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1012
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1012.0157_arXiv.txt
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We study the thermophysical properties of dense helium plasmas by using quantum molecular dynamics and orbital-free molecular dynamics simulations, where densities are considered from 400 to 800 g/cm$^{3}$ and temperatures up to 800 eV. Results are presented for the equation of state. From the Kubo-Greenwood formula, we derive the electrical conductivity and electronic thermal conductivity. In particular, with the increase in temperature, we discuss the change in the Lorenz number, which indicates a transition from strong coupling and degenerate state to moderate coupling and partial degeneracy regime for dense helium.
| 10 | 12 |
1012.0157
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1012
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1012.5128_arXiv.txt
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We study temporal variations in the amplitudes and widths of high-degree acoustic modes in the quiet and active Sun by applying ring-diagram technique to the GONG+ and MDI Dopplergrams during the declining phase of cycle 23. The increase in amplitudes and decrease in line-widths in the declining phase of the solar activity is in agreement with previous studies. A similar solar cycle trend in the mode parameters is also seen in the quiet-Sun regions but with a reduced magnitude. Moreover, the amplitudes obtained from GONG+ data show long-term variations on top of the solar cycle trend.
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The global and local analysis of solar acoustic modes has shown that the mode amplitudes and lifetimes are anti-correlated with the solar activity level, and depend strongly on the local magnetic flux \cite{Chaplin00,Rajaguru01,Howe04}. Most of the analyses have been done for average quantities without discrimination between quiet and active areas, while the solar cycle behavior of the quiet-Sun regions is not well understood. High-degree \textit{p}-mode lifetimes measured in the quiet-Sun regions using cross-correlation analysis show variations with the activity cycle \cite{Burtseva09a}. In the quiet Sun, no large-scale magnetic field concentrations, one of the potential damping sources, are present on the surface. Another explanation could be the activity-related variation of the convective properties near the solar surface. The studies of the solar cycle variations of the size of the solar granules so far arrive at contradictory conclusions \cite{Berrilli99,Saldana04}. An attempt to mask strong surface activity and analyze high-degree \textit{p}-mode amplitudes in the quiet-Sun regions at solar minimum and maximum indicated that the amplitude at solar minimum is higher than that at solar maximum \cite{Burtseva09b}, however, the effect introduced by the mask needs to be better understood. In this work we apply ring-diagram analysis to eight years of data to characterize the high-degree acoustic mode amplitudes and widths in the quiet and active Sun during declining phase of the activity cycle.
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The results for amplitudes and widths shown in this paper are for a multiplet $\ell$ = 440, $n$ = 2 ($\nu$ = 3.2 mHz). We find that the correlation of this multiplet with other multiplets in the 2.5$-$3.5 mHz frequency range over solar activity cycle is higher than 70$\%$. \subsection{Activity-related trend} The variations of amplitude and width as a function of time and MAI, computed from all patches and the quiet-Sun patches at the disk center using GONG data, are shown in Figure 1. According to the Mt.Wilson sunspot index data (see Figure 2), during most of the year 2008 the Sun was quiet with no or very few small sunspots appearing on the solar surface. Therefore, we took average of year 2008 consisting of 13 Carrington Rotations as a reference, and plotted the mode parameters relative to this mean value. The amplitudes increase by $\sim 10\%$ from the high solar activity period in 2001 till 2004. After this period, we find a long-term variation, but no clear association with activity cycle. It is interesting to note that the mode amplitude, derived from Variability of Solar Irradiance and Gravity Oscillations (VIRGO) and GONG data using global helioseismology technique, shows similar wiggles though they are small in amplitude relative to the extent of the solar cycle trend \cite{Salabert11}. If we ignore the unexplained long-term variations, we can conclude that the mode amplitudes obtained in our study increased from 2001 to 2008 by $\sim 22\%$ in all patches and by $\sim 16\%$ in the quiet-Sun regions. The widths decreased from 2001 to about 2008 by $\sim 9\%$, then started showing a rising trend. \begin{figure}[t] \begin{center} \includegraphics[width=0.95\linewidth]{fig11111.eps} \end{center} \caption{Variation of amplitude and width (relative to the mean value of 2008) as a function of time and MAI, computed in the quiet-Sun (closed blue circles) and all patches without MAI threshold (open green cicrcles). The mode parameters are derived from disk center patches using GONG data. Each point is an average over one Carrington rotation. Solid light-blue and dashed magenta curves represent the temporally smoothed values in the quiet-Sun and all patches, respectively. Red stars show the amplitude and width variation as a function of time, computed from MDI data in all patches at the disk center. In this case, each point is an average over one year. The typical error bars for the measurements are shown in the top panels.} \end{figure} \begin{figure}[t] \begin{center} \includegraphics[width=0.95\linewidth]{fig444.eps} \end{center} \caption{Normalized magnetic activity index, computed from MDI magnetograms in the quiet-Sun (dash-dot blue curve) and all regions (dashed green curve) at the disk center, plotted as a function of time. The solid red line represents the Mt.Wilson Sunspot Index. Note, that each curve is normalized to its maximum value to show the relative variation.} \end{figure} The variations in amplitude and width from MDI data (see Figure 1, top panels) show a solar cycle trend similar to that from GONG data. The long-term variations in the amplitudes obtained from GONG data are not seen in the MDI amplitudes. However, this could be due to the limited amount of data (only 135 days) that were analyzed with the MDI Dopplergrams. We plan to improve statistics on MDI data to confirm these results. The decrease in amplitudes and increase in widths with increase in MAI (bottom panels of Figure 1) are consistent with known results from global and local analysis. The linear correlation analysis shows that the correlation coefficient between amplitude and MAI is $-0.49$ for all patches and $-0.29$ for the quiet-Sun patches. The correlation coefficient between width and MAI is 0.73 for all patches and 0.48 for the quiet-Sun patches. Figure 2 shows the variation of the magnetic activity index over time. The MAI values of all patches at the disk center are well correlated with Mt. Wilson sunspot index. A weaker solar-cycle trend is also visible in the magnetic indices of the quiet-Sun patches. The solar cycle variation of the MAIs of the quiet-Sun regions computed from MDI magnetograms was also noted in \cite{Burtseva09a} and could be due solar cycle variation of the strong-field component of the quiet-Sun network \cite{Hagenaar03,Pauluhn03}. This suggests that the magnetic field plays a role in the activity related variations of the acoustic mode parameters in the quiet Sun. \subsection{Long-term variations} In an attempt to understand the long-term variations, on top of the solar cycle trend, seen in the amplitudes computed from GONG data, we restrict our analysis to 2 G $\le$ MAI $\le$ 3 G patches. We notice the long-term variation in the amplitude but no dependence on magnetic activity. This leads to the conclusion that the variations are probably not related to solar activity cycle. Moreover, according to our preliminary estimates from the patches at higher latitudes, the variations are present up to the regions centered at 52$^o$.5 in latitude. We plan to investigate this aspect in more detail after applying the necessary geometrical corrections. \ack This work utilizes data obtained by the Global Oscillation Network Group (GONG) Program, managed by the National Solar Observatory, which is operated by AURA, Inc. under a cooperative agreement with the National Science Foundation. The data were acquired by instruments operated by the Big Bear Solar Observatory, High Altitude Observatory, Learmonth Solar Observatory, Udaipur Solar Observatory, Instituto de Astrofisica de Canarias, and Cerro Tololo Interamerican Observatory. SOHO is a mission of international cooperation between ESA and NASA.
| 10 | 12 |
1012.5128
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1012
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1012.4012_arXiv.txt
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We present results of Hubble Space Telescope NICMOS H-band imaging of 73 of most luminous (i.e., log[$L_{IR}/L_{\odot}]>11.4$) Infrared Galaxies (LIRGs) in the Great Observatories All-sky LIRG Survey (GOALS). This dataset combines multi-wavelength imaging and spectroscopic data from space (Spitzer, HST, GALEX, and Chandra) and ground-based telescopes. In this paper we use the high-resolution near-infrared data to recover nuclear structure that is obscured by dust at optical wavelengths and measure the evolution in this structure along the merger sequence. A large fraction of all galaxies in our sample possess double nuclei ($\sim$63\%) or show evidence for triple nuclei ($\sim$6\%). Half of these double nuclei are not visible in the HST B-band images due to dust obscuration. The majority of interacting LIRGs have remaining merger timescales of 0.3 to 1.3 Gyrs, based on the projected nuclear separations and the mass ratio of nuclei. We find that the bulge luminosity surface density $\mathrm{L_{Bulge}/R_{Bulge}^2}$ increases significantly along the merger sequence (primarily due to a decrease of the bulge radius), while the bulge luminosity shows a small increase towards late merger stages. No significant increase of the bulge S\'{e}rsic index is found. LIRGs that show no interaction features have on average a significantly larger bulge luminosity, suggesting that non merging LIRGs have larger bulge masses than merging LIRGs. This may be related to the flux limited nature of the sample and the fact that mergers can significantly boost the IR luminosity of otherwise low luminosity galaxies. We find that the projected nuclear separation is significantly smaller for ULIRGs (median value of 1.2~kpc) than for LIRGs (mean value of 6.7~kpc), suggesting that the LIRG phase appears earlier in mergers than the ULIRG phase.
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Luminous IR Galaxies (LIRGs: i.e., $L_{IR} \geq 10^{11} L_{\odot}$) were discovered in the 1980s with the Infra Red Astronomical Satellite \citep[IRAS;][]{Soi84}, due to their strong emission at 12-100$\mu$m, and today are an important class of objects for understanding massive galaxy formation and evolution. Follow-up observations with ISO and Spitzer have shown that LIRGs comprise a significant fraction (50-70\%) of the cosmic infrared background and dominate the star-formation activity at $z\geq1$ \citep{Elb02, LeF05, Cap07, Bri07, Mag09}. Ultra Luminous IR Galaxies (ULIRGs: i.e., $L_{IR} \geq 10^{12} L_{\odot}$) are typically a mixture of disk galaxy pairs, interactions, or mergers \citep{Jos85, Arm87, San88b, Cle96,Mur96}. (U)LIRGs are also observed to be rich in molecular gas and dust \citep{San91, Sol97}; interaction drives the gas inward, resulting in enhanced star formation rates. In the local universe, ULIRGs are relatively rare, but are about a thousand times more frequent at $z\geq 2$ as shown by observations of the population of sub-mm galaxies \citep{Bla02, Cha05}.\par Recent discoveries have revealed that the central black hole mass is constrained by and closely related to properties of the host galaxy's bulge \citep[e.g. ][]{Mar03,Har04, Geb00}, which suggests a co-evolution of the central black hole and the galaxy itself. Given both the evidence that merging LIRGs likely evolve into massive elliptical and S0 galaxies \citep[e.g. ][]{Bar92, Gen01, Tac02} and that LIRGs are stellar nurseries \citep[see e.g.][]{San91, Sol97, Mag09, How10}, the study of nuclear regions of a complete sample of LIRGs spanning all merger and interaction stages allows us to investigate the coeval nature of both processes.\par The picture that mergers transform disc galaxies into elliptical galaxies is supported by numerical simulations \citep[e.g.][]{Bar90, Kor92}, potentially linking LIRG activity to a key step in the formation of the galaxy morphologies seen locally. Other recent simulations have shown evidence that this may be more complicated \citep[e.g.][]{Bou05, Rob06}, suggesting the importance of different merger-mass ratios, gas-fraction and feedback processes due to supernovae or AGN activity. Moreover, the exact number of merger progenitors is controversial; while the vast majority of interacting (U)LIRGs are likely two galaxy systems \citep[e.g.][]{Mur96,Vei02}, there may be also a small fraction of LIRGs that involve three galaxies and substantial star-formation between the merging nuclei \citep{Vai08}. While previous ground-based observations \citep[e.g.][]{Wri90, Jam99, Rot04} have revealed that a significant fraction of merger remnants seems to exhibit radial profiles similar to elliptical galaxies based on their large-scale appearance (at radii of several kpc), we are able to probe the first time the stellar light distribution in the central kpc for a large sample of nearby mergers. \par Merger models that incorporate hydrodynamics and star formation suggest that dissipation in mergers produces central starbursts which lead to a steepening of the central profile well above the extrapolation of a \cite{deV48} profile fit to the main body of the galaxy \citep{Her93, Mih94}. Indeed, recent observations have revealed such central ``extra light'' or ``cuspiness'' components for some elliptical galaxies \citep{Rot04, Fer06, Cot07, Kor09}. In particular, \cite{Hop08a} have shown that the observed extra light can be identified with the central density excess produced in simulations of gas-rich mergers and that such components are ubiquitous in the local cusp elliptical population \citep{Hop09a}, based on samples from \cite{Lau07} and \cite{Kor09}. However, although the idea that almost all the gas falls to the center to fuel a central starburst or AGN might be true for ULIRGs, this process may be less prevalent in LIRGs. In principle, a LIRG sample spanning all merger stages can be used to test the model of merger-induced cusp-building along the merger stage sequence. The Great Observatories All-sky LIRG Survey \citep[GOALS, see][]{Arm09} provides an excellent sample to compare to these models and observations. \par To reveal the detailed structure of the nuclear regions, where dust obscuration may mask star clusters, AGN and additional nuclei from optical view, we study 73 of most luminous (U)LIRGs in the GOALS sample using the Hubble Space Telescope (HST) Near Infrared Camera and Multi-Object Spectrometer (NICMOS). This high-resolution Near Infra-Red (NIR) imaging sample is unique not only in its completeness and sample size, but also in the proximity and brightness of the galaxies. Moreover, a wealth of additional information about the properties of these galaxies (e.g. mid-IR distribution, mid-IR line diagnostics, merger stage) is available, since the GOALS project combines multi-wavelength imaging and spectroscopic data from space (Spitzer, HST, GALEX, and Chandra) and ground-based telescopes. In particular, a detailed multi-wavelength study of the LIRG Markarian 266 \citep{Maz10} reveals nuclear and galactic-scale outflows, shock-heated gas and the presence of a dual AGN. \cite{Dia10} found that at least half of the MIR emission is extended for more than 30\% of local LIRGs, as well as an increasing compactness towards the final stage of major merger interaction. Furthermore, the relationship between the IR and UV properties \citep[e.g. the IRX-$\beta$ relation as a function of LIRG properties][]{How10} has been investigated. The optical GOALS HST images (ACS/WFC F435W and F814W filters) have clearly shown the inner spiral structure, dust lanes, extended filamentary emission, and star clusters in the nuclear regions and tidal tails of LIRGs \citep{Vav10, Kim10}.\par To understand the manner in which clusters, AGN, and nuclei evolve as a function of luminosity and merger stage, it is essential to first unveil the hidden nuclei and to model the nuclear structure. This requires both high spatial resolution ($\sim$100~pc) and the ability to penetrate the dust. Unfortunately, it is almost impossible to reveal the circumnuclear regions and to identify the nuclei in LIRGs with observations at optical wavelengths, simply because they are extremely dusty. On the other hand, NIR imaging is 1) less affected by dust extinction and hence provides an almost unobscured view on the nuclear region of galaxies, and 2) best suited to trace the old stellar population since NIR light is less biased by young blue stars that can dominate the optical light but contribute only a minor fraction by mass to the total stellar population. In particular, detailed information about the nuclear structure is crucial to reveal double or multiple nuclei to estimate merging time scales which can provide important constraints for the lifetime of the IR luminous phase and associated galaxy evolution scenarios. Therefore, HST NICMOS images are perfectly suited to observe more directly the true nuclear morphologies by combining the high resolution of the HST ($0.15\arcsec$ FWHM) which corresponds to 30 - 300~pc over the distance range 40 - 400~Mpc (with a median resolution of 106~pc at a median distance of 142~Mpc) at wavelengths where dust extinction is reduced by an order of magnitude compared to visual wavelengths.\par Our NICMOS imaging has revealed the presence of double and triple nuclei LIRG systems. Fitting the NIR light distribution allows us to extract the structural stellar components of the galaxies and to study their evolution along the merger stage sequence. Our NICMOS sample and the observations are described in \S~\ref{sec:obs}. The modeling of the stellar components using GALFIT \citep{Pen02} and their parameters are characterized in \S~\ref{sec:galfit}, and the results are presented in \S~\ref{sec:res} together with the numbers of double and multiple nuclei found in our sample. In \S~\ref{sec:dis} we discuss the number of nuclei that are obscured by dust and its implication for high-redshift studies, merger time scales, the evolution of the central stellar morphology along the merger stage sequence, and the role of AGN activity and young stellar populations on nuclear properties. For convenience, we refer to the full sample (log[$L_{IR}/L_{\odot}]>11.4$) as LIRGs despite the inclusion of a small number (17 systems, $\sim$23\%) of ULIRGs (log[$L_{IR}/L_{\odot}]>12.0$).
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\label{sec:dis} \subsection{Obscured Nuclei} Understanding the exact number of merger progenitors for LIRGs is important as these numbers provide information about merger timescales and subsequently LIRG activity timescales \citep[see][]{Mur96}. The study of the small and large-scale morphologies of LIRGs in the nearby universe allows us to understand the physical processes that drive galaxy evolution. While ULIRGs are predominantly merger systems, the fraction of merger systems and the merger timescales for LIRGs are more controversial: The imaging analysis of 30 local LIRGs \cite{Alo06} indicates that most of these LIRGs (log[L$_{IR}$/L$_\odot$]$= 11.0-11.9$) have prominent spiral patterns, and of those, a non negligible number are weakly interacting or even isolated systems. \cite{Bri07} find in an analysis of MIPS 24~$\mu$m detected and undetected mergers ($0.2 \leq z \leq 1.3$), that a larger fraction of LIRGs appear later, in the merger phase, than in the pre-merger, close pair phase. However, it is expected that the fraction of major mergers increases with IR luminosity. Thus far, the trigger mechanisms and end products of LIRGs have not been constrained as well as those of local ULIRGs.\par We find for our sample that the fraction of LIRG systems (log[L$_{IR}$/L$_\odot$]$= 11.4-12.5$) with at least two interacting nuclei is 63\% (see Fig.~\ref{nuclei_H-B}). Note that the fraction of major merger for the entire GOALS sample (log[L$_{IR}$/L$_\odot$]$>11$) is $\sim$50\% (based on Spitzer IRAC and HST data). Interestingly, the comparison with our HST ACS B-band images revealed that roughly half of the double nuclei seen in the HST NICMOS H-band are obscured by dust (not visible in the B-band). Therefore, NIR observations of LIRGs and ULIRGs at high-redshift (z$>$2), which correspond to the B-band rest-frame range, might require significant correction factors of $\sim$2 or higher to accurately estimate the true number of multiple nuclei systems.\par In particular, \citet[using HST NICMOS F160W]{Das08} found that the fraction of binary systems in ULIRGs at redshift 2 is roughly 50\%, which is significantly smaller than that of the local ULIRG population \citep[nearly 100\% of objects in the IRAS 1 Jy sample show signs of major merger interaction][]{Kim02, Vei02}. Also in studies of \citet[MIPS 24~$\mu m$ and HST ACS F850LP filter]{Bel05} and \citet[MIPS 24~$\mu m$ and HST ACS Viz band]{Mel05} major merger appear to make up a smaller fraction of z$\simeq0.7$ LIRGs compared to local systems. The strong wavelength dependence on the number of visible nuclei may explain some of the apparent discrepancy between the (U)LIRG population at local and high-redshift, particularly in identifying late stage mergers when tidal tails have faded and double nuclei are surrounded by dust. \subsection{Merger Time Scales} The space density of LIRGs combined with an estimate of the time-scales of their power sources (star formation and/or AGNs) provides important constraints for galaxy evolution scenarios that link the IR-luminous phase with the luminous, unobscured AGNs observed in Quasar Hosts \citep{San88b}. The merger time scale depends on the separation between the multiple nuclei observed in our NIR images (Fig.~\ref{hist_dist}) and we follow the approach by \cite{Mur96}, namely: In a first stage we adopt a constant velocity $v_r$ until the nuclei are 10~kpc apart. Because of the small range in nuclear separations, and the large uncertainties involved, we treat the radial velocity as constant. In this case, the time from observation until the nuclei are within 10~kpc is ($r-10$~kpc)/$v_r$ with $r$ as the projected nuclear separation. The time spent in the second stage is estimated by calculating the dynamical time scale times the mass ratio of the two nuclei, given as $(M_1/M_2)\cdot 2\pi r/v_c$ with the mass ratio $M_1/M_2$, the projected nuclear separation $r$ and the orbital velocity $v_c$ \citep[see][]{Mur96, Bin08}. As the nuclei spiral together, the relative velocity is similar to the circular velocity at that radius \citep{Bin08}. In the calculations that follow we adopt an average radial velocity $v_r$ and an average circular velocity $v_c$, of 200~km~s$^{-1}$ for all the double nucleus galaxies in our sample \citep{Mur96}. These average velocities are supported by simulations \citep[e.g. see Fig.~15 in][]{Xue08} and the distribution of the line-of-sight velocities differences between the galaxy pairs in our sample (typically between 60--180~km~s$^{-1}$; note that the radial velocity is on average a factor of 1.5 larger than the line-of-sight velocity component). It is important to note that for an individual galaxy all we can estimate from the images and these simple dynamical arguments is the time from the present until the time of final merger, when the nuclei of interacting galaxies coalesce. We must rely on the distribution of these time scales within the sample, along with the morphologies, to allow a calculation of the average remaining merger time scales of our sample. For our LIRG sample (log[L$_{IR}$/L$_\odot$]$>11.4$) we find a median remaining time until the nuclei merge ($t-t_m$) of $4.3 \cdot 10^8$~yrs. \begin{figure}[ht] \begin{center} \includegraphics[scale=0.6]{f16.eps} \end{center} \caption{Histogram of the remaining merger time scale, $t - t_{merg}$ (see text for details) with a median time scale of $4.2 \cdot 10^8$~yrs. Fourteen apparent single nuclei in galaxies that exhibit prominent interaction features, such as tidal tails, are also included in white, converting our angular resolution limit of $0.15\arcsec$ to the linear size at their corresponding distance and assuming a nuclei mass ratio of 1:1.} \label{hist_time} \end{figure} As shown in Fig.~\ref{hist_time}, we find two peaks in the numbers of LIRGs as function of the remaining merger time scale: The largest occurs at 0.3~Gyr$< [t - t_{merg}] < $1.3~Gyr, roughly representing the first passage of interacting galaxies, while the second peak occurs at the final coalescence of the nuclei (most at $[t - t_{merg}] < 10^7$~yr), including all galaxies with apparent single nuclei that show interaction features. Note that the second peak represents the same number of LIRGs and ULIRGs (each 6). The fact that we see two peaks is consistent with merger simulations \citep{Cox08, Hop08b}, which show two star formation rate (SFR) peaks as function of merger time. Furthermore, the time-scales of these two peaks are roughly consistent with the two SFR peaks as seen the in merger simulations, but we see slightly more LIRGs at $[t - t_{merg}] < 1$~Gyr. In detail, the peak at $[t - t_{merg}] \sim 1$~Gyr is a bit broader (0.3~Gyr$< [t - t_{merg}] < $1.3~Gyr) than we would expect from the merger models (0.8~Gyr$< [t - t_{merg}] < $1.3~Gyr), indicating that we observe more LIRGs at shorter timescales. The fraction of interacting LIRGs is larger for the peak at $\sim$1~Gyr (53\% of interacting LIRGs) than for the peak at the nuclear coalescence (26\%), which likely corresponds to the broader width ($\sim$0.7 Gyr) of the SFR peak at $\sim$1~Gyr than the narrow peak ($\sim$0.3 Gyr) at the nuclear coalescence \citep[see][]{Cox08, Hop08b}. As our GOALS sample is a flux-limited snapshot of galaxies passing through the LIRG and ULIRG phase, the data suggest that galaxies spend a longer time in the LIRG phase, which is more easily associated with the first peak in the model SFR at $[t - t_{merg}] < 1$~Gyr. Furthermore, the IR luminosity is slightly larger for LIRGs at $[t - t_{merg}] \simeq 0$ (median log[$L_{IR}/L_{\odot}] = 11.7$) than for LIRGs at $[t - t_{merg}] \simeq 1$ (median log[$L_{IR}/L_{\odot}] = 11.9$). This is consistent with our finding that LIRGs in late stages of merging have significantly larger total IR luminosities (roughly a factor of two) than pre- or non-merging LIRGs. We also discuss this point below (see Fig.~\ref{IR_dist} as well). \par To estimate the timescale that LIRGs spend in the final nuclear coalescence (post-merger stage), we have applied the following approach: In principle, the post-merger time scale can be estimated by multiplying the merger time-scale with the ratio of galaxies with single nuclei to double nuclei. However, this approach might be oversimplified since not all LIRGs might have undergone a merging process, and instead other mechanisms may be responsible for a significant fraction of LIRGs. In fact, 23\% of LIRGs in our HST NICMOS sample show no major interaction features in the H-, I-, or B-band (merger stage 0 and 1, see \S~\ref{subsec:res-GALFIT}) and hence it is unlikely that they have undergone a previous major merger. Therefore, we take into account only galaxies that are classified as mergers (merger stage 2---6, see \S~\ref{subsec:res-GALFIT}). The time spent in the pre-merger stage for a random galaxy is statistically roughly twice the average observed time required to complete the merging $t-t_m$ \citep[see][]{Mur96}, or about 10$^9$yrs for our sample. Thus, the time spent in the post-merger stage is the median merger time of LIRGs with double nuclei ($2 \times 0.43$ Gyr) times the ratio of the number of single nuclei with interaction features (15 LIRGs) to double nuclei (45 LIRGs). This results in an average post-merger time of $\sim 2.9\cdot 10^8$~yrs, which is very similar to the typical timescale of the SF burst of 0.3 Gyr as seen in major merger simulations \citep{Cox08, Hop08b}.\par Since we expect most ULIRGs to pass through a LIRG phase, we might expect to see a correlation of merger stage, nuclear separation and IR luminosity \citep{Mur01, Vei06}. A comparison of LIRGs and ULIRGs in our HST NICMOS sample, presented in Fig.~\ref{IR_dist}, reveals that the projected nuclear separation is significantly different between both populations (probability level of 98\% in KS test) with a mean (median) projected separation of 3.4~kpc (1.2~kpc) and 11.1~kpc (6.7~kpc) for ULIRGs and LIRGs, respectively. \par The results for our subsample of ULIRGs (log[$L_{IR}/L_{\odot}]>12.0$) are also in good agreement with a study for a larger sample of ULIRGs \citep{Mur96}, which is based on ground-based observations (seeing of $\Delta R \simeq 0.8\arcsec$ for K band images), and hence the lower spatial resolution limits the ability to resolve double nuclei. Including galaxies with apparent single nuclei (upper limits on the separation of double nuclei), most of the ULIRGs in the sample of \cite{Mur96} (more than 60\%) have nuclear separation between 1-4~kpc (median of $\sim$1.8~kpc) which is very similar to our results for the ULIRG population (median of $\sim$1.2~kpc). Note that we have an overlap of ten ULIRGs with the sample of \cite{Mur96} and we found for one ULIRG, namely IRAS~F19297-0406, two nuclei with a projected nuclear separation of $0.77\arcsec$ that was previously characterized as an apparent single nucleus by \cite{Mur96} due to limited spatial resolution of the R-band image (1.9\arcsec). However, only our entire sample covering a IR luminosity range of log[$L_{IR}/L_{\odot}]>11.4$, allows us to see a strong decrease of the projected nuclear separation as a function of IR luminosity, as shown in Fig.~\ref{IR_dist}. \begin{figure}[h] \begin{center} \includegraphics[scale=0.7]{f17.eps} \end{center} \caption{The projected nuclear separation as a function of IR luminosity for the observed double nuclei (plus marker), as well as apparent single nuclei (arrows down, converting our angular resolution limit of $0.15\arcsec$ to the linear size at their corresponding distance) whose host galaxies exhibit interaction features. The mean and median values of the projected nuclear separation are shown as solid and dashed horizontal lines, respectively.} \label{IR_dist} \end{figure} \clearpage \subsection{The Evolution of the Central Stellar Structure along the Merger Stage} \label{subsec:evo} Evidence suggests that ULIRGs may evolve into elliptical galaxies once the starburst subsides and the gas is either used up or expelled in a wind \citep{San88b, Hec90, Gen01, Vei06}. Therefore, the stellar surface brightness profiles of LIRGs should provide a glimpse of the process of bulge and black hole building. Unlike previous studies, which focused primarily on ULIRGs (log[$L_{IR}/L_{\odot}]>12.0$), the GOALS NICMOS sample targets the nuclear regions of all merger systems in the IRAS Bright Galaxy Sample \citep[RGBS;][]{San03} with log[$L_{IR}/L_{\odot}]>11.4$. Because of the low red-shift range of our sample ($0.01 < z < 0.05$), the galaxies are bright and well-resolved due to their large angular size.\par Although we find a significant increase in the bulge luminosity surface density (a factor of 5 to 60) along the merger sequence, the growth of the BH mass (bulge luminosity) toward later merger stages is only marginal: A factor of $\sim1.8$ in bulge luminosity from merger stage 3 to 5/6 with a probability level of $\sim$90\% that both merger stage populations are not drawn from the same parent population (KS and MWU test, see Table~\ref{tab:ks_test}). Two possibilities may explain why we do not see a a more significant increase of the bulge luminosity: First, as the merging proceeds, the bulge may get partially disrupted or stripped and may assemble again after the LIRG phase has passed. Thus, our measured bulge luminosities may miss a significant bulge fraction at the late merger stage, resulting in smaller estimated black hole masses. A simpler explanation could be that the flux limited nature of our HST LIRG sample (log[$L_{IR}/L_{\odot}]>11.4$) excludes galaxies with smaller bulge masses in single isolated galaxies and galaxies in the pre-merging stage. Indeed, we find evidence for this, as the LIRGs in our sample that show no evidence for interaction features (single isolated galaxies and merger class 1) have on average a significantly larger BH mass (a factor of two). Since LIRGs are drawn from the top end of the IR luminosity function and mergers boost IR luminosity, we might expect the non-interacting LIRGs to be more luminous and hence more massive than the progenitors of interacting LIRGs. The interacting LIRGs would tend to include more systems with smaller bulges and then grow their bulges during the merging process. If this is the case, the non-interacting LIRGs would tend to have more massive bulges - exactly as we observed in our sample. A detailed comparison of the bulge masses of non-interacting spiral galaxies with log[L$_{IR}$/L$_\odot$]$ < 11.4$ would be needed to test this hypothesis. \subsection{Merger-induced Building of Central Starbursts: Comparison to Models} \label{subsec:model} Dissipation in mergers can generate central starbursts, imprinting a central ``extra light'' component into the surface brightness profiles of merger remnants \citep{Her93}. Such an excess of the central light or ``cuspiness'' in the surface brightness profiles has been found for some elliptical galaxies \citep{Fer06, Cot07, Kor09, Hop09a}. This can be explained by the assumption that the envelopes of cusp ellipticals are established by violent relaxation in mergers acting on stars present in gas-rich progenitor disks, while their centers are structured by the relics of dissipational, compact starbursts \citep{Hop09b}. Given the evidence that ULIRGs likely evolve into massive elliptical and S0 galaxies \citep{Gen01, Tac02} and the fact that LIRGs and ULIRGs host powerful starbursts, the study of our LIRG sample as a function of merger stage allows us in principle to test the build-up process of nuclear starbursts. \par To test such a scenario we studied the following parameters as a function of merger stage (1---6, see \S.~\ref{subsec:res-GALFIT}): (1) the effective bulge radius $R_{Bulge}$, (2) the bulge luminosity $L_{Bulge}$, (3) the bulge S\'{e}rsic index $n_s$ which defines the steepness of the inner bulge profile (see Fig.~\ref{sersic_stage}), (4) the bulge surface brightness defined as $L_{Bulge}/R^2_{Bulge}$, and (5) the central light concentration given by the ratio of core excess luminosity to total bulge luminosity (see Fig.~\ref{res_stage}). We have compared our results with recent models of \cite{Hop09a}, which are based on N-body simulations and smoothed particle hydrodynamics, account for radiative cooling and for heating by a UV background, and incorporate a subresolution model of a multiphase interstellar medium (ISM) to describe star formation and supernova feedback. Since the model of \cite{Hop09a} focuses on local gas-rich mergers (unlike high-redshift simulations which have a factor of 2--5 larger gas content), it is ideally suited for a comparison to our data-set of local major mergers. \subsubsection{Bulge Radius and Surface Brightness} The merger model of \cite{Hop09a} predicts a decrease of the effective bulge radius along the merger process due to gas inflow: Tidal torques excited during major merger provide the fuel to power intense starbursts boosting the concentration and central phase-space density. As shown in \cite{Hop09a} a gas inflow of e.g $\sim$10\% shrinks the effective bulge radius to about half its previous size and subsequently enhances the bulge surface brightness. Our findings show the same trend as the model predictions (see Fig.~\ref{rad_stage} and Fig.~\ref{dens_stage}), given a decrease (increase) of the effective bulge radius (surface brightness) by a factor of $\sim 2-3$ ($5-20$), and hence provide evidence for the idea that gas inflow can dramatically reduce the apparent bulge radius during a merger. \subsubsection{Bulge S\'{e}rsic Index} Due to the violent relaxation of stars during a merger, the model of \cite{Hop08a} predicts an increase of the bulge S\'{e}rsic index $n_s$ towards late merger-stage (from $t - t_{merg}=0.5$~Gyr to $t - t_{merg}=0$~Gyr). We do not find a significant increase of $n_s$ (see Fig.~\ref{sersic_stage} and Table~\ref{tab:ks_test}). This discrepancy can have several reasons, such as that young stellar populations and age (metallicity) gradients can shift $n_s$ significantly and/or tidal interactions are changing $n_s$ during the more active/merging phases. \subsection{The Role of AGN Activity and Young Stellar Populations on Nuclear Properties} \label{subsec:dis-AGN} Throughout most of this paper, we have assumed that the near-infrared H-band data is an accurate tracer of the stellar mass, as it is less affected by dust extinction (than the UV or optical), and has less of a contribution from young, blue stars. However, LIRGs and ULIRGs are known to be powered by starbursts and AGN, which can, in some cases, significantly contribute to the near-infrared light. In the following sub-sections we discuss the possible contributions of central starbursts and AGN to the near-infrared light. \subsubsection{Young Stellar Populations} Evidence suggests that asymptotic giant branch (AGB) and red supergiant (RSG) stars can dominate the NIR in the center of some starburst galaxies \citep{Arm95, Rot10, Mel10}. Since high-spatial resolution spectroscopy is required to identify and resolve the young stellar populations, their spatial distribution is unknown for our sample. However, two scenarios can constrain the possible effects of young stellar populations on our measured nuclear properties: (1) If young stars form in the core of the galaxies at scales $\lesssim$ 300~pc, these populations would be taken into account by our measured PSF component and hence not affect our measured bulge properties. (2) In case young stars are more widely distributed on scales typical for stellar bulges (median effective bulge radius for our sample is 0.7~kpc), they could affect our measured properties. If young stars form in a disk, they would manifest an exponential radial profile, and a central GALFIT component with a S\'{e}rsic index $n_s\simeq1$ (exponential disk) would be expected, instead of $n_s>1$ (more typical for old stellar bulges), which we find for most of our galaxies (see Fig.~\ref{rad_stage}). If this is not the case, namely that the majority of young stars form at bulge scales ($>300$~pc, not in in the core), are distributed in a spheroid rather than in a disk, and dominate the NIR, then the H-band bulge luminosity may not be a reliable estimate of the central BH mass. However, this does not affect our comparison of the central stellar structure along the merger stage to merger models (\S.~\ref{subsec:model}), since young stellar contributions are already taken into account in the merger simulations \citep{Hop09a, Hop09b}. \subsubsection{AGN Activity} As described in \S~\ref{subsec:res-GALFIT} we find for some galaxies a large ratio of core excess to total bulge luminosity $L_\mathrm{excess}/L_\mathrm{Bulge}$. One possible explanation of such a bright central light emission could be that the LIRG activity is caused by AGN heating. In general, lower Equivalent Widths (EQWs) of the 6.2~$\mu$m PAH feature are associated with AGN activity \citep[see e.g.][]{Gen98, Stu00, Arm07, Des07} because the hot dust continuum increases and the hard AGN photons may destroy/ionize the PAH molecules. \cite{Pet10} estimated that AGN are responsible for $\sim$12\% of the total bolometric luminosity of local LIRGs based on several mid-IR line diagnostics measured with the Infrared Spectrograph on Spitzer. \par We find no significant trend between the ratio of core excess to bulge luminosity and the PAH EQW (see Fig.~\ref{EQW_res}) of the 6.2~$\mu$m PAH feature \citep{Pet10}. However, one example of a galaxy where a central AGN might dominate the NIR emission is AM~0702-601: A galaxy in our sample with a very small PAH EQW (0.037~$\mu$m) but also exhibits one of the most significant core excess luminosity fractions ($L_\mathrm{excess}/L_\mathrm{Bulge} = 1.5$). Counter examples, where the core excess light fraction in the H band and PAH EQW are not correlated: ESO~060-IG016 and NGC~3690 East, which have very small PAH EQWs ($<$ 0.27~$\mu$m) and hence likely large AGN contribution, but they exhibit no significant core excess light fraction ($L_\mathrm{excess}/L_\mathrm{Bulge} < 0.001$). One possible explanation might be that those AGN systems are so deeply embedded in dust that they don't show a central light excess in the H-band, in contrast to some ULIRGs which show a very large light excess in the center when an AGN is present \citep[see e.g.][]{Sur99}. We find also galaxies (ESO~239-IG002, WKK~2031, and NGC~1614) with a significant core light excess ($L_\mathrm{excess}/L_\mathrm{Bulge} > 0.05$) but relatively large PAH EQWs (0.4 - 0.7~$\mu$m), suggesting that those galaxies build up a concentrated stellar ``cusp'' in the center due to luminous nuclear starbursts. The sources with the largest core excess ($L_\mathrm{excess}/L_\mathrm{Bulge} > 0.05$) are not responsible for the decrease in the bulge radius along the merger sequence, as the central PSF (even if it is due to an AGN) is already taken into account by the GALFIT fitting process (see \S.~\ref{subsec:model}). We obtain roughly the same results for the bulge properties along the merger sequence if we exclude the galaxies with small 6.2~$\mu$m PAH EQW (EQW $<$ 0.27~$\mu$m) from our sample.\par \begin{figure}[ht] \begin{center} \includegraphics[scale=0.7]{f18.eps} \end{center} \caption{Equivalent Width (EQW) of the 6.2~$\mu$m PAH feature as function of the core excess luminosity to total bulge luminosity. A smaller EQW indicates a larger contribution of AGN generated emission to the mid-IR emission. No significant trend is visible between 6.2~$\mu$m PAH EQW and core excess luminosity (see text).} \label{EQW_res} \end{figure} AGN activity might also effect the observed NIR luminosity and hence one would expect a possible contribution to the measured bulge luminosity. This could subsequently explain the increase of the BH mass (bulge luminosity) with larger IR luminosity (see Fig.~\ref{bulge_IRlum} and Eq.~\ref{eq:IR_bulge}). However, a possible contribution of AGN activity to the bulge luminosity is not very likely since most of our bulges are resolved and have typical ranges of 0.5-2~kpc while the main dust heating by AGN (500-2000K) is generated on scales of (10-100)~pc \citep{Soi01}. Furthermore, the central light excess expected from an AGN contributes on average less than 5\% to the total bulge luminosity for our sample (see Fig.~\ref{EQW_res}).
| 10 | 12 |
1012.4012
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1012
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1012.0588_arXiv.txt
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{ The diffuse gamma ray emission from astrophysical backgrounds in our Galaxy and the signal due to the annihilation or decay of Dark Matter (DM) in the Galactic Halo are expected to have a substantially different morphology and spectral signatures. In order to exploit this feature we perform a full sky and spectral binned likelihood fit of both components, using data collected during the first 21 months of operation of the Fermi-LAT observatory. Preliminary constraints are presented on the DM annihilation cross section and decaying rate for various masses and annihilation/decay modes.} \FullConference{Identification of Dark Matter 2010-IDM2010\\ July 26-30, 2010\\ Montpellier France} \begin{document}
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The Dark Matter (DM) properties, despite of being one of the most widely investigated topics in contemporary fundamental physics, are still largely undetermined. However, the current and near future experiments are cutting more and more into the parameter space region predicted for the most popular type of DM candidates, i.e. Weakly Interacting Massive Particles (WIMPs) (for a review see \cite{Bertone:2004pz}). In particular, high energy gamma ray astronomy is an exciting probe into these fundamental questions. In fact, gamma rays are products of hadronization and radiative loss processes, and are therefore unavoidably emitted in DM annihilation and decay. Also, the propagation of gamma rays is unaffected by the Galaxy, and therefore the data contain the information on the morphology of the emission region, together with the spectra (unlike, i.e. cosmic rays). Not least, the Large Area Telescope (LAT), onboard the Fermi gamma-ray observatory \cite{Atwood:2009ez}, is now providing unprecedented high quality data and statistics which makes it timely to investigate DM in gamma-rays. We will focus in particular on DM signatures in the gamma-ray diffuse emission. Diffuse data contain about $90\%$ of the LAT photons, and is extremely rich with information. Its {\em Galactic component} encodes information on the propagation and origin of cosmic rays, distribution of cosmic ray sources, magnetic and starlight fields in our Galaxy, while its {\em extragalactic component} is a superposition of all unresolved sources emitting gamma rays in the Universe, providing a signature of energetic phenomena over cosmological scales. {\em Both} components are expected to receive a contribution from {\em dark matter} annihilation/decay: the Galactic signal from the smooth dark matter halo and Galactic substructures, while the extragalactic one from self-annihilation of dark matter in all halos at all redshifts. The dark matter searches with the diffuse data by the Fermi-LAT team is being pursued with the Extragalactic (Isotropic) Signal, by using the intensity and spectral shape of the signal \cite{Abdo:2010dk} and angular anisotropies \cite{jsg}. Here we will investigate the DM signal originating in our Halo (see also the analysis presented in \cite{ba}). Due to the bright sources present in the Galactic Center and the bright diffuse emission along the plane, it has been argued e.g. in \cite{Serpico:2008ga}, that the region of the inner Galaxy, extending $10-20~\deg$ away from the Galactic Plane, is the most promising in terms of a signal to background ratio $S/N$. As an additional advantage, the perspectives to constrain a DM signal in that region, become less sensitive to the unknown profile of the DM halo. In particular the $S/N$ of quasi-cored profiles is only a factor $\sim 2$ worse than for the NFW case (compared to an order of magnitude of uncertainty, when one considers the Galactic center region). The increase in $S/N$ ratio away from the plane is further emphasized for DM models in which DM annihilations result in a significant fraction of leptons in the final state. These leptons diffuse in the Galaxy and produce high energy gamma rays mainly through the Inverse Compton (IC) scattering on the Interstellar Radiation Field (ISRF). By diffusing away from the Galactic Center region, electrons produce an extended gamma ray signal which is easier to distinguish from the astrophysical signal (\cite{Borriello:2009fa}). In this work we build on this idea, and consider the whole sky data, while masking out the Galactic Plane. We test DM annihilation and decay signal in the smooth Galactic halo, by performing a spatial and energy fit. We also take into account the most up to date modeling of the diffuse signal of astrophysical origin \cite{paper2}. However, at present, this approach faces few limitations. The high statistics of the Fermi data is currently unmatched with our knowledge of the cosmic ray production and propagation mechanisms, as determined by the cosmic ray experiments and cosmic ray source observations. In other words, current diffuse models, which rely on the gamma ray independent measurements, have both significant energy and spatial residuals, on small and large scales. This poses a significant challenge to the statistical approach to constrain the presence of an additional (e.g. dark matter) component in the fit (see also \cite{ba}). In this talk, we present preliminary limits, based on the astrophysical background model, found to be one of the most conservative in terms of dark matter searches, among the models currently considered in analysis of the diffuse emission, by the Fermi-LAT collaboration \cite{paper2}.
| 10 | 12 |
1012.0588
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1012
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1012.2919_arXiv.txt
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The latest observation of a Shapiro delay of the binary millisecond pulsar J1614-2230 by Demorest et al. \cite{demorest2010} yielded the pulsar gravitational mass to be 1.97$\pm$0.04 M$_\odot$, the heaviest observed pulsar to-date. This result produces a stringent constraint on Equation(s) of State (EoS) of high density neutron star matter. One of the main conclusions of Demorest et al. was that their result makes the presence of non-nucleonic components in the neutron star matter unlikely. We compare the result with our recent work and conclude that hyperons in high-density matter are fully consistent with the observation and that their presence is a necessary consequence of general physical laws.
| 10 | 12 |
1012.2919
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1012
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1012.5014_arXiv.txt
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We study the cosmic microwave background temperature and polarisation spectra sourced by multi-tension cosmic superstring networks. First we obtain solutions for the characteristic length scales and velocities associated with the evolution of a network of F-D strings, allowing for the formation of junctions between strings of different tensions. We find two distinct regimes describing the resulting scaling distributions for the relative densities of the different types of strings, depending on the magnitude of the fundamental string coupling $g_s$. In one of them, corresponding to the value of the coupling being of order unity, the network's stress-energy power spectrum is dominated by populous light F and D strings, while the other regime, at smaller values of $g_s$, has the spectrum dominated by rare heavy D strings. These regimes are seen in the CMB anisotropies associated with the network. We focus on the dependence of the shape of the B-mode polarisation spectrum on $g_s$ and show that measuring the peak position of the B-mode spectrum can point to a particular value of the string coupling. Finally, we assess how this result, along with pulsar bounds on the production of gravitational waves from strings, can be used to constrain a combination of $g_s$ and the fundamental string tension $\mu_F$. Since CMB and pulsar bounds constrain different combinations of the string tensions and densities, they result in distinct shapes of bounding contours in the $(\mu_F, g_s)$ parameter plane, thus providing complementary constraints on the properties of cosmic superstrings.
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\label{intro} Although it has been established that a network of cosmic strings cannot source the majority of the observed cosmic microwave background (CMB) temperature anisotropy \cite{Albrecht:1997nt}, the CMB can still provide a distinctive signature of their presence through the specific primordial B-mode polarisation spectrum \cite{Seljak:1997ii,Battye:1998js,PogosianTye,SelSlo,Pogosian:2007gi,Bevis:2007qz,Urrestilla:2008jv,Mukherjee:2010ve}. The spectrum generated by strings is different from the one generically produced from tensor modes arising in inflationary scenarios, and future probes of the B-mode should be able to reveal the presence of cosmic strings, even if strings contribute as little as $0.1\%$ to the CMB temperature anisotropy \cite{SelSlo,Pogosian:2007gi,Bevis:2007qz,Urrestilla:2008jv,Mukherjee:2010ve}. Interest in cosmic strings has revived following the realisation that they can arise in superstring theory \cite{Sarangi:2002yt, Jones:2003da}, for example in models of brane inflation \cite{Dvali:1998pa,models,Kachru:2003sx,Burgess:2004kv} (for a review see \cite{Copeland:2009ga}). Cosmic superstrings can have small tensions ($10^{-12} \gsim G\mu \gsim 10^{-7}$ \cite{Sarangi:2002yt, Jones:2003da, Polchinski:2004ia}), can be effectively stable over cosmological timescales, and can stretch over cosmological distances \cite{Copeland:2003bj,Leblond,Dvali:2003zj}. Hence, they can have interesting cosmological implications. Furthermore, their intercommutation probabilities can be significantly less than unity \cite{Jones:2003da, Jackson:2004zg, Hanany:2005bc,Jackson:2007hn} and, because of the charges present on them, they can zip together to form Y-junctions (trilinear vertices), leading to more complicated networks than those usually considered in the case of `standard' Abelian cosmic strings. Understanding the imprint of such additional network features on observables, such as CMB temperature and polarisation, is a step that may lead to interesting new constraints on the basic parameters of string theory, such as the string coupling $g_s$ and the fundamental string tension $\mu_F$. Several approaches have been developed to model the evolution of cosmic string networks, and an interesting recent attempt to extend them to cosmic superstring networks -- which contain different types of string -- is due to Tye, Wasserman and Wyman \cite{TWW}. Their model, based on the ``velocity dependent one-scale'' model of Martins and Shellard \cite{VOS,VOSk}, describes evolution of a multiple tension string network (MTSN) under the assumption that all types of strings have the same correlation length and root-mean-square velocity. By studying the evolution of the number density of strings, they find that scaling is achieved when the energy associated to the formation of junctions is assumed to be radiated away. This model has been extended in \cite{NAVOS}, where the authors assigned a different correlation length and velocity to each string type, and enforced energy conservation at each junction. Scaling is again achieved (with different number densities), but not as generically as in \cite{TWW}. In a complementary approach, a number of authors have studied the kinematics of cosmic string collisions \cite{CopKibSteer1,CopKibSteer2}. When two Nambu-Goto (NG) strings (of generally different tensions) collide, rather than intercommuting in the standard way, they can form two junctions and a linking string of a third tension. Kinematically this can only occur if the relative orientation, velocity and string tensions lie in certain ranges. In \cite{CopFirKibSteer}, the authors extended their earlier studies to $(p,q)$-cosmic superstrings by modifying the NG equations to take into account the additional requirements of flux conservation. Once again the kinematic conditions required for the formation of Y-junctions were established, with results very similar to the ones obtained for NG strings. These kinematic constraints have been checked quite extensively with dynamical field theory simulations of strings collisions, and the agreement is (generally) good \cite{Sakellariadou:2008ay,Salmi,Urrestilla:2007yw,Bevis:2008hg,Bevis:2009az}. Recently, in \cite{Avgoustidis:2009ke} these constraints have been incorporated into the generalised MTSN velocity one-scale model of \cite{NAVOS}, leading to new conditions required for scaling and thereby providing the most complete model of cosmic superstring evolution to date. In this paper, we use the model of \cite{Avgoustidis:2009ke} to study the evolution of a cosmic superstring network for different values of the string coupling $g_s$ and different charges $(p,q)$ on the strings. We find that in all cases the three lightest strings, i.~e. the $(1,0)$, $(0,1)$ and $(1,1)$ strings, dominate the string {\it number} density. When the string coupling is large, $g_s \sim {\cal O}(1)$, most of the network {\it energy} density is in the lightest $(1,0)$ and $(0,1)$ strings (respectively F and D strings), whose tensions are approximately equal. At smaller values of $g_s \sim {\cal O}(10^{-2})$, the $(1,0)$ string becomes much lighter than both the $(0,1)$ and $(1,1)$ strings, and dominates the string {\it number} density. However, because of their much larger tension, the {\it energy} density of the network at small couplings can be dominated by the rarer $(0,1)$ and $(1,1)$ strings. The existence of these two distinct limiting scaling behaviours at large or small values of $g_s$ is quite generic, although the specific details are somewhat dependent on the model-dependent value of the effective volume of the compactified dimensions. In either of the two limiting regimes, the {\it energy} density of the multi-tension network is effectively dominated by strings of one tension. With the scaling solutions to hand we then focus on the CMB imprints of these networks, using a modified version of the publicly available code CMBACT \cite{Pogosian:1999np, cmbact}. In particular, we extend the Unconnected Segment Model (USM), first introduced in \cite{Albrecht:1997nt,ABR99}, to describe the MTSN of \cite{Avgoustidis:2009ke} and implement it in CMBACT to obtain the CMB temperature and polarisation spectra. We find that for sufficiently large values of the parameter $w$, which is inversely proportional to the effective volume of the compactified dimensions (see Eq.~(\ref{wpar}) below), the two limiting regimes, one with the network energy dominated by light populous strings and the second with it dominated by rare heavy strings, can each produce distinct shapes for their CMB B-mode polarisation. In particular, for $w \sim 1$, the position of the peak in the B-mode spectrum is at $\ell \approx 770 $ for $g_s=0.9$ and at $\ell \approx 610$ for $g_s=0.04$. This opens up the exciting possibility that upcoming observations may not only constrain the overall contribution of strings, but in fact rule out certain values of the string coupling. Namely, the combination of the normalisation and the peak position of the B-mode spectrum can point to a particular combination of $g_s$ and the fundamental string tension $\mu_F$. It is common to report constraints on standard cosmic strings in terms of bounds on the single dimensionless string tension $G\mu$. These bounds have an implicit assumption on the number density of strings corresponding to the usual Abelian Higgs model strings with intercommutation probability ${\cal P} =1$. However, in a more general situation such as that of cosmic superstrings considered here, there can be smaller intercommutation probabilities as well as strings of different tensions. As a result, each type of string will in principle have a different number density: the same fraction of CMB anisotropy can be sourced either with many light strings or with a few heavy ones. In general, each type of observational bound will constrain a different combination of the string tensions and densities (which, for cosmic superstrings, are derived from the fundamental string tension $\mu_F$ as well as $g_s$). In particular CMB and pulsar bounds, which we discuss in Section~\ref{sec:pulsars}, will lead to different shapes of bounding contours in the $(\mu_F, g_s)$ parameter plane. We show that combining these two constraints can lead to complementary constraints on properties of superstrings. The position of the peak in the B-mode spectrum can be used to further eliminate a large region of the $(\mu_F, g_s)$ parameter space. This paper is organised as follows. In Section \ref{sec:scaling} we summarise the extended VOS model which describes multi-tension networks with junctions, and for the case of cosmic superstrings we present the scaling solutions as a function of $g_s$. In Section \ref{sec:cmb} we determine the temperature and B-mode spectra for these scaling solutions using a generalised version of CMBACT. In Section \ref{sec:pulsars} pulsar constraints on gravitational waves from string networks are discussed, and we conclude in Section \ref{sec:conc}.
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\label{sec:conc} Using the MTSN scaling model developed in \cite{Avgoustidis:2009ke}, we have studied the evolution of a network of cosmic superstrings with the aim of identifying characteristic trends in their scaling properties at different values of the fundamental string coupling $g_s$. One particularly interesting trend that we have discovered is in the so-called power spectrum density, which controls the amplitude of the two-point function of the string stress-energy. As $g_s \rightarrow 1$ we see it is dominated by populous light F and D strings where as as $g_s$ decreases to smaller values rare heavy D strings come to dominate. Using the MTSN scaling model in CMBACT, we evaluated the contribution of the FD networks to the CMB temperature and polarisation spectra. We found that the difference between the scaling patterns at high and low values of $g_s$ are manifest as a shift in position of the peak in the B-mode spectrum. In the one-scale approximation adopted in this work, the correlation length is equal to the average inter-string distance, thus string networks of higher (lower) number density have smaller (larger) correlation lengths. The correlation lengths, along with the string velocities, determine the position of the peak. This points to the possibility of constraining $g_s$ with future CMB experiments, possibly providing us with the first opportunity to date to constrain this fundamental parameter of string theory with observations. Most observables, at least those that rely on averaged properties of string networks, constrain a combination of the string tension $\mu$ and their number density, determined by $\xi^{-2}$. Thus, in principle, it is quite hard to distinguish between the effect of many light strings vs a few heavy ones. Measuring a particular peak position in the B-mode spectrum would be one way to partially break this degeneracy, as we have discussed in this work. In addition, one can explore the fact that different observables constrain different combinations of $\mu$ and $\xi$. For instance, while the amplitude of the CMB spectra is determined by $(\mu / \xi)^2$, the energy density of strings is proportional to $\mu / \xi^2$. In this paper, we have shown how this difference can be explored in the case of FD networks to partially break the degeneracy between the fundamental string tension $\mu_F$ and the coupling $g_s$ by combining the CMB constraints with those from bounds on gravity waves (GW). The main trends we have identified in this paper are largely independent of many of the details of the underlying string theory model, as well as the assumptions that went into the CMB calculation and the predictions for GW. However, follow up studies are needed in several directions in order to make firmer quantitative predictions. For instance, we need to improve our understanding of the string interaction rates for different choices of $g_s$ and $w$, especially in the non-perturbative regime. At very small string couplings, the density of the dominant species becomes so high that the one-scale approximation is almost guaranteed to break down. In such cases, a more sophisticated model is needed to properly describe the scaling of the network and its prediction for the CMB spectra. Whether it will be possible to measure a peak at high $\ell$ in the B-mode spectrum is another interesting question which is addressed in an upcoming publication \cite{MossPogosian}. It will depend strongly on the resolution and sensitivity of the experiments, as well as our ability to clean the contribution from weak lensing. The GW bounds on FD strings depend on the presence of cusps \cite{Davis:2008kg}, and on the loop size distribution, which is not fully understood at present, and should be revisited in the light of additional GW signatures coming from $Y$-junctions \cite{Binetruy:2010cc}. Also, while in all string models considered so far the B-mode from the ordinary strings is sourced predominately by vector modes, the tensor modes (i.e. large scale GW) were never properly worked out (since it requires accounting for the backreaction). It is our hope that the potentially very exciting opportunity for testing fundamental theory based on the general trends identified in this work will serve as an additional motivation for pursuing the remaining open questions.
| 10 | 12 |
1012.5014
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1012
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1012.2128_arXiv.txt
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{ We use cosmological SPH simulations to study the cool, accreted gas in two Milky Way-size galaxies through cosmic time to $z=0$. We find that gas from mergers and cold flow accretion results in significant amounts of cool gas in galaxy halos. This cool circum-galactic component drops precipitously once the galaxies cross the critical mass to form stable shocks, $\Mvir = \Msh \sim 10^{12}\Msun$. Before reaching $\Msh$, the galaxies experience cold mode accretion ($T < 10^5$ K) and show moderately high covering fractions in accreted gas: $\CF \sim 30-50\%$ for $R<50$ co-moving kpc and $\NHI>10^{16}$ cm$^{-2}$. These values are considerably lower than observed covering fractions, suggesting that outflowing gas (not included here) is important in simulating galaxies with realistic gaseous halos. Within $\sim500$ Myr of crossing the $\Msh$ threshold, each galaxy transitions to hot mode gas accretion, and $\CF$ drops to $\sim5\%$. The sharp transition in covering fraction is primarily a function of halo mass, not redshift. This signature should be detectable in absorption system studies that target galaxies of varying host mass, and may provide a direct observational tracer of the transition from cold flow accretion to hot mode accretion in galaxies. }
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\label{Introduction} Recent advances in galaxy formation theory have emphasized the importance of cool gas accretion onto galaxies: gas that never shock-heats to the virial temperature of the halo \citep[e.g.,][]{BirnboimDekel03, DekelBirnboim06, Keres09}. Under this picture of galaxy formation, the cooling time of gas entering halos less massive than a critical threshold, $\Mvir=\Msh$, is too short to sustain compressive shocks, with $\Msh\sim10^{11.5-12}\Msun$. Galaxies within dark matter halos less massive than $\Msh$ experience ``cold mode'' accretion, since baryonic accretion onto these galaxies is dominated by cool gas. Galaxies above this transition experience ``hot mode'' accretion, with infalling gas shock-heating to the virial temperature. In practice, the precise value of $\Msh$ depends somewhat on definition. In what follows we are interested in the absolute shut-down mass, when almost no cool accreted gas reaches the central regions of the halo. Thus, we adopt $\Msh\sim10^{12} \Msun$, motivated by \cite{DekelBirnboim06} for shocks at large fractions of the virial radius for $Z\sim0.1$ metallicity gas\footnote{\cite{Keres09} find that less than half of a \emph{galaxy's} gas is accreted from cold mode accretion for $\Mvir\gtrsim2\times10^{11}\Msun$; this fraction drops to near zero at $10^{12}\Msun$. Above $\sim10^{12}\Msun$ cold mode accretion is negligible.}. Unfortunately, there are currently no definitive observational tests to detect cosmological cool gas accretion. On the contrary, numerous observational studies of cool halo gas around galaxies have emphasized the presence of gas \emph{outflows}, not inflows \citep[e.g.][]{Steidel96, Shapley03, Weiner09, Steidel10,Rubin11}. The stark contrast between theory and observations is understandable at high redshift, as gas inflow to galaxies at $z>2$ is expected to flow along dense filaments, resulting in small global covering fractions \citep{FGKeres10,Kimm10}. In addition, galaxies at $z\sim2$ are at the peak of cosmological star formation \citep{HopkinsBeacom06}; one might expect feedback processes to dominate any observational indicator of gas accretion at these epochs. At lower redshifts star formation rates decline and the gaseous halos of galaxies are observed as quasar absorption systems \citep[e.g.,][]{BergeronBoisse91,Bowen95,Churchill96, BartonCooke09, Chen10,Gauthier10}. These studies are primarily of metal lines (MgII, CIV, OVI, etc.) though with COS on HST we expect Lyman $\alpha$ observations to increase. In this letter, we utilize cosmological hydrodynamic simulations to study cool gas accretion and the possibility of detecting it as quasar absorption systems. To correctly produce metal lines in a simulation requires radiative transfer, metal diffusion and modeling of local ionizing sources; however, we are focusing on the qualitative behavior of halo gas, so we will instead give results in terms of HI column density calculated in the optically thin limit, without local sources. For column densities below the Lyman limit ($2 \times 10^{17}$ cm$^{-2}$) this should be fairly robust, but for higher column densities we expect a full treatment would lead to quantitative but not qualitative differences.
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\label{conclusion} We have used two high-resolution, cosmological SPH simulations as a tool for studying the cool gaseous halos of galaxies. By creating mock observation sight lines, representing cool gas detection via absorption, we present the covering fraction of neutral hydrogen in our galaxies over time. Our primary results are summarized as follows. \begin{enumerate} \item To first order, the covering fraction of \emph{accreted} cool gas is relatively stable at $\CF\sim 30-40\%$ ($R<50$ co-moving kpc, $\NHI>10^{16}$ cm$^{-2}$), as long as the galaxy continues to accrete cool gas from the cosmic web. \item As soon as our simulated galaxies cross the threshold between cold mode accretion and hot mode accretion (when the halo is massive enough to support stable shocks at a large fraction of the virial radius, $\Mvir\sim10^{12}\Msun$) the lack of cool accreted gas results in a suppression of cool halo gas. Within $\sim500$ Myr of reaching this threshold mass, the covering fraction drops from $30-50\%$ to $5-10\%$. A transition this sharp should be directly observable via metal line absorption system studies. \end{enumerate} We have used a feedback model without cool gas outflows here, focusing on galaxy halo properties that are a natural consequence of cosmological gas accretion in LCDM. However, observations have shown that galaxy outflows are an abundant phenomenon, and likely play an important role in shaping the properties of cool gaseous halos around galaxies. We believe future work in comparing observations to a variety of simulations with different feedback models would prove a valuable tool in testing theoretical models of galaxy formation, as well as understanding the underlying nature of galaxy halo observations.
| 10 | 12 |
1012.2128
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1012
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1012.2634_arXiv.txt
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We use the largest homogeneous sample of globular clusters (GCs), drawn from the ACS Virgo Cluster Survey (ACSVCS) and ACS Fornax Cluster Survey (ACSFCS), to investigate the color gradients of GC systems in 76 early-type galaxies. We find that most GC systems possess an obvious negative gradient in ($g$--$z$) color with radius (bluer outwards), which is consistent with previous work. For GC systems displaying color bimodality, both metal-rich and metal-poor GC subpopulations present shallower but significant color gradients {\it on average}, and the mean color gradients of these two subpopulations are of roughly equal strength. The field-of-view of ACS mainly restricts us to measuring the inner gradients of the studied GC systems. These gradients, however, can introduce an aperture bias when measuring the mean colors of GC subpopulations from relatively narrow central pointings. Inferred corrections to previous work imply a reduced significance for the relation between the mean color of metal-poor GCs and their host galaxy luminosity. The GC color gradients also show a dependence with host galaxy mass where the gradients are weakest at the ends of the mass spectrum---in massive galaxies and dwarf galaxies---and strongest in galaxies of intermediate mass, around a stellar mass of $\mstellar \approx 10^{10}\msun$. We also measure color gradients for field stars in the host galaxies. We find that GC color gradients are systematically steeper than field star color gradients, but the shape of the gradient--mass relation is the same for both. If gradients are caused by rapid dissipational collapse and weakened by merging, these color gradients support a picture where the inner GC systems of most intermediate-mass and massive galaxies formed early and rapidly with the most massive galaxies having experienced greater merging. The lack of strong gradients in the GC systems of dwarfs, which probably have not experienced many recent major mergers, suggests that low mass halos were inefficient at retaining and mixing metals during the epoch of GC formation.
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\label{sec:intro} Galactic radial gradients in stellar populations are a result of a galaxy's star formation, chemical enrichment, and merging histories, and thus can be an important discriminant of galaxy formation scenarios. Galaxies that form in a strong dissipative collapse are expected to have steep gradients in metallicity, as the central regions retain gas more effectively and form stars more efficiently. Thus, in isolation, higher mass galaxies formed in this way are expected to have steeper negative metallicity gradients due to their deeper potential wells (e.g., \citealt{Chiosi2002_MNRAS_335_335, Kawata2003_MNRAS_340_908}). By contrast, in galaxies where merging is a dominant process, radial gradients are expected to weaken due to radial mixing that occurs during mergers \citep{White1980_MNRAS_191_1, Bekki1999_ApJ_513_108, Kobayashi2004_MNRAS_347_740}. So if the most massive, quiescent galaxies are the ones most shaped by major merging (e.g., \citealt{van_der_wel2009ApJ_706_120}), one would expect their metallicity gradients to be relatively flat. Recently, however, \citet{Pipino2010_MNRAS_407_1347} argue that shallow gradients in massive galaxies can also result from lower star formation efficiency and do not necessarily require extensive merging. The existence of negative optical and near-infrared color gradients, where the outer regions are bluer, have been well-established in elliptical and disk galaxies (e.g., \citealt{Franx1989_AJ_98_538, Peletier1990_AJ_100_1091, Michard2005_A+A_441_451, Wu2005_ApJ_622_244, Liu2009_RAA_9_1119}), and have generally been interpreted as gradients in metallicity, or sometimes age (e.g., \citealt{Kobayashi1999_ApJ_527_573, Kuntschner2006_MNRAS_369_497, Rawle2008_MNRAS_389_1891}). In lower mass galaxies, however, gradients appear to be shallower, nonexistent, or even positive. This shows that gradient properties can be a function of galaxy mass and perhaps reflects the greater diversity in the star formation and evolutionary histories of low-mass galaxies. Recent results with large samples of galaxies show that while the most massive galaxies have shallow or flat color gradients, gradients get increasingly negative toward lower stellar mass until $\mstellar\sim 3\times10^{10} \msun$, at which point gradients again become shallower and even positive \citep{Spolaor2009_ApJ_691_138, Tortora2010_MNRAS_407_144}. For early-type galaxies in particular, this has been interpreted as an intrinsic correlation between gradient and galaxy mass---more negative at higher mass---modulated by dry merging at higher masses, especially for brightest cluster galaxies \citep{Roche2010_MNRAS_407_1231}. Nearly all previous studies of stellar population gradients are of the main stellar body (bulge or disk) of a galaxy. Given the complex star formation histories of galaxies, the effects of age and metallicity are often difficult to disentangle, and require multiband photometry and spectroscopy (e.g., \citealt{MacArthur2004_ApJS_152_175}). Moreover, these studies say little about the stellar halo, perhaps the oldest galactic component. We thus approach the issue of population gradients using a unique tool: globular clusters. Globular clusters (GCs) are among the oldest stellar populations in galaxies, and preserve information from the earliest epochs of star formation. Population gradients in GC systems have not received much attention, but one notable exception was the study of metallicity gradients in the Milky Way GC system by \citet{Searle1978_ApJ_225_357}. They showed that although the inner halo GCs had a negative gradient, the outer halo GCs had no gradient, leading them to suggest that the outer halo was accreted from dwarf-like fragments. In both the Milky Way and nearby galaxies, GCs are found to be nearly universally old, with ages greater than $\sim8$~Gyr (e.g., \citealt{Puzia2006_ApJ_648_383, Hansen2007_ApJ_671_380, Mar'in-Franch2009_ApJ_694_1498, Woodley2010_ApJ_708_1335}). Although in extragalactic systems we are mostly limited to broadband colors, the lack of any significant age spread in GCs, and the fact that they are generally simple stellar populations, allows us to interpret GC colors as largely representative of metallicity. The color distributions of GCs in massive galaxies are often bimodal, and usually interpreted as two metallicity subpopulations \citep[e.g., ][]{Gebhardt1999_AJ_118_1526, Larsen2001_AJ_121_2974, Kundu2001_AJ_121_2950, Peng2006_ApJ_639_95} (although there is still uncertainty in the transformation from color to metallicity, see Yoon, Yi \& Lee 2006). Metal-rich (red) GCs are found to have a more concentrated spatial distribution than the metal-poor (blue) GCs, which results in the total mean color of GCs becoming gradually bluer with projected radius \citep[e.g.,][] {Rhode2001_AJ_121_210, Jord'an2004_AJ_127_24, Tamura2006_MNRAS_373_601}. Many studies of massive galaxies have confirmed that GC systems taken as a whole have negative color and metallicity gradients \citep{ Geisler1996_AJ_111_1529, Rhode2001_AJ_121_210, Jord'an2004_AJ_127_24, Cantiello2007_ApJ_668_209}. The conventional wisdom, however, has been that individual metal-rich or metal-poor GC subpopulations have no color or metallicity gradients \citep{Lee1998_AJ_115_947, Rhode2001_AJ_121_210}. Additional studies of individual galaxies, however, have shown that GC subpopulations in M49, M87, NGC~1427, and NGC~1399, and nearby brightest cluster galaxies do have a slightly negative color gradients \citep{Geisler1996_AJ_111_1529, Forte2001_AJ_121_1992, Bassino2006_A+A_451_789, Harris2009_ApJ_703_939, Harris2009_ApJ_699_254}. Furthermore, very little is known about color gradients in the GC systems of dwarf galaxies, whose systems are dominated by metal-poor GCs. Similar to population gradient studies of the main bodies of galaxies, investigating the color or metallicity gradients of GC systems across a range of galaxy mass can provide direct constraints on the formation of GC systems and the merging history of their host galaxies. In this paper, we present the results from the first homogeneous study of color gradients in the GC systems of early-type galaxies. The ACS Virgo Cluster Survey (ACSVCS, \citealt{Cot'e2004_ApJS_153_223}) and ACS Fornax Cluster Survey (ACSFCS, \citealt{Jord'an2007_ApJS_169_213}) observed 100 galaxies in the Virgo Cluster and 43 galaxies in the Fornax Cluster using the Hubble Space Telescope Advanced Camera for Surveys (HST/ACS). All 143 objects are early-type galaxies and range in mass from dwarf to giant galaxies. One of the main goals of the surveys is the investigation of extragalactic GC systems, and previous studies have examined their color distributions \citep{Peng2006_ApJ_639_95}, size distributions \citep{Jord'an2005_ApJ_634_1002, Masters2010_ApJ_715_1419}, luminosity functions \citep{Jord'an2006_ApJ_651_25, Jord'an2007_ApJS_171_101, Villegas2010_ApJ_717_603}, formation efficiencies \citep{Peng2008_ApJ_681_197}, and color-magnitude relations \citep{Mieske2006_ApJ_653_193, Mieske2010_ApJ_710_1672}. Likewise, the surface photometry of the galaxies themselves have also been studied in detail \citep{Ferrarese2006_ApJS_164_334, Cot'e2007_ApJ_671_1456}, allowing us to perform a homogeneous comparison of the color gradients in the field stars with those in the GC systems. Another advantage of this sample is that distances to most galaxies have been determined using the method of surface brightness fluctuations \citep{Mei2007_ApJ_655_144, Blakeslee2009_ApJ_694_556}. Using this large and homogenous sample of extragalactic GCs \citep{Jord'an2009_ApJS_180_54}, we measure the color gradients of GC systems in the targeted galaxies within the field of view (FOV) of the ACS camera, except for four galaxies where we use multiple ACS fields. The high resolution and quality of the HST images allow us to measure the gradients of GC systems in dwarf galaxies as well as in individual GC subpopulations for systems showing color bimodality. This paper is organized as follows: In Section~2, we give a description of the GC selection and data analysis. The results and discussion are presented in Sections~3 and 4, respectively. Finally, we conclude in Section~5.
| 10 | 12 |
1012.2634
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1012
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1012.5222_arXiv.txt
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We present new ultraviolet spectra of the peculiar white dwarf (WD) \hh, obtained with \emph{COS} on \emph{HST}. \hh\ is the hottest known WD (\Teff\ = 200\,000\,K) and has an atmosphere mainly composed of C and O, augmented with high amounts of Ne and Mg. This object is unique and the origin of its surface chemistry is completely unclear. We probably see the naked core of either a C--O WD or even a O--Ne--Mg WD. In the latter case, this would be the first direct proof that such WDs can be the outcome of single-star evolution. The new observations were performed to shed light on the origin of this mysterious object.
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\hh\ is a faint blue star that was identified as the counterpart of a bright soft X-ray source \citep{nousek:86}. Spectroscopically, it is a member of the PG\,1159 class but, within this class \hh\ is an extraordinary object. It was shown that it is not only hydrogen-deficient but also helium-deficient. From optical spectra it was concluded that the atmosphere is primarily composed of carbon and oxygen, by equal amounts \citep{werner:91}. Strong neon lines were detected in soft X-ray spectra taken with the \emph{EUVE} satellite and in a Keck spectrum, and a high abundance of neon was derived \citep{werner:99}. The origin of this exotic surface chemistry (C = 49\%, O = 49\%, Ne = 2\%, mass fractions) is completely unclear. We have speculated that \hh\ represents the naked C--O core of a white dwarf. Another, even more exciting possibility is that we see the eroded C--O envelope of a O--Ne--Mg white dwarf. This is corroborated by our analysis of a \emph{Chandra} soft X-ray spectrum \citep{werner:04}, which allowed the detection of magnesium. \hh\ turns out to be the hottest single WD ever analyzed with model atmospheres (\Teff\ = 200\,000\,K $\pm$ 20\,000\,K, \logg\ = 8.0 $\pm$ 0.5), only rivaled by the hot DO \kpd\ (see Wassermann \etal in these proceedings, and \cite{was10}). Because of this extremely high temperature, a unique photospheric absorption-line spectrum can be observed with \emph{Chandra-LETG}. It shows a wealth of lines from highly ionized species (\ion{O}{vi}, \ion{Ne}{vi} -- \ion{Ne}{viii}, \ion{Mg}{v} -- \ion{Mg}{viii}). But the spectral analysis of the X-ray data is seriously hampered by heavy line blanketing of iron-group elements. In principle, we can account for this in our synthetic spectra, but two problems prevent a detailed quantitative analysis of relatively weak absorption lines from light metals within the Fe-group forest. First, accurate line positions are unknown for the majority of the Fe-group lines. Second, the bulk of the Fe-group lines from very high ionization stages (a single \ion{Fe}{x} UV line was recently discovered in a \emph{FUSE} spectrum, see Werner, Rauch, \& Kruk, these proceedings; and \cite{we10}) are completely unknown. Consequently, effects of Fe-group line blending on other weak metal lines cannot be calculated with sufficient precision. Nevertheless, we have roughly estimated the Mg abundance to about 1\%, an amount similar to the Ne abundance. If this strong overabundance (20 times solar) can be confirmed, then that would strongly support the idea that \hh\ is a O--Ne--Mg WD. This means that \hh\ might had been one of the ``heavy-weight'' intermediate-mass stars (8\,M$_{\odot}\ \lappr$ $ M\ \lappr$\ 10\,M$_{\odot}$) that form white dwarfs with electron-degener\-ate O--Ne--Mg cores. Evolutionary models \cite{iben:97} predict strong Ne and Mg overabundances in the C/O envelope. Another strong argument in favor of this idea would be the detection of sodium, which would be direct evidence for C-burning. The models predict that the $^{23}$Na abundance at the bottom of the C/O envelope is comparable to that of neon (main isotope $^{20}$Ne) and magnesium ($^{24,25,26}$Mg, see Fig.~34 in \cite{iben:97}). Unfortunately, we were not able to detect Na lines in the \emph{Chandra} spectrum beyond doubt because of, again, heavy metal-line blanketing.
| 10 | 12 |
1012.5222
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1012
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1012.3461_arXiv.txt
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We present All-Sky Automated Survey data starting 25 days before the discovery of the recent type IIn SN~2010jl, and we compare its light curve to other luminous IIn SNe, showing that it is a luminous ($M_I \approx -20.5$) event. Its host galaxy, UGC~5189, has a low gas-phase oxygen abundance ($\rm 12 + log(O/H) = 8.2\pm0.1$), which reinforces the emerging trend that over-luminous core-collapse supernovae are found in the low-metallicity tail of the galaxy distribution, similar to the known trend for the hosts of long GRBs. We compile oxygen abundances from the literature and from our own observations of UGC~5189, and we present an unpublished spectrum of the luminous type Ic SN~2010gx that we use to estimate its host metallicity. We discuss these in the context of host metallicity trends for different classes of core-collapse objects. The earliest generations of stars are known to be enhanced in [O/Fe] relative to the Solar mixture; it is therefore likely that the stellar progenitors of these overluminous supernovae are even more iron-poor than they are oxygen-poor. A number of mechanisms and massive star progenitor systems have been proposed to explain the most luminous core-collapse supernovae. Any successful theory that tries to explain these very luminous events will need to include the emerging trend that points towards low-metallicity for the massive progenitor stars. This trend for very luminous supernovae to strongly prefer low-metallicity galaxies should be taken into account when considering various aspects of the evolution of the metal-poor early universe, such as enrichment and reionization.
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\label{sec:intro} The bright SN~2010jl, discovered on UTC 2010 November 3.5 \citep{newton10} was classified as a type IIn supernova \citep{benetti10,yamanaka10}. A possible luminous blue progenitor has been identified in archival HST WFPC2 data \citep{smith10jl}. If this detection is of a massive compact cluster, the turnoff mass is constrained to be $> 30$~M$_\odot$. If it is a single star, it is either a massive $\eta$~Carinae-like star or something fainter that has been caught in an LBV-like eruption phase, possibly a precursor explosion. Spectropolarimetric observations indicate possible asymmetry in the explosion geometry and limit the amount of dust in the progenitor environment \citep{patat10}. \citet{prieto10} point out that the supernova's host galaxy, UGC~5189, has been shown to be metal-poor based on a spectrum obtained by the SDSS survey a few arcseconds away from the site of the supernova. It is included in the \citet{tremonti04} DR4 catalog of galaxy metallicities, which estimates $\rm 12 + log(O/H) = 8.15\pm 0.1$~dex based on strong recombination and forbidden emission lines in the spectrum. \citet{pilyugin07} estimate $\rm 12 + log(O/H) = 8.3$~dex from the same spectrum based on the ``direct'' photoionization-based abundance method using an estimate of the electron temperature ($T_e$) from the [\ion{O}{3}]$\lambda 4363$\AA{} auroral line. Throughout this paper, as above, we will adopt O/H as a proxy for the ``metallicity'' of the host galaxy, because only gas-phase oxygen abundance estimates are available. In Section~\ref{sec:discuss} we will discuss how this relates to total metal abundance and iron abundances. There is mounting evidence showing that the majority of the most optically luminous supernovae explode in low-luminosity, star-forming galaxies, which are likely to be metal-poor environments. Circumstantial evidence comes from the fact that most of these objects have been discovered in new rolling searches that are not targeted to bright galaxies, even though they were bright enough to be discovered in galaxy-targeted searches (the only exception is the recently discovered SN~2010jl). Some of the rolling searches that are discovering the most energetic supernovae include the Texas Supernova Search \citep[TSS;][]{quimbyTSS}, the Catalina Real-Time Transient Survey \citep[CRTS;][]{drake09}, the Palomar Transient Factory \citep[PTF;][]{rau09}, and the Panoramic Survey Telescope \& Rapid Response System (Pan-STARRS). \citet{neill10} find that luminous SNe occur predominantly in the faintest, bluest galaxies. \citet{li10} find that SNe IIn may preferentially occur in smaller, less-luminous, later-type galaxies than SNe II-P. The mass-metallicity relationship implies that such small, faint galaxies should tend to be low-metallicity \citep{tremonti04}. \citet{kozlowski10} presented the first host galaxy luminosity vs. oxygen abundance diagram with a small sample of three energetic type IIn and type Ic core-collapse supernovae (CCSNe) compiled from published work in the literature (two objects) and new data for SN~2007va. These initial results suggest that the host galaxy environments of the most energetic CCSNe are on average metal-poor (metallicities $\sim 0.2 - 0.5$~$Z_\odot$) compared to the bulk of star-forming galaxies in SDSS, and are similar in metallicity to the hosts of local GRBs \citep{stanek06}. It has now been well-established that long GRBs are accompanied by broad-line type Ic SNe \citep[e.g.][]{stanek03}, firmly connecting these very energetic explosions to the deaths of massive stars. In contrast, luminous type IIn SNe are not shown to be connected with local GRBs \citep[although see][]{germany00,rigon03}. The sample of luminous supernovae with measured abundances is small and incomplete because only some of the brightest host galaxies have been targeted, which hinders the interpretation and comparison of the results. In this paper we expand this sample with metallicities of the hosts of SN~2010gx and SN~2010jl and put all metallicity measurements on a common scale to confirm the emerging trend that these optically luminous core-collapse events appear to occur in low-metallicity or low-luminosity hosts.
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Recent observations of the environments of supernovae and GRBs show that metallicity is a key parameter in the lives and deaths of massive stars. % Our finding that the hosts of luminous SNe lie in the low-metallicity extremes of the distribution of star-forming galaxies supports this emerging picture. Verifying and constraining this emerging relationship with metallicity may be an important probe of the mechanisms of the most luminous supernovae, and this trend may be an important factor in various aspects of the evolution of the metal-poor early universe.
| 10 | 12 |
1012.3461
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1012
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1012.1464_arXiv.txt
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{ The decay of kinetic helicity is studied in numerical models of forced turbulence using either an externally imposed forcing function as an inhomogeneous term in the equations or, alternatively, a term linear in the velocity giving rise to a linear instability. The externally imposed forcing function injects energy at the largest scales, giving rise to a turbulent inertial range with nearly constant energy flux while for linearly forced turbulence the spectral energy is maximum near the dissipation wavenumber. Kinetic helicity is injected once a statistically steady state is reached, but it is found to decay on a turbulent time scale regardless of the nature of the forcing and the value of the Reynolds number. }
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\label{Introduction} The physical properties of astrophysical turbulence are often studied by solving the hydrodynamic equations in a periodic domain with an assumed forcing function. In particular, isotropic homogeneous turbulence is often studied as a proxy of turbulence in more complicated situations, where specific concepts and general aspects are harder to isolate. An important concept is that of the forward cascade of kinetic energy to smaller scales. This leads to a $k^{-5/3}$ energy spectrum, where $k$ is the wavenumber. Many flows of geophysical and astrophysical relevance are subject to rotation and stratification and can therefore attain kinetic helicity. Closure calculations (\cite{AL77}), direct numerical simulations (\cite{BO97}; \cite{BS05a}), as well as shell model calculations (\cite{Chkhetiani}; \cite{DG01},\cite{DG01b}; \cite{Stepanov}) show an approximate $k^{-5/3}$ scaling for the kinetic helicity, suggesting that kinetic helicity too is subject to a forward cascade toward smaller length scales. On the other hand, kinetic helicity is conserved by the quadratic interactions and might therefore play an important role in the inviscid limit. Although this is also the case in ideal magnetohydrodynamics (MHD), where magnetic helicity is also conserved by the quadratic interactions, there is a significant difference. In the ideal case of small magnetic diffusivity, magnetic helicity can only evolve on resistive time scales. This can have profound effects on the saturation behavior of large-scale dynamos (\cite{B01}). As mentioned above, in a polytropic flow, the quadratic interactions conserve kinetic helicity, $\bra{\oo\cdot\uu}$, where $\oo=\nab\times\uu$ is the vorticity, $\uu$ the velocity, and angular brackets denote volume averaging over a closed or periodic domain. This conservation property becomes evident when writing the Navier-Stokes equation in the form \EQ {\partial\uu\over\partial t}=\uu\times\oo-\nab P+\ff +\nu[\onethird\nab(\nab\cdot\uu)-\nab\times\nab\times\uu+\GG], \EN where $P=\half\uu^2+h$ is the sum of specific turbulent pressure, $\uu^2/2$, and specific enthalpy, $h=\cs^2\ln\rho$. Furthermore, $\rho=\const$ is the density, and $\ff$ is a forcing function. In the absence of forcing, $\ff=\bm{0}$, kinetic helicity is just subject to viscous decay, because the nonlinear term $\uu\times\oo$ is perpendicular to $\oo$, i.e., we have \EQ {\dd\over\dd t}\bra{\oo\cdot\uu}=-2\nu\bra{\qq\cdot\oo}, \EN where $\qq=\nab\times\oo$ is the curl of the vorticity. In the ideal case, $\nu=0$, we have $\bra{\oo\cdot\uu}=\const$. However, in fluid dynamics the ideal case is hardly representative of the limit of large Reynolds numbers, where $\nu\to0$. Indeed, for a self-similar decay of kinetic energy, the wavenumber of the energy-carrying eddies, $\kf$, decreases with time such that $\nu\kf^2t\approx\const$. This implies that the rate of energy decay, $\nu\bra{\oo^2}$, is essentially independent of $\nu$ and hence independent of the Reynolds number. Given that $\bra{\uu^2}$ is related to the kinetic energy, which is also independent of the Reynolds number, we expect that the ratio \EQ \bra{\oo^2}/\bra{\uu^2}\equiv\kT^2, \label{kT} \EN which is related to the Taylor micro-scale wavenumber $\kT$, should be proportional to $\Rey$, and therefore \EQ \kT\sim\Rey^{1/2}. \label{kscaling} \EN However, for helical flows the rate of kinetic helicity dissipation is proportional to $\nu\bra{\qq\cdot\oo}$. Thus, if we define \EQ \bra{\qq\cdot\oo}/\bra{\oo\cdot\uu}\equiv\ke^2, \label{ke} \EN we see that kinetic helicity dissipation is related to kinetic energy by altogether 3 wavenumber factors. If all these factors scale like in \Eq{kscaling}, we may expect the rate of kinetic helicity dissipation to diverge with decreasing $\nu$ like $\nu^{-1/2}$. This is in stark contrast to the related case of magneto-hydrodynamic turbulence where, following similar reasoning, the magnetic helicity dissipation converges to zero like $\eta^{1/2}$ as the magnetic dissipation $\eta$ goes to zero (\cite{BS05}); see also the appendix of \cite{BK07} for a clear exposition of these differences. A problem with the simple argument above is that in cases of practical relevance the forcing function $\ff$ usually breaks kinetic helicity conservation. This is particularly evident for the so-called linear forcing model of \cite{Lun03}, where \EQ \ff=A\uu \label{LinForcing} \EN is a positive multiple of the velocity vector. In that case we have \EQ {\dd\over\dd t}\bra{\oo\cdot\uu}=2A\bra{\oo\cdot\uu}-2\nu\bra{\qq\cdot\oo}, \label{ExpIncrease} \EN so that $\bra{\oo\cdot\uu}$ could even exhibit exponential growth. An aim of this paper is thus to investigate to what degree kinetic helicity is conserved in forced turbulence using both the linear forcing model and compare it with the more traditional stochastic forcing in a narrow wavenumber band. Another motivation is the fact that, by analogy, magnetic helicity turned out to be of crucial importance in understanding the saturation properties of helically forced dynamos in periodic domains (see, e.g., \cite{B01}). We study the kinetic helicity evolution by monitoring the response to adding a large-scale helical component to the flow for both types of forcing.
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The linear forcing model is based on the driving of turbulence by a linear instability instead of a forcing function that is independent of the flow. This type of forcing might be more physical, because it does not change abruptly and depends on the local flow properties. Nevertheless, both types of forcing are Galilean invariant. This would change if the flow-independent monochromatic forcing were no longer $\delta$ correlated in time. A disadvantage of linearly forced turbulence is that energy injection occurs at all wavenumbers. Our simulations show that the energy spectrum is perhaps slightly shallower than with the flow-independent monochromatic forcing function. Part of this is explained by the bottleneck effect (\cite{Fal94}; \cite{Dob03}), and that the energy spectra compensated by $k^{5/3}$ show a stronger uprise toward the dissipation wavenumber. It is unclear whether this result would persist at larger resolution. The linear forcing model has the interesting property of amplifying not only kinetic energy, but also kinetic helicity. Indeed, at early times, just after having injected kinetic helicity into the system, the coherent helical part continues to increase exponentially, but the flow soon breaks up into smaller eddies, giving rise to enhanced effective dissipation. For practical applications, it should be noted that the linear forcing model has the disadvantage that the initial exponential growth of kinetic energy and kinetic helicity might prevail for too long. This will be the case when the {\it initial} random perturbations are too weak to perturb the flow sufficiently. In that case, the kinetic energy would quickly increase to large values without producing three-dimensional turbulence. With regards to geophysical and astrophysical applications we can say that in a turbulent system, kinetic helicity is no longer a conserved quantity, even though it would be if $\nu=0$ were strictly true. The latter requirement is of course not really possible in a turbulent system, because kinetic energy would then accumulate at the smallest possible scale resolved within the hydrodynamics framework and kinetic energy would not be able to decay, which is unphysical. While this should not be surprising, it is important to remember that this is quite different in the case of magnetohydrodynamics, where magnetic helicity dissipation really does go to zero in a turbulent system -- even for finite (but small) values of the magnetic diffusivity. At the same time, magnetic energy dissipation does stay finite and is able to accomplish magnetic reconnection on the smallest resolved scales of the turbulent cascade (\cite{GN96}; \cite{LV99}). This paper has also shown that, regardless of the nature of the forcing, there are fairly strong helicity fluctuations. They appear to be coherent over many turbulent eddy timescales. One may wonder how generic such fluctuations are and if such fluctuations could be relevant for say the incoherent dynamo effects that have been investigated by several authors in recent years (\cite{Vishniac}; \cite{Sur}; \cite{Heinemann}; \cite{Mitra}; \cite{Proctor}). It will therefore be interesting the associate the kinetic helicity fluctuations with those of $\alpha$, which have already been determined in simulations of turbulent shear flows (\cite{BRRK}).
| 10 | 12 |
1012.1464
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1012
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1012.3711_arXiv.txt
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\noindent We have observed an area of $\approx 27$~deg$^{2}$ to an rms noise level of $\lessapprox 0.2~\mathrm{mJy}$ at 15.7~GHz, using the Arcminute Microkelvin Imager Large Array. These observations constitute the most sensitive radio-source survey of any extent ($\gtrsim 0.2$~deg$^{2}$) above 1.4~GHz. This paper presents the techniques employed for observing, mapping and source extraction. We have used a systematic procedure for extracting information and producing source catalogues, from maps with varying noise and \textit{uv}-coverage. We have performed simulations to test our mapping and source-extraction procedures, and developed methods for identifying extended, overlapping and spurious sources in noisy images. In an accompanying paper, \citet{davies2010}, the first results from the 10C survey, including the deep 15.7-GHz source count, are presented.
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The 9C survey \citep{waldram2003,waldram2010} mapped 29~deg$^{2}$ of sky to 5.5~mJy completeness at 15~GHz, in addition to several larger and shallower areas. The Ryle Telescope (RT), which carried out the survey, has subsequently been reconfigured and re-equipped to form the Large Array (LA) of the Arcminute Microkelvin Imager \citep[AMI;][]{zwart2008}. As part of this metamorphosis, three of the RT's eight antennas were moved to the north of the (almost) east-west line on which they originally stood, providing the telescope with north-south baselines. In addition, new front-end receivers and back-end electronics, including a new correlator, were installed. The result is a telescope with much larger bandwidth and improved flux-density sensitivity, allowing us to extend our 15-GHz-band survey work to much deeper flux-density levels. The 9C survey was conceived to provide information regarding the foreground radio sources that contaminated the Cosmic Microwave Background radiation observations of the Very Small Array \citep{watson2003}. Similarly, the 10C source survey has been designed to complement the other AMI science programmes, which also require knowledge of contaminating radio sources. This paper is focussed on the techniques employed for observing, mapping and source extraction in the 10C survey. In an accompanying paper \citep[hereafter Paper~II]{davies2010} the first results from the 10C survey, including the deep 15.7-GHz source count, are presented.
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In order to investigate the high-frequency-radio sky we have observed 10 survey fields, covering $\approx 27$~deg$^{2}$ to an rms noise level of $\lessapprox 0.2~\mathrm{mJy}$, using the AMI LA at 16~GHz. In an accompanying paper, \citet{davies2010}, we present the first results from the survey, including the deep 15.7-GHz source count. Here, we have concentrated on developing techniques for producing and analysing the survey raster maps. In particular, we have: \begin{enumerate} \item[(1)] developed systematic, automated methods for identifying and characterising sources in maps with varying noise levels and sythesised beams. \item[(2)] proposed a straightforward and robust method for distinguishing between point and extended sources over a wide range of SNRs. Our method has been tested using maps including simulated sources and noise, and has been shown to be successful in identifying extended emission. \item[(3)] applied our techniques to real sky maps and demonstrated that our automated techniques are useful for identifying and chracterising complex structure (see, for example, Fig.~\ref{fig:AMI001L_map_section}). \end{enumerate}
| 10 | 12 |
1012.3711
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1012
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1012.5301_arXiv.txt
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We report monitoring observations of the T~Tauri star EX~Lupi during its outburst in 2008 in the CO fundamental band at 4.6--5.0~$\mu$m. The observations were carried out at the {\it VLT} and the Subaru Telescope at six epochs from April to August 2008, covering the plateau of the outburst and the fading phase to a quiescent state. The line flux of CO emission declines with the visual brightness of the star and the continuum flux at 5~$\mu$m, but composed of two subcomponents that decay with different rates. The narrow line emission (50~km~s$^{-1}$ in FWHM) is near the systemic velocity of EX~Lupi. These emission lines appear exclusively in v=1-0. The line widths translate to a characteristic orbiting radius of 0.4~AU. The broad line component (FWZI $\sim$ 150~km~s$^{-1}$) is highly excited upto v$\leq$6. The line flux of the component decreases faster than the narrow line emission. Simple modeling of the line profiles implies that the broad-line emitting gas is orbiting around the star at 0.04--0.4~AU. The excitation state, the decay speed of the line flux, and the line profile, indicate that the broad-line emission component is physically distinct from the narrow-line emission component, and more tightly related to the outburst event.
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EX~Lupi, the prototype of the class of the eruptive pre-main sequence stars called EXors,\footnote{The group of low-mass stars with large variability is coined as EXors by \citet{her89} for its nature of the outburst similar to FU~Ori variables that are collectively called FUors.} underwent its largest outburst recorded in the past 50 years in early 2008 \citep[history of the observations of EX~Lupi and EXors:][]{her07,her08}. The observational definition of EXor variables relies strongly on the prototype EX~Lupi, which spectroscopically looks like a classical T~Tauri star in the quiescent phase \citep{her50,her01}, but flares by 1--3 magnitudes in the visible roughly once a decade, with each outbreak lasting about a year \citep{her07}. The spectroscopic behavior in the outburst is complicated, and sometimes related to an expanding stellar shell \citep[e.g.][]{her89}. It is increasingly accepted, however, that the circumstellar disk plays the primary role in the outbursts, in which mass accreted from the envelope is stalled somewhere in the disk, until it reaches a critical surface density and triggers a disk instability. \begin{figure*} \includegraphics[width=0.8\textheight,angle=0]{f1.eps} \caption{Light curve in the visible (crosses; from AAVSO http://www.aavso.org/.) with the ordinate to the left, and in $M$ band (diamonds for CRIRES, and a square for IRCS) to the right. The axis of $M$ band magnitude is arbitrarily scaled to that of $V$ band. The $M$ band photometry is performed with the spectroscopic data, with the continuum level of the spectroscopic standard star as the photometric reference. The uncertainty of $M$ band photometry is $\pm$0.5~mag (the error bar shown in the upper right corner). \label{lc} } \end{figure*} The idea that EXor outbursts arise inside the accretion disk originates in their observational similarity with FUors. The most compelling evidence that the FUor outburst has its origin in a disk is the split absorption lines in the visible to the near-infrared with wavelength separation of two absorption peaks gradually decreasing with the wavelength \citep{har04}. This is most naturally accounted for when the absorption lines at long wavelengths arise from the cool outer part of the disk, where the orbital velocity is slower \citep[e.g.][]{zhu09}. FUors and EXors are historically considered as different classes of eruptive pre-main sequence stars. However, both classes may represent a similar eruptive event, and differ only quantitatively in terms of flare amplitude, outburst duration, and outburst frequency. Recent outbursts of V1647~Ori \citep[the central source of ``McNeal's Nebula'', e.g.][]{asp08} are recurrent like EXors, yet the amplitude of the outburst is more compatible with FUors, which fills the gap between these two classes. The mass accreted in a single outburst and the frequency of the outbursts indicates that an EXor/FUor star acquires a substantial fraction of its mass solely during the outbursts. It lends circumstantial support to the theory that EXors/FUors are not a group of special variable stars, but a state of the accretion phase that stars commonly experience during their evolution \citep[e.g.][]{har96}. However, there are indeed qualitative differences between FUors and normal T~Tauri stars, such as the clustering nature, as FUors tend to be in isolation, while other T~Tauri stars are in clusters \citep{her03}. Moreover, the number of known FUors in the Orion Nebula Cluster falls short if all class~I sources go through the FUor phase with a finite duty cycle \citep{fed07}, although the number may increase with the inclusion of the ``FUor-like'' stars being identified recently \citep{gre08}. The current focus of EXor studies is to understand (1) what is the mechanism of the disk instability, (2) what is the relationship between EXors and FUors, and (3) what is the role of EXor/FUor events in the context of star formation. In this report, we will concentrate on the cause of the disk instability and the physical scale of the outburst, as it often gives a clue to the trigger of the instability. The most favored model to date, thermal instability in a disk, is induced by ionization, and subsequent opacity increase due to negative hydrogen ions \citep{bel91,cla89}. The thermal instability is therefore only operational within a limited disk radius, where atomic hydrogen is subject to thermal ionization. If the physical size of the outburst is larger, magnetohydrodynamic (MHD) instability may be favored, as it predicts a massive accretion in the magnetic dead zone at 1--10~AU \citep{arm01}. If the outburst involves the whole disk, the scenario of an unseen companions leading to global disk instability may also work \citep{bon92}. The vibrational transitions of CO are a unique probe of the inner regions of disks (0.01--1~AU) due to their high critical density ($n_{\rm cr}>10^{10}$~cm$^{-3}$) and high excitation energy ($\Delta E_{\rm v=1-0}>3000$~K). Measurement of the location of the emitting gas is relatively straightforward under the assumption that the gas is in Keplerian rotation \citep[e.g.][]{naj96}. We utilized multi-epoch observations of CO vibrational transitions in EX Lupi to constrain where in the disk the hot gas is located during the outburst, and how its physical property develops with time. The present study is part of the first multi-wavelength campaign of EX~Lupi from the optical \citep{sic10}, near-infrared \citep{kos10} to mid-infrared wavelengths \citep{abr09, juh10} during the 2008 eruption. We will use $d=155$~pc as the distance to EX~Lupi from \citet{lom08}, and the stellar parameters from \citet{gra05} following \citet{sip09} (Table~\ref{t1}).
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\subsection{Spectral Components} \subsubsection{Broad Line Emission} The disk parameters of the broad line emission that we modeled in the previous section, $i$, $R_i$, $N_{\rm CO}$, $T_{\rm x}$ and $T_{\rm v}$ are not independent, but interlocked in a non-trivial way. The line width of $v \sin{i} \sim 150$~km~s$^{-1}$ on Apr 28 testifies that the radius of the inner edge of the disk is not larger than 0.09~AU. On the other hand, the disk cannot be too small in order to account for the absolute line flux. In the cases when the line profile is not simple and the signal-to-noise ratio is not high, such as the present case, the photometric calibration is important because the absolute line flux effectively constrains the size of the emitting region and therefore the disk radius. If we take the photometry at 5~$\mu$m derived based on the spectroscopic data as face value, the single line flux ranges from 10$^{-14}$ to 10$^{-13}$~erg~cm$^{-2}$. The maximum flux density that we could retrieve from a disk at $d$ pc away is $f_\nu \approx \pi r^2 B_\nu(T_{\rm b})/d^2$, where $B_\nu$ is the Planck function. Given the blackbody temperature of $T_{\rm b}$ = 2000~K, the distance of $d =155$~pc, and the line width of 100~km~s$^{-1}$, we will get $f_\nu < 10^{-14}$~erg~s$^{-1}$cm$^{-2}$, if we assume a disk radius of 0.01~AU. This is too small to account for the observed line flux, therefore the radius of the inner hot disk cannot be smaller. The lower limit for the disk radius then calls for the inclination angle to be larger than $i >$30\degr~to match the observed line width. We first applied $T_{\rm x}$ and $N_{\rm CO}$ to the models as determined by the optically thick population diagrams, and optimized gas radius to fit the line profiles best. The vibrational temperatures $T_{\rm v}$ are set equal to $T_{\rm x}$. The results underestimate the absolute line flux, because the outer radius $R_o$ that is consistent with the kinematics in the line profiles is larger than the specific emitting radius $R_s$ we used to reproduce the population diagrams. The model spectra are normalized by the square of $R_o/R_s$ and shown in Fig.~\ref{sp1} in blue lines. The absolute line flux is reproduced without such an arbitrary scaling, by increasing the excitation temperature, allowing it to differ from $T_{\rm v}$. However, when the emitting gas is hotter than 3000~K, the line ratios of different rotational transitions diminish in a short wavelength interval. It becomes less straightforward to determine $T_{\rm x}$ uniquely. In the optically thick regime, increasing $N_{\rm CO}$ and $T_{\rm x}$ imposes similar effects on the output spectra. In order to cope with the degeneracy, and avoid increasing the optical depth any further, which our simple model cannot handle, the rotational temperature was set to 4500~K, as the thermal dissociation temperature of CO is 4500--4600~K \citep{tat66}. $T_{\rm v}$ is adjusted so that the model reproduce the absolute line flux under the condition of $T_{\rm x}=4500$~K. The second model with $T_{\rm x}=$4500~K is shown in Fig.~\ref{sp1} in red lines. The two solutions (Table~\ref{t3}) can be regarded as the two extreme cases of the spectral models, and provide a fair idea how the model is affected by the degeneracy of the parameters. The best-fit line profile of a disk in Keplerian rotation sets the inner and the outer radius at 0.04--0.05~AU ($=$5--7~$R_\ast$) and 0.2--0.4~AU, respectively, for all epochs with the fixed inclination angle of $i=$45\degr. These geometric parameters are commonly used in the two representative models discussed above, and less subjected to the degeneracy. The optical depth is $\tau_\nu=$0--1000 at the line center in the model with $T_{\rm x}=T_{\rm v}$, and $\tau_\nu=$0--100 for $T_{\rm x}=4500$~K, assuming the line is only thermal broadened. Although the heavy overlap of the lines makes it difficult to discern the shape clearly, the line profile looks double-peaked (Fig.~\ref{dp}). This is the first time that the line profile is found to be double peaked in EXors. It lends support to the underlying similarity of the two classes of variable stars, FUors and EXors, and the origin of EXor outbursts associated with the circumstellar disk. Hartmann and his school \citep[e.g.][]{har96,har04,zhu09} attribute the split absorption lines of FUors, unusually broad for a stellar photosphere, to the gas in Keplerian rotation in unstable disks. The interval of line splitting is smaller in the near-infrared than in the visible. This is naturally explained if the absorption lines at longer wavelengths arise from the cooler, more distant, slow-rotating region \citep{zhu09}. On the other hand, \citet{her03} and \citet{pet08} found no correlation between the line widths and the excitation potentials of the visible absorption lines of FU~Ori, which should be present if the lines arise from the disk, and the disk has a temperature gradient. They argued against the disk origin of the optical lines, and claimed that the split lines are better reproduced by a large cool spot at the polar region on a rapidly rotating star. Moreover, the absorption lines of FUors are not perfectly Keplerian, but have a rectangular shape at the bottom \citep[e.g.][]{har04}. It is hard to distinguish whether the unusual line shape represents the turbulence in the upper layer of the disk through which the absorption takes place, or a particular morphology and geometry of the cool spot on the surface of a star \citep[e.g.][]{pet08}. The line profile of EX~Lupi is found to be double peaked, but in emission. The double peaked emission lines would not be compatible with the stellar spot hypothesis, unless the stellar surface is extended to at least a few times the stellar radius. The line profiles are asymmetric between the blue and the red peaks. Figure~\ref{as} shows the blow-up view of the spectra near $R$(6) v=2-1 (4.6675~$\mu$m). There are minor contributions from $R$(14) v=3-2 and $R$(23) v=4-3 at 4.6665~$\mu$m and 4.6670~$\mu$m, respectively. The asymmetric line profile, however, cannot be fully explained by changing the vibrational temperature, thus the relative contribution of the three lines. The line profile is fitted better if we artificially introduce intrinsic asymmetry in the Keplerian line profile by modulating the azimuthal intensity distribution sinusoidally [$I_\nu^\prime(\theta)=I_\nu (\theta) (1+(1-a)\cdot\sin{\theta})$, where $a$ is the asymmetric factor] by down to $a=$0.6. Asymmetries in infrared CO line emission have been found by \citet{got06} in HD~141569~A, where the v=2-1 emission is three times brighter in the northern disk than in the southern disk in the spatially resolved spectroscopy. \citet{pon08} proposed that the transition disks around SR~21 and HD~135344B have azimuthal structures, as the spectroastrometric signal of CO v=1-0 is asymmetric, which is consistent with the millimeter dust continuum imaging \citep{bro09}. The line asymmetry of EX~Lupi varies with time. \subsubsection{Narrow Line Emission} The single peak profile of the narrow line emission with 50~km~s$^{-1}$ line-width in FWHM defies pure Keplerian rotation modeling, if the line intensity decreases as a power law of the radius. Instead, the line width is simply taken as the orbital velocity of the gas, and is translated to the disk radius of 0.4~AU, using the same inclination angle as the broad line emission. FWHM was used instead of FWZI, as the latter is hard to measure unambiguously because of the overlap with the broad line components. The disk radius above does not mean the inner truncation, but a characteristic radius where the narrow line emission arise from. The quiescent line emission comes exclusively from the v=1-0 transition with no obvious higher vibrational transitions. Except the emergence of absorption at the blue shoulder implying a disk wind, no significant changes are found either in the equivalent width or in the line shape (Fig.~\ref{ew}; Fig.~\ref{dp}, right). Note that the small equivalent width on July 21 is of less significance than other epochs, as the data suffer from the low spectral resolution and the limited signal-to-noise ratio, which makes the removal of the telluric CO absorption challenging. The systematic error brought in by the telluric correction is magnified by the subtraction of the outburst model. The broad line emission is less subject to such problems in the measurements of the equivalent widths. The equivalent width of the narrow line component looks constant in time. \subsubsection{Disk Wind} The model spectra of the broad line emission systematically overestimate the flux of v=1-0 lines in the blue shoulder, and underestimate the red. After subtracting the spectral model, therefore, the v=1-0 lines systematically show a blue depression between $-$40 and $-$100~km~s$^{-1}$ with excess emission at the red shoulder at $+$80~km~s$^{-1}$, which is similar to a P~Cyg profile overlaid on the sharp quiescent emission (Fig.~\ref{dp}, right). This blueshifted absorption is more prominent toward the later epoch. The absorption lines can be noted at $P$(21) to $P$(26) on August 21, implying that the absorbing cloud is warm ($T_{\rm x}\gg$100~K). The absorption is not likely photospheric, although EX~Lupi is an M-type star in its quiescence, as the amount of blueshift is too large. The emergence of blue shifted absorption is also observed in infrared CO spectra of V1647~Ori toward the end of an outburst \citep{bri07}. We infer that the blue absorption in the EX~Lupi spectra also represents the disk wind, similarly to the case of V1647~Ori. \begin{figure} \begin{center} \includegraphics[height=0.52\textheight,angle=0]{f5.eps} \end{center} \caption{Line profile of the $R$(6) v=2-1 transition at 4.6675~$\mu$m. The systemic velocity of EX~Lupi is close to 0~km~s$^{-1}$ \citep{her03}. The model spectra of the outburst with $T_{\rm x}=4500$~K are overlaid in gray lines. There are minor contributions from the higher vibrational transitions of $R$(14) v=3-2 and $R$(23) v=4-3 as shown in the right bottom panel on 21 August. The observed line profiles cannot be reproduced without introducing the asymmetry. The asymmetry flips three times in August, which is consistent with the orbital period at 0.04--0.06~AU. \label{as}} \end{figure} \subsection{Heating Mechanism} The excitation mechanism of the CO vibrational band is somewhat ambiguous. The elevated accretion energy can be directly converted to the kinetic energy of the gas through the viscous heating. In this case, the broad line emission that we locate at 0.04--0.4~AU traces the region where the accretion rate is high. On the other hand, the connection between the CO emission and the outburst could be indirect. The radiation from the central star, either from the hot accretion funnels or the footprints of the magnetic field lines, is first raised by the high disk accretion. The gas in the disk is heated afterwards by the enhanced irradiation. The primary difference of the two mechanisms is the vertical structure of the temperature that the heating imposes. In the case that the gas is viscously heated by the accretion, the disk is hottest near the mid-plane, and gradually cools down toward the surface. The emergent spectrum should be in absorption, as is the case for FUors \citep{har04}. When the irradiative heating is dominant, the temperature gradient is inverse with emission lines from the upper layer of the disk surface \citep{cal91a}. The irradiative heating has a few advantages to account for the CO line emission of EX~Lupi. First, most of the infrared CO emission lines observed today are thought to come from the hot surface of a disk, irradiatively heated \citep[e.g.][]{naj03}. This is fully consistent with the line flux of the narrow line emission closely following the continuum level. A critical problem of the direct heating of gas by the accretion is the cooling time of the outburst. The duration of the EX~Lupi outburst as is seen in the visible light curve, is roughly 8 months, from January to August 2008. The viscous time scale on which the gas accretes on the star is given by \[\frac{1}{t_{\rm vis}} = 3 \pi \, \alpha (H/r)^2 \cdot \frac{1}{t_{\rm orbit}},\] for standard $\alpha$ prescription of an accretion disk. The viscous time scale at 0.04~AU is 150 to 600 days, depending on $H/r$ from 0.05 to 0.025 at the radius, even for the fastest case with $\alpha$ being unity. The observed time scale of the outburst is uncomfortably short compared to the most optimistic case \citep{juh10}. On the other hand, some features of the EX~Lupi outburst are not fully explained by the irradiative heating either. The line flux of both broad and narrow components decline with the visual magnitude and the continuum flux at 5~$\mu$m (Fig.~\ref{ew}, upper panels), but on the slightly different decay rates. The narrow line emission is diminished by the factor of 4 from April to August (Table~\ref{t4}), which is the same factor with the decline of 5~$\mu$m continuum flux, while the broad line emission is decreased faster by the factor of 10. In order to contrast the decay rates, the equivalent widths are calculated in the same manner with the line flux. While the equivalent width of the narrow line emission maintains the level of 1.5$\times$10$^{-4}$~$\mu$m at the end of the outburst (Fig~\ref{ew}, bottom right), where the visible magnitude of EX~Lupi returned to $V=$12--13~mag, comparable to the level of the pre-outburst brightness in 2007 (Fig.~\ref{lc}), the equivalent width of the broad line emission declines, keeping pace with the continuum flux at the visual wavelength. The differential time scale of the two line components is not accounted for, if the broad and the narrow line emissions are a single disk component radiatively heated in the identical manner. In addition, the excitation state and the kinematics are distinct in two emission components, with no smooth transition between them. The characteristic radius where the narrow line emission arises from is 0.4~AU, and is constant as seen in the stable line width. The radius is immediately outside of the broad line emitting region at 0.04--0.4~AU (Fig.~\ref{car}). The location of this transition zone is close to the inner-rim of the dust disk found by the SED modeling by \citet{juh10}. With the vibrational excitation state upto v=6, and the faster decay of the line flux than the radiatively heated narrow line component, the broad line emission apparently traces the outburst event more directly. \begin{figure} \begin{center} \includegraphics[height=0.34\textheight,angle=0]{f6.eps} \end{center} \caption{Top: (Left) Line flux of the emission lines from vibrationally excited states (v$\ge$2, open circles), and the lowest transition v=1-0 (triangles) after subtraction of the outburst model of $T_{\rm x}=4500$~K shown as functions of the continuum flux of EX~Lupi at the visible wavelength. The line intensity is integrated between $-$120 and $+$120~km~s$^{-1}$ for the broad line component. The composite line profiles used in the calculations are shown in Fig.~\ref{dp}. (Right) Same with the left panel, but shown with the date in abscissa. Bottom: Same with the top panels but for the equivalent widths. The line flux of the v$\ge$2 and v=1-0 lines decrease with the time and the visual brightness of the star, but with different speed. The difference is clearly seen in the plots of the equivalent widths. The equivalent widths of the broader, highly excited lines decays quickly with the optical brightness of the star, while the equivalent widths of the v=1-0 narrow line emission stays constant during the outburst. \label{ew}} \end{figure} \setlength{\tabcolsep}{5pt} \begin{deluxetable*}{lcccccc} \tablecolumns{7} \tablewidth{0pc} \tablecaption{Line flux and equivalent width.\label{t4}} \tablehead{ \colhead{Epoch}& \colhead{JD}& \colhead{$F_\lambda(V)$ }& \multicolumn{2}{c}{Line flux}& \multicolumn{2}{c}{Equivalent width}\\ \colhead{}& \colhead{}& \colhead{[$\times$10$^{-12}$ Wm$^{-2}$$\mu$m$^{-1}$]}& \multicolumn{2}{c}{[$\times$10$^{-17}$ Wm$^{-2}$]}& \multicolumn{2}{c}{[$\times$10$^{-4}$ $\mu$m]}\\ \colhead{}& \colhead{}& \colhead{}& \colhead{v$\ge$2\tablenotemark{a}}& \colhead{v$=$1-0\tablenotemark{b}}& \colhead{v$\ge$2}& \colhead{v$=$1-0}} \startdata 28 Apr 2008 &2454585 & 4.5 & 9.7$\pm$ 1.8 & 8.3$\pm$ 0.8 & 1.6$\pm$ 0.3 & 1.4$\pm$ 0.1\\ 21 Jul 2008 &2454669 & 3.4 & 4.9$\pm$ 2.5 & 2.7$\pm$ 0.6 & 1.4$\pm$ 0.7 & 0.8$\pm$ 0.2\\ ~7 Aug 2008 &2454686 & 1.6 & 5.5$\pm$ 1.0 & 7.5$\pm$ 0.5 & 1.2$\pm$ 0.2 & 1.7$\pm$ 0.1\\ 14 Aug 2008 &2454693 & 1.2& 1.7$\pm$ 0.5 & 2.3$\pm$ 0.2 & 1.0$\pm$ 0.3 & 1.4$\pm$ 0.1\\ 17 Aug 2008 &2454696 & 0.4& 1.0$\pm$ 0.5 & 2.1$\pm$ 0.2 & 0.8$\pm$ 0.4 & 1.6$\pm$ 0.2\\ 21 Aug 2008 &2454700 & 0.2& 0.9$\pm$ 0.4 & 2.2$\pm$ 0.2 & 0.7$\pm$ 0.3 & 1.7$\pm$ 0.1\\ \enddata \tablenotetext{a}{Integrated over $\pm$120~km~s$^{-1}$.} \tablenotetext{b}{Integrated over $\pm$40~km~s$^{-1}$, after the spectral model with $T_{\rm x}=4500$~K is subtracted.} \end{deluxetable*} The heating mechanism of similar outburst variables, FUors, is most likely the elevated accretion near the disk-midplane, as the broad and double-peak CO lines in absorption testify \citep[e.g.][]{har04}. To the contrary, the CO vibrational lines of EX~Lupi are observed in emission in both overtone \citep{kos10} and fundamental bands. The critical difference that might explain the emergent spectra of the two classes of the outburst variables is the optical depth of the disk. The broad line emitting region of EX~Lupi is likely free of dust as is discussed above. The SED analysis by \citet{juh10} also excludes the optically thick continuum emission either by gas or by dust grains within 0.3~AU during the outburst. Moreover, the accretion rate during the outburst is a few orders magnitude smaller in EX~Lupi ($\sim 10^{-7}M_\odot$~yr$^{-1}$) than FUors ($\sim 10^{-4}M_\odot$~yr$^{-1}$), implying the lower disk surface density. The disk close to the midplane that is viscously heated by the accretion might be more exposed to the surface in the case of EX~Lupi, which possibly makes the CO vibrational band observed in emission. \begin{figure*} \begin{center} \includegraphics[height=0.2\textheight,angle=0]{f7.eps} \end{center} \caption{Schematic representation of the broad and the narrow line emitting regions discussed in the text. \label{car}} \end{figure*} \subsection{Origin of Outburst} Let us extend our speculation to the trigger of the outburst, taking the physical scale of it being traced by the broad line emission. The physical size of the outburst on the scale of 0.2--0.4~AU argues against mechanisms that predict a global disk instability. \citet{bon92} proposed that a binary system can tidally disrupt the disks of each other in a close encounter, and temporarily increase the accretion rate up to 10$^{-4}M_\odot$~yr$^{-1}$. The tidal disruption by a companion is attractive because it naturally explains the repetitive nature of EXor outbursts. The binary fraction among low-mass pre-main sequence stars is indeed high \citep[e.g.][]{ghe93}. FU~Ori, the prototype FUor, has a close companion \citep{wan04}, and there are even FUor-like binaries where both of the binary components are FUor-like, which is statistically unlikely if two are randomly paired \citep{rei04}. The gravitational disruption by a binary companion may work for the present case as a trigger of the disk instability combined with other mechanisms, but most likely does not explain the small radius of the high mass-accretion region by itself. No visible \citep{ghe97,bai98} or spectroscopic companion \citep{her07} to EX~Lupi has been identified as of today. \citet{vor06} argued that planetary cores -- which possibly have already formed in a gravitationally unstable disk -- can trigger an FUor outburst when they migrate inward and eventually fall into the central star. Protoplanetary clumps form at $>$10--50~AU away from the star. The migration from the outer disk over this distance does not explain the discrete boundary at 0.2--0.4~AU either. Their simulation ends at the inner radius of 0.5~AU. Further investigation how the infalling a protoplanetary clump behaves in the final 0.5~AU is awaited. When mass inflow from the outer disk is faster than what the inner disk can transport further in, the accretion slows down, and piles up material at the boundary between the inner and the outer disk. The elevated surface density triggers a local disk instability, and the accretion mode switches to a higher state. The mass accretion could slow down at a certain radius simply because the disk viscosity is too low, or the viscosity is locally too low because of the decoupling of the disk and the magnetic field, or if the disk has a gap, or if it is magnetically truncated. The disk viscosity is provided by the coupling of the magnetic field and slightly ionized medium in the MHD instability model. Deep inside the disk where the stellar radiation no longer penetrates, there is a low ionization region called the dead zone, where the viscosity is close to zero. The incoming mass piles up there, until the surface density becomes high enough to heat the disk to $\approx$1000~K to restore the coupling with the magnetic field. A dead zone starts at 0.1~AU, and extends beyond 1~AU in a disk around a solar mass star \citep{gam96}. The disk instability is triggered at the coldest region of the dead zone near the outer boundary \citep{arm01}. \citet{zhu07} constructed a radiation transfer model of the disk around FU~Ori in outburst, and calculated the size of the active region to be around 1~AU, which naturally fits with MHD disk instability. The observed transition region of EX~Lupi is a few times smaller than the typical size of a dead zone. This may not be surprising, as the luminosity of EX~Lupi in quiescence is small as well [$L_\ast<0.5L_\odot$ in Gras-Vel\'azquez \& Ray (2005), $L_\ast=0.75L_\odot$ in Sipos et al. (2009)]. \citet{wun05} discovered another type of instability near the inner edge of the dead zone in their numerical model. When a small perturbation is applied to the height of the dead zone, the mass inflow locally increases while the outgoing mass decreases. The height of the perturbed region grows, until it splits up from the dead zone and triggers a ring instability, or the inner edge of the dead zone starts oscillating radially \citep{wun06}. This instability works at the correct physical scale of 0.1--0.2~AU with small accretion rates $10^{-9}M_\odot$~yr$^{-1}$ which are appropriate for EX~Lupi. However, the predicted timescale of the outburst is 100~yrs or longer. This is close to the observed timescales of FUor outburst, but is too long for EX~Lupi ($\sim$1~yr). When a star is strongly magnetized, the gas infall stops at the radius where the Keplerian angular velocity is equal to the stellar rotation, as the orbital motion of the gas is tied to the star. The matter piles up at the co-rotation radius, until the surface density becomes so high that the thermal pressure of the gas overcomes the magnetic pressure \citep{dan10}. This mechanism requires the presence of a strong magnetic field. Among young eruptive stars, FU~Ori is the first where a magnetic field was discovered in its disk \citep{don05}. There are no attempts known to date, however, to measure the magnetic field of EX~Lupi. The inner rim of the high accretion region we measured (0.04--0.05~AU) translates to a period of 4--5 days assuming Keplerian rotation. The typical rotation time for T~Tauri stars is a few days to a week \citep[e.g.][]{rod09}. The rotation period of EX~Lupi itself is however not known yet. A self-regulated thermal instability theory assumes relatively high, constant inflow of mass \citep[$10^{-6}M_\odot$~yr$^{-1}$--$10^{-7}M_\odot$~yr$^{-1}$, ][]{cla89,bel94}. The constant mass transportation is hindered at a certain radius, where the viscosity is simply too low to drain the incoming mass in time. The material piles up until the surface density becomes high enough to thermally ionize the gas. Near the ionization front, the disk opacity increases with temperature due to negative hydrogen ions adding continuum opacity \citep{fau83}. Thermal instability is triggered by the runaway local heating of the disk. The transition front propagates outward until the instability is suppressed at the radius where the viscous heating is no longer effective to keep the disk warm. The radius at which the thermal instability is turned off is a function of the mass accretion. It is 0.1~AU for a fiducial star of 1~$M_\odot$ with 3~$R_\odot$ with a constant mass accretion of $3\times 10^{-6} M_\odot$ \citep{bel94}. Self-regulating thermal instability is well studied in the parameter space suited to FUors. The mechanism seems to work with accretion rates from $\dot{M}=10^{-7}M_\odot$~yr$^{-1}$ to $\dot{M}=10^{-5}M_\odot$~yr$^{-1}$, and is able to reproduce the observed flaring magnitude at optical wavelengths ($\approx$5~mag), the rise time to the maximum ($\sim$~1 year), and the duration of the outburst (decades to a century). These physical parameters are one order of magnitude smaller in the case of EX~Lupi, with the quiescent phase accretion rate being $\dot{M}=4\times 10^{-10}M_\odot$~yr$^{-1}$ \citep{sip09}, which goes up to $\dot{M}=2\times 10^{-7}M_\odot$~yr$^{-1}$ in the outburst \citep{juh10} with the rise time of a month, and duration of about a year. It is yet to be seen if the self-regulating thermal instability works for the small accretion rate with no external triggers to cooperate. An interesting possible triggering mechanism is an embedded planet. When a protoplanet opens a gap, the mass transportation stops at its outer edge. The mass is banked up until the thermal instability is triggered \citep{lod04,cla96}. With the additional degrees of freedom of the mass of the planet and its location, this model could cover a wider parameter space, and closely reproduce the observed properties of FUor outbursts. The triggering via embedded planet makes the instability propagate initially outside in, and qualitatively reduces the rise time. The smaller the mass of the planet, the smaller the outburst, triggered by the smaller accretion rate. It is still to be seen though, if the mechanism works for EX~Lupi, as the exact parameter space of EX~Lupi outburst is not covered in the models by \citet{lod04}. \subsection{Variation of Line Profile Asymmetry} Nevertheless, a possible planet in the disk is interesting in connection with the temporal variation of the line symmetry in CO emission spectra (Fig.~\ref{as}). The line asymmetry indicates that the receding side of the disk was brighter on August 7, but by August 14 the approaching side became brighter. The receding part is again brighter 3 days later, which again moved to the other side in 4 days. The modulation of the line profile on the timescale of 3--4 days was known in the visible spectra of FU~Ori and V1057~Cyg. \citet{her03} discussed two possible origins of the variation, that the modulation comes from the stellar spot and that it comes from the hot spot in the circumstellar disk, which may or may not be related to the hot accretion columns rising out of the disk. In the case of V1057~Cyg, the line modulation is not accompanied by photospheric variability, which makes the stellar spot hypothesis difficult. On the other hand, the line modulation in V1057~Cyg is stable over 3 yrs, which also makes the hot spot on a differentially rotating disk unlikely. The question of whether the hot spot resides on the photosphere or on the disk is therefore still open. The period of the spectroscopic modulation of EX~Lupi is hard to constrain with only six epochs of observations. If we simply take 7 and 3.5~days as half the period of the orbit in the first and the second half of August, the locations of the hot spot are 0.06 and 0.04~AU, respectively. This is within the span of the active disk that we derived from the outburst model, with the latter location close to the inner edge. Further spectroscopic monitoring with finer temporal sampling should be done in order to tell whether this represents real migration of a hot spot in the disk, as well as the theoretical study on the fate of protoplanteary clumps that fall into the star \citep{vor06}. A more general discussion on the line asymmetry caused by a planet embedded in the disk is found in \citet{reg10}.
| 10 | 12 |
1012.5301
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1012
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1012.1591_arXiv.txt
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\small{ A large number of cosmological parameters have been suggested for obtaining information on the nature of dark energy. In this work, we study the efficacy of these different parameters in discriminating theoretical models of dark energy, using both currently available supernova (SNe) data, and simulations of future observations. We find that the current data does not put strong constraints on the nature of dark energy, irrespective of the cosmological parameter used. For future data, we find that the although deceleration parameter can accurately reconstruct some dark energy models, it is unable to discriminate between different models of dark energy, therefore limiting its usefulness. Physical parameters such as the equation of state of dark energy, or the dark energy density do a good job of both reconstruction and discrimination {\it if} the matter density is known to high accuracy. However, uncertainty in matter density reduces the efficacy of these parameters. A recently proposed parameter, ${\rm Om}$, constructed from the first derivative of the SNe data, works very well in discriminating different theoretical models of dark energy, and has the added advantage of not being dependent on the value of matter density. Thus we find that a cosmological parameter constructed from the first derivative of the data, for which the theoretical models of dark energy are sufficiently distant from each other, and which is independent of the matter density, performs the best in reconstructing dark energy from SNe data. }
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The nature of dark energy is one of the most tantalizing mysteries in cosmology today. Observations of high redshift type Ia supernovae tell us that the expansion of the universe is accelerating at present \cite{union2}, which can not be satisfactorily explained in the standard cold dark matter (CDM) scenario. Independent observations of the cosmic microwave background \cite{wmap7} and large scale structure \cite{sdss} tell us that two-thirds of the present density of the universe is composed of some unknown component. The study of this unknown ``\de'' is of great interest among cosmologists today. Various theoretical models for \de~have been suggested, the simplest being the \cc~model with constant dark energy density and equation of state $\w = -1$. Other models of dark energy include physically motivated models like the scalar field quintessence and Chaplygin gas models, as well as geometrically motivated models like scalar-tensor theories and higher dimensional braneworld models (see \cite{de} and references therein). The growing number of theoretical \de~models has inspired a complementary, data-driven approach in which the properties of \de~are reconstructed from the data by studying the cosmological parameters characterizing \de. Two primary methods are used for reconstructing cosmological parameters. In the first approach, known as parametric reconstruction, a sufficiently general fitting function is used to represent the parameter in the analysis. This suffers from the possibility of bias, depending on the form chosen for the parameter. The second method is that of non-parametric reconstruction, in which no specific form is assumed for the parameter. The difficulty with this is that the parameters of interest are usually obtained by taking the first or second derivative of the data, therefore, direct reconstruction involving differentiation of noisy data can lead to large errors. Many different cosmological parameters have been suggested both these reconstruction methods (see \cite{parm} and references therein). In this work, we attempt to study the relative efficacy of the different cosmological parameters in reconstructing \de~from observations, and discriminating between different theoretical \de~models, using the parametric reconstruction formalism. The paper is arranged as follows-- section~\ref{meth} contains a description of the data and methods used in the analysis, section~\ref{res} outlines the results, and section~\ref{concl} presents the conclusions.
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We reconstruct four different cosmological parameters to examine their potential in extracting information from observations about \de. We find that the parameter $\omm$, which is constructed from the Hubble parameter, has the narrowest errors, and gives the most accurate reconstruction. We also see that results for the physical parameters $\omde, \w$ are extremely dependent on knowledge of the matter density, and without very tight bounds on matter density, physical parameters of \de~may not perform well in reconstruction \de. The deceleration parameter reconstructs quite well, but has low discriminatory powers, since different models of \de~have very similar values for this parameter. Overall it appears that the parameters constructed out of the first derivative of the data, \ie the Hubble parameter, are somewhat better poised to give information about the nature of dark energy than the second derivatives, and geometric parameters of \de, such as $\omm$, are better at reconstructing \de~since they are not biased by lack of knowledge about the matter density, provided they are constructed in such a way that they can discriminate between different models of \de. To obtain information about \de~from the physical parameters of \de, it is important to have independent sources of observation of the matter density. Thus, for a parameter to obtain maximum information out of the data, it works better if it is constructed out of the first derivative of the data or less, does not depend on other physical parameters such as the matter density, and is constructed such that different \de~models can be easily discriminated from each other with the parameter. We note here that the reconstruction methods considered here are that of parametric reconstruction, non-parametric reconstruction of cosmological parameters would suffer less from biases in reconstruction such as those Model C. However, typically these methods have higher errors, and also the issues of degeneracy with other, non-\de~parameters exist for these methods as well. Therefore if a simple cosmological parameter, like $\omm$, can be constructed which performs well for a large class of models, both in reconstructing and in discriminating between models, and is also independent of other non-\de~parameters such as $\omt$, it would be extremely useful for understanding the nature of \de.
| 10 | 12 |
1012.1591
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1012
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1012.3077_arXiv.txt
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{} {We applied a decomposition method to the energy dependent pulse profiles of the accreting binary pulsar A\,0535+26, in order to identify the contribution of the two magnetic poles of the neutron star and to obtain constraints on the geometry of the system and on the beam pattern. } {We analyzed pulse profiles obtained from RXTE observations in the X-ray regime. Basic assumptions of the method are that the asymmetry observed in the pulse profiles is caused by non-antipodal magnetic poles and that the emission regions have axisymmetric beam patterns. } {Constraints on the geometry of the pulsar and a possible solution of the beam pattern are given. We interpreted the reconstructed beam pattern in terms of a geometrical model of a hollow column plus a halo of scattered radiation on the neutron star surface, which includes relativistic light deflection.} {}
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The Be/X-ray binary A\,0535+26 was discovered by \textsl{Ariel V} in 1975 \citep{rosenberg75} during a giant outburst. The system consists of a neutron star orbiting the optical companion HDE\,245770 on an eccentric orbit ($e=0.47\pm0.02$) of orbital period $P_{\mathrm{orb}}=111.1\pm 0.3$\,d \citep{finger06}. The source presents quiescent X-ray emission with a luminosity of $L_{\mathrm{X}}\lesssim10^{35-36}\,\mathrm{erg}\,\mathrm{s}^{-1}$, sometimes interrupted by ``normal'' outbursts ($L_{\mathrm{X}}\approx10^{37}\,\mathrm{erg}\,\mathrm{s}^{-1}$) linked to the periastron passages of the neutron star, and less frequent ``giant'' outbursts ($L_{\mathrm{X}}>10^{37}\,\mathrm{erg}\,\mathrm{s}^{-1}$) of longer duration and less clearly related to the orbital phase \citep[see, e.g.,][]{giovanelli92,kend94,finger96}. The system presents two cyclotron resonant scattering features at $E\sim45$\,keV and $E\sim100$\,keV, from which a magnetic field strength of $B\sim 4\times10^{12}$\,G is inferred (\citealt{kend94,grove95,kretschmar05,caballero07}). Pulsations are observed to have a period of $P_{\mathrm{spin}}\sim$103\,s, typically with spin-up during stronger outbursts and spin-down during quiescent periods\footnote{see, e.g., the results of \textsl{Fermi}-GBM monitoring at\newline \url{http://gammaray.nsstc.nasa.gov/gbm/science/pulsars/}}. The pulse profile evolves from a complex profile at lower energies to a simpler, two-peaked structure at higher energies. This behavior is observed in several accreting X-ray pulsars. Similar to other sources \citep[e.g.,][]{staubert80}, individual pulses show strong pulse-to-pulse variations, while the average pulse profile is rather stable, with slower variations over the course of an outburst \citep{caballero08a}. The basic concept of pulsed emission is well understood. Pulsed emission originates in regions close to the magnetic poles of a rotating neutron star with the magnetic axis misaligned with respect to the rotation axis. In contrast, physical modeling of the pulsed emission turns out to be a complex task. Many processes are in fact involved in modeling pulse profiles, from the modeling of the emission regions and their local emission pattern to the formation of the pulse profiles seen by a distant observer. Comparison of model calculations with observations has been performed for instance by \citet{wang81}, \citet{meszaros85}, and \cite{leahy91}. A proper model calculation should include relativistic light deflection, which has a significant effect on the pulse shape\footnote{The importance of relativistic light deflection in model calculations can be visualized in \url{http://www.spacetimetravel.org/xpulsar06/xpulsar06.html}} \citep{riffert88}. For slowly rotating neutron stars, the metric around a neutron star can be approximated by the Schwarzschild metric (see, e.g., \citealt{pechenick83}). Due to the strong gravitational field around the neutron star, the X-rays will be observed at red-shifted energies. Geometrical models of filled and hollow accretion columns of accreting neutron stars, including relativistic light deflection, were computed in \cite{kraus01} and \cite{kraus03}. These models give the beam pattern or energy-dependent flux of one emission region as a function of the angle, as seen by a distant observer. Introducing the rotation of the pulsar and its geometry, i.e., the orientation of the rotation axis with respect to the direction of observation and the location of the two poles, the pulsed emission from each of the two poles (single-pole pulse profiles) observed by a distant observer can be modeled. The sum of the single-pole contributions gives the total pulse profile. An alternative method of analyzing pulse profiles is to start from the observed pulse profiles and, based on symmetry considerations, decompose the pulse profile into single-pole contributions. This is then transformed into the visible section of the beam pattern. This method has been successfully applied to the accreting X-ray binary pulsars Cen\,X-3\, Her\,X-1 and EXO\,2030+375 (\citealt{kraus96}, \citealt{blum00}, \citealt{sasaki10} respectively) and is applied in this work to A\,0535+26. Preliminary results were presented in \cite{caballero08c}.
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\label{sec:summary} In this work, a decomposition analysis was applied to the A\,0535+26 energy-dependent pulse profiles. A dipole magnetic field is assumed with axisymmetric emission regions. The asymmetry in the total pulse profiles is explained with a small offset from one of the magnetic poles from the antipodal position. We find a physically acceptable decomposition of the pulse profiles that allows us to extract information on the geometry of the pulsar. We obtain $\Theta_{1}\approx50^{\circ}$ and $\Theta_{2}\approx130^{\circ}$ for the position of the magnetic poles, and an offset of $\delta\approx25^{\circ}$. The visible section of the beam pattern was reconstructed. A characteristic feature of the reconstructed beam pattern at all energies is a minimum observed in the flux between $\theta\approx30^{\circ}-40^{\circ}$, where $\theta$ is the angle between the direction of observation and the magnetic axis. This was interpreted in terms of a simple geometrical model that includes relativistic light deflection. The model includes a hollow column emitting isotropically black body radiation, plus a thermal halo created on the neutron star surface around the column from scattered radiation emitted from the column walls. Another characteristic feature of the reconstructed beam pattern is a steep increase in flux at high values of $\theta$ ($\theta>120^{\circ}$). This could come from gravitational light bending, which produces a similar feature in model calculations. We performed model calculations for different column thicknesses and opening angles, and we found the best estimates of the half-opening angle and column thickness to be $\alpha_{\mathrm{o}}=0.2\,\mathrm{rad}$, $\alpha_{\mathrm{o}}-\alpha_{\mathrm{i}}=0.06\,\mathrm{rad}$. We would like to stress, however, that this model is simplified, and we do not claim it is true in all details, but it does reproduce the basic shape of the energy-dependent reconstructed beam pattern of A\,0535+26 for values of $\theta<40$\,\degr well. Computation of beam patterns at different energies has revealed a weak dependence of the minimum and its depth with the energy, suggesting that the minimum is mainly an effect of the geometry of the system, produced when the observer looks directly onto the accretion stream. This weak energy dependence on the minimum is also found in the reconstructed beam patterns of A\,0535+26.
| 10 | 12 |
1012.3077
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1012
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1012.3900_arXiv.txt
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{Absorption line indices are widely used to determine the stellar population parameters such as age and metallicity of galaxies, but it is not easy to obtain the line indices of some distant galaxies that have colours available. } {This paper investigates the correlations between absorption line indices and colours.} {A few statistical fitting methods are mainly used, via both the observational data of Sloan Digital Sky Survey and a widely used theoretical stellar population model.} {Some correlations between widely used absorption line indices and $ugriz$ colours are found from both observational data of early-type galaxies and a theoretical simple stellar population model. In particular, good correlations between colours and widely used absorption line indices such as D$_{\rm n}$(4000), H$\gamma_{\rm A}$, H$\gamma_{\rm F}$, H$\delta_{\rm A}$, Mg$_{\rm 1}$, Mg$_{\rm 2}$, and Mg$_{\rm b}$, are shown in this paper.} {Some important absorption line indices of early-type galaxies can be estimated from their colours using correlations between absorption line indices and colours. For example, age-sensitive absorption line indices can be estimated from $(u-r)$ or $(g-r)$ colours and metallicity-sensitive ones from $(u-z)$ or $(g-z)$. This is useful for studying the stellar populations of distant galaxies, especially for statistical investigations.}
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The determination of stellar population properties (e.g., age, metallicity, and element abundances) of galaxies has long been an important subject in galaxy studies \citep{Tantalo:2004}. The evolutionary population synthesis technique has been widely used in such works \citep [see, e.g., ][]{Renzini:2006}. Many works confirm that some spectral features, which are called absorption line indices (the most famous ones are the well-known Lick indices), can reliably estimate stellar population properties \cite [e.g., ][]{Worthey:1994, Kong:2001, Bruzual:2003, Thomas:2003, Gallazzi:2005}. The definitions of widely used line strength indices can be seen in the papers of \citet {Bruzual:1983}, \citet{Worthey:1994}, \citet{Worthey:1997}, \citet{Huchra:1996}, \citet{Diaz:1989}, \citet{Balogh:1999}, and \citet{Maraston:2009} among others. Although absorption line indices can determine the stellar population properties of galaxies well, this method cannot be used to study very distant (e.g., $z > 0.3$) galaxies because of the difficulty obtaining reliable spectral line indices. Meanwhile, some colours of such distant galaxies can be measured well. If some estimations of absorption line indices of galaxies can be derived from their colours, it will be able to investigate the stellar population properties of galaxies better, and then the formation and evolution of galaxies. This work proposes to study the correlations between colours and absorption line indices, and then presents a new method for estimating absorption line indices from the colours of galaxies. The organization of the paper is as follows. In Sect. 2, we briefly introduce the observational data and theoretical stellar population model used in this work. In Sect. 3, we study the correlations between colours and absorption line indices using the data of some early-type galaxies and theoretical stellar populations. Then the estimation of absorption line indices of galaxies from colours are discussed. Finally, we give our conclusions in Sect. 4.
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We have investigated the correlations between absorption line indices and $ugriz$ colours using observational data relating to SDSS and a widely used stellar population model. Both observational data of early-type galaxies and a theoretical stellar population model show good correlations between some line indices and $ugriz$ colours of stellar populations. These correlations can be fitted as linear, exponential, or polynomial relations. The fitting uncertainties are found to be similar to, but somewhat larger than, the metrical uncertainties of line indices. In particular, it was found that some widely used metallicity and age-sensitive line indices, such as Mg$_{\rm b}$, Fe$_{\rm {5270}}$, Fe$_{\rm {5335}}$, H$\delta_{A}$, H$\gamma_{\rm A}$, H$\gamma_{\rm F}$, and D$_{\rm n}$(4000), are well correlated to some $ugriz$ colours. The relative uncertainties caused by the fitted correlations are small when transforming colours into absorption line indices. The results suggest that some estimates of absorption line indices can be obtained from the colours of galaxies. Our results are possibly useful for estimating the absorption line indices of galaxies that have no reliable spectral data available, and then for studying the stellar metallicities, ages, and element abundances.
| 10 | 12 |
1012.3900
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1012
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1012.2137_arXiv.txt
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We present the results of time-integrated searches for astrophysical neutrino sources in both the northern and southern skies. Data were collected using the partially-completed IceCube detector in the 40-string configuration recorded between 2008 April 5 and 2009 May 20, totaling 375.5 days livetime. An unbinned maximum likelihood ratio method is used to search for astrophysical signals. The data sample contains 36,900 events: 14,121 from the northern sky, mostly muons induced by atmospheric neutrinos and 22,779 from the southern sky, mostly high energy atmospheric muons. The analysis includes searches for individual point sources and targeted searches for specific stacked source classes and spatially extended sources. While this analysis is sensitive to TeV--PeV energy neutrinos in the northern sky, it is primarily sensitive to neutrinos with energy greater than about 1~PeV in the southern sky. No evidence for a signal is found in any of the searches. Limits are set for neutrino fluxes from astrophysical sources over the entire sky and compared to predictions. The sensitivity is at least a factor of two better than previous searches (depending on declination), with 90\% confidence level muon neutrino flux upper limits being between $E^{2} dN/dE \sim 2-200 \times 10^{-12}~\mathrm{TeV \, cm^{-2} \, s^{-1}}$ in the northern sky and between $3-700 \times 10^{-12}~\mathrm{TeV \, cm^{-2}\, s^{-1}}$ in the southern sky. The stacked source searches provide the best limits to specific source classes. The full IceCube detector is expected to improve the sensitivity to $E^{-2}$ sources by another factor of two in the first year of operation.
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\label{sec1} Neutrino astronomy is tightly connected to cosmic ray (CR) and gamma ray astronomy, since neutrinos likely share their origins with these other messengers. With a possible exception at the highest observed energies, CRs propagate diffusively losing directional information due to magnetic fields, and both CRs and gamma rays at high energies are absorbed due to interactions on photon backgrounds. Neutrinos, on the other hand, are practically unabsorbed {\it en route} and travel directly from cosmological sources to the Earth. Neutrinos are therefore fundamental to understanding CR acceleration processes up to the highest energies, and the detection of astrophysical neutrino sources could unveil the origins of hadronic CR acceleration. Whether or not gamma ray energy spectra above about 10~TeV can be accounted for by only Inverse Compton processes is still an open question. Some observations suggest contributions from hadronic acceleration processes \citep{Morlino:2009ci,boettcher2009}. Acceleration of CRs is thought to take place in shocks in supernova remnants (SNRs) or in jets produced in the vicinity of accretion disks by processes which are not fully understood. Black holes in active galactic nuclei (AGN), galactic micro-quasars and magnetars, or disruptive phenomena such as collapsing stars or binary mergers leading to gamma ray bursts (GRBs), all characterized by relativistic outflows, could also be powerful accelerators. The canonical model for acceleration of CRs is the Fermi model \citep{PhysRev.75.1169}, called first-order Fermi acceleration when applied to non-relativistic shock fronts. This model naturally gives a CR energy spectrum with spectral index similar to $E^{-2}$ at the source. The neutrinos, originating in CR interactions near the source, are expected to follow a similar energy spectrum. More recently, models such as those in \citet{Caprioli:2009fv} can yield significantly harder source spectra. In the framework of these models it is possible to account for galactic CR acceleration to energies up to the knee, at about $Z \times 4\times10^{15}$~eV, where $Z$ is the atomic number of the CR. Extragalactic sources, on the other hand, are believed to be responsible for ultra-high energy CRs observed up to about $10^{20}$~eV. The concept of a neutrino telescope as a 3-dimensional matrix of photomultiplier tubes (PMTs) was originally proposed by \citet{Markov1961385}. These sensors detect the Cherenkov light induced by relativistic charged particles passing through a transparent and dark medium such as deep water or the Antarctic ice sheet. The depth of these detectors helps to filter out the large number of atmospheric muons, making it possible to detect the rarer neutrino events. The direction and energy of particles are reconstructed using the arrival time and number of the Cherenkov photons. High energy muon neutrino interactions produce muons that can travel many kilometers and are almost collinear to the neutrinos above a few TeV. The first cubic-kilometer neutrino telescope, IceCube, is being completed at the South Pole. IceCube has a large target mass. This gives it excellent sensitivity to astrophysical neutrinos, enabling it to test many theoretical predictions. Reviews on neutrino sources and telescopes can be found in \citet{Anchordoqui:2009nf, Chiarusi:2009ng, Becker:2007sv, Lipari:2006uw, Bednarek:2004ky, Halzen:2002pg, Learned:2000sw, Gaisser:1994yf}. Recent results on searches for neutrino sources have been published by IceCube in the 22-string configuration \citep{Abbasi:2009cv,Abbasi:2009iv}, AMANDA-II \citep{AMANDA7YR_Collaboration:2008ih}, Super-Kamiokande \citep{Thrane:2009tw}, and MACRO \citep{MACRO_Ambrosio:2000yx}. This paper is structured as follows: Sec.~\ref{sec2} describes the detector. The data sample and cut parameters are discussed in Sec.~\ref{sec2andahalf}, along with the simulation. In Sec.~\ref{sec3} the detector performance is characterized for source searches. Sec.~\ref{sec4} describes the unbinned maximum likelihood search method, and in Sec.~\ref{sec5} the point-source and stacking searches are discussed. After discussing the systematic errors in Sec.~\ref{sec6}, the results are presented in Sec.~\ref{sec7}. Sec.~\ref{sec8} discusses the impact of our results on various possible neutrino emission models, and Sec.~\ref{sec9} offers some conclusions.
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\label{sec9} A search for sources of high energy neutrinos has been performed using data taken during 2008--09 with the 40-string configuration of the IceCube Neutrino Observatory. Five searches were performed: 1) a scan of the entire sky for point sources, 2) a predefined list of 39 interesting source candidates, 3) stacking 16 sources of TeV gamma rays observed by Milagro and Fermi, along with an unconfirmed hot spot (17 total sources), 4) stacking 127 starburst galaxies, and 5) stacking five nearby CGs, testing four different models for the CR distribution. The most significant result of the five searches came from the all-sky scan with a p-value of 18\%. The cumulative binomial probability of obtaining at least one result of this significance or higher in five searches is 63\%. This result is consistent with the null hypothesis of background only. The sensitivity of this search using 375.5 days of 40-string data already improves upon previous point-source searches in the TeV--PeV energy range by at least a factor of two, depending on declination. The searches were performed using a data set of up-going atmospheric neutrinos (northern sky) and higher energy down-going muons (southern sky) in a unified manner. During 2010--2011, 79 strings of IceCube are operating and detector construction should finish during the austral summer of the same years. The full IceCube detector should improve existing limits by at least another factor of two with one year of operation. Additional improvement is foreseeable in the down-going region by developing sophisticated veto techniques and at lower energies by using the new dense sub-array, DeepCore, to its fullest potential.
| 10 | 12 |
1012.2137
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1012
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1012.5521_arXiv.txt
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The discovery of neutrino masses suggests the likely existence of gauge singlet fermions (right-handed neutrinos) that participate in the neutrino mass generation [1]. A sterile neutrino is a hypothetical neutrino that does not interact via any of the fundamental interactions of the Standard Model except gravity. It is a right-handed neutrino or a left-handed anti-neutrino. Such a particle belongs to a singlet representation with respect to the strong interaction and the weak interaction and has zero weak hypercharge, zero weak isospin and zero electric charge. Sterile neutrinos would still interact via gravity, so if they are heavy enough, they could explain cold dark matter or warm dark matter. The X-ray observations make use of the radiative decay of a sterile neutrino [2, 3], can yield a non-negligible flux from concentrations of dark matter in astrophysical systems, such as, e.g., galaxies, clusters, and dwarf spheroidal galaxies [1, 4]. The photons emitted from decays sterile neutrinos can affect the formation of the first stars. Their production in a supernova can also explain the pulsar kicks and they have many other implications in astrophysics and cosmology. It is of interest, therefore to study the interactions of sterile neutrinos in matter with the purpose of possibly using them to inform the direction of current and future experimental searches. \\
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Sterile neutrinos have important implications in cosmology and astrophysics which can be employed to inform the direction of current and future experimental searches. In this paper we have studied the interaction of sterile neutrinos with atoms and their role on ionization of atoms in MeV and GeV energy scale. We have also studied the interaction of sterile neutrinos with nuclei in the MeV and GeV energy scale. We obtained the relevant cross sections for both these two interactions. we have compared our results with the results of keV energy range. Although it seems difficult to detect sterile neutrinos, but the experimental approach should not be hopeless in the long run.\\
| 10 | 12 |
1012.5521
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1012
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1012.0242_arXiv.txt
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{Planetary nebulae (PN) are an excellent laboratory to investigate the nucleosynthesis and chemical evolution of intermediate mass stars. In these objects accurate abundances can be obtained for several chemical elements that are manufactured or contaminated by the PN progenitor stars, such as He, N, C, and also elements that were originally produced by more massive stars of previous generations, namely O, Ne, Ar, and S. Some of these elements are difficult to study in stars, so that PN can be used in order to complement results obtained from stellar data. In the past few years, we have obtained a large sample of PN with accurately derived abundances, including objects of different populations, namely the solar neighbourhood, the galactic disk and anticentre, the galactic bulge and the Magellanic Clouds. In this work, we present the results of our recent analysis of the chemical abundances of He, O, N, S, Ar and Ne in galactic and Magellanic Cloud PN. Average abundances and abundance distributions of all elements are determined, as well as distance-independent correlations. These correlations are particularly important, as they can be directly compared with the predictions of recent theoretical evolutionary models for intermediate mass stars.} \FullConference{11th Symposium on Nuclei in the Cosmos\\ 19-23 July, 2010 \\ Heidelberg, Germany.} \begin{document}
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Planetary nebulae (PN) are an excellent laboratory to investigate the nucleosynthesis and chemical evolution of intermediate mass stars. Accurate abundances can be obtained for several chemical elements, including (i) those elements that are manufactured by the PN intermediate-mass progenitor stars (He, N, C), and (ii) also elements that were originally produced by more massive stars of previous generations (O, Ne, Ar, S). The abundances of the first class of elements measured in PN include the original content previous to the formation of the progenitor stars and the contamination effects during the nuclear processes in these objects. As a consequence, PN can be used to study the nucleosynthetic processes in intermediate mass stars. On the other hand, elements such as O, Ne, etc. reveal the interstellar abundances at the time and place the progenitor stars were formed, so that the determination of their chemical abundances produces important observational constraints to the chemical evolution models for the galaxies hosting these objects. In the past few years, we have obtained a large sample of PN of different galactic populations with accurately derived abundances (cf. Maciel et al. \cite{mci09}, \cite{mci10}, and references therein). In this work we present average abundances and abundance distributions of several elements, as well as distance-independent abundance correlations that can be directly compared with the predictions of recent theoretical evolutionary models for intermediate mass stars.
| 10 | 12 |
1012.0242
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1012
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1012.0851_arXiv.txt
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We present high angular resolution Submillimeter Array (SMA) and Combined Array for Research in Millimeter-wave Astronomy (CARMA) observations of two GLIMPSE Extended Green Objects (EGOs)--massive young stellar object (MYSO) outflow candidates identified based on their extended 4.5 \um\/ emission in \emph{Spitzer} images. The mm observations reveal bipolar molecular outflows, traced by high-velocity \co(2-1) and \hco(1-0) emission, coincident with the 4.5 \um\/ lobes in both sources. SiO(2-1) emission confirms that the extended 4.5 \um\/ emission traces active outflows. A single dominant outflow is identified in each EGO, with tentative evidence for multiple flows in one source (G11.92$-$0.61). The outflow driving sources are compact millimeter continuum cores, which exhibit hot-core spectral line emission and are associated with 6.7 GHz Class II \meth\/ masers. G11.92$-$0.61 is associated with at least three compact cores: the outflow driving source, and two cores that are largely devoid of line emission. In contrast, G19.01$-$0.03 appears as a single MYSO. The difference in multiplicity, the comparative weakness of its hot core emission, and the dominance of its extended envelope of molecular gas all suggest that G19.01$-$0.03 may be in an earlier evolutionary stage than G11.92$-$0.61. Modeling of the G19.01$-$0.03 spectral energy distribution suggests that a central (proto)star (M \q10 \msun) has formed in the compact mm core (M$_{gas}$ \q 12-16\msun), and that accretion is ongoing at a rate of \q10$^{-3}$ \msun\/ year$^{-1}$. Our observations confirm that these EGOs are young MYSOs driving massive bipolar molecular outflows, and demonstrate that considerable chemical and evolutionary diversity are present within the EGO sample.
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Massive star formation remains a poorly understood phenomenon, largely due to the difficulty of identifying and studying massive young stellar objects (MYSOs)\footnote{We define MYSOs as young stellar objects (YSOs) that will become main sequence stars of M$>$8 \msun\/ (O or early-B type ZAMS stars).} in the crucial early active accretion and outflow phase. During the earliest stages of their evolution, young MYSOs remain deeply embedded in their natal clouds. Most massive star-forming regions are also distant ($>$ 1kpc) and crowded, with massive stars forming in close proximity to other MYSOs and to large numbers of lower-mass YSOs. Studying the early stages of massive star formation thus requires high angular resolution observations (to resolve individual objects in crowded regions) at long wavelengths unaffected by extinction. Large-scale \emph{Spitzer} surveys of the Galactic Plane have yielded a promising new sample of young MYSOs with \emph{active outflows}, which may be inferred to be actively accreting. Identified based on their extended 4.5 \um\/ emission in \emph{Spitzer} images, these sources are known as ``Extended Green Objects (EGOs)'' \citep{egocat,maserpap} or ``green fuzzies'' \citep{Chambers09} from the common coding of the 4.5 \um\/ band as green in 3-color IRAC images. In active protostellar outflows, the \emph{Spitzer} 4.5 \um\/ broadband flux can be dominated by emission from shock-excited molecular lines \citep[predominantly \h:][]{SmithRosen05,Smith06,Davis07,Ybarra09,Ybarra10,DeBuizer10}. The resolution of \emph{Spitzer} at 4.5 \um\/ ($\sim$2\pp) is sufficient to resolve the extended emission from outflows in massive star forming regions nearer than $\sim$7 kpc. Over 300 EGOs have been cataloged in the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE-I) survey area by \citet{egocat}. The mid-infrared (MIR) colors of EGOs are consistent with those of young protostars still embedded in infalling envelopes \citep{egocat}. A majority of EGOs are also associated with infrared dark clouds (IRDCs), identified by recent studies as sites of the earliest stages of massive star and cluster formation \citep[e.g.][]{Rathborne07,Chambers09}. Remarkably high detection rates for two diagnostic types of \meth\/ masers in high-resolution Very Large Array (VLA) surveys provide strong evidence that GLIMPSE EGOs are indeed \emph{massive} YSOs with active outflows \citep{maserpap}. There are two Classes of \meth\/ masers, both associated with star formation, but excited under different conditions by different mechanisms. Class II 6.7 GHz \meth\/ masers are radiatively pumped by IR emission from warm dust \citep[e.g.][and references therein]{Cragg05} and are associated exclusively with massive YSOs \citep[e.g.][]{Minier03,Bourke05,Xu08,Pandian08}. Class I 44 GHz \meth\/ masers are collisionally excited in molecular outflows, and particularly at interfaces between outflows and the surrounding ambient cloud \citep[e.g.][]{PlambeckMenten90, Kurtz04}. Of a sample of 28 EGOs, $>$64\% have 6.7 GHz Class II \meth\/ masers (nearly double the detection rate of surveys using other MYSO selection criteria), and of these 6.7 GHz maser sources, $\sim$89\% also have 44 GHz masers \citep{maserpap}. A complementary James Clerk Maxwell Telescope (JCMT; resolution \q20\pp) molecular line survey towards EGOs with 6.7 GHz \meth\/ maser detections found SiO(5-4) emission and \hco(3-2) line profiles consistent with the presence of active molecular outflows \citep{maserpap}. SiO is particularly well-suited to tracing \emph{active} outflows, as it persists in the gas phase for only $\sim$10$^{4}$ years after being released by shocks \citep[e.g.][]{pdf97}. A single-dish (resolution \q80\pp) 3 mm spectral line survey of all EGOs visible from the northern hemisphere by \citet{Chen10} found associated gas/dust clumps of mass 69-29000 \msun, consistent with the identification of EGOs as MYSOs. The nature of the driving sources of the 4.5 \um\/ outflows is only loosely constrained by the survey results. Bright ultracompact (UC) HII regions are, in most cases, ruled out as powering sources by the lack of VLA 44 GHz continuum detections \citep{maserpap}. A high detection rate (83\%) for thermal \meth\/ emission in the \citet{maserpap} JCMT survey indicates the presence of warm dense gas, and possible hot core line emission. Further understanding of the nature of EGOs, and their implications for the mode(s) of high-mass star formation, requires identifying the driving source(s) and characterizing their physical properties, as well as those of the outflows associated with EGOs. Interferometric millimeter-wavelength line and continuum observations provide access to \emph{direct} tracers of molecular outflows and dense, compact gas and dust cores, including a wealth of chemical diagnostics. In this paper, we present Submillimeter Array (SMA)\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} and Combined Array for Research in Millimeter-wave Astronomy (CARMA)\footnote{Support for CARMA construction was derived from the Gordon and Betty Moore Foundation, the Kenneth T. and Eileen L. Norris Foundation, the James S. McDonnell Foundation, the Associates of the California Institute of Technology, the University of Chicago, the states of California, Illinois, and Maryland, and the National Science Foundation. Ongoing CARMA development and operations are supported by the National Science Foundation under a cooperative agreement, and by the CARMA partner universities.} observations at 1 and 3 mm of two EGOs from the \citet{maserpap} sample: G11.92$-$0.61 and G19.01$-$0.03. The targets were chosen to have bipolar (and in some cases quadrupolar) 4.5 \um\/ morphology, associated 24 \um\/ emission, associated (sub)mm continuum emission in single-dish surveys, and 6.7 Class II and 44 GHz Class I \methanol\/ maser detections in the \citet{maserpap} survey. The promise of extended 4.5 \um\/ emission as a MYSO diagnostic lies largely in its ability to identify very young sources with ongoing accretion and outflow that are missed by other sample selection methods. These sources had not been targeted for study prior to their identification as EGOs and inclusion in the \citet{maserpap} sample, and very little is known about them beyond the results of that survey (see also \S\ref{g11_previous} and \S\ref{g19_previous}). In \S\ref{obs} we describe our observations, and in \S\ref{results} we present our results. In \S\ref{discussion} we discuss the physical properties of the compact cores and outflows associated with our target EGOs, and in \S\ref{conclusions} we summarize our conclusions.
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Our high-resolution millimeter observations of two EGOs unambiguously show that they are young MYSOs driving massive bipolar outflows. The spatial coincidence of high velocity \co(2-1) and \hco(1-0) emission with the extended 4.5 \um\/ lobes supports the outflow hypothesis for the 4.5 \um\/ emission. A single dominant outflow is identified in each EGO, with tentative evidence for multiple outflows in one source (G11.92$-$0.61). The morphology and kinematics of the SiO(2-1) emission differ from the other outflow tracers in that some of the strongest red and blueshifted features are offset from the extended 4.5 \um\/ emission, and may trace the impact of outflow shocks on dense gas in the surrounding cloud. The morphology of the high-velocity gas with respect to 44 GHz Class I \meth\/ maser emission further solidifies the association of this type of maser with outflows. Anomalously intense and narrow components of 229.759 GHz \meth\/ emission are also detected in the outflow lobes from both objects, suggesting additional Class I maser activity. The outflow driving sources appear as compact cores of millimeter continuum emission and dense gas, including the hot core molecules \meth, \methcyn\/ and OCS. Coincident with 22 GHz water maser emisson, G11.92$-$0.61-MM1 shows considerably richer and stronger hot core line emission than G19.01$-$0.03-MM1, consistent with its warmer temperature derived from the multi-transition analysis of the \methcyn\/ and \meth\/ emission (166$\pm$20 v. 114$\pm$15 K). Both hot cores exhibit 24 \um\/ and 70 \um\/ emission in MIPSGAL images and contain 6.7 GHz Class II \meth\/ masers, all consistent with their identification as MYSOs. Our observations also reveal considerable diversity within the EGO sample. Although observed at the same spatial resolution, G19.01$-$0.03 appears as a single MYSO while G11.92$-$0.61 resolves into a cluster of three compact dust cores. In addition to the difference in multiplicity, several other factors point to G19.01$-$0.03 being in a earlier evolutionary stage: SED modeling, the relative weakness of its hot core emission, and the dominance of the extended envelope of molecular gas. In contrast, G11.92$-$0.61 appears to have already formed a protocluster whose members span a range of ages -- one is a hot core and two are almost entirely devoid of line emission. These initial results demonstrate the potential of the EGO sample for probing the importance of protostellar feedback in the formation of massive stars and star clusters. The future capabilities of the EVLA and ALMA will enable uniform surveys of a statistically meaningful number of regions which will enable the relative importance of outflows, photoionization, and radiative feedback to be assessed.
| 10 | 12 |
1012.0851
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1012
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1012.0583_arXiv.txt
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The dense concentration of stars and high velocity dispersions in the Galactic centre imply that stellar collisions frequently occur. Stellar collisions could therefore result in significant mass loss rates. We calculate the amount of stellar mass lost due to indirect and direct stellar collisions and find its dependence on the present-day mass function of stars. We find that the total mass loss rate in the Galactic centre due to stellar collisions is sensitive to the present-day mass function adopted. We use the observed x-ray luminosity in the Galactic centre to preclude any present-day mass functions that result in mass loss rates $> 10^{-5} \mathrm{M_{\odot} yr^{-1}}$ in the vicinity of $\sim 1''$. For present-day mass functions of the form, $dN/dM \propto M^{-\alpha}$, we constrain the present-day mass function to have a minimum stellar mass $\lesssim 7\mathrm{M_{\odot}}$ and a power law slope $\gtrsim 1.25$. We also use this result to constrain the initial mass function in the Galactic centre by considering different star formation scenarios.
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The dense stellar core at the Galactic centre has a radius of $\sim 0.15 - 0.4 \mathrm{pc}$, a stellar density $> 10^6\mathrm{M_{\odot} pc^{-3}}$ \citep{genzel:1996, eckart:1993, genzel:2003, schodel:2007} high velocity dispersions ($\ge 100\mathrm{km} \, \mathrm{s^{-1}}$), and Sgr A*, the central supermassive black hole with a mass $\approx 4 \times 10^{6}M_{\odot}$ \citep{eckart:2002, schodel:2002, schodel:2003, ghez:2003, ghez:2008}. Due to the extreme number densities and velocities, stellar collisions are believed to play an important role in shaping the stellar structure around the Galactic centre, and in disrupting the evolution of its stars. \citet{genzel:1996} found a paucity of the brightest giants in the galactic center, and proposed that collisions with main sequence (MS) stars could be the culprit. This hypothesis was found to be plausible by \citet{alexander:1999}. Other investigations of collisions between giants and MS, white dwarf and neutron stars \citep{bailey:1999} and collisions between giants and binary MS and neutron stars \citep{davies:1998} could not account for the dearth of observed giants. The contradictory results were resolved by \citet{dale:2009}, who concluded that the lack of the faintest giants (but not the brightest giants) could be explained by collisions between giants and stellar mass black holes. Significant mass loss in the giants' envelopes after a collision would prevent the giants from becoming bright enough to be observed. The above studies concentrated on collisions involving particular stellar species with particular stellar masses. To examine the cumulative effect of collisions amongst an entire ensemble of a stellar species with a spectrum of masses, one must specify the present-day stellar mass function (PDMF) for that species. The PDMF gives the current number of stars per unit stellar mass up to a normalization constant. Given a certain star formation history, the PDMF can be used to determine the initial mass function of stars (IMF), the mass function with which the stars were born. There is currently no consensus as to whether the IMF in the Galactic centre deviates from the canonical IMF \citep{bastian:2010}. First described by Salpeter more than 50 years ago \citep{salpeter:1955}, the canonical IMF is an empirical function which has been found to be universal \citep{kroupa:2001}, with the Galactic centre as perhaps the sole exception. \citet{maness:2007} found that models with a top-heavy IMF were most consistent with observations of the central parsec of the Galaxy. \citet{paumard:2006}, and subsequently \citet{bartko:2010} found observational evidence for a flat IMF for the young OB-stars in the Galactic centre. On the other hand, \citet{lockmann:2010} concluded that models of constant star formation with a canonical IMF could explain observations of the Galactic centre. In this work we use calculated mass loss rates due to stellar collisions as a method to constrain the PDMF for main sequence stars in the Galactic centre. We construct a simple model to estimate the actual mass loss rate in the Galactic centre based on observed x-ray emission. PDMFs that predict mass loss rates from stellar collisions greater than the observed rate are precluded. This method allows us to place conservative constraints on the PDMF, because we do not include the contribution to the mass loss rate from stellar winds from massive evolved stars \citep{baganoff:2003}. Specifically, this method allows us to place a lower limit on the power-law slope and an upper limit on the minimum stellar mass of the PDMF in the Galactic centre (see \S~\ref{sec:constrain}). Inclusion of the mass loss rate from stellar winds (or other sources) could further constrain the PDMF of the Galactic centre. We present novel, analytical models to calculate the amount of stellar mass lost due to stellar collisions between main sequence stars in \S~\ref{sec:condition} through \S~\ref{sec:mass_loss_direct}. In \S~\ref{sec:coll_rates} we develop the formalism for calculating collision rates in the Galactic center. We utilize our calculations of the mass loss per collision, and the collision rate as a function of Galactic radius to find the radial profile of the mass loss rate in \S~\ref{sec:mass_loss_rates}. Since the amount of mass lost is dependent on the masses of the colliding stars, the mass loss rate in the Galactic centre is sensitive to the underlying PDMF. By comparing our calculations to mass loss rates obtained from the x-ray luminosity measured by \textit{Chandra}, in \S~\ref{sec:constrain} we constrain the PDMF of the Galactic centre. We derive analytic solutions of the PDMF as a function of an adopted IMF for different star formation scenarios, which allows us to place constraints on the IMF in \S 6. In \S~7, we estimate the contribution to the mass loss rate from collisions involving red giant (RG) stars.
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We have have derived novel, analytical methods for calculating the amount of mass loss from indirect and direct stellar collisions in the Galactic centre. Our methods compares very well to hydrodynamic simulations, and do not require costly amounts of computation time. We have also computed the total mass loss rate in the Galactic centre due to stellar collisions. Mass loss from direct collisions dominates at Galactic radii below $\sim0.1\mathrm{pc}$, and thereafter indirect collisions dominate the total mass loss rate. Since the amount of stellar material lost in the collision depends upon the masses of the colliding stars, the total mass loss rate depends upon the PDMF. We find that the calculated mass loss rate is sensitive to the PDMF used, and can therefore be used to constrain the PDMF in the Galactic centre. As summarized by Fig.~\ref{fig:param_space}, our calculations rule out $\alpha \lesssim 1.25$ and $M_{\min} \gtrsim 7 \mathrm{M_{\odot}}$ in the $M_{min}-\alpha$ parameter space. Finally, we have used our constraints on the PDMF in the Galactic centre to constrain the IMF to have a power-law slope $\gtrsim $ 0.4 to 0.9 depending on the star formation history of the Galactic centre.
| 10 | 12 |
1012.0583
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1012
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1012.0310_arXiv.txt
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{Differences in the stellar populations of galaxies can be used to quantify the effect of environment on the star formation history. We target a sample of early-type galaxies from the Sloan Digital Sky Survey in two different environmental regimes: close pairs and a general sample where environment is measured by the mass of their host dark matter halo. We apply a blind source separation technique based on principal component analysis, from which we define two parameters that correlate, respectively, with the average stellar age ($\eta$) and with the presence of recent star formation ($\zeta$) from the spectral energy distribution of the galaxy. We find that environment leaves a second order imprint on the spectra, whereas local properties -- such as internal velocity dispersion -- obey a much stronger correlation with the stellar age distribution.}
| 10 | 12 |
1012.0310
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1012
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1012.5703_arXiv.txt
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We report our 110 ks Chandra observations of the supernova remnant (SNR) 0104-72.3 in the Small Magellanic Cloud (SMC). The X-ray morphology shows two prominent lobes along the northwest-southeast direction and a soft faint arc in the east. Previous low resolution X-ray images attributed the unresolved emission from the southeastern lobe to a Be/X-ray star. Our high resolution Chandra data clearly shows that this emission is diffuse, shock-heated plasma, with negligible X-ray emission from the Be star. The eastern arc is positionally coincident with a filament seen in optical and infrared observations. Its X-ray spectrum is well fit by plasma of normal SMC abundances, suggesting that it is from shocked ambient gas. The X-ray spectra of the lobes show overabundant Fe, which is interpreted as emission from the reverse-shocked Fe-rich ejecta. The overall spectral characteristics of the lobes and the arc are similar to those of Type Ia SNRs, and we propose that \theSNR\ is the first case for a robust candidate Type Ia SNR in the SMC. On the other hand, the remnant appears to be interacting with dense clouds toward the east and to be associated with a nearby star-forming region. These features are unusual for a standard Type Ia SNR. Our results suggest an intriguing possibility that the progenitor of \theSNR\ might have been a white dwarf of a relatively young population.
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Among about a dozen supernova remnants (SNRs) detected in the Small Magellanic Cloud (SMC), only three X-ray brightest SNRs (0102-72.3, 0103-72.6, and 0049-73.6) have been relatively well studied. Because all these three SNRs most likely originate from a core-collapse SN, revealing the type of other SNRs is of great importance to study the star-forming history, SN rates, and chemical evolution of the SMC. \theSNR\ is the fourth brightest X-ray SNR in the SMC, whose origin has been controversial. Its optical emission was suggested to be Balmer-dominated \citep{1984ApJS...55..189M}, indicating Type Ia origin. On the other hand, with the ROSAT data, \citet{1994AJ....107.1363H} identified an unresolved X-ray source in the southern part of the SNR as a candidate Be/X-ray star which appeared to be co-spatial with the SNR. Based on this plus other evidence for a Population I environment, they suggested a core-collapse origin for \theSNR. With the XMM-Newton data, \citet{2004A&A...421.1031V} found evidence for enhanced Fe L line emission from the remnant, with which they reinstated Type Ia origin. Recently, with Akari IRC observations, \citet{2007PASJ...59S.455K} discovered bright infrared shells which are positionally coincident with the optical \halpha\ filaments surrounding the X-ray emission \citep{1994AJ....107.1363H}. IR color distributions indicated shock interaction with ambient molecular clouds, which may support a core-collapse explosion of a massive progenitor star for the origin of \theSNR. We report here on the results from our observations of \theSNR\ with the \emph{Chandra} X-Ray observatory. We note that the remnant has been serendipitously detected a number of times by \emph{Chandra} during the ACIS calibration observations of SNR 0102-7219. However, \theSNR\ was detected mostly in FI CCDs at large off-axis angles. Due to the low sensitivity and the poor spatial resolution, these archival data were not suitable for our analysis and we only report results based on our new observations.
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\label{sec:discussion} \subsection{The SN type of \theSNR} Based on the ROSAT HRI observation, \citet{1994AJ....107.1363H} attributed the unresolved X-ray emission feature at $\alpha_{2000}$, $\delta_{2000}$ = \hms{01}{06}{18}, \dms{-72}{06}{00} to emission associated with a candidate Be star, interpreting it as a Be/X-ray star. The arcsecond resolution and good photon statistics of our deep Chandra data reveal that this X-ray emission feature is extended (Figure~\ref{fig:chand}), and no point-like source is found within a few arcsecond from the position of the Be star candidate (on-axis astrometric uncertainties of the ACIS is $\sim 0.6\arcsec$). We examined the RGB composite image and various narrow band images of different energy ranges and their ratio images, but could not find any point-like feature that is positionally associated with the Be candidate. The overall X-ray spectrum of this region is also similar to X-ray spectra in other regions. Thus we conclude that the X-ray emission in this region originates from the shocked SNR gas, and not from the point source. Although the Be star that \citet{1994AJ....107.1363H} identified is unlikely to be a significant X-ray emitter, the possibility that \theSNR\ and the Be star belong to the same OB association, as they also suggested, remains valid, but would then argue against a typical Type Ia remnant. This will be further discussed in \S~\ref{sec:env}. \label{sec:comp-ir} The overabundance of Fe in the bar region suggests that \theSNR\ is likely the remnant of Type Ia SN. Indeed, the spectra of the bar regions are quite similar to ejecta spectra found in the well-known Type Ia SNR DEM~L71 \citep{2003ApJ...582L..95H} in the Large Magellanic Cloud (LMC) and several other relatively old Type Ia LMC SNRs \citep[e.g.,][and references therein]{2006ApJ...652.1259B}. DEM~L71's central X-ray emission, which is well described by shocked Fe-rich ejecta, is similar to that of \theSNR. DEM~L71 also shows faint soft emission of normal LMC abundances surrounding the central emission, similar to the arc in \theSNR. Therefore, based on the overabundant Fe from our spectral fit and the spectral similarity of the bar and the arc region to other Type Ia SNRs, we propose that \theSNR\ is the first solid candidate for Type Ia SNR in the SMC. \MYNOTE{ Another element expected to be a signature of Type Ia SNR is Si \citep[e.g.,][]{2007ApJ...662..472B}. Allowing the Si abundance to vary in the fit yields a slightly larger value, but still comparable to the SMC abundance within the fit uncertainty. } The Type Ia origin is supported by non-detection of a pulsar and/or pulsar wind nebula, although they could be too faint to be detected as in other core-collapse SNRs in the SMC \citep[e.g.,][]{2003ApJ...598L..95P}. \subsection{Type Ia SNR in Unusual Environment?} \label{sec:env} Our deep Chandra observations find that the X-ray emission from the bar region is prominently from Fe-enriched SN ejecta of a Type Ia SNR. Based on limited sample, the X-ray morphologies of young Type Ia SNRs have been claimed to be statistically more symmetric than core-collapse SNRs \citep{2009ApJ...706L.106L}. Although a quantitative analysis of the asymmetry in the morphology of \theSNR\ should be performed (which is beyond the scope of this work), the highly elongated morphology of \theSNR\ appears to be inconsistent with the results by \citet{2009ApJ...706L.106L}. Such an elongated Fe-rich ejecta feature is similar to that found in the Galactic SNR W49B. Although the type is uncertain for W49B \citep[e.g.,][]{2006A&A...453..567M}, the highly elongated morphology of ejecta may suggest an asymmetric explosion. Alternatively, given the possibility that the remnant is interacting with dense material toward the east (see below), the atypical morphology could be due to highly inhomogeneous ambient environment of the remnant. The environment around \theSNR\ is also unusual for a Type Ia SNR. The relatively high absorbing column density toward \theSNR\ indicates that the remnant could be located at dense environment. The eastern X-ray arc is positionally coincident with an optical and infrared shell (Figure~\ref{fig:ir-comp}). The optical emission shows dominant \halpha\ line with weak lines of other ions such as [\ion{O}{3}] and [\ion{S}{2}], and \citet{2007MNRAS.376.1793P} found that their line ratios are consistent with radiative shocks in other SNRs. \citet{2007PASJ...59S.455K} reported infrared emission associated with \theSNR\ from their \emph{Akari} observations. The observed infrared colors were consistent with those from shocked molecular clouds, and \citet{2007PASJ...59S.455K} suggested that \theSNR\ is probably interacting with molecular clouds. In this regard, the radio continuum emission of \theSNR\ could be mostly from the shocked ambient gas. Although the radio data quality is poor, Figure~\ref{fig:chand}(b) shows that the extent of the radio continuum covers the eastern arc region, and also the central region of the bar where the Fe overabundance appears to be less prominent. Thus, % observations seem to provide a consistent picture that the arc is where the remnant is interacting with dense ambient (possibly molecular) clouds. On the other hand, the NW and SE parts of the SNR, where only the X-ray emitting ejecta are visible and there is no optical or IR emission, probably corresponds to low density regions where the forward shock in the ISM is still relatively fast and radiative shocks have not formed. In fact, the projected location of \theSNR\ is at the eastern boundary of the superbubble DEM S124 \citep{1976MmRAS..81...89D} (see Figure~\ref{fig:ir-comp}). In Figure~\ref{fig:ir-comp}(c), we show the \halpha\ emission complex toward the direction of \theSNR\ and overlay locations of known SNRs in the SMC. It appears that \theSNR\ and SNR~0102-7219, a well-known core-collapse SNR, belong to the same star-forming region. The Type~Ia origin inferred from the Fe-rich ejecta and the association of the remnant with the star-forming region are not consistent with conventional scenarios of standard Type Ia, which involve the explosion of an old C/O white dwarf in a binary. Recent studies of a large sample of Type Ia SNe suggested that progenitors of some Type Ia SNe (``prompt'' type) could have been relatively young massive stars with high metallicity \citep[e.g.,][]{2005ApJ...629L..85S,2008A&A...492..631A}. \MYNOTE{ More recently, from the systematic study of the SNRs in the Magellanic Clouds, \citet{2010MNRAS.407.1314M} suggested that some of the SNRs could be remnants of the prompt Type Ia SNe. } While we cannot completely rule out the possibility of mere coincidence, the unusual nature of \theSNR\ (i.e., Type Ia characteristics of Fe-rich X-ray ejecta in the SNR and the SNR's spatial association with an environment with a strong star-forming activity) suggests an intriguing possibility that \theSNR\ may be a candidate SNR of a prompt SN Ia.
| 10 | 12 |
1012.5703
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1012
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1012.1120_arXiv.txt
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Active galactic nuclei (AGNs) exhibit variability across the entire electromagnetic spectrum with distinct flaring episodes at different frequencies. The high sensitivity and nearly uniform sky coverage of the Large Area Telescope on board the {\it Fermi}\, satellite make it a powerful tool for monitoring a large number of AGNs over long timescales. This allowed us to detect several flaring AGNs in $\gamma$ rays, triggering dedicated multifrequency campaigns on these sources from radio to TeV energies. We discuss the results for two different types of flaring AGN: the flat spectrum radio quasar 3C 279, in particular the coincidence of a $\gamma$-ray flare from this source with the drastic change of the optical polarization angle, and the first $\gamma$-ray flare from a radio-loud narrow-line Seyfert 1, PMN J0948$+$0022.
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Since its launch on 2008 June 11, the {\it Fermi} Gamma-ray Space Telescope has opened a new era in high-energy astrophysics. The primary instrument on board {\it Fermi}, the Large Area Telescope (LAT), is a pair-conversion telescope covering the energy range $\sim$ 20 MeV to 300 GeV with unprecedented sensitivity and effective area~\cite{Atwood}. The combination of deep and fairly uniform exposure over two orbits, very good angular resolution, and stable response of the LAT has allowed it to produce the most sensitive, best-resolved survey of the $\gamma$-ray sky, and to efficiently find episodes of flaring from $\gamma$-ray sources of different nature. One of the major scientific goals of the {\it Fermi} mission is to investigate the high-energy emission in Active Galactic Nuclei (AGNs) in order to understand the mechanisms by which the particles are accelerated and the precise site of the $\gamma$-ray emission, and investigate on long timescales the AGN variability and the $\gamma$-ray duty cycle. With respect to previous $\gamma$-ray instrument, such as EGRET~\cite{Thompson} and AGILE~\cite{Tavani}, the LAT provides opportunities to investigate more in detail the behaviour of flaring $\gamma$-ray AGNs. Best examples are the extraordinary outbursts of 3C 454.3 in December 2009 and April 2010 (see~\cite{Escande} for details). Together with simultaneous multiwavelength observations collected over the entire electromagnetic spectrum, LAT measurements allow us to reach a deeper insight on the jet structure and the emission mechanisms at work in AGNs. In the following we focus on two different types of AGN: the Flat Spectrum Radio Quasar (FSRQ) 3C 279, and the radio-loud Narrow-Line Seyfert 1 (RL-NLS1) PMN J0948$+$0022.
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The $\gamma$-ray all-sky monitoring by {\it Fermi}-LAT, combined with other simultaneous ground- and space-based observations, provide us new insights into the relativistic jets and broad emission models of AGNs. We presented the results from the multifrequency campaign of the FSRQ 3C 279, including the discovery of a $\gamma$-ray flare associated with a drastic change of the optical polarization angle as well as the detection of an ``orphan'' X-ray flare. In addition, the discovery by {\it Fermi}-LAT of $\gamma$-ray emission from 4 RL-NLS1s provided evidence for relativistic jets in these systems. The $\gamma$-ray flare of PMN J0948$+$0022 in July 2010 confirmed that extreme power can be produced also in this class of AGNs.
| 10 | 12 |
1012.1120
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1012
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1012.4705_arXiv.txt
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We study luminosities of millisecond pulsars in globular clusters by fitting the observed luminosity distribution with single and double power laws. We use simulations to model the observed distribution as the brighter part of some parent distribution for Terzan 5 and try to find a model which simultaneously agrees with the observed diffuse radio flux, total predicted number of pulsars and observed luminosity distribution. We find that wide ranges of parameters for log-normal and power-law distributions give such good models. No clear difference between the luminosity distributions of millisecond pulsars in globular clusters and normal disk pulsars was seen.
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Millisecond pulsars (MSPs) in globular clusters (GCs) can be used as a tool to understand the properties of GCs as well as the recycling process in the dense stellar environments inside GCs \citep{cr05}. As luminosity is a fundamental property of pulsars, one necessary step to achieve this goal is to understand the luminosity distribution of MSPs in GCs. A full dynamical approach, where one models evolution of pulsars and observational limits following appropriate choice of birth distributions of pulsar parameters, can not be adopted for GC pulsars where it is difficult to model the effects of stellar encounters and the cluster potential. The simplest way is to use a snapshot approach where one models pulsar luminosities as observed. We first adopt this method and then, using Monte Carlo simulations, try to find a good model which not only fits the observed luminosity distribution, but also agrees with the total radio and $\gamma$-ray fluxes for Terzan 5.
| 10 | 12 |
1012.4705
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1012
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1012.1857_arXiv.txt
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Two aspects in studies of X-ray emission from massive stars attract most attention: {\em i)} how X-rays are generated in massive stars, and {\em ii)} how X-ray emission can be used in analyzing stellar winds. In the basic concept, the wind has two components: a general cool wind with temperature of $T_{\rm w}\sim 10$\,kK which contains nearly all the wind mass, and a hot tenuous component with $\Tx\sim$few\,MK where the X-rays originate. The X-ray photons suffer continuum K-shell absorption in the cool wind and, in turn, can affect the wind ionization via the Auger process. In this review we concentrate on X-ray emission from single stars. {\changed This is thermal emission from gases heated in the stellar wind shocks or in magnetically confined wind regions. Cassinelli \& Olson (1979) proposed X-radiation from a base coronal zone plus Auger ionization in the surrounding cool wind to explain the superionzation (e.g.\ N\,{\sc v}, O\,{\sc vi}) that was observed to be present in Copernicus UV spectra of OB stars. From the analysis of {\em Einstein} spectra of OB-stars, Cassinelli \& Swank (1983) concluded that the base corona idea was not correct since soft X-rays were observed. The Si\,{\sc xiii} and S\,{\sc xv} line emission was detected in the SSS spectrum of the O-star $\zeta$ Ori. These ions correspond to high temperature and are located at a energy where the wind would be thin to X-rays. This led to a conclusion about two sources of X-ray emission, X-rays that arise from fragmented shocks in the wind and X-rays from very hot, probably magnetically confined loops, near the base of the wind. Furthermore since X-ray variability was already known to be less than about 1\%, Cassinelli \& Swank (1983) suggested that there had to be thousands of shock fragments in the wind. }Radiation hydrodynamic simulations of the nonlinear evolution of instabilities in stellar winds were performed by Owocki, Castor, \& Rybicki (1988). They demonstrated that the X-ray can originate from plasma heated by strong reverse shocks, which arise when a high-speed, rarefied flow impacts on slower material that has been compressed into dense shells. Feldmeier, Puls, \& Pauldrach (1997) assumed a turbulent seed perturbation at the base of the stellar wind and found that the shocks arising when the shells collide are capable of explaining the observed X-ray flux. These 1D hydrodynamical models predict plasma with temperatures 1--10\,MK which is permeated with the cool wind. X-rays suffer absorption as they propagate outwards through the ensemble of dense, radially compressed shells. Waldron (1984) calculated the opacity of O-star winds for the X-ray radiation. The absorption of X-rays in Wolf-Rayet (WR) star winds was investigated by Baum \etal\ (1992). They employed detailed non-LTE stellar atmosphere models and showed that since the WR wind opacity is very high, the observed X-rays must emerge from the far outer wind region. Hillier \etal\ (1993) computed the wind opacity of the O5Ia star $\zeta$\,Pup. They found that the high opacity of the stellar wind would completely block the soft X-rays ($<0.5$\,keV) unless some significant fraction of hot plasma is located far out in the wind, at distances exceeding 100\,\Rstar. The shape of X-ray emission line profiles was predicted by MacFarlane \etal\ (1991). They considered the effect of wind absorption on the emission from an expanding shell of hot gas. When the cool wind absorption is small, the line is broad and has a box-like shape. For stronger wind absorption, the line becomes more skewed (see Fig.\,7 in MacFarlane \etal\ 1991). The line shape is largely determined by a parameter $\tau_0$: \begin{equation} \tau_0=\kappa_\lambda R_\ast = \rho_{\rm w}\chi_{\lambda} R_\ast \label{eq:t0} \end{equation} where $R_\ast$\,[cm] is the stellar radius, and the atomic opacity $\kappa_\lambda$ is the product of the mass absorption coefficient $\chi_{\lambda}$ [cm$^2$\,g$^{-1}$] and the density of the cool wind ($\rho_{\rm w}$) as defined from the continuity equation {\changed{ $\dot{M}=4\pi\rho_{\rm w} v(r)r^2R_\ast^2$, where $r$ is the radial distance in units of $R_\ast$, and $v(r)$ is the velocity law, that can be prescribed by the formula $v(r)=v_\infty(1-1/r)^\beta$.}} MacFarlane \etal\ notice that when $\tau_0$ increases, the red-shifted part of the line ($\Delta\lambda>0$) becomes significantly more attenuated than the blue-shifted part. They suggested that evaluating the line shape can be used to determine $\tau_0$. {\changed{ The K-shell opacity varies with wavelength with the power between 2 and 3 (Hillier \etal\ 1993), therefore in the X-ray band $\tau_0$ should change by orders of magnitude. Consequently, the X-ray emission line shape at shorter and longer wavelengths should be different.}} Waldron \& Cassinelli (2001) expanded the single-shock model of MacFarlane \etal\ and considered emission from spherically symmetric shocks equally distributed between 0.4\vinf\ and 0.97\vinf\, with temperatures ranging from 2 to 10\,MK. In similar spirit, Ignace (2001) provided a formalism that accounts for the emission from a flow that is embedded with zones of X-ray emitting gas. Owocki \& Cohen (2001) calculated model X-ray line profiles for various combinations of the parameters $\beta$, $\tau_0$, and onset radii for X-ray emission.
| 10 | 12 |
1012.1857
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1012
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1012.3549_arXiv.txt
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We impose constraints on the topology of the Universe determined from a search for matched circles in the temperature anisotropy patterns of the 7-year \emph{WMAP} data. We pay special attention to the sensitivity of the method to residual foreground contamination of the sky maps, and show that for a full sky estimate of the CMB signal (the ILC map) such residuals introduce a non-negligible effect on the statistics of matched circles. In order to reduce this effect, we perform the analysis on maps for which the most contaminated regions have been removed. A search for pairs of matched back-to-back circles in the higher resolution \emph{WMAP} W-band map allows tighter constraints to be imposed on topology. Our results rule out universes with topologies that predict pairs of such circles with radii larger than $\alpha_{\rm min} \approx 10^\circ$. This places a lower bound on the size of the fundamental domain for a flat universe of about 27.9 Gpc. This bound is close to the upper limit on the size of Universe possible to detect by the method of matched circles, i.e.\ the diameter of the observable Universe is 28.3 Gpc.
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According to General Relativity, the pseudo-Riemannian manifold with signature (3,1) is a mathematical model of spacetime. The local properties of spacetime geometry are described by the Einstein gravitational field equations. However, they do not specify the global spatial geometry of the universe, i.e.\ its topology. This can only be constrained by observations. The concordance cosmological model assumes that the universe possesses a simply-connected topology, yet recently detected anomalies on large angular scales in the cosmic microwave background (CMB) anisotropy suggest that it may be multiply-connected. Evidence of such anomalies comes from the suppression of the quadrupole moment, alignment of the quadrupole and octopole and asymmetry in the temperature anisotropy observed in two hemispheres on the sky \citep{de Oliveira-Costa:2004, copi:2004, eriksen:2004, hansen:2004, schwarz:2004}. The present constraints on topology were placed by studying two signatures of multi-connectedness in the CMB maps: the large scale damping of the power in the direction of a small dimension of the domain, which causes a breakdown of statistical isotropy \citep{de Oliveira-Costa:2004,kunz:2006, kunz:2008}, and the distribution of matched patterns \citep{cornish:2004, aurich:2005, aurich:2006, then:2006, key:2007}. In this work, we constrain the topology of the Universe using the method of matched circles proposed by \citet{cornish:1998} and apply it to the 7-year \emph{WMAP} data \citep{jarosik:2010}. In contrast to the majority of previous studies, we will pay special attention to the impact of Galactic foreground residuals on the constraints. Some consideration of this problem was made by \citet{then:2006} in their analysis of the first year \emph{WMAP} data release. We will use also Monte Carlo (MC) simulations for the estimation of the false detection level of the statistic. The method is applied to higher resolution maps than previously, which implies a lower level of false detection and therefore tighter constraints on the size of the Universe. As a result of computational limitations, we will restrict the analysis to a search for back-to-back circle pairs\footnote{pairs of circles centred around antipodal points}. In the following two sections we describe data used in analysis and simulations of the CMB maps for a flat universe with the topology of a 3-torus which were used to test the reliability of our implementation. The statistic adopted in our studies is presented in \sect\ref{sec:statistic}. The results and conclusions are described in the last two sections.
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\label{sec:conclusions} We have studied constraints on the topology of the Universe using the method of matched circles as applied to the 7-year \emph{WMAP} ILC map and the foreground-reduced W-band map. We paid special attention to three aspects of the analysis that have been neglected in previous studies -- the application of a mask, the use of high resolution data and the estimation of the false detection level on the basis of detailed MC simulations of the sky maps. The necessity for the application of a mask is due to the presence of residual Galactic foreground emission present even in the ILC map. These introduce a non-negligible effect on the matched circles statistic that is used for constraining topology. However, the possibility to apply the analysis to masked maps yields the opportunity to more tightly constrain topology by using higher resolution, foreground corrected \emph{WMAP} W-band data. Constraints on the topology depend significantly on the threshold for a significant match of the circle pairs. In order to estimate this correctly, we used 100 MC simulations of the ILC and W-band maps assuming a simply-connected universe. The level of false detection calibrated in this manner is slightly higher than that derived in \citet{cornish:2004} using a method in which the analyzed maps are resampled by shuffling their spherical harmonic coefficients. Although the difference is not very big, it is worth noting that the constraints on the size of Universe are overestimated if one uses a lower level of false detection. The analysis of the \emph{WMAP} W-band map, after correction using templates of Galactic foreground emission, did not reveal any significant correlations for pairs of back-to-back circles with a radius greater than $\sim 10^\circ$. This substantially extends the previous constraint on the minimum radius of detectable matched circles given by \citet{key:2007} of $20^\circ$. It also places a lower bound on the size of fundamental domain of about 27.9 Gpc for a flat universe described by the best-fit 7-year \emph{WMAP} cosmological parameters. Although this constraint concerns only those universes with such dimensions and orientation of the fundamental domain with respect to the mask that allow the detection of pairs of matched circles, the probability of overlooking circle pairs is rather low for the KQ85y7 mask that removes only a relatively small fraction of the sky. Of course, observations of the CMB with higher angular resolution and significantly lower noise level by the \emph{Planck} satellite may yield even tighter constraints on the topology of the Universe. However, one should bear in mind that the possible improvement in the lower bound on the size of the Universe will not be substantial. The current constraint is not much smaller than the diameter of the observable Universe $2\, R_{\rm LSS} = 28.3\ {\rm Gpc}$, which imposes an upper limit on the size of the fundamental domain that it is possible to detect using the method of matched circles. The only significant improvement might be related to improved modeling of the Galactic emission allowing the application of a smaller mask approaching closer to the Galactic plane. This would minimize the probability of overlooking topologies with matched circles that are hidden within the masked region of the sky. Finally, as in \citet{cornish:2004}, the studies could also be extended to search for nearly back-to-back circle pairs. However, the much higher computational requirements for the analysis of high resolution maps make such studies difficult and extremely time-consuming. Nevertheless, we can hope that the steadily increasing speed of processors and availability of larger computational resources will make such computations feasible in the coming years, thus allowing a final resolution of the problem of the topology of our Universe.
| 10 | 12 |
1012.3549
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1012
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1012.4369_arXiv.txt
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Carr and Hawking showed that the proper size of a spherical overdense region surrounded by a flat FRW universe cannot be arbitrarily large as otherwise the region would close up on itself and become a separate universe. From this result they derived a condition connecting size and density of the overdense region ensuring that it is part of our universe. Carr used this condition to obtain an upper bound for the density fluctuation amplitude with the property that for smaller amplitudes the formation of a primordial black hole is possible, while larger ones indicate a separate universe. In contrast, we find that the appearance of a maximum is not a consequence of avoiding separate universes but arises naturally from the geometry of the chosen slicing. Using instead of density a volume fluctuation variable reveals that a fluctuation is a separate universe iff this variable diverges on superhorizon scales. Hence Carr's and Hawking's condition does not pose a physical constraint on density fluctuations. The dynamics of primordial black hole formation with an initial curvature fluctuation amplitude larger than the one corresponding to the maximum density fluctuation amplitude was previously not considered in detail and so we compare it to the well-known case where the amplitude is smaller by presenting embedding and conformal diagrams of both types in dust spacetimes.
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Einstein's general relativity (GR) permits the fascinating possibility that matter and fields interact with geometry in such a way that compact regions form from which temporarily no information can escape. Furthermore, assuming reasonable properties of matter and fields, the appearance of such a closed horizon (classically) implies the existence of a curvature singularity in the interior and that information can never escape \cite{HE73}. These configurations are black holes (BHs). In a spherically symmetric spacetime, which we assume throughout the paper, a horizon and hence a BH forms as soon as the areal radius $R$ of a region which determines the measurable surface area $A= 4 \pi R^2$, decreases and becomes smaller than twice the enclosed mass $M$ \footnote{We use natural units: $c=G=1$}. Achieving this matter configuration requires strong compression of the material and in most circumstances pressure forces will overcome the gravitational forces preventing the BH formation. A dying star is a well-known exception. Depending on its mass gravity can overcome the forces exerted by the neutrons in the collapsed stellar core. In this way the quantum physics of neutrons determines the minimal mass of a stellar BH to be larger than the Tolman-Oppenheimer-Volkoff limit, which is a few solar masses. Another possibility is that BHs form in the very early universe from primordial density perturbations \cite{ZN66, H71,CH74,C75}. These BHs are called primordial black holes (PBHs) and are the subject of this paper (we do not consider the formation involving the collapse or collision of defects, see for instance \cite{CS82,H89, RKS00, K10}). The formation process is bottom-up like in the ordinary structure formation scenario since only perturbations smaller than the Hubble radius $R<H^{-1}$ can collapse. The maximum size of a PBH forming in the early universe is approximately given by the Hubble radius $H^{-1}_\hc$ at the time $t_\hc$ when the perturbation crosses the horizon (see Sec.\,\ref{derivnosu}). As a consequence, PBH sizes range in principle between the Planck length and today's Hubble radius. The actual mass spectrum of PBHs then depends on the spectrum of primordial density perturbations and the equation of state. Numerical simulations of PBH formation in a radiation filled universe \cite{NNP78, BH79, SS99, HS02, MMR05, PM07} confirmed Carr's estimate \cite{C75}, that the initial fluctuation amplitude of a density perturbation must be very high in order not to disperse but to form a PBH. PBHs would open up windows into many interesting aspects of GR. For instance, the formation process in a radiation-filled universe is a critical phenomenon \cite{C93,NJ99,HS02,MMR05, MMP08}. The mass spectrum of PBHs strongly depends on the probability distribution of metric fluctuations \cite{C75, ZMN82, GLMS04}, which depends in turn on the mechanism that generates the primordial density perturbations \cite{INN94, GLW96, BP97, I98, Y98, Y99}, such that a (non-)observation of PBHs constrains these mechanisms. And most importantly, while Hawking radiation \cite{H74} of stellar BHs is much colder than the CMB and hence undetectable, PBHs could have already evaporated or evaporate at the current time \cite{C75,C76, BCL91}, which could give rise to a gravitational wave background \cite{P76, BR04, AEG09}. Additionally, quantum gravity effects may alter the evaporation process at its final stage such that there could have remained a Planck mass remnant for each evaporated PBH \cite{M87,BCL92}, while PBH evaporation could have created the baryons \cite{BCKL91,DNN00,AM07}. Many more reasons for studying PBH and more references can be found in \cite{C05}. While there is still no direct observation of a BH, there is convincing evidence for the existence of stellar BHs \cite{MR03} and supermassive BHs in the centers of galaxies \cite{R84, A93}, in particular our own Galaxy \cite{GETAG09,BLN09}. In contrast there is no trace of PBHs and their existence is strongly constrained by observations \cite{JGM09,CKSY10}. However it is still possible that both dark matter and supermassive BHs may be the results of PBH evaporation \cite{M87} and PBH clustering \cite{C06,KRS05}. The intention of the first part of this paper is to clarify a certain aspect of PBH formation which originated from the earliest treatments of PBHs \cite{H71, CH74,C75} and can still be found in recent works on PBH formation, e.g. in \cite{BP97,NJ98, HC05, MMP08, PM07, HP09} and probably many more. Carr and Hawking considered the collapse of an initially superhorizon-sized homogeneous spherical overdensity surrounded by an otherwise flat FRW universe \cite{CH74,C75} (see Fig.\,\ref{sphconfig}). They estimated the necessary amplitude of the fluctuation for PBH formation using an energy argument: the gravitational energy at the moment $t_\m$ of maximal expansion, where the kinetic energy and the expansion rate is zero in the interior, has to overcompensate the internal energy density that will cause the pressure forces during the subsequent contraction. For a radiation-filled universe, they found that this requires the overdensity to have a size $R_\m \gtrsim \rho_\m^{-1/2}$, where $\rho_\m=\rho(t_\m)$ is the density inside the fluctuation at maximal expansion $t_\m$. On the other hand, it was noted that the space inside the fluctuation at $t_\m$ can only have positive curvature determined by $\rho_\m$. This leads to a maximal proper size $L_\m$ of the overdensity since the region would otherwise close up on itself and represent a separate universe (SU). They concluded that having no SU requires $R_\m \lesssim \rho_\m^{-1/2}$ \cite{CH74}. Hence Harada and Carr concluded that a fluctuation must be finely tuned in order to become a black hole but not a separate universe \cite{HC05}. One aim of this paper is to show that the no-SU condition actually does not pose a constraint on density fluctuations, and hence fine tuning is not necessary. We see this from several points of view in Secs. \ref{derivnosu} -- \ref{densfluc}. Carr's and Hawking's no-SU condition will be quickly rederived in Sec.\,\ref{derivnosu} keeping numerical factors in order to show why it does not constrain fluctuations. Defining the amplitude of a fluctuation in terms of a curvature fluctuation variable $\zeta$ in Sec.\,\ref{curvfluc}, we will conclude that a SU corresponds to $\zeta = \infty$ on superhorizon scales and hence SUs are physically impossible and irrelevant for PBH formation. In contrast to a result in \cite{C75}, we find in Sec.\,\ref{densfluc} that the density fluctuation $\delta$ features a maximal value $\delta_\ma$ which is independent of the no-SU condition and instead is a consequence of the spatial geometry of the fluctuation determined by the slicing and the Hamiltonian constraint. Furthermore it will be clear that this maximum arises naturally if one expresses $\delta$ in terms of $\zeta$. To every $\delta$ (except $\delta_\ma$) belong two different $\zeta$'s; one is larger and the other smaller than $\zeta_\ma\equiv\zeta(\delta_\ma)$. We will point out the relevance of the rather conceptual findings concerning SUs for the existing PBH literature in Sec.\,\ref{results}. \begin{figure}[t] \centering \includegraphics[width=0.25\textwidth]{sphconfig} \caption{Sketch of an idealized fluctuation: flat FRW for $r> r_b$ and closed FRW for $r<r_a$} \label{sphconfig} \end{figure} The second part of the paper is devoted to gaining intuition about PBH formation: in Sec.\,\ref{ltbcoll}, we compare the collapse of two spherical overdensities with the same $\delta$ but different $\zeta$ patched to a flat dust FRW universe using the LTB metric. The comparison will be carried out with the help of embedding diagrams and conformal diagrams, which enables us to intuitively understand the differences and similarities of both types of fluctuations. Throughout the paper we consider highly idealized models of density fluctuations. This is because they are analytically tractable and at the same time exhibit all the necessary features to clarify and show some new aspects of PBH formation.
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\label{concl} Type II ($\zeta>\zeta_\ma$) and marginal ($\delta=\delta_\ma$) fluctuations sometimes misinterpreted as separate universes do form primordial black holes. The formation process looks completely different compared to the type I class ($\zeta<\zeta_\ma$) in synchronous gauge. This is only a gauge artifact of the synchronous and related ``natural'' slicings and can be avoided in the quite unnatural $\tilde{\eta}$-slicing used in the conformal diagram. The appearance of $\delta_\ma$ is independent of the no-SU condition formulated by Carr and Hawking \eqref{rhomaxhaw}. It is impossible to violate this condition due to \eqref{rhomaxhar}. The actual limiting value to form a SU obtained by closing a type II is $\delta_\su\rightarrow 0$ corresponding to $\zeta\rightarrow\infty$. The nonmonotonicity of $\delta(\zeta)$ suggests not to use the density fluctuation variable $\delta$ but instead the curvature fluctuation variable $\zeta$ if one also wants to consider larger fluctuations, for instance in a probability distribution function of fluctuation amplitudes. For homogeneous overdensities every type I fluctuation has a $\delta$-twin in the type II class related by $\chi_\ti=\pi-\chi_\tii$. If there is no pressure, both fluctuations form a PBH with exactly the same mass although the formation process looks very different for a type II. This is related to the fact that the gravitating mass energy that one can probe by using Kepler's law is not the integral of the density over the proper volume but over the volume of the would-be flat space using the areal radius $R$ (see Eq.\,\eqref{MSmass}). A difference between type I and II arises if there is pressure present. We expect the existence of a shape and pressure dependent threshold $\zeta_\mathrm{nc}$ with the property that for $\zeta\geq\zeta_\mathrm{nc}>\zeta_\mathrm{max}$ a non-central singularity forms, while for $\zeta_\mathrm{nc}>\zeta\geq \zeta_\mathrm{max}$ the type II fluctuation relaxes via pressure driven matter loss within the horizon to a centrally collapsing fluctuation. To find this threshold it is necessary to perform a numerical simulation. \appendix
| 10 | 12 |
1012.4369
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1012
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1012.1250_arXiv.txt
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The new technique of Speckle Stabilization \citep[][]{E08,K08,K10} has great potential to provide optical imaging data at the highest angular resolutions from the ground. While Speckle Stabilization was initially conceived for integral field spectroscopic analyses, the technique shares many similarities with speckle imaging (specifically shift-and-add and Lucky Imaging). Therefore, it is worth comparing the two for imaging applications. We have modeled observations on a 2.5-meter class telescope to assess the strengths and weaknesses of the two techniques. While the differences are relatively minor, we find that Speckle Stabilization is a viable competitor to current Lucky Imaging systems. Specifically, we find that Speckle Stabilization is 3.35 times more efficient (where efficiency is defined as signal-to-noise per observing interval) than shift-and-add and able to detect targets 1.42 magnitudes fainter when using a standard system. If we employ a high-speed shutter to compare to Lucky Imaging at 1\% image selection, Speckle Stabilization is 1.28 times more efficient and 0.31 magnitudes more sensitive. However, when we incorporate potential modifications to Lucky Imaging systems we find the advantages are significantly mitigated and even reversed in the 1\% frame selection cases. In particular, we find that in the limiting case of Optimal Lucky Imaging, that is zero read noise {\it and} photon counting, we find Lucky Imaging is 1.80 times more efficient and 0.96 magnitudes more sensitive than Speckle Stabilization. For the cases in between, we find there is a gradation in advantages to the different techniques depending on target magnitude, fraction of frames used and system modifications. Overall, however, we find that the real strength of Lucky Imaging is in observations of the brightest targets at all frame selection levels and in observations of faint targets at the 1\% level. For targets in the middle, we find that Speckle Stabilization regularly achieves higher $S/N$ ratios.
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Speckle Stabilization (SS) is a new technique that has the potential to achieve diffraction-limited spatial resolutions in the optical regime from ground-based telescopes \citep[][]{E08,K08,K10}. The technique is based on a relatively simple idea. On short timescales, atmospheric turbulence gets frozen out into a speckle pattern. Each one of these moderate-to-low Strehl speckles is at the diffraction-limited spatial resolution of the telescope. By tracking the speckle patterns in real time with a high speed camera (such as an electron-multiplying CCD, EMCCD), it is possible to identify the location of the brightest speckle and stabilize it onto a traditional ``low-speed'' science camera through the use of a fast steering mirror. Over time the system tracks and stabilizes bright speckles as they appear and the result is a diffraction-limited core surrounded by a diffuse halo of scattered light. This was demonstrated through simulation by \citet{K08}. In many ways, SS is simply a real-time implementation of the shift-and-add (SAA) technique developed by \citet{BC80} who showed that by stacking speckle images on top of one another based on the location of the best quality speckle, one could extract much higher spatial resolutions. Recently, on-sky tests by \citet{K10} using a SS proof-of-concept instrument have produced z' images with spatial resolutions of $\approx 3\lambda /D$ on the 4.2-meter William Herschel Telescope. The instrument consists of a Andor DU-860 EMCCD for speckle sensing, an Optics-in-Motion fast steering mirror for guiding and an SBIG ST-237 as a science camera with light picked off via a beam splitter. Using this instrument, they were also able to resolve the components of the binary star system WDS 14411+1344 which were blended in the seeing-limit. This instrument is the first implementation of the Stabilized sPeckle Integral Field Spectrograph (SPIFS) envisioned by \citet{E08}. SPIFS will be a system capable of achieving diffraction-limited angular resolutions at optical wavelengths. What is unique about this instrument over other techniques is that it would exploit these high angular resolutions with an integral field spectrograph (IFS). An IFS behind a SS system highlights a primary advantage of SS over Lucky Imaging and SAA: long exposures are possible in the science field. This enables both faint target imaging and spectroscopy. As a result, SPIFS could be coupled to an 8- or 10-meter class telescope and produce resolutions as fine as 10 milliarcseconds complete with associated spectral information (although at low Strehl values $\approx2\%$). As diffraction-limited angular resolutions are currently elusive in the optical from the ground for IFU work, SPIFS will be able to probe into astrophysical structures at unprecedented resolutions. SS will enable supermassive black holes to be measured in many more galaxies via the calcium triplet line, young stellar objects will have more of their structure probed, and even dense stellar fields will be spectroscopically classified. A particular example of the SPIFS potential was given in \citet{E08} describing how SPIFS will be able to resolve the optical jets produced by the micro-quasar SS 433. In the simulations, the jets were assumed to be 30 mas from the source as per the model developed by \citet{E01}. At this distance, a broad band image of the system would merely reveal elongation. However, when integral field spectroscopy is performed, the distinct locations of the jet and counter-jet become apparent. These observations would be invaluable to further our understanding of jet astrophysics as well as accretion processes. The ability to access spectroscopic information is one of the reasons why SPIFS is being developed with 8 or 10-meter class telescopes in mind and is the primary advantage of a speckle-stabilization system over any other type of speckle imaging system. In this paper, we also examine Lucky Imaging-- a descendant of the SAA technique. For Lucky Imaging, the driving principle is that random fluctuations in atmospheric turbulence will occasionally result in high-quality, high-Strehl images. This technique was first proposed by \citet{F78} who calculated the fraction of time one would expect high Strehl images. Because these fluctuations are not necessarily long lived, a Lucky Imager needs to take thousands of images at high speeds and selects only the highest quality ones for analysis \citep[][]{Bald01,T02,Law06,Law07}. The resulting PSF from Lucky imaging observations is similar to a stabilized speckle PSF. That is, there is a diffuse halo with a distinct, sharp core. The primary difference is the fraction of light present in the diffraction-limited core, i.e. the Strehl ratio. The key point, however, is that both Lucky imaging and SS produce \emph{the same} spatial resolutions in the core. Lucky Imaging has produced interesting science in the fields of high-resolution companion searches where it has helped constrain the binary fraction of M dwarfs and M subdwarfs (Law et al. 2006, Lodieu et al. 2009, Bergfors et al. 2010) and has also been helpful in determining if exoplanet systems also have multiple stellar components \citep[][]{D09}. Lucky Imaging, by itself, acquires a usable fraction ($\approx 5-20\%$) of high-Strehl speckle images up to telescopes 2.5-m in diameter, however much beyond that size the fraction of time when high-Strehl images are present is effectively zero \citep[][]{F67,F78,Law07,S09}. Because of this drop off in high-Strehl images, Lucky Imaging produces data nearly identical to SAA on large telescope. As a way to increase the potential of Lucky Imaging on large telescopes, \citet{Law09} have put a Lucky Imager behind an adaptive optics system at Palomar and were able to acquire near diffraction-limited spatial resolutions at Strehl ratios of $\approx 0.10$. The purpose of this paper, however, is to compare Speckle Stabilization to SAA and Lucky Imaging purely in imaging mode. We note that there are inherent differences between SS and speckle techniques and our goal is to determine if these differences result in a significant advantage to SS. One of the differences we are interested in is the number of read outs each technique requires. SS only needs to read out the science camera a few times during observations, whereas Lucky Imaging has to read out thousands of times. Even though the read noise on the EMCCDs employed by speckle teams is typically very low (on the order of $0.2e^-$/pix), it is still present and we hypothesize it can become a significant noise term when observing faint targets. Another significant term we think might have an effect on sensitivity is the additional $\sqrt{2}$ noise term present in images taken with an EMCCD. This stochastic noise is introduced in the readout process and can only be worked around in photon-counting mode. Therefore, we present simulations to determine if a Speckle Stabilization imager has the capability to overcome these features of Lucky Imaging systems in a way that is advantageous to an observer.
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We find that Speckle Stabilization is a viable competitor to current Lucky Imaging systems when used solely for imaging in certain circumstances. The results from all of our models and simulations are presented in Table 1. Both SS and Lucky Imaging have their own strengths. In general, SS is more sensitive and efficient than speckle imaging at the faintest magnitudes with normal frame selection meaning it could be well-employed for faint object work. Additionally, for most mid-magnitude targets SS and SS+S is a factor of $\sqrt{2}$ times more efficient owing to the lack of the additional noise term. As a counter, there are several cases in which speckle techniques are superior. We find that the Lucky Imaging techniques are the only way to observe the brightest stars with any usable \emph{S/N}. Our work has also revealed that simple modifications to traditional Lucky Imaging systems can greatly improve their performance and completely mitigate the advantage of a SS system at the faint end. The most effective alterations are approaching zero read noise and using photon counting techniques beyond a particular threshold. When these features are implemented, we found the Lucky Imaging was both more sensitive and efficient than Speckle Stabilization at the faint end. We highlight again that the main advantage to Speckle Stabilization is long exposures for IFU work, but this paper has revealed SS is also useful from an imaging perspective. Overall the differences are fairly minor between the output products, but as telescope time is a valuable commodity it is useful to have instruments in place that are able to perform efficient observations. We find that speckle stabilization is one way to achieve this aim, but similar goals can be met by modifying existing Lucky Imaging systems. While SS is still in its infancy, instruments like SPIFS will be able to reveal some of the potential of this technique and help solve outstanding issues in astrophysics.
| 10 | 12 |
1012.1250
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1012
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1012.2968_arXiv.txt
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We present our analysis aimed at inferring average magnetic fields in slowly-rotating solar-like stars. Using the spectral line inversion code \texttt{SPINOR}, we perform high-accuracy line profile fitting and investigate whether Zeeman broadening can be reliably detected in optical data of unprecedented quality. We argue that our usage of both high- and low-$g_{\rm{eff}}$ lines does provide a certain sensitivity to magnetic fields that may, indeed, be detected. However, the measurement is subject to a model dependence and prone to ambiguities, e.g. due to spectral blends. Hence, while a field may be successfully recovered, the quantification of this field is subject to large uncertainties, even for the highest-quality optical data.
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Direct detections of magnetic fields in slowly-rotating ($v\sin i < 10\,\rm{km\,s}^{-1}$) solar-like stars (SRSLS) are somewhat rare. Since the first application of the Zeeman broadening technique to cool stars by \cite{1980ApJ...239..961R}, much effort was made to infer their field strengths, e.g. in order to constrain stellar dynamo theory. The more recent advances, however, come from different directions, such as direct detections for very late-type stars, e.g. \cite{2007ApJ...656.1121R}, and detections of magnetic field geometry reversals for faster rotators \citep{2009MNRAS.398.1383F}. The rotation-activity relationship and the absence of accurate and recent measurements of SRSLS magnetic fields thus suggest that the corresponding field strengths may be too low to enable clear and robust detections. We aim to shed some light on the question of Zeeman broadening detectability in optical Stokes\,I of SRSLS, and quantify the fields found. To this end, we apply the state-of-the-art line inversion code \texttt{SPINOR} \citep[][]{2005A&A...444..549F} to data of unprecedented quality and search for signs of Zeeman broadening, see \cite{2010A&A...522A..81A}[from hereon PaperI].
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Our analysis to infer average magnetic fields in optical spectra takes us close to the detectability limit of Zeeman broadening for SRSLS. Overall, we find that even the highest quality data do not yield clear and robust detections. The analysis is subject to a model dependence related to the inital choice of temperature configuration for the stellar surface, and may further be influenced by the approximative treatment of convection by analytical macroturbulence profiles (see PaperI). Also, blends (and their reproducability) gain in importance at the present level of accuracy and therefore constitute a limiting factor. However, there does remain a sensitivity to the Zeeman broadening signature, cf. the case of 59\,Vir, thanks to the simultaneous use of both high- and low-$g_{\rm{eff}}$ lines. It is therefore the numerical amount of magnetic flux detected that is questioned by the OC-TC-related model dependence, not the presence of a magnetic field on the star.
| 10 | 12 |
1012.2968
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1012
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1012.5854_arXiv.txt
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{ {The observed acceleration of the universe, explained through dark energy, could alternatively be explained through a modification of gravity that would also induce modifications in the evolution of cosmological perturbations.} {We use new weak lensing data from the COSMOS survey to test for deviations from General Relativity. The departure from GR is parametrized in a model-independent way that consistently parametrizes the two-point cosmic shear amplitude and growth.} {Using CMB priors, we perform a likelihood analysis. We find constraints on the amplitude of the signal that do not indicate a deviation from General Relativity.} }
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\label{sec1} The $\Lambda$CDM model is currently the best fit to the available cosmological data. In this framework, the period of accelerated expansion of the Universe is due to dark energy, a dominant component of the energy density with a negative equation of state. A possible alternative, worth to investigate observationally, is to replace the standard gravitational theory of General Relativity (GR) by a modification that admits self-accelerating solutions. Such gravity theory must converge to GR both at high redshift, when the expansion is not accelerated, and on small scales, since GR has consistently passed local tests of gravity (\cite{will}). The various classes of modified gravity and dark energy models (\cite{uzan}) cannot be distinguished based only on observations of the cosmological background. They are only distinguishable due to their different predictions for cosmological structure formation (\cite{zhang}). The evolution of cosmological background and perturbations have been studied in various classes of physically motivated modified gravity theories. Most prominent today are theories where the modifications arise from extra dimensions, with matter fields confined to a 4-dimensional brane embedded in a higher dimensional bulk, such as the Dvali-Gabadadze-Porrati (DGP) model (\cite{dvali}; \cite{deffayet}; \cite{lue}), and also the so-called f(R) theories (\cite{carroll}; \cite{songfr}). In the latter, the field equations are modified because the corresponding Einstein action no longer depends linearly on the Ricci scalar (R) but there are extra terms on R. The standard way of studying cosmological structure formation is to introduce scalar perturbations in the Friedmann-Robertson-Walker (FRW) spacetime metric. In the Newtonian or longitudinal gauge (\cite{mukhanov}), where by construction the metric remains diagonal, there are only two scalar perturbations: the scalar fields $\Psi$ (the Newtonian potential) and $\Phi$ (the spatial curvature potential). They are sourced by the cosmological fluid perturbations contained in the stress-energy tensor. Conservation of the stress-energy tensor imposes evolution equations for its components. In addition, the equations of the theory of gravitation, relating the metric components with the energy components, bring extra evolution and constraint equations. In modified gravity, structure formation has been studied with various parametric and non-parametric approaches (\cite{bertschinger}; \cite{silvestri}). Parametric approaches include direct parametrization of the linear growth (\cite{linder}; \cite{lindercahn}); the so-called parametrized post-Friedmann framework (PPF) (\cite{husawicki}) and several related phenomenological post-GR parametrizations, which give a model-independent description of deviations to GR in a way that is convenient for the testing of cosmological predictions (\cite{caldwell}; \cite{amendola}; \cite{daniel08}; \cite{daniel10}; \cite{pogosian}); and phenomenological applications to parametrizations of particular theories of modified gravity (\cite{acquaviva}; \cite{zhaoa}; \cite{songbd}). The various metric and energy perturbation variables in modified gravity scenarios may be measured by different probes (\cite{huterer}; \cite{jainz}; \cite{schmidt}; \cite{song}; \cite{guzik}; \cite{beynon}). In particular, weak lensing by large-scale structure (cosmic shear), which probes the sum $\Phi + \Psi$, is well suited for such studies (\cite{uzanb}; \cite{knox}). In cosmic shear measurements, the gravitational information is contained in the equations relating the measured matter density perturbations with the two potentials, or equivalently in one of those equations (a generalized Poisson equation) and an equation relating the two potentials (the anisotropy equation). A parametrization of those two relations is thus well suited for cosmic shear measurements, which will provide a degenerate estimation of those two quantities. This is the approach we follow in this work. We compute cosmic shear two-point correlation models in tomographic bins, using standard cosmological parameters plus two modified gravity parameters, and analyse them against the new COSMOS tomographic cosmic shear measurements of Schrabback \etal (2010). Other cosmic shear data have already been used in recent studies of modified gravity, using similar or different parametrizations (\cite{dore}; \cite{thomas}; \cite{daniel09}; \cite{bean}; \cite{zhaotom}; \cite{daniellinder}). The standard parameters are constrained by using WMAP7-year CMB data (\cite{komatsu}). Our goal is to study the impact of this phenomenological model in the degeneracy between the amplitude of the cosmic shear signal and the linear growth of structure. In other words, to investigate if, given the strong cosmological priors imposed by the WMAP7 results, there is room to compensate a possible alteration of the predicted redshift-dependent amplitude of the cosmic shear effect with a modification of the growth of matter perturbations, which, phenomenologically, would keep open the possibility of an alternate gravity theory as an origin for an effective dark energy behaviour. In the next section, we present the parametrization of modified gravity used. The new COSMOS dataset is described in Sect.~3. The statistical analysis and its results are presented and discussed in Sects.~4 and 5. We conclude in Sect.~6 with a reminder of the main assumptions made.
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\label{sec6} We used the new COSMOS cosmic shear data to test possible deviations from General Relativity. The departure from GR is parametrized in a phenomenological model-independent way extensively used in the literature. The two phenomenological parameters allow to derive two other parameters that consistently parametrize the modifications to two-point cosmic shear amplitude and growth, and are scale-independent and have a scaling redshift dependence with the $\Lambda$CDM background. A more realistic parametrization should take into account three regimes of modified gravity (\cite{husawicki}), allowing for compatibility with the background expansion on large scales and agreement with GR on sub-galactic scales. Hu \& Sawicki (2007) give a fitting formula for the non-linear matter power spectrum across the regimes. That formula was confirmed for particular classes of theories by perturbation theory (\cite{koyama}) and N-body simulations (\cite{oyaizu}). In particular for f(R), according to Koyama \etal (2009) it shows systematic deviations of at least 10\% from the formula of Smith \etal (2003), for scales smaller than $k\approx 0.8h\,{\rm M_{pc}}^{-1}$. COSMOS has an average source redshift of $z\approx 1.5$ and is mainly sensitive to lensing structures at around half comoving distance to the source. Therefore, this limiting scale corresponds to $\theta\approx 20\,{\rm arcmin}$ and the COSMOS range is affected. We do not apply this result to our model-independent study, but in general a systematic increase in the amplitudes of the models implies that cosmic shear becomes more sensitive to $\Sigma$, leading to a more precise determination of this parameter for the same data. Our results use WMAP priors obtained for $\Lambda$CDM models. This is consistent with using a $\Lambda$CDM background evolution. The CMB power spectrum is affected by the Integrated Sachs-Wolfe effect essentially at the recent epoch when the modified gravity effects become important and the potentials change significantly over time. This affects only the scales large enough compared with the photons crossing time. Since the effect is at low redshift those large physical scales correspond to very large angular scales. In this way, the peak structure is not changed and the parameters of modified gravity are not degenerate with the standard parameters at perturbation level, allowing to use $\Lambda$CDM priors. We did not find evidence for inconsistency with GR, in agreement with other cosmological tests in the linear regime (see \cite{jain} for a review). In particular, we do not find the evidence for deviation reported by Daniel \& Linder (2010) when using small-scale CFHTLS data. Our result using the new tomographic correlations of COSMOS corroborates the results obtained by Daniel \& Linder (2010) using a different COSMOS catalogue (\cite{massey07}). We cannot however make a direct comparison of the constraints, since we assume a different redshift-dependence and different priors. Our analysis relied on the tomographic amplitude of 2-pt cosmic shear, and did not allow to break the $\Sigma-\gamma$ degeneracy. The scale-dependence of the modifications, which is more distinctive on scales larger than those probed by COSMOS, should be explored with future wide high-precision weak lensing surveys. Such an analysis will have a better chance of reaching our original goal, i.e., to investigate if given strong cosmological priors there would be room to compensate a modification of the predicted amplitude of the cosmic shear signal with a modification of the growth of matter perturbations.
| 10 | 12 |
1012.5854
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1012
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1012.2190_arXiv.txt
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\noindent \noindent \noindent We study the structure of electromagnetic field of slowly rotating magnetized star in a Randall-Sundrum II type braneworld. The star is modeled as a sphere consisting of perfect highly magnetized fluid with infinite conductivity and frozen-in dipolar magnetic field. Maxwell's equations for the external magnetic field of the star in the braneworld are analytically solved in approximation of small distance from the surface of the star. We have also found numerical solution for the electric field outside the rotating magnetized neutron star in the braneworld in dependence on brane tension. The influence of brane tension on the electromagnetic energy losses of the rotating magnetized star is underlined. Obtained "brane" corrections are shown to be relevant and have non-negligible values. In comparison with astrophysical observations on pulsars spindown data they may provide an evidence for the brane tension and, thus, serve as a test for the braneworld model of the Universe.
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The study of magnetic and electric fields around the compact objects is an important task for several reasons. First is that we obtain information about such stars through their observable characteristics, which are closely connected with electromagnetic fields inside and outside the stars. Magnetic fields play an important role in the life history of majority astrophysical objects especially of compact relativistic stars which posses surface magnetic fields of $10^{12}G$ and $\sim 10^{14}G$ in the exceptional cases for magnetars. The strength of compact star's magnetic field is one of the main quantities determining their observability, for example as pulsars through magneto-dipolar radiation. Electric field surrounding the star determines energy losses from the star and therefore may be related with such observable parameters as period of pulsar and it's time derivation. The second reason is that we may test various theories of gravitation through the study of compact objects for which general relativity effects are especially strong. Considering different metrics of space-time one may investigate the effect of the different phenomena on evolution and behavior of stellar interior and exterior magnetic fields. Then these models can be checked through comparison of theoretical results with observational data. The third reason may be seen in influence of stellar magnetic and electric field on different physical phenomena around the star, such as gravitational lensing and motion of test particles. In the Newtonian framework the exterior electromagnetic fields of magnetized and rotating sphere are given in the classical paper of \citet{d55} and interior fields are studied by many authors, for example, in the paper of \citet{rt73}. In the general-relativistic approach the study of the magnetic field structure outside magnetized compact gravitational objects has been initiated by the pioneering work of \citet{go64} and have been further extended by number of authors (\citet{ac70}, \citet{p74}, \citet{ws83}, \citet{mh97}, \citet{mt92}, \citet{ram01a}, \citet{ram01b}, \citet{japan}), while in some papers (\citet{getal98}, \citet{pg97}, \citet{gpz00}, \citet{pgz00}, \citet{zr}) the work has been completed by considering magnetic fields interior relativistic star for the different models of stellar matter. General-relativistic treatment for the structure of external and internal stellar magnetic fields including numerical results has shown that the magnetic field is amplified by the monopolar part of gravitational field depending on the compactness of the relativistic star. We are interested in study of stellar electric and magnetic fields in frames of recently popular model of the braneworld first proposed in the work of \citet{RaSu99}. According to this model our four-dimensional space-time is just a slice of five-dimensional bulk and only gravity is the force which can freely propagate between our space-time and bulk while other fields are confined to four-dimensional Universe. The review of braneworld models is given in the work of \citet{maar04} and some cosmological and astrophysical implications of the braneworld theories may be found in works \citet{maar00}, \citet{cs01}, \citet{lan01}, \citet{hm03}, \citet{ger06}, \citet{kg08}, \citet{MaMu05}. For astrophysical interests, static and spherically symmetric exterior vacuum solutions of the braneworld models were initially proposed by \citet{dmp00} which have the mathematical form of the Reissner-Nordstr\"{0}m solution, in which a tidal Weyl parameter $Q^\ast$ plays the role of the electric charge squared of the general relativistic solution. It should be noted that besides this solution was pioneering and there are many different vacuum braneworld solutions at the moment, this solution still stays interesting and actual and, for example, was recently applied to the solar system tests in the paper of \citet{brhl10}. Observational possibilities of testing the braneworld black hole models at an astrophysical scale have been intensively discussed in the literature during the last several years, for example, through the gravitational lensing, the motion of test particles, and the classical tests of general relativity (perihelion precession, deflection of light, and the radar echo delay) in the Solar System (see \citet{lobo08}). In the paper of \citet{pkh08} the energy flux, the emission spectrum, and accretion efficiency from the accretion disks around several classes of static and rotating braneworld black holes have been obtained. The complete set of analytical solutions of the geodesic equation of massive test particles in higher dimensional spacetimes which can be applied to braneworld models is provided in the recent paper of \citet{Lam08}. The relativistic quantum interference effects in the spacetime of slowly rotating object in braneworld and phase shift effect of interfering particle in neutron interferometer have been studied in the recent paper of \citet{mht10}. The influence of the tidal charge onto profiled spectral lines generated by radiating tori orbiting in the vicinity of a rotating black hole has been studied in the paper of \citet{sstuch09}. Authors showed that with lowering the negative tidal charge of the black hole, the profiled line becomes flatter and wider, keeping their standard character with flux stronger at the blue edge of the profiled line. The role of the tidal charge in the orbital resonance model of quasiperiodic oscillations in black hole systems has been investigated in the paper of \citet{stuch09}. The influence of the tidal charge parameter of the braneworld models on some optical phenomena in rotating black hole spacetimes has been extensively studied in the paper of \citet{ssstuch09}. A braneworld corrections to the charged rotating black holes and to the perturbations in the electromagnetic potential around black holes are studied, for example, in works of \citet{ag05} and \citet{aa10}. Our preceding paper, \citet{af08}, was devoted to the stellar magnetic field configurations of relativistic stars in dependence on brane tension. The present paper extends the paper of \citet{af08} to the case of rotating relativistic star. In our paper we will consider rotating spherically symmetric star in the braneworld endowed with strong magnetic fields. We assume that the star has dipolar magnetic field and the field energy is not strong enough to affect the spacetime geometry, so we consider the effects of the gravitational field of the star in the braneworld on the magnetic and electric field structure without feedback. The motion of test particles near black holes immersed in an asymptotically uniform magnetic field and some gravity surrounding structure, which provides the magnetic field has been intensively studied in the paper of \citet{kon06}. The author has calculated the binding energy for spinning particles on circular orbits. The bound states of the massive scalar field around a rotating black hole immersed in the asymptotically uniform magnetic field are considered in the paper of \citet{kon07}. The uniform magnetic field in the background of a five dimensional black hole has been extensively studied in the work of \citet{alfr04}. In particular, authors presented exact expressions for two forms of an electromagnetic tensor and the electrostatic potential difference between the event horizon of a five dimensional black hole and the infinity. The paper is organized as follows. In section \ref{meq} we present a set of Maxwell's equations in the space-time of spherically symmetric rotating relativistic compact star in the braneworld. Section \ref{ss} is devoted for solutions of Maxwell's equations. In subsection \ref{toy} we consider the solution for "toy model" - monopolar structure of magnetic field of the star. This solution is not realistic but it can be used to obtain first estimates of the influence of brane tension on the electromagnetic field of the star. In subsection \ref{srst} we are looking for analytical solution of the Maxwell's equations for the exterior magnetic field of the star. We obtain approximate solution of the differential equation for magnetic field in the near vicinity of the surface of the star. In subsection \ref{electr} we get the differential equation for the electric field outside the star and solve them numerically. We show that both magnetic and electric fields will be essentially modified by five-dimensional gravity effects. In subsection \ref{application} we investigate the astrophysical application of obtained result, namely, calculate energy losses from the slowly rotating magnetized neutron star in the braneworld. The last section is devoted to the conclusions of the research done. Throughout, we use a space-like signature $(-,+,+,+)$ and a system of units in which $G = 1 = c$. Greek indices are taken to run from 0 to 3 and Latin indices from 1 to 3; covariant derivatives are denoted with a semi-colon and partial derivatives with a comma.
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In the present paper we considered modifications of electromagnetic field of a rotating magnetized neutron stars in the braneworld. We formulated Maxwell's equations for the case of slowly rotating magnetized compact star with non-zero brane tension and considered the particular case of monopolar magnetic field. Despite this configuration may not be considered as realistic it helped us to see the strong influence of brane tension on the electric field of the gravitating object. As the analytical solution is always more valuable for further calculations we attempted to find analytical solution for the dipolar magnetic field configuration. We have derived an approximate analytical expression for the magnetic field just near the surface of the star as a solution of II type hypergeometric equation. This region of the magnetosphere is very important because exactly in this region the processes of plasma generation responsible for the radio emission take place. We have got equations for the electric field of the slowly rotating magnetized neutron star on branes and solved one of them numerically for different values of brane tension. It is shown that the effect of brane tension on the electromagnetic field of the star is non-negligible (may have the order of tens percents of the initial value) and may help in future in testing the braneworld model. As an important application of the obtained results we have calculated energy losses of slowly rotating magnetized neutron star in the braneworld and found that the star with non-zero brane parameter will lose more energy than typical rotating neutron star in general relativity. The obtained dependence may be combined with the astrophysical data on pulsar period slowdowns and be useful in further investigations of the possible detection/estimation of the brane parameter.
| 10 | 12 |
1012.2190
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1012
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1012.2229_arXiv.txt
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My presentation was an overview of what we know about the Local Group of galaxies, primarily from optical imaging and spectroscopy. AGB stars are on the whole a very sparse and unrepresentative stellar population in most Local Group galaxies. However, more detailed studies of star formation histories and chemical evolution properties of populations, like Main Sequence dwarf stars and Red Giant Branch stars, allow a better understanding of the evolutionary context in which AGB stars can be observed. There are a variety of galaxy types in the Local Group which range from predominantly metal poor (e.g., Leo~A) to metal-rich (e.g., M~32). Dwarf galaxies are the most numerous type of galaxy in the Local Group, and provide the opportunity to study a relatively simple, typically metal-poor, environment that is likely similar to the conditions in the early history of all galaxies. Hopefully the range of star formation histories, peak star formation rates and metallicities will provide enough information to properly calibrate observations of AGB stars in more distant systems, and indeed in integrated spectra. Here I will summarise what we know about the star formation histories of nearby galaxies and their chemical evolution histories and then attempt to make a connection to their AGB star properties.
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Within the Local Universe galaxies can be studied in great detail star by star. The Colour-Magnitude Diagram synthesis analysis method is well established, at optical wavelengths, as the most accurate way to determine the detailed star formation history of galaxies going back to the earliest times \citep[e.g.,][]{Tolstoy09}. This approach has benefited enormously from the exceptional data sets that wide field CCD imagers on the ground and the Hubble Space Telescope can provide. Spectroscopic studies using large ground-based telescopes have allowed the determination of abundances and kinematics for significant samples of red giant branch (RGB) stars and also more massive O, B and A stars in several nearby galaxies \citep[e.g.,][and references therein]{Tolstoy09}. These studies have shown directly how properties can vary spatially and temporally, and how this information can give important constraints to theories of galaxy formation and evolution. Dwarf galaxies are commonly used as probes of a simple ``single cell'' star forming environment. They cover a range of mass and metallicity, and are considered to be representative of how galaxies in the early universe may have looked. A working definition of dwarf galaxies includes all galaxies that are fainter than M$_B \le -16$ (M$_V \le -17$) and more spatially extended than globular clusters \citep[e.g.,][]{Tammann94}, see Figures~\ref{bing}~\&~\ref{bel}. Although these limits were not physically motivated, they are broadly consistent with the limit of mass and concentration at which gas outflows are likely to start to significantly effect the baryonic mass of a galaxy. This includes a number of different types: early-type dwarf spheroidals (dSphs); late-type star-forming dwarf irregulars (dIs); the recently discovered very low surface brightness, ultra-faint, dwarfs (uFd); as well as centrally concentrated actively star-forming blue compact dwarf galaxies (BCDs). The newly discovered, even more extreme, so-called ultra-compact dwarfs (UCDs) are identified as dwarf galaxies from spectra but are of a similar compactness to globular clusters (see purple crosses in Figure~\ref{bel}). The dIs, BCDs, dSphs, late-type and spheroidal galaxies tend to overlap with each other in global properties in Figures~\ref{bing}~\&~\ref{bel}. These overlapping properties of early and late-type dwarfs have long been assumed to be convincing evidence that early-type dwarfs are late-type systems that have been stripped of, or otherwise used up, their gas \citep[e.g.,][]{Kormendy85}. Thus, like larger systems, the global properties of dwarf galaxies correlate closely with luminosity, half-light radius and surface brightness, over a large range. Dwarf galaxies thus allow us to study specific aspects of galaxy formation and evolution on a small scale. \begin{figure} \hskip 1.5cm \resizebox{0.85\columnwidth}{!}{\includegraphics{etolstoy-fig1.eps}} \vskip -0.7cm \caption{ The relationship between the structural properties (absolute magnitude, M$_V$ and central surface brightness $\mu_V$) for a range of different galaxy types. The dotted line is the classical maximum luminosity of the dwarf galaxy class, from Tammann (1994). Local Group galaxies are plotted as open pentagons, with the colour depending upon their gas content. The Sloan discovered ultra-faint systems as plotted as star symbols. Blue Compact Dwarf galaxies are squares, Ultra-compact systems as crosses and Galactic globular clusters as dots. See \citet{Tolstoy09} and \citet{Binggeli94} for more details. } \label{bing} % \end{figure}
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It is clear the that AGB stars can play a very significant role in the chemical evolution of a galaxy, especially a dwarf galaxy. A dwarf galaxy with an extended star formation history will likely be highly sensitive to the chemical enrichment created by the relatively slow and steady stellar winds from AGB stars. In small galaxies Supernovae may drive mass and metals entirely out of the galaxy, but stellar winds from AGB stars probably will not. The effect of AGB stars is likely to be dependent upon the time scale over which star formation occurred. The products of these stellar winds must be returned to the ISM on a time frame consistent with the subsequent star formation episodes in a galaxy to have an impact on the chemical evolution. From the lack of AGB stars in very metal poor systems it also seems likely that the [Fe/H] plays a role in the evolution of AGB star populations. It seems to be more difficult to produce metal poor AGB stars, and also to measure any effect in the abundance ratios that may come from them. In this review I have just touched upon the connections that can be made between the AGB star properties of nearby galaxies and their star formation histories and metallicities. These results are likely to be placed on much more quantitative basis in the coming years as more and more wide-field near-IR and optical imaging and spectroscopic surveys are carried out for both nearby and more distant galaxies. It is clear that to sort out the complex and intertwined effects of star formation, stellar winds, supernovae explosions and their effect on the ISM we need to use information from a variety of sources that are sensitive to different time scales, and different physical processes. This means that we need to combine information from optical imaging (SFHs) and spectroscopy (abundances) with IR imaging and spectroscopy to get the full story.
| 10 | 12 |
1012.2229
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1012
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1012.4125_arXiv.txt
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We present a new approach to the sky subtraction for long-slit spectra suitable for low-surface brightness objects based on the controlled reconstruction of the night sky spectrum in the Fourier space using twilight or arc-line frames as references. It can be easily adopted for FLAMINGOS-type multi-slit data. Compared to existing sky subtraction algorithms, our technique is taking into account variations of the spectral line spread along the slit thus qualitatively improving the sky subtraction quality for extended targets. As an example, we show how the stellar metallicity and stellar velocity dispersion profiles in the outer disc of the spiral galaxy NGC5440 are affected by the sky subtraction quality. Our technique is used in the survey of early-type galaxies carried out at the Russian 6-m telescope, and it strongly increases the scientific potential of large amounts of long-slit data for nearby galaxies available in major data archives.
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Low-surface brightness ($\mu_B>23$ mag/arcsec$^2$) outer regions of galaxies contain crucially important information for understanding the properties of their extended discs and dark matter haloes. Brightness profiles of dwarf early-type galaxies whose mean surface brightness is correlated with the luminosity, can be entirely in the low-surface brightness regime. Analysis of absorption line spectra at such surface brightness levels is often hampered by systematic errors of the sky subtraction that sometimes may lead to wrong astrophysical conclusions. Therefore, in order to analyse deep spectral data, it is important to improve the sky subtraction technique. Here we present a new approach to the sky subtraction for long-slit spectra based on the controlled reconstruction of the night sky spectrum in the Fourier space using twilight or arc line frames as references. \articlefigurethree{lsf_Ypos_variation_4877.ps}{lsf_Ypos_variation_5155.ps}{lsf_Ypos_variation_5448.ps} {fig1_lsf_shape}{An example of the LSF shape of the SCORPIO reconstructed from the twilight frame at different wavelength and slit positions. We used the Gauss-Hermite LSF representation. One can see that the profile asymmetry increases towards the outer slit regions. There is also a notable change of the overall spectral resolution from blue to red.} Due to optical distortions, the shape of the spectral line spread function (LSF) in a long-slit spectrograph varies along the wavelength range as well as along the slit. In Fig.~\ref{fig1_lsf_shape}, we provide an example of the LSF shape of the SCORPIO \citep{AM05} universal spectrograph at the Russian 6-m telescope reconstructed from the twilight frame (i.e. the Solar spectrum). The LSF is slightly asymmetrical and cannot be described by the Gaussian function, a usual parametrization in most data reduction packages. Here we use the Gauss-Hermite representation \citep{vdMF93} up-to the 4th order moment that allows one to describe first-order differences of the line profile from the Gaussian shape. These LSF variations affect the night sky spectrum which is subtracted from science frames during the data reduction. On the Fig.~\ref{fig2_objspec} we show a reduced long-slit spectrum of the spiral galaxy NGC 5440 before the sky subtraction step. \articlefigure{obj_sky_region_for_substruction_small.ps}{fig2_objspec}{A reduced long-slit spectrum of the spiral galaxy NGC~5440 before the sky subtraction step. Yellow areas denote a region of the frame used to construct the night sky spectrum used for the sky subtraction. Taking into account the profile variation shown in Fig.~\ref{fig1_lsf_shape}, it is clear that the intrinsic LSF shape in these region will differ from that in regions of the galaxy closed to the slit centre.}
| 10 | 12 |
1012.4125
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1012
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1012.3019_arXiv.txt
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In spite of recent detections of magnetic fields in a number of $\beta$\,Cephei and slowly pulsating B (SPB) stars, their impact on stellar rotation, pulsations, and element diffusion is not sufficiently studied yet. The reason for this is the lack of knowledge of rotation periods, the magnetic field strength distribution and temporal variability, and the field geometry. New longitudinal field measurements of four $\beta$\,Cephei and candidate $\beta$\,Cephei stars, and two SPB stars were acquired with FORS\,2 at the VLT. These measurements allowed us to carry out a search for rotation periods and to constrain the magnetic field geometry for four stars in our sample.
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\begin{table} \centering \caption{ The observed $\beta$\,Cephei and SPB stars. } \label{tab:table1a} \begin{tabular}{rccc} \hline \hline \multicolumn{1}{c}{HD} & \multicolumn{1}{c}{Other} & \multicolumn{1}{c}{Spectral} & \multicolumn{1}{c}{Comments} \\ \multicolumn{1}{c}{} & \multicolumn{1}{c}{Identifier} & \multicolumn{1}{c}{Type} & \multicolumn{1}{c}{} \\ \hline 24587 & 33\,Eri & B5V & SPB, SB1\\ 46328 & $\xi^1$\,CMa & B1III & $\beta$\,Cep \\ 50707 & 15\,CMa & B1Ib & $\beta$\,Cep \\ $\ast$ 74575 & $\alpha$\,Pyx & B1.5III & $\beta$\,Cep \\ 74560 & HY\,Vel & B3IV & SPB, SB1\\ $\ast$ 136504 & $\epsilon$\,Lup & B2IV-V & $\beta$\,Cep, SB \\ \hline \end{tabular} \end{table} For several years, a magnetic field survey of main-sequence pulsating B-type stars, namely the slowly pulsating B (SPB) stars and $\beta$\,Cephei stars, has been undertaken by our team with FORS\,1 in spectropolarimetric mode at the VLT, allowing us to detect in four $\beta$ Cephei stars and in 16 slowly pulsating B stars, for the first time, longitudinal magnetic fields of the order of a few hundred Gauss \citep{Hubrig2006,Hubrig2009}. $\beta$~Cephei variables have spectral types B0--B2 and pulsate in low-order pressure and gravity modes with periods between 2 and 6\,hours. Slowly pulsating B (SPB) stars are mid-B type (B3--B9) objects pulsating in high-order gravity modes with periods in the range of 0.5--3\,days. Pulsating stars are currently considered as promising targets for asteroseismic analysis (e.g.\ \citealt{ShibahashiAerts2000}), which requires as input the observed parameters of the magnetic field topology. Early magnetic field searches of $\beta$\,Cephei stars were mostly unsuccessful due to low precision (see \citealt{Babcock1958,RudyKemp1978}). Before we started our systematic search for magnetic fields in pulsating B-type stars, a weak magnetic field was detected in two $\beta$\,Cephei stars, in the prototype of the class, $\beta$\,Cep itself, by \citet{Henrichs2000} and in V2052\,Oph by \citet{Neiner2003a}. The first detection of a weak magnetic field in the SPB star $\zeta$\,Cas was reported by \citet{Neiner2003b}. The detected magnetic objects for which we gathered several magnetic field measurements showed a field that varies in time, but no periodicity could be derived yet due to the limited amount of VLT observing time. Among these targets with a detected magnetic field, we selected two slowly pulsating B stars, two $\beta$ Cephei stars, and two candidate $\beta$ Cephei stars with suitable coordinates, for successive VLT multi-epoch magnetic measurements. The list of the selected targets is presented in Table~\ref{tab:table1a}. In the four columns we list the HD number, another identifier, the spectral type retrieved from the SIMBAD database, the pulsating type, and membership in a spectroscopic binary system. An asterisk in front of the HD number denotes candidate $\beta$\,Cephei stars (cf.\ \citealt{StankovHandler2005}). This most recent study aimed at the determination of magnetic field properties for these stars, such as field strength, field geometry, and time variability. Here we present the results of 62 new magnetic field measurements of the six selected stars and discuss the obtained results on their rotation periods and magnetic field geometry.
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Using FORS\,1/2 and SOFIN longitudinal magnetic field measurements collected in our recent studies, we were able to determine rotation periods and constrain the field geometry of two $\beta$\,Cephei stars, one candidate $\beta$\,Cephei star, and one SPB star. The dipole model provides a satisfactory fit to the data and among the very few presently known magnetic $\beta$\,Cephei stars, $\xi^1$\,CMa and $\alpha$\,Pyx possess the largest magnetic fields, with a dipole strength of several kG. \citet{Briquet2007} discussed the position of SPB and chemically peculiar Bp stars in the H-R diagram. indicating that the group of Bp stars is significantly younger than the group of SPB stars. A similar conclusion was deduced by \citet{Hubrig2007}, who studied the evolution of magnetic fields in Ap and Bp stars with definitely determined magnetic field geometries across the main sequence. The vast majority of Bp stars exhibits a smooth, single-wave variation of the longitudinal magnetic field during the stellar rotation cycle. These observations are considered as evidence for a dominant dipolar contribution to the magnetic field topology. It is of interest that the study of \citet{Hubrig2007} indicates the prevalence of larger obliquities $\beta$, namely $\beta$$>$60$^{\circ}$, in more massive stars. The magnetic field models for the three $\beta$\,Cephei stars and the one SPB star presented in this work confirm this trend. The insufficient knowledge of the strength, geometry, and time variability of magnetic fields in hot pulsating stars prevented until now important theoretical studies on the impact of magnetic fields on stellar rotation, pulsations, and element diffusion. Although it is expected that the magnetic field can distort the frequency patterns (e.g.\ \citealt{Hasan2005}), such a perturbation is not yet detected in hot pulsating stars. Splitting of non-radial pulsation modes was observed for 15\,CMa \citep{Shobbrook2006}, but the identification of these modes is still pending. The magnetic $\beta$\,Cephei star sample indicates that they all share common properties: they are N-rich targets (e.g., \citealt{Morel2008}) and, as discussed by \citet{Hubrig2009}, their pulsations are dominated by a non-linear dominant radial mode (see also \citealt{Saesen2006} for $\xi^1$\,CMa). The presence of a magnetic field might consequently play an important role to explain such a distinct behaviour of magnetic $\beta$\,Cephei stars. More precisely, chemical abundance anomalies are commonly believed to be due to radiatively-driven microscopic diffusion in stars rotating sufficiently slowly to allow such a process to be effective. However, we need an additional clue to account for the fact that both normal and nitrogen-enriched slowly rotating stars are observed. The presence of a magnetic field is a very plausible explanation, as it can add to the stability of the atmosphere, allowing diffusion processes to occur \citep{Michaud1970}. On the other hand, among the studied stars, apart from the star 15\,CMa with rather low $v_{\rm eq} = 40\pm 6$\,km\,s$^{-1}$, the other three magnetic pulsating stars rotate much faster up to $v_{\rm eq} = 165\pm 21$\,km\,s$^{-1}$ for $\xi^1$\,CMa with the strongest magnetic field, indicating that these stars are not truly slowly rotating stars, but seen close to pole-on. Obviously, the topic of mixing signatures is not understood theoretically yet and more computational work as well as future additional observational validation of our results are needed to understand the link between the presence of a magnetic field, rotation, pulsating characteristics, and abundance peculiarities.
| 10 | 12 |
1012.3019
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1012
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1012.1470_arXiv.txt
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% {} {We study the influence of the environment on the evolution of galaxies by investigating the luminosity function (LF) of galaxies of different morphological types and colours at different environmental density levels.} {We construct the LFs separately for galaxies of different morphology (spiral and elliptical) and of different colours (red and blue) using data from the Sloan Digital Sky Survey (SDSS), correcting the luminosities for the intrinsic absorption. We use the global luminosity density field to define different environments, and analyse the environmental dependence of galaxy morphology and colour. The smoothed bootstrap method is used to calculate confidence regions of the derived luminosity functions.} {We find a strong environmental dependency for the LF of elliptical galaxies. The LF of spiral galaxies is almost environment independent, suggesting that spiral galaxy formation mechanisms are similar in different environments. Absorption by the intrinsic dust influences the bright-end of the LF of spiral galaxies. After attenuation correction, the brightest spiral galaxies are still about 0.5\,mag less luminous than the brightest elliptical galaxies, except in the least dense environment, where spiral galaxies dominate the LF at every luminosity. Despite the extent of the SDSS survey, the influence of single rich superclusters is present in the galactic LF of the densest environment.} {}
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% Understanding the formation and evolution of galaxies is one of the biggest challenges of observational cosmology. The luminosity function (LF) is in this respect one of the most fundamental of all cosmological observables, helping us to describe the global properties of galaxy populations and to study the evolution of galaxies. The dependence of the LF on cosmic time, galaxy type, and environmental properties gives insight into the physical processes that govern the assembly of the stellar content of galaxies. The first determinations of the galaxy LF were made several decades ago \citep{Kiang:61,Christensen:75,Kirshner:79}; in the following studies, the number of galaxies used to calculate the LF has increased continuously \citep{Tully:88,Efstathiou:88,Loveday:92}. The Las Campanas Redshift Survey measured the general LF of galaxies with a good accuracy \citep{Lin:96,Bromley:98,Christlein:00}. Our current understanding of the general LF owes much to the 2dFGRS \citep{Norberg:02} and SDSS surveys \citep{Blanton:03,Montero-Dorta:09}. These new samples of galaxies make it possible to study the dependence of the LF on a large number of different galaxy properties, as galaxy morphology, colours, star formation rate, local and global density environment, etc. The morphology of a given galaxy is a reflection of its merger history. In studies of the LF, galaxy morphology has been determined either by its colours \citep{Yang:09}, spectra \citep{Folkes:99,Madgwick:02,de-Lapparent:03} or photometric profile \citep{Bell:03,Driver:07a}; some studies use artificial neural networks for morphological classification \citep{Ball:06}. The most accurate, but by far the most time consuming approach is to use visual classification \citep{Marzke:94a,Marzke:98,Kochanek:01,Cuesta-Bolao:03,Nakamura:03}. For the SDSS survey, visual classification has become possible thanks to the Galaxy Zoo project \citep{Lintott:08} that will help us to study the morphology and the LF in detail in the future. Recently \citet{Huertas-Company:11} have published an automated morphological classification based on the Galaxy Zoo data. In all these studies, the classification of early-type and late-type galaxies is different, but all studies agree that later-type galaxies have a fainter characteristic magnitude and a steeper faint-end slope of the LF. The biggest differences in previous studies are found at the faint-end of the LF, where classification is less certain than for brighter galaxies. To understand how galaxies form, we also need to understand where galaxies are located; it is essential to study the LF dependence on the environment. It is well known from the halo occupation distribution models that the local environment is crucial for the galaxy distribution \citep[e.g.][]{Zandivarez:06,Park:07}: luminous galaxies tend to occupy high mass haloes and low luminosity galaxies reside mainly in low mass haloes. This motivates the study of the LF in galaxy groups \citep{Xia:06,Zandivarez:06,Hansen:09,Yang:09}. A likewise important, but not so well understood factor is the global environment where the galaxy is located -- its place in the supercluster-void network. In \citet{Tempel:09} we have found that the global environment has an important role in determining galaxy properties. Some studies have been dedicated only to special regions: e.g., \citet{Mercurio:06} investigate the Shapley supercluster. The dependence on the environment has been also studied numerically \citep{Mo:04} and using semi-analytical models \citep{Benson:03a,Khochfar:07}. These semi-analytical models allow us to study morphological evolution: how the morphology of a galaxy changes in time. To compare these models with the real Universe, we need to know the observed LF in detail. The influence of the global environment on the LF has been investigated by \citet{Hoyle:05}, using the SDSS data, and by \citet{Croton:05}, using the 2dFGRS data. These results show strong environmental trends: galaxies in higher density regions tend to be redder, of earlier type, have a lower star formation rate, and are more strongly clustered. Some of these trends can be explained with the well known morphology-density relation \citep{Einasto:74,Dressler:80,Postman:84} and luminosity-density relation \citep{Hamilton:88}. It is less well known how far these trends extend when moving toward extreme environments, into deep voids or superclusters. Recent studies have shown that dust plays an important role in galaxy evolution and it may significantly influence the luminosities and colours of galaxies \citep{Pierini:04,Tuffs:04,Driver:07,Rocha:08,Tempel:10}, especially for late-type galaxies. Thus, in order to study intrinsic properties of galaxies, it is necessary to take dust extinction into account. Using the SDSS data, \citet{Shao:07} have studied the influence of dust on the LF. In general, dust is important for late-type spiral galaxies; nearly edge-on galaxies are most affected. In the present paper we use the SDSS data to study the LF in different global environments and for different types of galaxies, taking the effect of dust attenuation into account. The LF dependency on group properties will be analysed in a forthcoming paper. Throughout this paper we assume a Friedmann-Robertson-Walker cosmological model with the total matter density $\Omega_\mathrm{m}=0.27$, dark energy density $\Omega_\Lambda=0.73$, and the Hubble constant $H_0=100\,h\,\mathrm{km\,s^{-1}Mpc^{-1}}$. Magnitudes are quoted in the AB~system.
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% We have used the SDSS data to derive the LF of spiral galaxies, elliptical galaxies, and all galaxies together in various environments. We have taken special care to correct the galaxy luminosities for the intrinsic attenuation. The principal results of our study are the following: \begin{itemize} \item The LF of elliptical galaxies depends strongly on the environment; this suggests that global environmental density is an important driving force (via merging history) of elliptical galaxy formation. Density environment is more important for red elliptical galaxies than for blue elliptical galaxies. \item The evolution of spiral galaxies (the LF of spiral galaxies) is almost independent of environment, especially for blue and red spirals separately, showing that spiral galaxy formation has to be similar regardless of the surrounding global density. \item The highest global density regions (superclusters) are significantly different from other regions, as stated also by \citet{Tempel:09}. Notably more faint galaxies are found in high density regions than in other environments. \item The brightest galaxies are absent from the void regions. After correcting for the intrinsic absorption, spiral galaxies dominate the LF of void regions at every luminosity. \item The faint-end of the LF is determined by spiral galaxies and the bright end by elliptical galaxies. The faint end includes mostly blue galaxies and the bright end mostly red galaxies. \item Detailed studies of LFs require galaxy luminosities to be corrected for the intrinsic absorption by dust. Dust absorption affects mostly the bright-end of the LF. For the full LF, including all galaxies, the characteristic luminosity increases after attenuation correction. The faint-end slope of the LF is practically independent on dust attenuation. \end{itemize} A comparison of these results with predictions of numerical simulations and/or semianalytical models would provide stringent constraints on the driving factors of the formation and evolution of galaxies in dark matter haloes.
| 10 | 12 |
1012.1470
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1012
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1012.3475_arXiv.txt
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In a planetary system with two or more well-spaced, eccentric, inclined planets, secular interactions among these planets may lead to chaos. The innermost planet may gradually become increasingly eccentric and/or inclined, as a result of the secular degrees of freedom drifting towards equipartition of AMD (angular momentum deficit). This ``secular chaos'' is known to be responsible for the eventual destabilization of Mercury in our own Solar System. Here we focus on systems with three giant planets. We characterize the secular chaos and demonstrate the criterion for it to occur, but leave a detailed understanding of secular chaos to a companion paper \citep{yoram}. After an extended period of eccentricity diffusion, the inner planet's pericentre can approach the star to within a few stellar radii. Strong tidal interactions and ensuing tidal dissipation extracts orbital energy from the planet and pulls it inward, creating a hot Jupiter. In contrast to other proposed channels for the production of hot Jupiters, such a scenario (which we term ``secular migration'') provides an explanation for a range of observations: the pile-up of hot Jupiters at 3-day orbital periods, the fact that hot Jupiters are in general less massive than other RV planets, that they may have misaligned inclinations with respect to stellar spin, and that they have few easily detectable companions (but may have giant companions in distant orbits). Secular migration can also explain close-in planets as low in mass as Neptune; { and an aborted secular migration can explain the ``warm Jupiters'' at intermediate distances. In addition,} the frequency of hot Jupiters formed via secular migration increases with stellar age. We further suggest that secular chaos may be responsible for the observed eccentricities of giant planets at larger distances, and that these planets could exhibit significant spin-orbit misalignment.
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\label{sec:intro} While around $10\%$ of sun-like stars surveyed harbor Jovian-mass planets, only $\sim 1\%$ are orbited by so-called hot-Jupiters with periods short-ward of $\sim 10 $ days \citep[see reviews by][]{Marcy,Udry}. There appears to be a pile-up of hot Jupiters around $3$ day orbital periods. This excess is genuine and has been confirmed by both radial velocity and transit surveys \citep{Gaudi,Butler,Cumming,Fressin}. Outward of hot Jupiters, there appears a deficit of gas giants with periods of $10$ to $100$ days \citep[the ``period valley;''][]{UdryMayorSantos,Wittenmyer}. According to current theories of planet formation, hot Jupiters could not have formed in situ, given the large stellar tidal field, high gas temperature, and low disk mass to be found so close to the star. Instead, hot Jupiters most likely formed beyond a few AU and then migrated inward. Candidate migration scenarios that have been proposed include protoplanetary disks, Kozai migration by binary or planetary companions, and scattering with other planets in the system. While each of these mechanisms may have contributed to the hot Jupiter population to some degree, the question remains as to which is the dominant one. The dominant mechanism has to explain a variety of observed correlations. In \S \ref{sec:critical}, we review some of these correlations and provide a critical assessment of the above three mechanisms. In this work, we propose a fourth channel for producing hot Jupiters, namely planet migration by secular chaos. Secular chaos may arise in planetary systems that are well-spaced and are dominated by long-range secular interactions. A system of two non-coplanar planets can be chaotic, but only if their eccentricities and inclinations are of order unity \citep{Libert,Migas09,Smadar}. So in this contribution we focus on systems with three planets. The criterion for secular chaos is less stringent, and the character of secular chaos is more diffusive, differing from that of the two planet case. This diffusive type of secular chaos promotes energy equipartition between different secular degrees of freedom. The physics behind secular chaos is analyzed in detail in a companion paper \citep{yoram}, where we show that Mercury, the innermost planet in our Solar system, experiences a similar type of secular chaos. Mercury may consequently be removed from the Solar system \citep{Laskar08,Batygin,LaskarGastineau}. Secular chaos tends to removes angular momentum in the inner most planet gradually, causing its pericenter to approach the star. Tidal dissipation may then remove orbital energy from this planet, turning it into a hot Jupiter. Hot Saturns or hot Neptunes may also be produced similarly. Such a migration mechanism, which we term ``secular migration,'' can reproduce a range of observations. It also predicts that in systems with hot planets, there are other giant planets roaming at larger distances.
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\label{sec:conclusion} Hot Jupiters, while representing only a small fraction of all known extra-solar planets, demand special attention. They are most at odds with planet formation theory; they are detected disproportionately in radial velocity and transit surveys; and they are most accessible to characterization. Their rarity may indicate that their formation requires extreme circumstances. However, they may teach us much about the general conditions of planetary systems. Hot Jupiters are piled up around 3-day orbital periods, with rapid cut-offs both inward and outward of this distance; they tend to be less massive than more distant planets; many of them have orbits that are misaligned relative to the stellar spin; and they are remarkably anti-social in having few detected companions. In this work, we show that most of these characteristics can be explained if hot Jupiters are produced by secular chaos. These planets, originally located at $\gtrsim 1$ AU, acquire AMD from planets that are further out in the system. The outer planets can be mildly eccentric and/or inclined. But the same AMD produces much greater eccentricities and inclinations when it is transported to an inner planet, especially if the inner planet is less massive than the outer ones. The extreme high eccentricity allows the inner planet to reach inward of a few stellar radii and be tidally ensnared by the central star into a hot Jupiter.\footnote{Kozai migration is another process that may give rise to large eccentricity to Jovian planets. Such planets may be similarly captured into hot Jupiters. Since Kozai oscillation is also a secular forcing, we propose to name both secular chaos and Kozai migration generically as ``Secular Migration''.} We find that the criterion for hot Jupiter production is a sufficient amount of AMD. Only 5 hot Jupiters have known planetary companions, and these companions are typically eccentric and may have contributed to secular chaos. We note that the hot Jupiter (at $0.06$ AU) in the Ups And system, where two {other} massive planets orbit at $0.8$ and $2.5$ AU, with eccentricities $0.2$ and $0.3$ respectively, may well be produced by the secular chaos presented here. This possibility is further boosted by the recent finding that the outer two planets' {inclinations} are misaligned by $\sim 30 \deg$ from each other \citep{McArthur}. Similarly, the retrograde hot Jupiter in WASP-8b \citep{Queloz} is in a stellar system including an M-star companion at $\sim 600$ AU. Moreover, the radial velocity trend indicates a companion at distance $> 1 AU$ that is more massive than $2 M_J$. This could also be a hot Jupiter produced by secular chaos, with the M-star acting as the third planet. Hot Neptunes may also be formed via secular migration. However, hot Earths are different . It is not that Earth-like planets could not undergo secular chaos, but that once they do, they cannot be stalled at a safe enough distance to avoid being swallowed by the central star. Secular chaos has also been found to be responsible for instability in the inner Solar System \citep{Laskar08,yoram}. We speculate that secular chaos may be a frequent phenomenon in planetary systems. It may help to excite inner planets to higher eccentricities or inclinations. If these planets are removed, the remaining planetary systems may be stabilized for a time comparable to the system age. The success of this theory depends on two unknown factors. One is the amount of initial AMD in the system. The other is the typical configuration of planetary systems when emerging out of the protoplanetary disk. It also needs to be demonstrated solidly that secular chaos can lead to a large fraction of retrograde hot Jupiters. Observationally, if not only hot planets, but also warm or cold planets can be shown to have significant orbital inclinations relative to the spin of their host stars, this would boost the case for the ubiquity of secular chaos. Future Rossiter-McLaughlin measurements should be extended to transiting planets at large distances.
| 10 | 12 |
1012.3475
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1012
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1012.0277_arXiv.txt
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In the past two decades, secular evolution has emerged as an important new paradigm for the formation and evolution of the Hubble sequence of galaxies. A new dynamical mechanism was identified through which density waves in galaxies, in the forms of nonlinear and global spiral and bar {\em modes}, induce important collective dissipation effects previously unknown in traditional studies. These effects lead to the evolution of the {\it basic state} of the galactic disk, consistent with the gradual transformation of a typical galaxy's morphological type from a late to an early Hubble type. In this paper, we review the theoretical framework and highlight our recent result which showed that there are significant qualitative and quantitative differences between the secular evolution rates predicted by the new theory compared with those predicted by the classical approach of Lynden-Bell and Kalnajs. These differences are the outward manifestation of the dominant role played by collisionless shocks in disk galaxies hosting quasi-stationary, extremely non-linear density-wave modes.
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The possibility that galaxy morphologies can transform significantly over their lifetime, not only through violent episodes such as merger or satellite accretion, but also through slow but steady internal secular dynamical processes, is a notion that is gaining acceptance in the recent decades. In the past, the work on secular evolution has been focused on gas accretion in barred galaxies and the growth of pseudo bulges (Kormendy \& Kennicutt 2004 and the references therein). This is partly due to the long-held notion that gas is the only mass component capable of dissipation, and the stellar component is adiabatic and generally does not lose or gain energy and angular momentum as they orbit around the center of a galaxy. The first indication that there is the possibility for significant {\it stellar mass redistribution} in galaxies originates from the seminal work of Lynden-Bell and Kalnajs (1972, hereafter LBK), who showed that a trailing spiral density wave possesses a gravitational torque that over time can transport angular momentum outward. LBK were interested in the angular momentum transport phenomenon because they were seeking a generating mechanism for the spiral density waves, thought to be short-lived wave trains constantly being amplified out of noise and subsequently absorbed at the inner Lindblad resonance. Since the density wave is considered to possess negative energy and angular momentum inside corotation relative to the basic state (i.e. the axisymmetric disk), an outward angular momentum transport would encourage the spontaneous growth of the wave trains. LBK at that time was not interested in the secular morphological evolution of the basic state of the disks. In fact, in the same paper, they showed that for WKBJ (tightly wrapped) waves, the long-term energy and angular momentum exchange between the wave and the basic state is zero away from the wave-particle resonances. This is possible in the presence of the outward angular momentum transport by gravitational torque couple because they showed that there is a second type of torque couple, the so-called advective torque couple (due to lorry transport), that opposes the gravitational torque couple, and the sum of the two types of torque couples is a constant independent of the galactic radii. The total torque couple, which is equal to the rate of total radial angular momentum flux, is thus a constant during the outward angular momentum transport, and there is no interaction of the wave and basic state except at the wave-particle resonances (i.e., they thought the wave picks up angular momentum from the basic state at the inner Lindblad resonance and dumps it at the outer Lindblad resonance, and en route of this radial transport the total angular momentum flux remains constant). Zhang (1996,1998,1999) showed that the classical theory of LBK ignored an important collective dissipation process present in the gravitational N-body disks possessing {\em self-organized, or spontaneously-formed, density wave modes}. This process is mediated by collisionless shocks at the density wave crest, which breaks the adiabaticity or the conservation of the Jacobi integral condition -- a condition which is shown to be valid only for a {\em passive} orbit under an {\em applied} spiral or bar potential, and is now shown not to be obeyed by orbits undergoing collective dissipation. The overall manifestation of the collective dissipation process is an azimuthal phase shift between the potential and the density distribution of the density wave pattern, and for a self-sustained mode this phase shift is positive inside corotation, and negative outside. The presence of the phase shift means that for every annulus of the galaxy, there is a secular torque applied by the density wave on the disk matter in the annulus, and the associated energy and angular momentum exchange between the wave and the basic state of the disk. As a result the disk matter inside corotation (both stars and gas) loses energy and angular momentum to the wave, and spirals inward, and the disk matter outside corotation gains energy and angular momentum from the wave and spirals outward. This energy and angular momentum exchange between the wave and the basic state of the disk thus becomes the ultimate driving mechanism for the secular evolution of the mass distribution of the basic state of galaxy disks. The energy and angular momentum received by the wave from the basic state, incidentally, serve as a damping mechanism for the spontaneously growing unstable mode, allowing it to achieve quasi-steady state at sufficiently nonlinear amplitude. In Zhang (1998), a set of analytical expressions for the secular mass accretion/excretion rate was derived, and was confirmed quantitatively in the N-body simulations presented in the same paper. However, due to the 2D nature of these simulations, where the bulge and halo were assumed to be spherical and inert, the simulated wave has an average density contrast of 20\% and potential contrast of 5\%, both much lower than the average observed density wave contrast in physical galaxies, so the simulated disk did not evolve a lot (Zhang 1999), despite the fact that these low evolution rates conform exactly to the analytical formula's prediction for the corresponding wave amplitude (Zhang 1998). Zhang \& Buta (2007) and Buta \& Zhang (2009) used near-frared images of observed galaxies to derive the radial distribution of the azimuthal potential-density phase shifts, and to use the positive-to-negative zero crossings of the phase shift curve to determine the corotation radii (CRs) for galaxies possessing spontaneously-formed density wave modes. This approach works because the alternating positive and negative humps of phase shift distribution lead to the correct sense of energy and angular momentum exchange between the wave mode and the disk matter to encourage the spontaneous emergence of the mode. In Zhang \& Buta (2007) and Buta \& Zhang (2009) we have found good correspondence between the predicted CRs using the potential-density phase shift approach with the resonance features present in galaxy images, and also with results from other reliable CR determination methods within the range of validity of these methods. Beside CR determination, an initial test case for mass flow rate calculation, for galaxy NGC 1530, was also carried out in Zhang \& Buta (2007) using the same volume-torque-integration/potential-density phase shift approach, and there we found that since this galaxy has exceptionally large surface density and density-wave arm-to-interarm contrast, mass flow rate more than 100 solar mass per year were obtained across much of the galactic radii for this galaxy. This level of mass flow rate is more than sufficient to transform the Hubble type of a late type galaxy to an early type within a Hubble time. Other galaxies we have tested have significantly lower mass flow rates but still are sufficient to lead to significant mass redistribution over a Hubble time. Recently, we have applied the potential-density phase shift/volume-torque method to a larger sample of galaxies in order to estimate their mass flow rates. In related earlier works, Gnedin et al. (1995) and Foyle et al. (2010) have applied LBK type gravitational torque integral to the calculation of angular momentum redistribution rate in a number of galaxies. In our own studies, we found that such earlier work using the gravitational torque couple alone had significantly under-estimated the (implied) total mass flow rates in these galaxies. As it turns out, the advective torque couple, which opposes the gravitational torque couple in the LBK classical theory, becomes to have the same sense of angular momentum transport direction for spontaneous-formed density wave modes at the nonlinear regime. Furthermore, at the extremely non-linear wave amplitudes usually found for observed galaxies, the contribution of advective torques due to the collisionless shocks far exceeds the contribution from gravitational torques, and becomes the dominant driver for the secular evolution of galaxy mass redistribution. The sum of both types of (surface) torque couples turns out to be equal to the (integral of the) volume-type torque we used in this work, which is first proved in Zhang (1998, 1999). The past calculations of gas mass accretion near the central region of galaxies (e.g. Haan et al. 2009) are likely to have significantly underestimated the gas mass flow rate for the same reason. Our current work also has other important implications for the fundamental questions of galactic dynamics. For example, on the modal versus transient nature of the density wave patterns in galaxies (see, e.g. Sellwood 2010 and the references therein). The bell-shaped total angular momentum flux or torque coupling integral, which is equivalent to the two-humped phase shift or volume torque distribution, that we have found to be overwhelmingly present in observed galaxies, has no explanation in the classical LBK theory of transient waves, which predicts constant total angular momentum flux for a wave train between the inner and outer Lindblad resonances (LBK; Binney \& Tremaine 2008), but is a natural consequence of the spontaneously-formed intrinsic modes of disk galaxies, as first demonstrated in Zhang (1998). Also, the result of CR determination for the majority of the more than 150 galaxies analyzed in Buta \& Zhang (2009) using the potential-density phase shift method also supports the modal view, since for transient waves one should not be able to use the Poisson equation alone to predict a partially-kinematic quantity such as the corotation radius. For the successful application of the potential-density phase shift method, the Poisson equation and the equations of motion must have achieved a good degree of mutual consistency to allow a quasi-steady state to form, a condition naturally met by self-sustained modes, and in general not expected to hold for transient waves.
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The correct treatment of gravitational many-body systems containing self-organized global patterns, such as density wave modes in disk galaxies, requires a re-examination of classical dynamical approaches and assumptions. Our experience so far has shown that entirely new qualitative and quantitative results can emerge from the collective interactions of the many particles in a complex dynamical system. Formerly sacred laws (such as the differential form of the Poisson equation) can break down at the crest of collisionless shocks, and new meta-laws (such as the equality of the volume torque integral with the derivative of the sum of gravitational and advective surface torque coupling integrals) appear as emergent laws. Such emergent behavior is the low-energy Newtonian dynamical analogy of high energy physics' spontaneous breaking of gauge symmetry, a well known pathway for forming new meta laws when traversing the hierarchy of organizations.
| 10 | 12 |
1012.0277
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1012
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1012.5387_arXiv.txt
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Neutrino energy losses through neutral weak currents in the triplet-spin superfluid neutron liquid are studied for the case of condensate involving several magnetic quantum numbers. Low-energy excitations of the multicomponent condensate in the timelike domain of the energy and momentum are analyzed. Along with the well-known excitations in the form of broken Cooper pairs, the theoretical analysis predicts the existence of collective waves of spin density at very low energy. Because of a rather small excitation energy of spin waves, their decay leads to a substantial neutrino emission at the lowest temperatures, when all other mechanisms of neutrino energy loss are killed by a superfluidity. Neutrino energy losses caused by the pair recombination and spin-wave decays are examined in all of the multicomponent phases that might represent the ground state of the condensate, according to modern theories, and for the case when a phase transition occurs in the condensate at some temperature. Our estimate predicts a sharp increase in the neutrino energy losses followed by a decrease, along with a decrease in the temperature that takes place more rapidly than it would without the phase transition. We demonstrate the important role of the neutrino radiation caused by the decay of spin waves in the cooling of neutron stars.
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Usually neutron stars consist mostly of a superdense neutron matter which is in $\beta$ equilibrium with a small fraction of protons and contains the triplet-correlated superfluid condensate of neutrons below some critical temperature \cite{Tamagaki}- \cite{Elg}. For a long time, it has been generally accepted that the pair condensation in the superdense neutron matter occurs into the $^{3}P_{2}$ state (with a small admixture of $^{3}F_{2}$) with a preferred magnetic quantum number $m_{j}=0$. This model has been conventionally used for estimates of neutrino energy losses in the minimal cooling scenarios of neutron stars \cite{Page04}, \cite{Page09}. During the last decade, considerable work has been done with the most realistic nuclear potentials to determine the magnitude of the energy gap in the triplet superfluid neutron matter for different densities \cite{Khodel}-\cite{0203046}. Sophisticated calculations have shown that, besides the above one-component state, there are also multicomponent states involving several magnetic quantum numbers that compete in energy and represent various phase states of the condensate dependent on the temperature. Whether the phase transitions modify the spectrum of low-energy excitations and the intensity of neutrino emission from the volume of neutron stars is the question we try to answer in this paper. Theoretical investigation of low-energy excitations responsible for the neutrino emission by the neutron triplet superfluid liquid is conducted first. Until recently, the only known excitations able to decay into neutrino pairs were the broken pairs. It is well known that the neutrino emission caused by the pair-recombination processes in the neutron triplet superfluid liquid can dominate in the long-term cooling of neutron stars \cite{YKL}. We will consider also the collective excitations in the timelike domain of energies and momenta, which can also be responsible for the intense neutrino emission. Since the neutrino emission in the vector channel of weak interactions is strongly suppressed \cite{L10a} we will focus on the collective spin-density oscillations that can decay into neutrino pairs through neutral weak currents. Previously spin modes have been studied in the $p$-wave superfluid liquid $^{3}He$ \cite{Maki}-\cite{Wolfle}. The pairing interaction in $^{3}He$ is invariant with respect to rotation of spin and orbital coordinates separately. In this case, the spin fluctuations are independent of the orbital coordinates. In contrast, the triplet-spin neutron condensate arises in high-density neutron matter owing mostly to spin-orbit interactions that do not possess the above symmetry. Therefore the results obtained for liquid $^{3}He$ cannot be applied directly to the superfluid neutron liquid. Recently spin waves with the excitation energy smaller than the superfluid energy gap were predicted to exist in the $^{3}P_{2}$ superfluid condensate of neutrons \cite{L10a}. The neutrino decay of such spin waves \cite{L10b} is important for thermal evolution of neutron stars with the conventional one-component ground state with $m_{j}=0$. In this paper, we consider spin-density excitations for the other superfluid phases, which can be preferred at some temperatures. We will not consider the spin oscillations of the normal component. These soundlike waves that transfer into the ordinary spin waves in the normal Fermi liquid above the critical temperature cannot kinematically decay into neutrino pairs. Instead, we will focus on the spin excitations of the order parameter, which are separated by some energy interval from the ground state and are kinematically able to decay into neutrino pairs. The dispersion equation for such waves in the $^{3}P_{2}$ superfluid one-component condensate with $m_{j}=0$ was derived in Ref. \cite{L10a} in the BCS approximation. In this paper we study the collective spin excitations in multicomponent phases of the condensate, while taking into account the Fermi-liquid interactions. This paper is organized as follows. Section II contains some preliminary notes and outlines some of the important properties of the Green functions and the one-loop integrals used below. In Sec. III we discuss the renormalization procedure which transforms the standard gap equation to a very simple form valid near the Fermi surface. In Sec. IV we derive the effective ordinary and anomalous three-point vertices responsible for the interaction of the multicomponent neutron superfluid liquid with an external axial-vector field. We analyze the poles of anomalous vertices in order to derive the dispersion of spin-density oscillations in the condensate. In Sec. V we derive the linear response of the multicomponent superfluid neutron liquid onto an external axial-vector field. In Sec. VI we briefly discuss the general expression that relates the neutrino energy losses through neutral weak currents to the imaginary part of response functions. We derive the neutrino losses caused by recombination of broken Cooper pairs and by decay of spin waves. Finally, in Sec. VII, we evaluate neutrino energy losses in the multicomponent superfluid neutron liquid undergoing the phase transition. Section VIII contains a short summary of our findings and the conclusion. Throughout this paper, we use the standard model of weak interactions, the system of units $\hbar=c=1$ and the Boltzmann constant $k_{B}=1$.
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Let us summarize our results. We have studied the linear response of the superfluid neutron liquid to an external axial-vector field. The calculation is made for the case of a multicomponent condensate involving several magnetic quantum numbers and allows us to consider various phases of superfluid neutron liquid. In order to estimate the neutrino energy losses, while taking into account possible phase transitions, we have considered the low-energy excitations of the multicomponent condensate. Along with the well-known excitations in the form of broken Cooper pairs, we consider the collective waves of spin density, which are known to exist in the one-component condensate at very low energy \cite{L10b}. Our theoretical analysis predicts the existence of such waves in all of the multicomponent phases we have considered. We found that the excitation energy of spin waves is identical for all of the phases and is independent of the Fermi-liquid interactions. In the angle-average approximation, the energy of spin-density oscillations is estimated as $\omega_{s}\left( q=0\right) \simeq\Delta /\sqrt{5}$. Neutrino energy losses caused by the pair recombination and spin-wave decays are given by Eqs. (\ref{PBF}) and (\ref{SWD}), respectively. Because of a rather small excitation energy, the decay of spin waves leads to a substantial neutrino emission at the lowest temperatures $T\ll T_{c}$, when all other mechanisms of the neutrino energy losses are killed by a superfluidity. We have evaluated the neutrino energy losses for all of the multicomponent phases that might represent the ground state of the condensate according to modern theories. Finally we have evaluated the temperature dependence of neutrino energy losses from the superfluid neutron liquid in the case when the phase transition occurs in the condensate at the temperature $T=0.7T_{c}$. Our estimate predicts a sharp increase of the neutrino energy losses followed by a decrease, along with a decrease of the temperature that takes place more rapidly than it would without the phase transition. Since the neutron triplet-spin pairing occurs in the core which contains more than 90\% of the neutron star volume, the neutrino processes discussed here could influence the evolution of neutron stars.
| 10 | 12 |
1012.5387
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1012
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1012.2272_arXiv.txt
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We observe a large excess of power in the statistical clustering of Luminous Red Galaxies in the photometric SDSS galaxy sample called MegaZ DR7. This is seen over the lowest multipoles in the angular power spectra $C_{\ell}$ in four equally spaced redshift bins between $0.45 \le z \le 0.65$. However, it is most prominent in the highest redshift band at $\sim 4\sigma$ and it emerges at an effective scale $k \lesssim 0.01 \mathrm{h \, Mpc^{-1}}$. Given that MegaZ DR7 is the largest cosmic volume galaxy survey to date (3.3 (Gpc $h^{-1}$)$^3$) this implies an anomaly on the largest physical scales probed by galaxies. Alternatively, this signature could be a consequence of it appearing at the most systematically susceptible redshift. There are several explanations for this excess power that range from systematics to new physics. This could have important consequences for the next generation of galaxy surveys or the $\Lambda$CDM model. We test the survey, data and excess power, as well as possible origins.
| 10 | 12 |
1012.2272
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1012
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1012.2791_arXiv.txt
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Observations of the Sun and of Sun-like stars provide access to different aspects of stellar magnetic activity that, when combined, help us piece together a more comprehensive picture than can be achieved from only the solar or the stellar perspective. Where the Sun provides us with decent spatial resolution of, e.g., magnetic bipoles and the overlying dynamic, hot atmosphere, the ensemble of stars enables us to see rare events on at least some occasions. Where the Sun shows us how flux emergence, dispersal, and disappearance occur in the complex mix of polarities on the surface, only stellar observations can show us the activity of the ancient or future Sun. In this review, I focus on a comparison of statistical properties, from bipolar-region emergence to flare energies, and from heliospheric events to solar energetic particle impacts on Earth. In doing so, I point out some intriguing correspondences as well as areas where our knowledge falls short of reaching unambiguous conclusions on, for example, the most extreme space-weather events that we can expect from the present-day Sun. The difficulties of interpreting stellar coronal light curves in terms of energetic events are illustrated with some examples provided by the SDO, STEREO, and GOES spacecraft.
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Magnetic activity of Sun and Sun-like or ``cool'' stars results in a rich variety of observable phenomena that range from the asterospheres that surround these stars down to the stellar surfaces and --~with rapid advances expected in asteroseismology~-- below. Many of these phenomena are directly observable on Sun and stars alike, including the outer-atmospheric phenomena of persistent chromospheres and coronae, as well as their perturbations in the form of short-lived light-curve perturbations that are the signature of energetic events. It is on the latter that I focus here. The proximity of the Sun enables us to see details in the evolving magnetic field and associated atmospheric phenomena that are simply impossible to infer from stellar observations. Small wonder, then, that many of the names for stellar phenomena are taken from the solar dictionary: active regions, spots, flares, eruptions, and even the processes such as differential rotation and meridional advection that are part of the equivalent of the magnetic activity cycle, all the way to the loss of angular momentum associated with the gusty outflow of hot, magnetized plasma. Ensembles of observations taken over periods of years to centuries are revealing the statistical properties of some of these phenomena, even as state-of-the art observatories in space and on the ground are revealing physical processes and the interconnectedness of the global outer atmosphere. But recent observations of the Sun only provide a very limited view of what its magnetic activity has on offer, for at least two reasons. First, our Sun has a magnetic activity cycle that is relatively long compared to researchers' careers as well as to the era of advanced technology that has aided us in our observations. Consequently, we can expect even the 'present-day Sun' to have surprises in store for us that we have not yet observed simply because we have not been looking long enough. Some of these surprises may lie hidden in records such as polar ice sheets, while others lie embedded in rocks from outer space. But the lessons that can be learned from these invaluable and, as yet, under-explored archives are limited by access to these resources, by the limited temporal resolution of such records, and by the long chains of processes that sit between a solar phenomenon like the sunspot cycle or its largest flares and the 'recording physics' for the archive from which we are attempting to learn about them. In addition to learning about the Sun from such 'geological' records of its activity, one can also perform an ensemble study of states of infrequent extreme solar activity by looking at a sample of stars like it. This can provide us with a large enough sample of Sun-like stars that we can begin to assess how frequently the Sun may subject us to rare but high-impact events such as dangerous superflares and disruptive geomagnetic storms: although rare, the damage that may be inflicted to our global society and its safety and economy by extreme events is of such a magnitude that in-depth study of their properties and likelihood is prudent.\footnote{See the NRC report on ``Severe Space Weather Events -- Understanding Societal and Economic Impacts" at \hbox{http://www.nap.edu/catalog.php?record$\_$id=12507}.} A second reason why stellar studies are crucial to understanding of solar activity is that only stellar observations allow us to explore what the Sun's activity has been in the very distant past or what it will be in the very distant future (measured on time scales up to billions of years) by selecting stars of a wide range of ages. In this review, I discuss a sampling of the results coming out of the study of what has been termed the solar-stellar connection. I focus, in particular, on lessons that we are learning about what could be called 'space climate', i.e., the characteristic state of activity of a star like the Sun including the fluctuations about the mean in the form of energetic events like flares and coronal mass ejections (CMEs). Consequently, one of the topics selected for this review is a comparison of frequency distributions of bipolar regions, flares, CMEs, and solar energetic particle (SEP) events, the possible relationships between them, and the lessons learned by combining solar and stellar observations in our quest to establish the 'laws' of astro-magnetohydrodynamics. Another topic is that of light curves, which touches on the need to have pan-chromatic knowledge to guide our interpretation of stellar observations as well as on the long-standing concept of 'sympathetic events' in stellar magnetic activity.
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The sample observations discussed above demonstrate that the combination of geophysical, heliophysical, and astrophysical data can teach us much about the Sun's magnetic climate, up to the most energetic of events. The interpretations outlined are, of course, to be tested and alternatives, that doubtlessly exist, are to be explored. Despite the speculative nature of the scenarios sketched above, it appears that we are close to having the material available to learn where the solar flare-energy spectrum drops below a solar-stellar power law: the combination of the study of archives in ice and of stellar flare statistics should be able to provide us with an answer. On a less positive note, the solar observations discussed above demonstrate that measuring the energies involved in explosive and eruptive events is difficult, that separating events based on lightcurves alone is an ambiguous exercise, and that broad wavelength coverage is essential to both of these objectives: to learn about the most severe space weather, we have to accept that long-duration, large-sample, pan-chromatic (and thus often multi-observatory) stellar observations are needed because they can provide crucial information that can otherwise only be gathered by observing the Sun for a very long time and undergoing the detrimental effects of extreme magnetic storms on the Sun and around the Earth.
| 10 | 12 |
1012.2791
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1012
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1012.0331_arXiv.txt
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In contrast to planets with masses similar to that of Jupiter and higher, the bulk compositions of planets in the so-called super-Earth regime ($M_{p}$ = 2 -- 10\,M$_{\oplus}$) cannot be uniquely determined from a mass and radius measurement alone. For these planets, there is a degeneracy between the mass and composition of the interior and a possible atmosphere in theoretical models\cite{adams08,rogers10a}. The recently discovered transiting super-Earth GJ\,1214b is one example of this problem\cite{charbonneau09}. Three distinct models for the planet that are consistent with its mass and radius have been suggested\cite{rogers10b}, and breaking the degeneracy between these models requires obtaining constraints on the planet's atmospheric composition\cite{millerricci09,millerricci10}. Here we report a ground-based measurement of the transmission spectrum of GJ\,1214b between 780 and 1000\,nm. The lack of features in this spectrum rules out cloud-free atmospheres composed primarily of hydrogen at 4.9\,$\sigma$ confidence. If the planet's atmosphere is hydrogen-dominated, then it must contain clouds or hazes that are optically thick at the observed wavelengths at pressures less than 200\,mbar. Alternatively, the featureless transmission spectrum is also consistent with the presence of a dense water vapor atmosphere.
| 10 | 12 |
1012.0331
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1012
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1012.1876_arXiv.txt
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We have computed the fate of exoplanet companions around main sequence stars to explore the frequency of planet ingestion by their host stars during the red giant branch evolution. Using published properties of exoplanetary systems combined with stellar evolution models and Zahn's theory of tidal friction, we modeled the tidal decay of the planets' orbits as their host stars evolve. Most planets currently orbiting within 2~AU of their star are expected to be ingested by the end of their stars' red giant branch ascent. Our models confirm that many transiting planets are sufficiently close to their parent star that they will be accreted during the main sequence lifetime of the star. We also find that planet accretion may play an important role in explaining the mysterious red giant rapid rotators, although appropriate planetary systems do not seem to be plentiful enough to account for all such rapid rotators. We compare our modeled rapid rotators and surviving planetary systems to their real-life counterparts and discuss the implications of this work to the broader field of exoplanets.
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As the number of known exoplanetary systems grows, we gain an ever more complete picture of the angular momentum reservoir stored in the exoplanetary orbits. This reservoir can become important as the star evolves and begins to expand. At some point, many of the known exoplanets will be near enough to their host stars for their gravity to raise tides on the star, which will distort the stellar shape and introduce a torque into the star-planet system. As long as the planets' orbital periods are shorter than the stellar rotation periods, the torque will act in the sense that will ``spin-up'' the star. Tidal dissipation of energy in the convective envelopes of these now-red giant stars allows angular momentum to be transfered from the planetary orbit to the stars. As the angular momentum is drained from the planetary orbit, the planet moves closer to the star, increasing the tidal distortion and accelerating the rate of the transfer. As a result, the planet rapidly spirals into the star, dumping its angular momentum in the process. The result of this planetary demise may help us understand the unusual class of rapidly rotating red giants. Because stars should spin down as they evolve and expand, red giant stars are expected to have slow rotation speeds. This expectation has been verified empirically by studies that find that most red giants are characterized by $v\sin i$~$\approx$~2~km~s$^{-1}$ \citep{gray81,gray82,demed96}. A small fraction of red giants, around a few percent \citep[see, e.g.,][]{demed99,massarotti08a,carlberg10b}, deviate from this general rule. They have $v\sin i$\ in excess of 10~km~s$^{-1}$ and sometimes significantly higher. Many of these stars have no known stellar companions with which to interact, and planet accretion is a simple explanation that may account for these rapid rotators. However, this explanation raises a number of questions for which answers are needed to verify planet accretion as the underlying cause. Does the number of rapid rotators predicted from modeling the future evolution of exoplanet systems match the number actually observed? If not, what does this imply about the occurrence of planets around the progenitors of the red giant rapid rotators? Can chemical abundances distinguish rapid rotators created by planet accretion from those created in some other way?
| 10 | 12 |
1012.1876
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1012
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1012.2980_arXiv.txt
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We present the light curves and spectral data of two exceptionally luminous gamma-ray outburts observed by the Large Area Telescope (LAT) experiment on board {\it Fermi} Gamma-ray Space Telescope from 3C~273 in September 2009. During these flares, having a duration of a few days, the source reached its highest $\gamma$-ray flux ever measured. This allowed us to study in some details their spectral and temporal structures. The rise and decay are asymmetric on timescales of 6 hours, and the spectral index was significantly harder during the flares than during the preceding 11 months. We also found that short, very intense flares put out the same time-integrated energy as long, less intense flares like that observed in August 2009.
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The LAT experiment, on board the {\it Fermi Gamma-ray Space Telescope} satellite, observes the entire sky, in the 0.02 -- $>$300 \GeV~ band, once every $\sim$3 hours. It is providing the first collection of well sampled gamma-ray light curves of several blazars useful to study their variability on time scales from day to several months (see e.g. Abdo et al. 2010c). Daily light curves can be obtained for several blazars and for those exceptionally bright it is possible to observe significant occasional variations on timescales shorter than a day. We have now observed such events in 3C273, the nearest quasar. \tc, the first quasar discovered by Schimdt (1963) and the first extragalactic source detected by {\it COS}-B in the gamma-ray band (Swanenburg et al. 1978), is one of the most extensively studied AGN across the entire electromagnetic spectrum. It is classified as a flat spectrum radio quasar (FSRQ) and has a redshift $z=0.158$ (see e.g., Strauss et al., 1992). It was observed by EGRET (3EG~J1229$+$0210 in Hartman et al. 1999) at an average flux of $0.18$ \latfluxvi with a peak flux of $1.27$\latfluxvi (Nandikotkur et al. 2007). Despite the huge amount of data collected to now (see for instance Soldi et al. 2008), its behaviour is still surprising and raises new challenges for physical models. It was detected by the LAT experiment since the beginning of the observation in 2008 (Marelli, 2008) and is identified with the $gamma$-ray source 1FGL~J1229.1$+$0203 in the First {\it Fermi} LAT catalog (Abdo et al. 2010a). For about one year its behavior was characterized by a rather stable flux of $\sim$0.3\latfluxvi, with some flares superposed reaching peak level higher by about an order of magnitude. Beginning at the end of July 2009, \tct started a brightening phase during which some very bright outbursts were observed. The first event lasted about 10 days in August 2009 (Bastieri, 2009) and was characterized by fast rise and decay times ($< 1$ day) and a relatively stable `plateau'. In contrast the second and third otbursts were sharply peaked from September 15 to 17 and from September 20 to 23, 2009 (Hill, 2009), when \tct reached a peak photon flux above $10^{-5}$ \latflux. Before the very recent (December 2009) flare observed in 3C 454.3 (Escande \& Tanaka, 2009; Stirani et al., 2009), it was, therefore, the brightest extragalactic source, non Gamma-Ray Burst, observed by {\it Fermi}. EGRET detected only a few blazars (e.g. 3C 279 and PKS 1622$-$297) with a flux above $10^{-5}$ \latflux. Thanks to the high flux is was possible to obtain light curves with a good S/N ratio with a time binning of only six hours (corresponding to two scans of sky) and therefore we can describe the evolution of this outburst with a level of detail never reached before. In this letter we present the results of the LAT observations.
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The large outburst of \tct observed by {\it Fermi} in September 2009 revealed that blazars can reach very high brightness levels for quite short short time intervals of 1-10 days. The fluences of these events, however, are comparable to that of longer but less intense flares. We measured significant variations of the flux over time scales as short as six hours, which occurred with only a mild change of the spectral shape. In contrast to previous low intensity flares (Abdo et al., 2010c), the very strong events of September have light curves characterized by decay times longer than the rise time. Light curves at different energies provide some indication, in particular for the second flare, that the decay rate above 0.4 \GeV~ is shorter than below. The most natural interpretation of this difference is that it is due to the radiative cooling of the high energy electrons responsible for the $\gamma$-ray emission. We can therefore to estimate that this time, as observed in the Earth's frame, is of the order of 0.5--1 day. Such short lifetimes explain why longer flares are generally symmetric. In fact, it is possible that they are structured in a series of short subflares with typical durations of two-three days, much shorter than the total flare length. The apparent rise and decay times of long events are then likely due to the superposition of quite shorter flares and any information on the radiative lifetimes is practically lost. In September 2009 the apparent position of \tct was rather close to the Sun, making it was practically impossible to perform observations in other frequency ranges. Therefore we do not have data to study the evolution of the broad band Spectral Energy Distribution. The typical SEDs of FSRQs (Abdo et al. 2010b) peak at rather low $\gamma$-ray energies, around 0.1 GeV and frequently lower. In the case of \tc, which was the most significant AGN detected by COMPTEL in the 1-30 MeV range, the peak of the inverse Compton component was in the range 1--10 MeV (Collmar et al 2000). Our data show that the photon index was always steeper than 2, and this indicates that the peak energy remained below 100 MeV also during the flares. $\gamma$-ray activity has been found to be related to changes of the blazars' radio structure observed with VLBI. In this respect \tct is one of the most interesting sources because of the flux and of the low redshift that allows a fine spatial resolution. Jorstad et al. (2001), on the basis of a large data set at 22 and 43 GHz, found an association between the ejection of superluminal radio knots and high states of $\gamma$-ray luminosity in ten blazars in EGRET observations, including 3C~273. They concluded that both the radio and high energy events are originating from the same shocked region of a relativistic jet. Similar correspondences were already reported for the two much more distant FSRQs S5 0836+710 (Otterbein et al. 1998) and PKS 0528+134 (Britzen et al. 1999). More recent results on the MOJAVE VLBA sample and Fermi-LAT observations (Lister et al. 2009, Savolainen et al. 2009a) provided evidence that $\gamma$-ray loud blazars have a Doppler factor higher than non LAT-detected sources. A detailed plot of the kinematics of the various components in the radio jet of \tct can be found in Lister et al. (2009b): the resulting mean superluminal velocity is $\beta_{app}$=13.4 with an estimated Doppler factor $\delta$=16.8. It will be very interesting to verify if the exceptional outbursts of September 2009 will or will not be associated with the ejection of new superluminal knots, possibly with even higher velocity and Doppler factor. Moreover, the discovery of a possible connection between the peak intensity and rise time of the $\gamma$-ray outbursts with the VLBI parameters can be very useful to constrain the modelling and the energetics of perturbations in the jet.
| 10 | 12 |
1012.2980
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1012
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1012.5758_arXiv.txt
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We have investigated the stellar population properties in the central regions of a sample of lenticular galaxies with bars and single-exponential outer stellar disks using the data from the SAURON integral-field spectrograph retrieved from the open Isaac Newton Group Archive. We have detected chemically decoupled compact stellar nuclei with a metallicity twice that of the stellar population in the bulges in seven of the eight galaxies. A starburst is currently going on at the center of the eighth galaxy and we have failed to determine the stellar population properties from its spectrum. The mean stellar ages in the chemically decoupled nuclei found range from 1 to 11 Gyr. The scenarios for the origin of both decoupled nuclei and lenticular galaxies as a whole are discussed.
| 10 | 12 |
1012.5758
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1012
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1012.5791_arXiv.txt
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{Tidal dissipation in late-type stars is presently poorly understood and the study of planetary systems hosting hot Jupiters can provide new observational constraints to test proposed theories.}{We focus on systems with F-type main-sequence stars and find that the recently discovered system CoRoT-11 is presently the best suited for such a kind of investigation.}{A classic constant tidal lag model is applied to reproduce the evolution of the system from a plausible nearly synchronous state on the ZAMS to the present state, thus putting constraints on the average modified tidal quality factor $\langle Q_{\rm s}^{\prime} \rangle$ of its F6V star. {{ Initial conditions with the stellar rotation period longer than the orbital period of the planet can be excluded on the basis of the presently observed state in which the star spins faster than the planet orbit.}}}{{{ It is found that $4 \times 10^{6} \ltsim \langle Q_{\rm s}^{\prime} \rangle \ltsim 2 \times 10^{7}$, if the system started its evolution on the ZAMS close to synchronization, with an uncertainty related to the constant tidal lag hypothesis and the estimated stellar magnetic braking within a factor of $\approx 5-6$.}} {{ For a non-synchronous initial state of the system, $\langle Q_{\rm s}^{\prime} \rangle \ltsim 4 \times 10^{6}$ implies an age younger than $\sim 1$~Gyr, while $\langle Q_{\rm s}^{\prime} \rangle \gtsim 2 \times 10^{7}$ may be tested by comparing the theoretically derived initial orbital and stellar rotation periods with those of a sample of observed systems. Moreover,}} we discuss how the present value of $Q_{\rm s}^{\prime}$ can be measured by a timing of the mid-epoch and duration of the transits as well as of the planetary eclipses to be observed in the infrared with an accuracy of $\sim 0.5-1$~s over a time baseline of $\sim 25$~yr. }{CoRoT-11 is an highly interesting system potentially allowing us a direct measure of the tidal dissipation in an F-type star as well as the detection of the precession of the orbital plane of the planet that provides us with an accurate upper limit for the obliquity of the stellar equator. If the planetary orbit has a significant eccentricity ($e \gtsim 0.05$), it will be possible to detect also the precession of the line of the apsides and derive information on the Love number of the planet and its tidal quality factor.}
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\label{intro} \subsection{Tidal dissipation theories} Tidal dissipation in close binary systems with late-type components is generally constrained by the ranges of orbital periods corresponding to circular orbits as observed in clusters of different ages. \citet{OgilvieLin07} review recent observations and conclude that the equilibrium tide theory is insufficient to explain binary circularization by at least two orders of magnitude. Therefore, in addition to the dissipation of the kinetic energy of the flow associated with the tidal bulge, which is considered in the equilibrium tide theory \citep[e.g., ][]{Zahn77,Zahn89}, other effects must be included. The dynamical tide theory treats the dissipation of waves excited by the oscillating tidal potential in the stellar interior whose kinetic energy is ultimately extracted from the orbital motion. For simplicity, we shall assume that stars are rotating rigidly. We consider a reference frame rotating with the stellar angular velocity $\Omega= 2\pi/P_{\rm rot}$, where $P_{\rm rot}$ is the stellar rotation period. In that frame, the tidal potential experienced by the star can be written as a sum of rigidly rotating components proportional to the spherical harmonics $Y_{l m} (\theta, \phi)$, viz., ${\rm Re}\, [\Psi_{l m} r^{l} Y_{m}^{l} (\theta, \phi) \exp(-i \hat{\omega}_{lm} t) ]$, where $(r, \theta, \phi)$ are spherical polar coordinates with the origin in the centre of the star, $\hat{\omega}_{lm}$ is the tidal frequency in that frame, $\Psi_{l m}$ the amplitude of the component of degree $l$ and azimuthal order $m$, and $t$ the time. The tidal frequency is given by $\hat{\omega}_{lm} = l n - m \Omega $, where $n = 2\pi/P_{\rm orb}$ is the mean motion of the binary, $P_{\rm orb}$ being its orbital period. Waves are expected to be excited with the different frequencies $\hat{\omega}_{lm}$ corresponding to the various components of the tidal potential and their amplitudes will depend on that of the exciting component and on the response of the stellar interior. For a nearly circular orbit, the $l=m=2$ component is the dominant one and is responsible for the synchronization of the stellar rotation with the orbital motion \citep[e.g., ][]{OgilvieLin04}. The efficiency of tidal dissipation is usually parameterized by a dimensionless quality factor $Q$ proportional to the ratio of the total kinetic energy of the tidal distortion to the energy dissipated in one tidal period $2\pi/\hat{\omega}$ \citep[e.g., ][]{Zahn08}. In the theory, $Q$ always appears in the combination $Q^{\prime} \equiv (3/2) (Q/k_{2})$, where $k_{2}$ is the Love number of the star that measures its density stratification\footnote{Note that $k_{2}$ is twice the apsidal motion constant of the star, often indicated with the same symbol, as in, e.g., \citet{Claret95}.}. Therefore, the smaller the value of $Q^{\prime}$, the stronger the tidal dissipation. In general, $Q^{\prime}$ depends on $l$, $m$, and the tidal frequency $\hat{\omega}$, thus a rigorous treatment of the tidal dissipation should consider the sum of the effects associated with the different tidal components each having its specific $Q^{\prime}$. In practise, we adopt a single value of $Q^{\prime}$ which represents an average of the contributions of the different components. Moreover, we also average on the tidal frequency, which means averaging along the evolution of a given system because the tidal frequency decreases with time and goes to zero when tidal dissipation has circularized and synchronized the binary. The observations reviewed by \citet{OgilvieLin07} indicate that an average $Q^{\prime}$ ranging between $5 \times 10^{5}$ and $2 \times 10^{6}$ is adequate to account for the circularization of late-type main-sequence binaries. Such low values require an efficient tidal dissipation mechanism that \citet{OgilvieLin07}, moving along the lines of previous work, propose to be the damping of inertial waves in the stellar interior. These waves have the Coriolis force as their restoring force and are excited provided that the tidal frequency $\hat{\omega}$ satisfies the relationship: \begin{equation} |\hat{\omega} | \leq 2 \Omega. \label{exc_cond} \end{equation} The corresponding $Q^{\prime}$ has a remarkable dependence on the tidal and stellar rotation frequencies owing to the complex details of wave excitation and dissipation that are still poorly understood \citep[cf. also][]{GoodmanLackner09}. The main point, which can be regarded as firmly established, is that $Q^{\prime}$ decreases by $2-4$ orders of magnitude when $|\hat{\omega}|/\Omega \leq 2$ with respect to the case when $|\hat{\omega}|/\Omega > 2$, because the excitation of inertial waves is forbidden in the latter case and only the damping of the equilibrium tide contributes to the dissipation. In view of the present uncertainties in the dynamical tide theory, a precise determination of the tidal dissipation in binary systems with well known parameters is highly desirable. The case of F-type main-sequence stars is particularly challenging from a theoretical viewpoint because their internal structure consists of a thin outer convective zone and a radiative interior hosting a small convective core at the centre of the star. Since the propagation and dissipation of inertial waves are remarkably different in the convective and radiative zones, the study of F-type stars provides a critical test for the theory. As the mass of the outer convection zone decreases rapidly with increasing stellar mass between $1.2$ and $1.5$~M$_{\odot}$, the value of $Q^{\prime}$ is expected to increase by $3-4$ orders of magnitude within this mass range \citep{BarkerOgilvie09}. \subsection{Testing tidal theory with planetary systems} A new opportunity to test the tidal theory comes from the star-planet systems, in particular those containing hot Jupiters. Systems with transiting planets have the best determined stellar and planetary parameters and are particularly suited to study tidal dissipation { \citep[see, e.g., ][]{CaronePatzold07}}. F-type host stars having a mass $M \geq 1.2-1.5$~M$_{\odot}$ evolve quite rapidly during their main-sequence lifetime, thus improving significantly their age estimate from model isochrone fitting in comparison with lower mass stars. A good age estimate is important to constrain the average value of $Q^{\prime}$ by modelling the tidal evolution of a particular system (cf. Sect.~\ref{back_integ}). In Table~\ref{hot-Jupiters}, we list the presently known transiting systems with a star having an effective temperature $T_{\rm eff} \geq 6250$~K which corresponds to a spectral type earlier or equal to F8V. \begin{table*} \begin{tabular}{lllllllllll} Name & $T_{\rm eff}$ & $M$ & $R$ & $P_{\rm rot}$ & $P_{\rm orb}$ & $M_{\rm p}$ & $R_{\rm p}$ & $\tau_{\rm syn}$ & Refs. \\ & (K) & (M$_{\odot}$) & (R$_{\odot}$) & (days) & (days) & (M$_{\rm J}$) & (R$_{\rm J})$ & (Gyr) & \\ & & & & & & & & & \\ CoRoT-11 & 6440 $\pm $ 120 & 1.27 $\pm $ 0.05 & 1.37 $\pm $ 0.03 & 1.73 $\pm $ 0.26 & 2.994 & 2.33 $\pm $ 0.34 & 1.43 $\pm $ 0.03 & 8.673 & Ga10 \\ HAT-P-06 & 6570 $\pm $ 80 & 1.29 $\pm $ 0.06 & 1.46 $\pm $ 0.06 & 8.49 $\pm $ 1.34 & 3.853 & 1.06 $\pm $ 0.12 & 1.33 $\pm $ 0.06 & 28.38 & To08, No08 \\ HAT-P-07 & 6350 $\pm $ 80 & 1.47 $\pm $ 0.07 & 1.84 $\pm $ 0.17 & 24.51 $\pm $ 5.59 & 2.205 & 1.78 $\pm $ 0.01 & 1.36 $\pm $ 0.14 & 0.166 & Na09 \\ HAT-P-09 & 6350 $\pm $ 150 & 1.28 $\pm $ 0.13 & 1.32 $\pm $ 0.07 & 5.61 $\pm $ 0.78 & 3.923 & 0.78 $\pm $ 0.09 & 1.40 $\pm $ 0.06 & 205.8 & Am09, Sh09 \\ HAT-P-14 & 6600 $\pm $ 90 & 1.30 $\pm $ 0.03 & 1.47 $\pm $ 0.05 & 8.85 $\pm $ 0.85 & 4.628 & 2.20 $\pm $ 0.04 & 1.20 $\pm $ 0.58 & 15.21 & To10, Si10 \\ HAT-P-24 & 6373 $\pm $ 80 & 1.19 $\pm $ 0.04 & 1.32 $\pm $ 0.07 & 6.67 $\pm $ 0.68 & 3.355 & 0.69 $\pm $ 0.03 & 1.24 $\pm $ 0.07 & 58.49 & Ki10 \\ HD147506 & 6290 $\pm $ 110 & 1.32 $\pm $ 0.08 & 1.42 $\pm $ 0.06 & 3.14 $\pm $ 0.30 & 5.633 & 8.62 $\pm $ 0.55 & 0.98 $\pm $ 0.04 & 4.379 & Ba07, Lo08 \\ HD15082 & 7430 $\pm $ 100 & 1.50 $\pm $ 0.03 & 1.44 $\pm $ 0.03 & 0.81 $\pm $ 0.11 & 1.220 & 4.10 $\pm $ 4.00 & 1.50 $\pm $ 0.05 & 0.347 & Ca10 \\ HD197286 & 6400 $\pm $ 100 & 1.28 $\pm $ 0.16 & 1.24 $\pm $ 0.05 & 3.68 $\pm $ 0.59 & 4.955 & 0.96 $\pm $ 0.20 & 0.93 $\pm $ 0.04 & 553.7 & He09a \\ Kepler-05 & 6297 $\pm $ 60 & 1.37 $\pm $ 0.06 & 1.79 $\pm $ 0.06 & 18.91 $\pm $ 4.80 & 3.548 & 2.11 $\pm $ 0.06 & 1.43 $\pm $ 0.05 & 1.009 & Ko10 \\ Kepler-08 & 6213 $\pm $ 150 & 1.21 $\pm $ 0.07 & 1.49 $\pm $ 0.06 & 7.16 $\pm $ 0.78 & 3.523 & 0.60 $\pm $ 0.19 & 1.42 $\pm $ 0.06 & 61.17 & Je10 \\ OGLE-TR-L9 & 6933 $\pm $ 60 & 1.52 $\pm $ 0.08 & 1.53 $\pm $ 0.04 & 1.97 $\pm $ 0.07 & 2.486 & 4.50 $\pm $ 1.50 & 1.61 $\pm $ 0.04 & 0.927 & Sn09 \\ WASP-03 & 6400 $\pm $ 100 & 1.24 $\pm $ 0.09 & 1.31 $\pm $ 0.09 & 4.95 $\pm $ 0.91 & 1.847 & 1.76 $\pm $ 0.09 & 1.29 $\pm $ 0.09 & 1.002 & Gi08, Po08 \\ WASP-12 & 6250 $\pm $ 150 & 1.35 $\pm $ 0.14 & 1.57 $\pm $ 0.07 & 36.12 $\pm $ 49.03 & 1.091 & 1.41 $\pm $ 0.09 & 1.79 $\pm $ 0.09 & 0.012 & He09 \\ WASP-14 & 6475 $\pm $ 100 & 1.21 $\pm $ 0.12 & 1.31 $\pm $ 0.07 & 23.68 $\pm $ 6.35 & 2.244 & 7.34 $\pm $ 0.50 & 1.28 $\pm $ 0.08 & 0.017 & Jos08, Jo09 \\ WASP-15 & 6300 $\pm $ 100 & 1.18 $\pm $ 0.12 & 1.48 $\pm $ 0.07 & 18.69 $\pm $ 13.64 & 3.752 & 0.54 $\pm $ 0.05 & 1.43 $\pm $ 0.08 & 22.41 & We09 \\ WASP-17 & 6550 $\pm $ 100 & 1.20 $\pm $ 0.12 & 1.38 $\pm $ 0.20 & 7.76 $\pm $ 2.49 & 3.735 & 0.49 $\pm $ 0.06 & 1.74 $\pm $ 0.24 & 128.7 & An09 \\ WASP-18 & 6400 $\pm $ 100 & 1.25 $\pm $ 0.13 & 1.22 $\pm $ 0.07 & 5.60 $\pm $ 1.11 & 0.941 & 10.30 $\pm $ 0.69 & 1.11 $\pm $ 0.06 & 0.002 & He09b \\ XO-3 & 6429 $\pm $ 100 & 1.21 $\pm $ 0.07 & 1.38 $\pm $ 0.08 & 3.81 $\pm $ 0.23 & 3.192 & 11.79 $\pm $ 0.59 & 1.22 $\pm $ 0.07 & 0.814 & Joh08, Wi09 \\ XO-4 & 6397 $\pm $ 70 & 1.32 $\pm $ 0.02 & 1.56 $\pm $ 0.05 & 8.97 $\pm $ 0.80 & 4.125 & 1.78 $\pm $ 0.08 & 1.34 $\pm $ 0.05 & 11.32 & Mc08, Na10 \\ \end{tabular} \caption{Parameters of the transiting planetary systems having stars with $T_{\rm eff} \geq 6250$~K; M$_{\rm J}=1.90 \times 10^{27}$~kg and R$_{\rm J}=7.15 \times 10^{7}$~m indicate the mass and the radius of Jupiter. {References}: Am09: \cite{Ammleretal09}; An09: \cite{Andersonetal10}; Ba07: \cite{Bakosetal07}; Ca10: \cite{Cameronetal10}; Ga08: \cite{Gandolfietal10}; Gi08: \cite{Gibsonetal08}; He09a : \cite{hellieretal09a}; He09b: \cite {hellieretal09b}; He09: \cite{Hebbetal09}; Je10: \cite{Jenkinsetal10}; Jos08: \cite{Joshietal08}; Joh08: \cite{JohnsKrulletal08}; Jo09: \cite{Johnsonetal09}; Ki10: \cite{Kippingetal10}; Ko10 : \cite{Kochetal10}; Lo08: \cite{Loeilletetal08}; Mc08: \cite{McCulloughetal08}; Na09: \cite{Naritaetal09}; Na10: \cite{Naritaetal10}; No08: \cite{Noyesetal08}; Po08: \cite{Pollaccoetal08}; Sh09: \cite{Shporeretal09}; Si10: \cite{Simpsonetal10}; Sn09: \cite{Snellenetal09}; To08 : \cite{Torresetal08}; To10: \cite{Torresetal10}; We09: \cite{Wessetal09}; Wi09: \cite{Winnetal09b}} \label{hot-Jupiters} \end{table*} The columns from left to right list the name of the system, the effective temperature $T_{\rm eff}$, the mass $M$ and the radius $R$ of the star, its rotation period $P_{\rm rot}$, as derived from the observed spectroscopic rotation broadening $v \sin i$ and the estimated stellar radius { assuming a equator-on view of the star}, the orbital period $P_{\rm orb}$, the mass $M_{\rm p}$ and radius $R_{\rm p}$ of the planet, % the timescale for the synchronization of the stellar rotation $\tau_{\rm syn}$, and the references. To compute the synchronization time, we assume that the entire star is synchronized as is customary in tidal theory and is suggested by the tidal evolution of close binaries observed in stellar clusters of different ages. The synchronization timescale is then a measure of the strength of tidal dissipation in the star and is computed according to the formula: \begin{equation} \tau_{\rm syn}^{-1} \equiv \frac{1}{\Omega} \left| \frac{d\Omega}{dt} \right| = \frac{9}{2} \frac{1}{\gamma^{2} Q^{\prime}_{\rm s}} \left( \frac{M_{\rm p}}{M} \right)^{2} \left( \frac{R}{a} \right)^{9/2} \left| 1 - \left( \frac{n}{\Omega} \right) \right| \sqrt{\frac{G M}{R^{3}}}, \label{tidal_time} \end{equation} where $\gamma R \simeq 0.22 \, R$ is the gyration radius of the star \citep{Siessetal00}, $Q_{\rm s}^{\prime}$ its modified tidal quality factor, here assumed to be $Q^{\prime}_{\rm s} = 10^{6}$, $a$ the semimajor axis of the orbit, and $G$ the gravitation constant \citep[see ][]{MardlingLin02}. Equation~(\ref{tidal_time}) is valid for circular orbits and when the spin axis is aligned with the orbital angular momentum. In this regard, the values given here must be considered as estimations and for illustration purpose only, as high eccentricity and/or obliquity have been measured for some of these systems. According to the dynamical tide theory by \citet{OgilvieLin07}, the stars experiencing the strongest tidal interaction are those with $n/\Omega = P_{\rm rot}/P_{\rm orb} \leq 2$ because they have $|\hat{\omega}|/\Omega \leq 2$ when the $l=m=2$ component of the tidal potential is considered \citep{BarkerOgilvie09}. For such stars, the orbital decay (or expansion) due to the tidal interaction can in principle be observed. Among those systems, the best candidate is CoRoT-11 \citep{Gandolfietal10} because it has a tidal synchronization timescale in between those derived for systems like HD~197286/WASP-7 or HAT-P-09, i.e., much longer than the main-sequence lifetime of the system, and those of, e.g., OGLE-TR-9 or WASP-3 that are shorter than the expected ages of the systems, implying that the synchronous final state has possibly already been reached. Other systems, e.g., HD15082/WASP-33, have a star so massive that a $Q^{\prime}_{\rm s}$ too large to be measurable is expected, or show a remarkable misalignment between the stellar spin and the orbital angular momentum that makes the derivation of the rotation period from the $v \sin i$ quite uncertain, as in the case of XO-4. In view of the peculiar characteristics of { the CoRoT-11 system}, we shall consider it for a detailed study of the tidal evolution. We shall derive constraints on the average $Q^{\prime}_{\rm s}$ value of its F6V star by considering a possible initial state for the system when its star settled on the zero-age main-sequence (hereafter ZAMS; see Sect.~\ref{initial_status}). Moreover, we shall demonstrate how the present value of $Q^{\prime}_{\rm s}$ can be directly measured with suitable transit observations extended on a time interval of a few decades.
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{ Assuming an initial state of the CoRoT-11 system close to synchronization between the stellar spin and the orbital period of the planet, we can put constraints on the average modified tidal quality factor of its F-type star, finding $4 \times 10^{6} \ltsim Q_{\rm s}^{\prime} \ltsim 2 \times 10^{7}$. Rigorously speaking, we should have expressed these constraints in terms of the average tidal time lag $\Delta t$ because our tidal equations are valid for a constant $\Delta t$ \citep[cf. Sect.~\ref{tidal_model} and ][]{Leconteetal10}. However, in view of the uncertainty on the above constraints, we prefer to give them in terms of $Q_{\rm s}^{\prime}$ that varies by a factor up to $\approx 4-5$ during the evolution, if $\Delta t$ is assumed to be constant (cf. Sects.~\ref{tidal_model} and \ref{back_integ}). An initial state close to synchronization is not the only possible one. However, the minimum estimated age of the system of $\sim 1$~Gyr implies $Q_{\rm s}^{\prime} \gtsim 4 \times 10^{6}$, otherwise the evolution from any initial state with $P_{\rm orb}> P_{\rm rot}$ would be too fast, leading to a present $P_{\rm orb}$ longer than observed. An upper limit on $Q_{\rm s}^{\prime}$ can be set when a larger statistics of the values of the $n/\Omega = P_{\rm rot}/P_{\rm orb}$ ratio will be available because for $Q_{\rm s}^{\prime} \gtsim 5 \times 10^{7}$ the initial value of $n/\Omega$ is predicted to be $\sim 0.25$ that is smaller than the presently observed minimum $n/\Omega \sim 0.5$ (see Sect.~\ref{back_integ}). The above limits on $Q_{\rm s}^{\prime}$ are only a factor of $2-15$ smaller than the average value predicted by \citet{BarkerOgilvie09} which can be considered a success of the dynamic tide theory of \citet{OgilvieLin07} given the current uncertainty on the tidal dissipation processes occurring in stellar convection zones. } We find that CoRoT-11 is one of the best candidates to look for orbital period variations related to tidal evolution by monitoring its transits and secondary eclipses in the optical and in the infrared passbands. In contrast to several systems hosting very hot Jupiters, e.g., WASP-12 or WASP-18, whose orbital decay may be too slow to be measurable \citep{OgilvieLin07,Lietal10}, the parameters of CoRoT-11 appear to be ideal for a detection by measuring the times of mid-transit and mid-eclipse with an accuracy of at least $1-5$~s along a time baseline of $\sim 25$~yr. This is a consequence of the peculiar ratio of the rotation period of the star to the orbital period of the planet that is presently the smallest among systems having a radial velocity curve compatible with zero eccentricity. { The test is particularly sensitive to values of $Q_{\rm s}^{\prime}$ between $10^{5}$ and $10^{6}$ leading to a fast tidal evolution of the system, as discussed in Sect.~\ref{back_integ}. } The CoRoT-11 system is also particular because the precession of the orbital plane due to a non-zero obliquity of the stellar spin can be measured through the variation of the duration of the transit on a timescale as short as $5-10$~yr by means of ground-based observations with telescopes of the $6-8$~m class. Note that even an obliquity as small as $5^{\circ}$ can be detected with this method. On the other hand, the eccentricity of the planetary orbit can be well constrained by measuring the times of the secondary eclipses in the infrared and their duration. If the system turns out to have an eccentric orbit, this puts a constraint also on the minimum value of the planet quality factor $Q_{\rm p}^{\prime}$ if we assume that the eccentricity is of primordial origin, or is a clear indication of the presence of a perturbing body that excites it \citep[e.g., ][]{TakedaRasio05}. The light-time effect due to a distant third body is indeed the only phenomenon that can seriously hamper the detection of the transit time variations expected from the tidal orbital evolution. The detection of a distant companion may be particularly difficult, but this limitation is present also in any other system candidate for a direct measurement of the tidal dissipation.
| 10 | 12 |
1012.5791
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1012
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1012.0661_arXiv.txt
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With gravity, ionization, and radiation being considered, we perform 2.5D compressible resistive MHD simulations of chromospheric magnetic reconnection using the CIP-MOCCT scheme. The temperature distribution of the quiet-Sun atmospheric model VALC and the helium abundance ($10\%$) are adopted. Our 2.5D MHD simulation reproduces qualitatively the temperature enhancement observed in chromospheric microflares. The temperature enhancement $\Delta T$ is demonstrated to be sensitive to the background magnetic field, whereas the total evolution time $\Delta t$ is sensitive to the magnitude of the anomalous resistivity. Moveover, we found a scaling law, which is described as $\Delta T/\Delta t \sim {n_H}^{-1.5} B^{2.1} {\eta_0}^{0.88}$. Our results also indicate that the velocity of the upward jet is much greater than that of the downward jet and the X-point may move up or down.
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Magnetic reconnection plays a very important role in solar flares, corona mass ejections, and other solar activities. During the solar minimum, many authors focus their attention on the solar small-scale activities such as microflares \citep{Qiu2004, Fang2006a, Ning2008, Brosius2009}, Ellerman bombs \citep{Fang2006b, Watanabe2008}, chromospheric jets \citep{Nishizuka2008}, and so on. Microflares, or subflares, or bright points, are small-scale and short-lifetime solar activities. The size of microflares ranges from several arcsecs to about 20 arcsecs, the duration and the total energy can be $10-30$ minutes and $10^{26}-10^{29}$ ergs, respectively \citep{Shimizu2002, Fang2006a}. The most distinctive feature in their visible spectra is the faint emission at the center and the brightening at the wing of some chromospheric lines such as H$\alpha$ line. Soft X-ray \citep{Golub1974, Golub1977}, hard X-ray \citep[HXR][]{Lin1984}, EUV \citep{Porter1984, Emslie1978} and microwave \citep{Gary1988} emissions have also been observed in some microflares. It is an interesting fact that there are some correlations among the emissions at different wavelengths. For instance, \cite{Qiu2004} found that about 40\% of microflares show HXR emissions at $>$10 keV and microwave emissions at $\sim$10 GHz. Recently, \cite{Ning2008} found that there is a correlation between the power-law index of the HXR spectrum and the emission measure of microflares. \cite{Brosius2009} found that the studied microflare is bright first in the chromospheric and transition region spectral lines rather than in the corona, which is consistent with the chromospheric heating by nonthermal electron beams. All these evidence indicates that microflares are related to the nonthermal processes driven by magnetic reconnection. Some microflares might result from the magnetic reconnection in the corona. However, there is accumulating evidence suggesting that some other microflares are due to the reconnection in the chromosphere \citep{Liu2004, Xia2007}. For instance, \cite{Tang2000} found that in many cases emerging flux occurred about $5-30$ minutes before the microflares, where the emerging flux may collide with the pre-existing magnetic field at the chromospheric height. \cite{Brosius2009} also stated that magnetic reconnection in the chromosphere could be a plausible mechanism for triggering the chromospheric microflare. \cite{Chen2001} made 2.5-dimensional (2.5D) numerical simulations of chromospheric magnetic reconnection in order to study Ellerman bombs and type II white-light flares. However, in the simulations, they made some simplifications, such as the omissionof the gravity. In this paper, in order to improve their results, we perform 2.5D magnetohydrodynamic (MHD) simulations using the CIP-MOCCT scheme, with gravity being included, and try to make some comparisons with the observations of chromospheric microflares. In the next section, the numerical method is given in detail. Our results, which include the dynamic process, parameter dependence, a scaling law and a comparison with a semi-empirical model, are described in Section \ref{NUMERICAL RESULTS}. Discussion and summary are given in Section \ref{Disc}.
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\label{Disc} More and more theoretical works have indicated that magnetic reconnection in the solar lower atmosphere can produce chromospheric microflares \citep{Tand1988, Liu2004, Fang2006a}, as well as other small activities like Ellerman bombs \citep{Fang2006b, Ding1998, Watanabe2008}. With gravity, ionization and radiation being considered, we performed 2.5D MHD simulations. Since we have used the temperature distribution of the quiet-Sun atmospheric model VALC \citep{Vernazza1981} and considered the helium abundance ($10\%$), our simulations are realistic to some extent. In our simulation, we calculate the magnetic reconnection rate by using the definition $R = v_{in}/v_A$, rather than the formula $R = d \psi / dt$ used by \cite{Chen1999, Chen2001} and $R = |\eta J_z|$ used by \cite{Yokoyama2001}, where $\psi$ means the magnetic flux function and $J_z$ stands for the $z$-component of current density. If we use the formula $R = d \psi / dt$ or $R=|\eta J_z|$, the reconnection rate has a large pulse, which can be greater than 1.0 at the beginning of the evolution. It is non-physical, as mentioned by \cite{Yokoyama2001}. However, it should be noted that all of these methods can give a similar result after enough time of evolution. In the subsections~\ref{Dynamic Process} and~\ref{Compare}, we studied the magnetic reconnection with the X-point at different heights, e.g., 500 km and 1000 km. Our results, in the case of 500 km, can well reproduce the temperature enhancement in the semi-empirical model \citep{Fang2006a}. At the heights of 500 km and 1000 km , the temperature enhancement is 1000-2000 K and 2000-3000 K, respectively. In the case of 1000 km, we found that the height of X-point can move up by tens of kilometers. This is different from the results of the case of 500 km, where the X-point moves down. It is noted that the temperature enhancement we got in our simulations is several thousand Kelvin, and our results can only represent the rise phase of chromospheric microflares since our simulations were performed only until the upward reconnection jet reaches the bottom of the transition region. Of course, those microflares with strong EUV, X-ray, and microwave emissions may occur at the lower corona. In this case, the emissions in X-rays, EUV and microwave can be naturally explained. It is worth noting that our simulations have some limitations. Even though we limited the computational height below $2000$ km, the plasma $\beta$ at the bottom boundary is still six orders of magnitude larger than that at the top boundary. It causes some numerical instability in our simulations. Besides, our initial magnetic configuration is relatively simple. Moreover, as it is well known, the current sheet should be very thin in the real situation, e.g., less than hundreds of meters, while in our simulation the minimum grid sizes are $\triangle x = 1.25$ km and $\triangle y = 5$ km, which are too large to simulate the current sheet. All these contribute to the fact that our simulations can not reproduce the real observations in details. Further improvement of this work is expected. In summary, we give the conclusions as follows: 1. Our 2.5D MHD simulations can reproduce the temperature enhancement in chromospheric microflares qualitatively. The temperature increase in the cases when the reconnection point is at the height of 500 km and 1000 km can reach 1000-2000 K and 2000-3000 K, respectively. 2. The free parameters in our 2.5D simulation are the background magnetic field ($B_0$) and the anomalous resistivity ($\eta_0$). We have found that the temperature enhancement is sensitive to the background magnetic field, while the total evolution time is sensitive to the magnitude of anomalous resistivity. 3. Our simulation results indicate that the velocity of the upward jet is much larger than that of the downward jet, and the X-point may exhibit downward or upward motions. 4. We have performed a parameter survey, and found that the temperature enhancement and the total evolution time in the chromospheric reconnection follow the scaling law: \begin{equation} \frac{\triangle T}{\triangle t} \approx 3.7 (\frac{n_H}{3.4 \times 10^{19} \ {\rm m}^{-3}}) ^{-1.5} (\frac{B_0}{25 \ {\rm G}})^{2.1} (\frac{\eta_0}{17.7 \ {\rm \Omega~m}})^{0.88} \ {\rm K~s}^{-1} \,\, . \end{equation}
| 10 | 12 |
1012.0661
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1012
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1012.1503_arXiv.txt
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Spectropolarimetry results for the starburst galaxy M82 are presented. The optical emission lines of the filaments in the energetic outflow (``superwind'') from the nuclear starburst region of M82 are substantially polarized. The H$\alpha$ polarization degrees and angles measured by our study are consistent with previous narrowband imaging polarimetry data. The polarized emission lines are redshifted with respect to the emission lines in the total light and systemic motion of the galaxy. The emission line intensity ratios [N~{\sc ii}]/H$\alpha$ and [S~{\sc ii}]/H$\alpha$ in the polarized light are similar to those of the nuclear star-forming region. In addition, the electron density $N_{\rm e}$ derived from the [S~{\sc ii}]$\lambda$6731/$\lambda$6717 line ratio of the polarized light is $\sim 600 - 1000$ cm$^{-3}$ at a distance of more than 1 kpc from the nucleus, whereas the $N_{\rm e}$ derived from the total light are less than 300 cm$^{-3}$. These facts strongly suggest that the emission from the nuclear starburst of M82 is scattered by dust grains entrained and transported outward by the superwind. A simple hollow biconical outflow model shows that the velocity of the outflowing dust grains, $v_{\rm d}$, ranges from 100 to 200 km~s$^{-1}$ near the nucleus, decreases monotonically with the distance from the nucleus, and reaches $\sim 10$ km~s$^{-1}$ at around 1 kpc. The motion of the dust is substantially slower than that of both ionized gas ($v_{\rm H\alpha} \sim 600$ km~s$^{-1}$) and molecular gas ($v_{\rm CO} \sim 200$ km~s$^{-1}$) at the same distance from the nucleus of M82. This indicates that dust grains in the superwind are kinematically decoupled from both gas components at large radii. Since the dust velocity $v_{\rm d}$ is much less than the escape velocity of M82 ($v_{\rm esc} \approx 170$ km~s$^{-1}$ at 1.5 kpc from the nucleus), most of the dust entrained by the superwind cannot escape to intergalactic space, and may fall back into the galaxy disk without any additional acceleration mechanisms (such as radiation pressure).
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Active star formation in starburst galaxies plays one of the most important roles in galaxy evolution. In spite of its short duration ($\sim 10^{7-8}$ yrs), a starburst reforms a substantial amount of the interstellar medium (ISM) of a galaxy into stars. In addition, a starburst creates an enormous hot outflow from a galaxy via the collective effect of supernovae explosions and the stellar winds of massive stars. This energetic outflow, known as a ``superwind,'' is a ubiquitous phenomenon in starburst galaxies (e.g., \cite{heckman03}; \cite{veil05}). Theoretical studies estimate that the terminal velocity of the hot gas of a superwind reaches $\sim 10^3$ km s$^{-1}$ (e.g \cite{chev85}), far exceeding the typical escape velocity of a galaxy. Metal-rich gas from the starburst region and ambient disk gas are expelled from the galaxy disk and pollute the galaxy halo and intergalactic space. Rapid consumption of the interstellar medium by a superwind will cause star formation in a galaxy to cease abruptly, although part of the expelled gas returns to the disk and induces further star formation. These negative and positive feedbacks greatly affect the chemical evolution of a galaxy. A significant amount of dust is associated with superwinds. Submillimeter (sub-mm), far-infrared (IR), and mid-IR maps, optical color maps, and polarization studies of several nearby starburst galaxies have exhibited extended dust emission along the galaxy minor axis (\cite{alton99}; \cite{leeuw09}; \cite{ichi94}; \cite{sca91}; \cite{engel06}; \cite{kaneda10}). Sub-mm observations of some starburst galaxies suggest that the mass of the dust may reach $\sim 10^{6-7}$ M$_{\odot}$ (\cite{alton99}; \cite{leeuw09}). Although such a large-scale dust outflow has been inferred to play an important role in the evolution and metal enrichment of the intergalactic medium and halo gas, as well as the evolution of the host galaxy itself, the fate of this dust has thus far remained uncertain. This is because kinematic information on dust outflow is quite difficult to obtain, since dust produces no sharp emission or absorption lines by which its radial velocity can be measured. One of the most promising techniques for probing the motion of dust in and around superwinds is optical spectropolarimetry. Outflowing dust grains entrained by a superwind scatter and polarize the continuum light and emission lines emanating from the nuclear starburst region. In other words, the dust grains act as ``moving mirrors'' for nuclear light. Hence, the velocity measured by the polarized emission lines of a superwind must reflect the motion of the dust with respect to the nucleus. Motivated by this idea, we carried out deep optical spectropolarimetry of the prototypical starburst galaxy M82 to reveal the dust kinematics of its superwind. M82 is distinguished by its very bright, kpc-scale superwind (\cite{nakai87}; \cite{sea01}; \cite{matsu05}; \cite{bierao08}; \cite{bland88}; \cite{shop98}; \cite{ohyama02}; \cite{mut07}; \cite{breg95}; \cite{tsuru07}; \cite{rana08}; \cite{strick07}). The superwind of M82 is accompanied by large extraplanar dust filaments, as well as hot ionized gas (\cite{ichi94}; \cite{alton99}; \cite{thuma00}). \citet{alton99} found a huge dust envelope extending along the minor axis of M82 using the 850-$\mu$m sub-mm observation. Recently, \citet{leeuw09} detected a much fainter dust emission as far as 1.5 kpc from the galaxy disk. They found that the sub-mm morphology has a north--south asymmetry, which is consistent with the H$\alpha$ and X-ray morphologies. Mid- and far-IR maps recently obtained using IR space telescopes reveal a complex structure in the kpc-scale filaments of polycyclic aromatic hydrocarbon (PAH) dust extending along the minor axis of M82 (\cite{engel06}; \cite{kaneda10}; \cite{roussel10}). The highly polarized nature of the optical continuum and H$\alpha$ emission of the outer region of M82 also indicates the presence of a vast quantity of dust in the superwind (\cite{schmidt76}; \cite{bing76}; \cite{vis72}; \cite{sca91}). The polarization degree of the H$\alpha$ emission reaches 30\%\ in some areas, and the polarization angle is almost perpendicular to the radial direction drawn from the nucleus of the galaxy \citep{sca91}. Therefore, M82 is an ideal object for studying dust kinematics in a starburst superwind via optical spectropolarimetry. In this paper, we present the results of deep optical spectropolarimetric observations of the superwind of M82. This is the first attempt to ascertain the spatial structure of the dust motion in a starburst superwind via spectropolarimetry. We adopted 3.89 Mpc as the distance to M82 \citep{sakai99}, which yields a linear scale of 18.9 pc~arcsec$^{-1}$ for the galaxy.
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We conducted optical spectropolarimetry observations of the starburst superwind of the prototypical starburst galaxy M82 to reveal the motion of the dust entrained by the superwind. The H$\alpha$ polarization degrees ($\sim$5 -- 15\%) and angles measured by our study are consistent with previous narrowband imaging polarimetry maps. The polarized emission lines are redshifted relative to the systemic motion of the galaxy. The emission line intensity ratios [N~{\sc ii}]/H$\alpha$ and [S~{\sc ii}]/H$\alpha$ in the polarized light are similar to those of the nuclear star-forming region. The electron densities $N_{\rm e}$ derived from the polarized [S~{\sc ii}] line ratio are much higher than those derived from the total light. These facts strongly suggest that the emission from a nuclear starburst is scattered by dust grains entrained and then transported outward by the starburst superwind. We derived the outflow velocity of the dust grains, $v_{\rm d}$, using a simple hollow biconical outflow model. The outflow velocity $v_{\rm d}$ is on the order of a few hundred km~s$^{-1}$ near the nucleus and decreases monotonically with the distance from the nucleus. The dust motion revealed by this study is substantially slower than the motion of the other components of the superwind (ionized gas and molecular gas). The outflow velocity of the dust is also much less than the escape velocity of M82. In the absence of any additional effective acceleration mechanisms (such as radiation pressure), the dust expelled by the superwind would fall back into the galaxy disk within several times $10^7$ yr. \bigskip We are grateful to the Subaru Telescope staff for their kind assistance with the observations. We also thank the anonymous referee for his/her helpful comments. Part of this study was carried out using the facilities of the Astronomical Data Center, National Astronomical Observatory of Japan. This research made use of NASA's Astrophysics Data System Abstract Service. This work was financially supported in part by the Japan Society for the Promotion of Science (Grant-in-Aid for Scientific Research No. 18340055) and the Ministry of Education, Culture, Sports, Science \& Technology, Japan (Grant-in-Aid for Scientific Research on Priority Areas No. 19047003).
| 10 | 12 |
1012.1503
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1012
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1012.1656.txt
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The existence of concentric low variance circles in CMB sky, generated by black-hole encounters in an aeon preceding our big bang, is a prediction of the Conformal Cyclic Cosmology. Detection of three families of such circles in WMAP data was recently reported by Gurzadyan \& Penrose (2010). We reassess the statistical significance of the low variance circles detected by Gurzadyan \& Penrose by comparing with Monte Carlo simulations of the CMB sky with realistic modeling of the anisotropic noise in WMAP data. We find that all three groups are consistent at 3$\sigma$ with a Gaussian CMB sky as predicted by inflationary cosmology model.
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% ====================== In a recent paper \cite{Gurzadyan:2010da} reported a high significance detection of concentric circles in the Cosmic Microwave Background (CMB) maps with anomalously low variance. The existence of these circles, if true, pose a serious challenge to our understanding of the CMB as being a Gaussian random field withing the framework of inflationary cosmology. The above authors used data from the the Wilkinson Microwave Anisotropy Probe (WMAP) to look for concentric low variance circles in the CMB sky. They examined 10885 choices of center in the CMB sky after masking the galactic plain by excluding the $|b|<20^\circ$ region from the maps. For each choice of center, they computed the variance of the temperature fluctuations in successively larger concentric rings of $0.5^\circ$, at increasing radii. They found three groups of rings of low variance at various radii. In this paper we compare the variance of the above low-variance-circles with the average variance of Monte Carlo simulations of the CMB sky to assess the statistical significance of these circles. % ======================
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By comparing with Monte Carlo simulations of the CMB sky, we find that the low variance circles of \cite{Gurzadyan:2010da} are not anomalous. They can naturally occur in a Gaussian CMB sky consistent with the predictions of the inflationary cosmology.
| 10 | 12 |
1012.1656
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1012
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1012.0396_arXiv.txt
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We report the timing and broad-band spectral properties of the Be transient high mass X-ray binary pulsar \gro~ using a $Suzaku$ observation in the declining phase of its 2007 November-December outburst. Pulsations with a period of 93.737 s were clearly detected in the light curves of the pulsar up to the 80-100 keV energy band. The pulse profile was found to be strongly energy dependent, a double peaked profile at soft X-ray energy bands ($<$8 keV) and a single peaked smooth profile at hard X-rays. The broad-band energy spectrum of the pulsar, reported for the first instance in this paper, is well described with three different continuum models viz. (i) a high energy cut-off power-law, (ii) a Negative and Positive power-law with EXponential cut-off (NPEX), and (iii) a partial covering power-law with high energy cut-off. Inspite of large value of absorption column density in the direction of the pulsar, a blackbody component of temperature $\sim$0.17 keV for the soft excess was required for the first two continuum models. A narrow iron K$_\alpha$ emission line was detected in the pulsar spectrum. The partial covering model, however, is found to explain the phase averaged and phase resolved spectra well. The dip like feature in the pulse profile can be explained by the presence of an additional absorption component with high column density and covering fraction at the same pulse phase. The details of the results are described in the paper.
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High mass X-ray binary (HMXB) systems are strong X-ray emitters via the accretion of matter from the OB companion. The majority of the HMXBs are known to be Be X-ray binaries. The mass donor in the Be binary systems is generally a B star that is still on the main sequence and lying well inside the Roche surface. In these Be binary systems, the compact object (a neutron star) is typically in a wide orbit with moderate eccentricity with orbital period in the range of 16 - 400 days (Coe 2000). The neutron star in these Be systems spends most of the time far away from the circumstellar disk surrounding the Be companion. Mass transfer from the Be companion to the neutron star takes place through the circumstellar disk. Strong X-ray outbursts are normally seen when the neutron star (pulsar) passes through the circumstellar disk or during the periastron passage (Okazaki \& Negueruela 2001). The outbursts in the Be/X-ray binaries are also theoritically explained by invoking the truncation of the circumstellar disk (Okazaki et al. 2002). According to the model, the neutron star exerts a negative torque outside some critical radius resulting in the truncation of the Be disk. The disk matter then accumulates in the outer rings of the disk until the truncation is overcome by the effects of global one-armed oscillations, disk warping, etc. The subsequent sudden infall of the high-density disk matter onto the neutron star causes X-ray outbusts in these systems. With the exception of a very few peculiar cases like LS~I+61303 (Massi et al. 2004 and references therein), all of the Be/X-ray binary systems appear to be accretion powered X-ray pulsars. The pulse period of the X-ray pulsars in Be binary systems is in the wide range of seconds to hundreds of seconds. There is a strong correlation between the pulse period and the orbital period (Corbet 1986) suggesting effective transfer of angular momentum from the accreted material. On the longer time-scale (months to years), the variability is observed in optical and infrared bands that is attributed to the structural changes in the circumstellar disk (Reig et al. 2001 and references therein). The X-ray spectra of Be/X-ray binary pulsars are usually hard. A fluorescent iron emission line at 6.4 keV is observed in the spectrum of most of the X-ray pulsars. It is possible that most of these systems have a soft X-ray excess above the power-law continuum component. However, detection of the the soft excess depends on the value of absorption column density (Paul et al. 2002; Naik \& Paul 2004a, 2004b and references therein). The transient X-ray pulsar \gro~ was discovered on 1993 July 14 by the BATSE experiment onboard the {\it Compton Gamma Ray Observatory (CGRO)} (Stollberg et al. 1993). X-ray pulsations of 93.587 s were detected in the 20-120 keV energy range of BATSE. From ASCA observation, the X-ray pulse profile of the pulsar was found to have a double-peak structure with a well-defined, sharp intensity minimum and a less prominent secondary minimum (Tanaka et al. 1993).The BATSE spectrum was described by an optically thin thermal bremsstrahlung model with {\it kT} = 25 keV. Following the discovery, the optical and infrared observations of the optical counterpart to \gro~ revealed the presence of strong Balmer emission lines and infrared excess (Coe et al. 1994). Based on these results, the system was classified as a massive binary system consisting of a neutron star as the compact object and a Be or a supergiant primary. After 260 days of this outburst, a second outburst was detected by BATSE (Finger et al. 1994). Assuming this 260 d as the orbital period of the pulsar, Finger et al. (1994) estimated the mass of the binary companion to be 3-8 $M_\odot$ indicating the system as a high mass X-ray binary. The ROSAT PSPC observation of the pulsar, in the declining phase of the discovery outburst by BATSE in 1993, clearly detected the 93.4 s pulsation with a double-peaked pulse profile in 0.1-2.4 keV light curve (Petre \& Gehrels 1994). A search in the archive of EXOSAT/Medium-Energy Experiment (ME) observation, centered on HD~88661, revealed the presence of the pulsed emission at the same period during the 1993 outburst (Macomb, Shrader, \& Schultz 1994). The X-ray spectrum (0.8--10 keV range) was found to be highly absorbed ($N_H = 0.7\times10^{22} atoms cm^{-2}$) and described by a hard power-law with a photon index of $\sim$1.2. A combined analysis of data from the CGRO and ASCA observations, though non-simultaneous, shortly after the peak of the discovery outburst, reported that the broadband spectrum of the pulsar can be well approximated by a power-law with an exponential cutoff and a 6.4 keV iron emission line (Shrader et al. 1999). The pulse profile was also found to be energy dependent, a double-peaked profile detected by ASCA that evolved into a single-peaked profile as detected by BATSE. Analyzing the BATSE and $Rossi X-ray Timing Explorer (RXTE)$/ASM flux histories, Shrader et al. (1999) suggested the orbital period of the system to be $\sim$135 days. However, Levine \& Corbet (2006) detected a 248.9 day periodicity in the $RXTE$ All-sky Monitor (ASM) X-ray light curve by analyzing data accumulated over nearly 10 years. This periodicity was found by visual identification of periodically occurring outbursts in the ASM light curve. An independent analysis of pulse period variations during outbursts, using BATSE data, estimated the orbital period precisely to be 247.8 d (Coe et al. 2007) which is good agreement with the orbital period determined from the recurrence of the X-ray outbursts in ASM light curve. Following the detection of an intense outburst from \gro~ with the Burst Alert Telescope (BAT) on Swift on 2007 November 17 (Krimm et al.\ 2007), the accreting X-ray pulsar was observed with various X-ray observatories. The RXTE observations detected the pulsar up to $\sim$70 keV along with the regular $\sim$93.75 s pulsations (Wilms et al. 2007). Suzaku performed a TOO observation of the pulsar on 2007 November 30. The results obtained from the analysis of the Suzaku observation are presented in this paper. \begin{figure} \centering \includegraphics[height=3.4in, width=3in, angle=-90]{Fg1} \caption{The RXTE-ASM and Swift/BAT light curves of \gro\ in 1.5-12 keV and 15-50 keV energy bands, from 2007 September 29 (MJD 54372) to 2008 January 20 (MJD 54485). The region between the vertical lines shows the duration of the Suzaku observation of the source.} \label{asm} \end{figure}
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\subsection{Pulse Profile} The temporal and spectral properties of the Be transient X-ray pulsar \gro~ have been reported only in a couple of occasions since its discovery. We detected X-ray pulsations in GRO~J1008-57 as high as the 80-100 keV energy band, which has not been reported earlier. The pulse profile of \gro~ is found to be strongly energy dependent i.e. a double-peaked profile in the soft X-ray energy band ($<$ 10 keV) and a single-peaked smooth profile in hard X-rays. The double-peaked profile at soft X-rays, as shown here in this paper, is found to be different from that of the single-peaked profile (in 1--4 keV energy range) obtained from the EXOSAT observation of \gro~ in 1985 (Macomb et al. 1994). Similar changes in the shape of pulse profiles are also seen in the Be transient X-ray pulsar A0535+262 (Naik et al. 2008), a single peaked profile in the quiescence (Mukherjee \& Paul 2005) and a double-peaked profile during outbursts (at high luminosity). The double-peaked pulse profiles seen in \gro, agrees with the luminosity dependence of the pulse profiles as seen in other pulsars. The presence of dip like structures in the pulse profiles of these X-ray pulsars is described as due to the obscuration of matter to the radiation. As we showed here, the dip like feature in the pulse profile of \gro~ is probably due to the additional absorption (other than the Galactic column density) at the pulse phase. \subsection{Spectroscopy} The broad-band X-ray spectrum of \gro~ has been described here for the first time in detail. Shrader et al. (1999) tried to explain the pulsar spectrum obtained from the OSSE data by a thermal bremsstrahlung model with a characteristic temperature $kT$ = 19 keV. However, when the spectral fitting was extended towards the low energy (ASCA energy range), the fitting was inconsistent. The statistics was improved marginally because of the addition of an absorption component at $\sim$88 keV for the possible cyclotron features. Shrader et al. (1999), when fitted the ASCA, BATSE and OSSE data simultaneously, found that the broad-band spectrum can be described by an absorbed power-law with high energy cutoff and a Gaussian component for the Iron emission line. However, the ASCA, BATSE, and OSSE observations of the pulsar were not simultaneous. The midpoints of ASCA and OSSE observations used were separated by 4 days during which time the source intensity was halved. As the spectrum of Be transient pulsars can differ at different luminosity states, the broad-band spectral analysis of the Suzaku observations of the pulsar can give better information on the properties. Selection of an appropriate continuum model is important to investigate the presence of several features in the pulsar spectrum, such as the soft excess represented by a blackbody component, emission lines, cyclotron absorption features etc. Most of the transient Be X-ray binary pulsars undergo periodic outbursts due to the enhanced mass accretion when the neutron star passes through the dense regions of the circumstellar disk or periastron passage when the outer edge of the disk is stripped off resulting in sudden accretion of matter onto the neutron star. During the passage, the value of the absorption column density increases compared to the value of the Galactic column density in the source direction. The spectral fitting to the data obtained from ASCA, BATSE, OSSE observations of the transient pulsar \gro~ during the 1993 August outburst, yielded the value of $N_H$ in the range of 0.8--1.73 $\times$ 10$^{22}$ atoms cm$^{-2}$ (Table~1; Shrader et al. 1999). In case of data obtained from the Suzaku observation of the pulsar, it is found that the 0.8-70.0 keV broad-band spectra can be well described by three different continuum models with similar statistical parameters. The high energy cutoff power-law model and NPEX continuum model yielded higher value of $N_H$ than that of the Galactic value in the direction of the pulsar. It is interesting to note that, inspite of a high value of column density, a blackbody component of temperature $kT$ $\sim$0.2 keV was also required for these two continuum models to describe the broad-band spectrum of the pulsar. In these two models, it is estimated that the absorption corrected flux of the soft X-ray excess (blackbody) in \gro~ is about 2\% of the unabsorbed source flux in 0.8-70 keV energy range. The third model i.e. the partial covering model, however, fits the pulsar spectrum comparatively better than the previous two continuum models. Based on our results from the phase resolved spectroscopy, the earlier two continuum models were not preferred to describe the properties of the pulsar. In the partial covering model, N$_{H1}$ is considered as the Galactic hydrogen column density, and N$_{H2}$ is interpreted as the column density of the material that is local to the neutron star. The value of N$_{H2}$ is maximum during the primary dip that is interpreted as due to the accretion column. The high value of N$_{H2}$ and the covering fraction at 0.4--0.5 pulse phase range explain the dip like feature in the pulse profile. The broad-band spectroscopy of \gro~ also shows the presence of a narrow iron K$_\alpha$ emission line at 6.4 keV. The iron emission line is generally interpreted as due to the fluorescent line from the cold matter in the surrounding region of the neutron star.
| 10 | 12 |
1012.0396
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1012
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1012.2446_arXiv.txt
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{Young Stellar Objects (YSOs) in the early evolutionary stages are very embedded, and thus they emit most of their energy at long wavelengths such as far-infrared (FIR) and submillimeter (Submm). Therefore, the FIR observational data are very important to classify the accurate evolutionary stages of these embedded YSOs, and to better constrain their physical parameters in the dust continuum modeling. We selected 28 YSOs, which were detected in the AKARI Far-Infrared Surveyor (FIS), from the Spitzer c2d legacy YSO catalogs to test the effect of FIR fluxes on the classification of their evolutionary stages and on the constraining of envelope properties, internal luminosity, and UV strength of the Interstellar Radiation Field (ISRF). According to our test, one can mis-classify the evolutionary stages of YSOs, especially the very embedded ones if the FIR fluxes are not included. In addition, the total amount of heating of YSOs can be underestimated without the FIR observational data.}
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Stars are born in their parental molecular cloud cores through gravitational contraction. After the gravitational collapse begins, a protostar with its circumstellar disk forms at the center of the collapsing core. Although the protostar/disk system is embedded in the dense infalling envelope at the early phase of evolution of star, the envelope gradually disappears as the protostar evolves to the main-sequence star. The Spectral Energy Distribution (SED) of a Young Stellar Object (YSO) is contributed by the protostar at Near-Infrared wavelengths whose temperature ranges from 2000 K to 4000 K, the disk at Mid-Infrared (MIR) wavelength which has a temperature gradient with radius, and the envelope at Far-Infrared (FIR, 40 $<$ $\lambda$ $<$ 350 $\micron$) or Submillimeter (Submm, $\lambda$ $\geq$ 350 $\micron$) wavelengths. Therefore, the analysis of the SED of a YSO can provide a way to identify the evolutionary stage of the YSO. In the SED analysis, however, we need to understand the effect of the interstellar radiation field (ISRF) since the envelope is heated by the ISRF as well as the internal source luminosity. There is a commonly known scheme classifying the evolutionary stages of YSOs based on the shape of the SED of a YSO (Lada 1987; Andr\'{e} et al. 1993). In this proposed scheme, the Class 0 object is in the earliest stage after the collapse begins. The Class 0 object already harbors a central heating source, but most of the mass is still in the dense envelope of gas and dust. Therefore, all emission from the central source is reprocessed through the envelope, which is almost isothermal with T$\sim$30 K (Young et al. 2003). Class I object is still heavily embedded, but emission from the central protostar and disk escapes through the outflow cavities. As a result, the peak of the SED shows at MIR wavelengths. As the envelope material accretes to the central star through the disk and/or is cleared out by the outflows, the protostar/disk system becomes the dominant mass reservoir. Therefore, Class II objects are less embedded and have the emission excesses at near and mid-infrared wavelengths by dusty disk. The Class III objects emit energy mainly from central protostars and have only a small amount of gas in disks. The Class II and the Class III are often classified as classical T Tauri Stars and weak-line T Tauri stars, respectively (White et al. 2007). As mentioned, in the early embedded phases, the radiation from central source undergoes the absorption/re-emission processes by the heavy envelope. As a result, significant energy is emitted at FIR wavelengths. Andr\'{e} et al. (2000) noted that deeply embedded Class 0, whose energy peak is located at beyond 100 $\micron$, were often missed by IRAS observations unlike the Class I objects, which reveal the peak energy emission at the wavelength from 12 to 100 $\micron$. Therefore, the FIR data are crucial to study the evolutionary stages of embedded sources. The envelope is also heated by the ISRF so that the SED at FIR and Submm is affected by the energy input from the ISRF as well as the central source. Since the heating from the ISRF also affects the total luminosity of a YSO and the shape of the dust temperature profile, it is very important to constrain the strength of the ISRF, and thus the internal luminosity of the central heating source. The external heating by the ISRF compared to the internal heating is not negligible especially in the early embedded phase and in low luminosity sources (Dunham et al. 2008). The AKARI/Far-Infrared Surveyor (FIS) All-Sky Survey Bright Source Catalog combined with precedented catalogs such as those by the Spitzer observations provides a useful tool to investigate the importance of the FIR fluxes in the calculations of the evolutionary indicators of YSOs, such as $T_{bol}$ and $L_{bol}$/$L_{smm}$. In this paper, AKARI/FIS All Sky Survey Bright Source Catalog is used to investigate the importance of the FIR fluxes on the evolutionary indicators of YSOs, $T_{bol}$ and $L_{bol}$/$L_{smm}$. In addition, we model the SED of one YSO to study how the FIR fluxes can constrain the internal luminosity and the ISRF strength around the YSO. This paper is structured as follows. In section 2 and 3, the catalogs used in this study and the criteria adopted to select our sample YSOs are introduced. The evolutionary indicators of YSOs are summarized in section 4. We calculate those evolutionary indicators of our samples and present the results in section 5. The dust continuum modeling of the SED of one sample YSO is presented in section 6, and the summary and conclusions are followed in section 7.
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We tested the effect of FIR fluxes in classification of the evolutionary stages of YSOs and in calculation of the internal luminosity as well as the UV strength of the external heating source, the ISRF. Twenty eight sample YSOs, which have their counterparts only in the AKARI/FIS Bright Source Catalog, were selected from the c2d YSO catalogs. First, we calculated an evolutionary indicator, $T_{bol}$, and found that 6 objects, which are classified into class I when only Spitzer fluxes were used, were reclassified into class 0 when the AKARI FIR fluxes were included. Two objects, which were classified into class II when only the Spitzer data were included, were also reclassified into class I. This result shows that the FIR observations are very important in the classification of the evolutionary stages of the early embedded YSOs since they emit most of their energy at the long wavelength regime. In order to test how the FIR fluxes constrain the physical structure of the envelope, the internal luminosity, and the UV strength of the ISRF, we modeled the SED of one of our sources. In the early evolutionary phase, the external heating by the ISRF can not be ignored compared to the internal heating. According to our dust continuum modeling, when the FIR fluxes are included, the density profile of the envelope is steeper (p=1.9 compared to p=1.6) indicating most of envelope material is possibly static, that is, the infall radius is very small in the sense of Shu's inside-out collapse model. The estimated internal luminosity is constrained to 4.7 $L_{\sun}$ when the FIR fluxes are included (compared to 3.8 $L_{\sun}$ obtained when the FIR fluxes are not included). The UV strength of the ISRF is $G_{0}$=0.4 (attenuated by $A_V$ = 0.5 mag) in the model with the FIR fluxes, but in the model without the FIR fluxes, the ISRF needs to be completely attenuated ($G_{0}$=0). Our modeling result also indicates that fluxes at FIR and even longer wavelengths are crucial to better constrain the physical parameters of embedded YSOs. \vspace{0.5cm}
| 10 | 12 |
1012.2446
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1012
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1012.4395_arXiv.txt
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We use two-dimensional hydrodynamic simulations of viscous disks to examine whether dynamically-interacting multiple giant planets can explain the large gaps (spanning over one order of magnitude in radius) inferred for the transitional and pre-transitional disks around T Tauri stars. In the absence of inner disk dust depletion, we find that it requires three to four giant planets to open up large enough gaps to be consistent with inferences from spectral energy distributions, because the gap width is limited by the tendency of the planets to be driven together into 2:1 resonances. With very strong tidal torques and/or rapid planetary accretion, fewer planets can also generate a large cavity interior to the locally formed gap(s) by preventing outer disk material from moving in. In these cases, however, the reduction of surface density produces a corresponding reduction in the inner disk accretion rate onto the star; this makes it difficult to explain the observed accretion rates of the pre/transitional disks. We find that even with four planets in disks, additional substantial dust depletion is required to explain observed disk gaps/holes. Substantial dust settling and growth, with consequent significant reductions in optical depths, is inferred for typical T Tauri disks in any case, and an earlier history of dust growth is consistent with the hypothesis that pre/transitional disks are explained by the presence of giant planets. We conclude that the depths and widths of gaps, and disk accretion rates in pre/transitional disks cannot be reproduced by a planet-induced gap opening scenario alone. Significant dust depletion is also required within the gaps/holes. Order of magnitude estimates suggest the mass of small dust particles ($\lesssim 1 \mu$m) relative to the gas must be depleted to 10$^{-5}$ -- 10$^{-2}$ of the interstellar medium value, implying a very efficient mechanism of small dust removal or dust growth.
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The transitional and pre-transitional disks around young stars exhibit strong dust emission at wavelengths $\gtrsim 10 \mu$m, while showing significantly reduced fluxes relative to typical T Tauri disks at shorter wavelengths (e.g., Calvet \etal 2002, 2005; D'Alessio \etal 2005; Espaillat \etal 2007, 2008). In the pre-transitional disks, there is evidence for emission from warm, optically-thick dust near the star (Espaillat \etal 2007, 2008, 2010), while in the transitional disks the emission at $\lesssim 10 \mu$m appears to be due entirely to optically-thin dust (Calvet \etal 2002, 2005; Espaillat \etal 2010). The depletion of near- to mid-infrared emission is generally interpreted as being due to evacuation of the disk interior to scales $\sim 5$ to $\sim 50$~AU (Marsh \& Mahoney 1992; Calvet \etal 2002, 2005; Rice \etal 2003; Schneider \etal 2003; Espaillat \etal 2007, 2008, 2010; Hughes \etal 2009), an interpretation confirmed in some cases via direct sub-mm imaging (e.g., Pietu \etal 2006; Brown \etal 2007, 2009; Hughes \etal 2009; Andrews \etal 2009). One proposed mechanism for clearing inner disks while leaving the outer disk relatively undisturbed is the formation of giant planets, which can open gaps in disks (e.g., Lin \& Papaloizou 1986; Marsh \& Mahoney 1992; Nelson \etal 2000; Calvet \etal 2002; Rice \etal 2003), and which are expected to form initially at relatively small disk radii due to the shorter evolutionary timescales compared with the outermost disk. However, there are three observational challenges, as discussed further in \S 2, which theoretical explanations of these systems must confront. First, the transitional and pre-transitional disk systems exhibit average gas accretion rates close to T Tauri disk accretion rates ($\sim 10^{-8}$$\msunyr$; Hartmann et al. 1998) onto their central stars (e.g. Calvet \etal 2002, 2005; Espaillat \etal 2007, 2008; Najita et al. 2007). Maintaining this accretion requires either a significant mass reservoir interior to the disk-clearing planets, or some way of allowing mass from the outer disk to move past the gap-clearing planets. Second, the cleared regions in these disks are {\em large} (e.g. Espaillat \etal 2010). In the case of the transitional disks, the optically-thin region must extend from radii as large as tens of AU all the way in to the central star. Even the pre-transitional disks, which have evidence for optically-thick dust emission in the innermost regions, must have large disk gaps. Furthermore, the spectral energy distribution (SED) modeling suggests that the detected optically thick region may only extend to radii $\sim 1$~AU, with an extremely dust free region beyond, until reaching the outer optically-thick disk (Espaillat \etal 2010; \S 2). Third, the requirement that the gap/hole be optically thin implies that the mass of dust in sizes of order a micron or less must be extremely small. Thus, either the planet-induced gap is very deep and is effectively cleared of gas {\it and} dust, or the dust abundance is reduced by many orders of magnitude from abundances in the diffuse interstellar medium. These three conditions must be fulfilled simultaneously. Photoevaporation has been proposed as one gap/hole clearing mechanism (Alexander \& Armitage 2007, 2009), but requires the mass loss rate by the photoevaporation to be comparable to the disk accretion rate. Initial estimates of photoevaporative mass loss suggested values of $\sim$ 4$\times$10$^{-10}$$\msunyr$ (Clarke et al. 2001), which would not be rapid enough to counteract accretion. More recent estimates suggest higher values, perhaps as much as $\sim$10$^{-8}$$\msunyr$ due to the inclusion of X-rays (Gorti \& Hollenbach 2009; Owen et al. 2010); however, the highest values are problematic, because then it becomes difficult to understand why so many T Tauri disks last for several Myr with accretion rates smaller than 10$^{-8}$$\msunyr$ (Hartmann et al. 1998). Najita, Strom, \& Muzerolle (2007) suggested that giant planets might explain the transitional disks by creating a gap while still maintaining accretion onto the central star. Najita \etal argued that there is some evidence of reduction in the accretion rates in transitional disks relative to the so-called ``primordial'' disks, perhaps by an order of magnitude, and that this was consistent with simulations of a single giant planet in the studies by Lubow \& D'Angelo (2006) and Varniere \etal (2006). However, as discussed in \S 2, the accretion rates of the pre/transitional disks are not very low in absolute magnitude. Moreover, while Varniere \etal (2006) were able to produce a large cleared inner region with a single planet, the study of Crida \& Morbidelli (2007) calls this result into question (see also \S 3). In this paper we consider whether disk gap formation by a multiple planet system can satisfy the following requirements: 1) creation of a gap extending over a large range in radius; 2) maintenance of the inner disk accretion rate by flow through the planet-created gap; and 3) sufficient reduction in the disk surface density within the gap such that extreme depletion of disk dust is not essential. We examine the problem in the context of viscous disks, with improvements over some previous simulations, including a more realistic temperature distribution and an improved inner boundary condition. We find that while multiple giant planets can indeed open large gaps, it is difficult to explain the inferred properties of the pre/transitional disks without invoking substantial depletion of small dust. We offer a speculative scenario in which combined dust depletion and planet formation explains the observations.
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Returning to the questions posed at the end of the Introduction, our findings are as follows: 1) As seen in previous simulations, a single planet opens up a small gap. Multiple planets can open wider common gaps. However, in order to explain the pre-transitional disk gaps spanning almost an order of magnitude in radius without dust depletion (\S 2), we need as many as three to four giant planets, given the tendency of the viscously evolving disk to drive the planets into 2:1 resonances. This is a large number of such planets, given current exoplanet statistics which do not include examples of systems containing more than two planets in resonance. One possibility is that resonant planet systems which contain more than two planets are able to remain stable over significant time periods ($\sim 1$ Myr) in the presence of the gas disk, whose contribution to eccentricity damping helps maintain dynamical stability. Once the gas disk dissipates, however, these resonant planetary systems may become unstable, leading eventually to ejection of some members and the formation of more sparsely populated systems of planets on eccentric orbits, similar to those observed. According to this hypothesis, large systems of planets help explain the presence of the large disk gaps in pre-transitional disks, and also explain the population of eccentric giant exoplanets. In principle, the gaps could be widened for fewer planets if the orbits were eccentric; however, with large $\alpha$=0.01, we do not see high eccentricities in stable systems. 2) A high accretion rate past the planetary gap can be maintained if accretion onto the planets is sufficiently slow. Planetary accretion rates of $f=10$ still permit substantial accretion ($>$10\% accretion rate outside the gap) past the planets to the inner disk. However, the reduction in surface density interior to the gap is in proportion to the reduction in mass accretion rate. This differs from the results of Varniere \etal (2006) but is consistent with the results of Crida \etal (2007), and is almost certainly the result of the improved treatment of the inner boundary condition. 3) Carefully comparing our simulations with observations indicates that the pre/transitional disk systems require substantial reduction in dust opacities within the gaps (and inner disk if present) since the reductions in surface density there are not sufficient to explain the very small amount of dust required to fit the observations. For example, if we adopt the ISM opacity at 10 $\mu$m (the assumption used by Espaillat \etal (2010) to estimate optically-thin dust masses) of $\kappa \sim$ 10 cm$^{2}$ g$^{-1}$, even the gaps with surface densities $\Sigma \gtrsim 0.1 {\rm g\, cm^{-2}}$ are still optically thick, which is the case for all but the $f=1$ simulations. $f=1$ simulations still have optical depth $\sim$0.1 and they are the cases with an extreme upper limit on the expected planetary accretion rate, which leads to low accretion rate onto the star (esp. with multiple planets). Furthermore, the dust depletion must be larger in the inner disk than in outer regions. This follows from our adoption of a realistic temperature distribution and improved inner boundary condition, which results in a significantly higher surface density at small radii than at large radii; thus the inner disk will not be optically thin unless the outer disk is also optically thin, which disagrees strongly with transitional disk observations. One alternative is to place planets sufficiently close to the central star that the gap essentially extends either up to the dust destruction radius, or near to it; however, in that case it is difficult to extend the gap all the way out to 20 - 70 AU as observed (Table 1) when using only four planets. In principle it may be possible to accommodate additional planets at smaller orbital radii which truncate the inner disk to a smaller radius. The other alternative is to allow the planets to accrete more mass, depleting the inner disk; but then the mass accretion rate onto the central star is reduced by unacceptably large values. Protoplanetary disks are probably not fully viscous (e.g., Gammie 1996), but the existence of ``dead zones'' and/or reduced viscosity regions is unlikely to improve the situation, as they will lead to higher surface densities for a given mass accretion rate, and may reduce the stability of systems of more than two planets in resonance. \subsection{Gap structure and dust depletion} A multi-giant planet model coupled with substantial dust depletion in the gap and inner disk does have an attractive feature. As pointed out in \S 2, there is some evidence - both from SED modeling and from near-infrared interferometry - that the pre-transitional and transitional disks exhibit structure within the optically-thin disk gap or hole, such that the innermost gap region is much less cleared of small dust than the outer gap region. Schematically, we might identify the highly-cleared region with the true planet-driven gap, and the inner, optically-thin dust region with accreting gas that has been strongly depleted in dust but not so depleted in gas surface density. The idea is illustrated in Fig.\ref{fig:tau}, which compares the inferred structure of one of the pre-transitional disk systems with the azimuthally-averaged dust surface density for the P4A10 case. In Fig. 9, the disk optical depth at 10 $\mu$m can be estimated by multiplying the dust surface density (dust refers to small dust particles which contribute to 10 $\mu$m opacity) by 1000 cm$^{2}$/g (which is estimated using the ISM opacity: a factor of 100 comes from the gas-to-dust mass ratio and 10 cm$^{2}$/g is the ISM opacity at 10 $\mu$m). Thus, disk regions with dust surface density $>$10$^{-3}$ g/cm$^{2}$ are optically thick. The dotted curve represents the dust surface density resulting from the model P4A10 assuming there is no dust depletion, which is obtained by dividing the gas surface density from the P4A10 simulation by the nominal gas-to-dust mass ratio $\sim$100. In this case even the gap is optically thick. The solid curve represents the case where we notionally deplete the dust content inside the gap by a factor of 100. This can explain the pre-transitional disks whose inner region within $\sim 1$ AU changes from being optically thick to optically thin, with the region beyond being almost dust free. This comparison with pre-transitional disks set a lower limit on the dust depletion factor, while the transitional disks can set a higher limit. If we assume all the dust inside the transitional disk gap has been detected (which is reasonable since transitional disks are optically thin), we can estimate the dust surface density by solving \begin{equation} \int_{0}^{R_{out,thin}}\Sigma_{d}(R)2\pi R dR=M_{d}\,, \end{equation} where R$_{out,thin}$ and M$_{d}$ are given in Table 1, and is assumed to be $\Sigma_{d}$(AU)(R/AU)$^{-1}$. The derived $\Sigma_{d}$ for GM Aur ($<$1 AU) using Table 1 is shown as the dashed curve in Fig. 9. By comparing with the dotted curve, we find the dust needs to deplete by a factor of $\sim$10$^{5}$. Thus we estimate the dust-to-gas mass ratio in the inner disk is between 10$^{-5}$ and 10$^{-2}$ of the ISM value (0.01) for our model to be consistent with pre/transitional disks. More broadly, a gap-opening perturbing body or bodies still remains a plausible explanation of the pre/transitional disks. This type of model naturally predicts a sharp transition in surface density at the outer gap edge, which is needed to explain the strong mid- and far-infrared emission of these systems. While radially-dependent dust depletion might also mimic this effect, most T Tauri disks do not show this behavior. There are a significant number of T Tauri systems in which the disk is optically-thick at a wide range of radii, but which are much more geometrically flat based on their SEDs (e.g., the ``group E'' systems shown in Figure 7 of Furlan \etal 2006). These objects suggest that dust growth and settling can occur roughly simultaneously over a wide range of radii, without resulting in the abrupt change between optically-thin and optically-thick regions that occurs in the pre/transitional disks. In addition, the large outer disk gap radii (as much as 20-70 AU) pose a challenge for pure coagulation models at an age of $\sim 1-2$~Myr (Tanaka \etal 2005; Dullemond \& Dominik 2005). The requirement of significant dust depletion is not completely surprising, given the evidence for small dust depletion in many T Tauri disks without obvious gaps or holes (e.g., D'Alessio \etal 2001; Furlan \etal 2006). In addition, formation of giant planets via core accretion clearly requires dust growth. Indeed, an outstanding theoretical problem has been to avoid clearing inner disks by ages of 1 Myr or less (e.g., Dullemond \& Dominik 2005; Tanaka \etal 2005), resulting in suggestions that small dust must be replenished to some extent by fragmentation as a result of collisions of larger particles. A hypothesis which explains many features of pre/transitional disks using both giant planets and dust depletion can be outlined as follows. Dust coagulation and growth in the disk interior to $\sim 20$ AU ensues, leading to the formation of planetesimals, and eventually to the formation of a system of numerous giant planets. This system of giant planets forms a large common gap which covers radii from $< 1$ AU out to $\sim 50$ AU. Remnant planetesimals within the gap region, whose collisions can act as a secondary source of dust, are dynamically cleared out - preventing in situ secondary dust formation. After a viscous time scale corresponding to the size scale of the planetary system, the gas present in the planet-induced gap and the inner disk within 1 AU has originated largely from that part of the disk which lies out beyond the planetary system. This gas may be substantially depleted of small dust because of significant grain growth at large radii, combined with filtration of dust at the outer edge of the gap (Paardekooper \& Mellema 2006; Rice \etal 2006). Furthermore, small dust particles may manage to pass through the gap, but they can quickly coagulate to big dust particles (Dullemond \& Dominik 2005). In this scenario, gas which accretes through the system at late times is strongly depleted of dust, can sustain a significant accretion rate onto the central star, and can provide a dust-rich wall at the outer edge of the gap required by SED modeling. We will explore this hypothesis in a forthcoming paper in which we include the effects of dust filtration. \subsection{Numerical limitations} Our simulations are highly simplified, and this needs to be borne in mind when interpreting our results. The resolution we adopt is insufficient to resolve the gas flow within the planet Hill sphere, and so we are forced to adopt a simplified approach to simulating gas accretion onto the planet, instead of simulating the accretion process directly. As such the detailed evolution of the gas flow in and around the planet Hill sphere may not be modeled with a high degree of fidelity in our simulations. The simulations are two dimensional, which may not be a good approximation early on when the planet masses are low, but improves as the planet masses become large and gap formation occurs. We have neglected a detailed treatment of the gas thermodynamics, adopting instead a locally isothermal equation of state. A proper treatment would allow the local disk temperature to be determined by a balance between viscous and stellar heating, and radiative cooling, and this might affect details of the gap structure due to the changing optical depth of the gas and its thermal evolution there (although viscous heating is generally much smaller than stellar heating). Ayliffe \& Bate (2010) suggest radiation due to circumplanetary disk accretion tends to suppress the spiral shocks and leads to a shallower gap (their Figure 18), which only increases the need for dust depletion within the gaps. Finally, we do not explicitly simulate the MHD turbulence that is believed to provide the effective viscous stresses in protoplanetary disks (Balbus \& Hawley 1991). But it appears from previous simulations that adopting the usual $\alpha$ prescription gives results in broad agreement with MHD simulations when considering gap formation and giant planet migration (Nelson \& Papaloizou 2003; Winters, Balbus \& Hawley 2003; Papaloizou, Nelson \& Snellgrove 2004). As such, we do not expect that the qualitative nature of our results have been compromised by the neglect of the above physical processes.
| 10 | 12 |
1012.4395
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1012
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1012.1323_arXiv.txt
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{The abundance patterns of metal-poor stars provide us a wealth of chemical information about various stages of cosmic chemical evolution. In particular, these stars allow us to study the formation and evolution of the elements, and the involved nucleosynthesis processes. This knowledge is invaluable for our understanding of the nature and condition of the early Universe, and the associated processes of early star- and galaxy formation. This proceeding summarizes the astrophysical topics and questions that can be addressed with metal-poor stars. For the full version of the review, the reader is referred to Frebel 2010.} \FullConference{11th Symposium on Nuclei in the Cosmos, NIC XI\\ July 19-23, 2010\\ Heidelberg, Germany} \begin{document}
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After the Big Bang, the first stars that formed from the pristine gas were very massive, of the order of 100\,M$_{\odot}$ (e.g., \citealt{bromm01, yoshida08}). After a very short life time these co-called Population\,III stars exploded as supernovae, which then provided the first metals to the still primordial interstellar medium. All subsequent generations of stars, Pop\,II, formed from chemically enriched material. The most metal-poor stars are the earliest and most extreme Population\,II objects and belong to the stellar generations that formed from the non-zero metallicity gas left behind by the first stars. In their atmospheres these old objects preserve details of the chemical composition of their birth gas cloud. They thus provide stellar archaeological evidence of the earliest times of the Universe. In particular, the chemical abundance patterns provide information about the formation and evolution of the elements and the involved nucleosynthesis processes and sites. By extension, metal-poor stars provide constraints on the nature of the first stars, the initial mass function, and the chemical yields of first/early SNe. This knowledge is invaluable for our understanding of the cosmic chemical evolution and the onset of star- and galaxy formation processes including the formation of the galactic halo. In summary, galactic metal-poor stars are the local equivalent of the high-redshift Universe, enabling observational constraints on the nature of the first stars and supernovae, and more generally, on various theoretical works on the early Universe. Due to their low masses ($\sim0.8$\,M$_{\odot}$) metal-poor stars have extremely long lifetimes that exceed the current age of the Universe of $\sim14$\,Gyr \citep{WMAP}. Hence, these stellar ``fossils'' of the early Universe are still observable. However, the most metal-poor stars (e.g., stars with\footnote{The main indicator used to determine stellar metallicity is the iron abundance, [Fe/H], which is defined as \mbox{[A/B]}$ = \log_{10}(N_{\rm A}/N_{\rm B})_\star - \log_{10}(N_{\rm A}/N_{\rm B})_\odot$ for the number N of atoms of elements A and B, and $\odot$ refers to the Sun. With few exceptions, [Fe/H] traces the overall metallicity of the objects fairly well.} $\mbox{[Fe/H]}<-5.0$; \citealt{HE1327_Nature}) are extremely rare \citep{schoerck}, and hence difficult to identify. The most promising way forward is to survey large volumes far out into the Galactic halo. Past surveys include the HK survey and the Hamburg/ESO survey \citep{ARAA} which have been very successful in producing large samples of extremely metal-poor stars (with $\mbox{[Fe/H]}<-3.0$). It was also shown that there are many different types of abundance patterns that arise as the result of specific nucleosynthesis processes. The fact that the nuclear physics details of these processes can be probed with the help of stars, means that stellar astrophysics becomes ``nuclear astrophysics''. This is a very complementary approach to experimental nuclear physics that is often limited in its attempts to create the most exotic nuclei or extreme processes, like the r-process, in the laboratory. Over the past few years, a number of extensive reviews have been published on the various roles and applications of metal-poor stars to different astrophysical topics. \citet{ARAA} report on the discovery history, search techniques and results, and the different abundance ``classes'' of metal-poor stars. \citet{sneden_araa} reviewed the evolution of neutron-capture elements in the Galaxy, including the s- and r-process stars. Recent advances regarding the stellar contents of different types of dwarf galaxies are presented in \citet{tolstoy_araa} and \citet{koch_biermann}. Finally, \citet{frebel10} summarized the role of metal-poor stars in the cosmological context, and how early star- and galaxy evolution can be studied with them. This proceeding is based on \citet{frebel10}, and thus only briefly summarizes a few specific aspects of what is described in more detail in the above review. \begin{figure*} [!] \begin{center} \includegraphics[width=13cm,clip=true,bbllx=65,bblly=423,bburx=528,bbury=655]{spectral_comparison_with_lines.ps} \caption{\label{spec_comp} Spectral comparison of stars in the main-sequence turn-off region with different metallicities. Several atomic absorption lines are marked. The variations in line strength reflect the different metallicities. From top to bottom: Sun with $\mbox{[Fe/H]}=0.0$, G66-30 with $\mbox{[Fe/H]}=-1.6$ \citep{norris_emp3}, G64-12 with $\mbox{[Fe/H]}=-3.2$ \citep{Aokihe1327}, and HE1327-2326 with $\mbox{[Fe/H]}=-5.4$ \citep{HE1327_Nature}. Reproduced from \citet{frebel10}.} \end{center} \end{figure*}
| 10 | 12 |
1012.1323
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1012
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1012.3326_arXiv.txt
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{In this review, I provide an overview of theoretical aspects related to the evolution of galaxies as a function of environment. I discuss the main physical processes at play, their characteristic time-scales and environmental dependency, and comment on their treatment in the framework of hierarchical galaxy formation models. I briefly summarize recent results and the main open issues.}
| 10 | 12 |
1012.3326
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1012
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1012.0325_arXiv.txt
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{In--correlator techniques offer the possibility of identifying and/or excising radio frequency interference (RFI) from interferometric observations at much higher time and/or frequency resolution than is generally possible with the final visibility dataset. Due to the considerable computational requirements of the correlation procedure, cross--correlators have most commonly been implemented using high--speed digital signal processing boards, which typically require long development times and are difficult to alter once complete. ``Software" correlators, on the other hand, make use of commodity server machines and a correlation algorithm coded in a high--level language. They are inherently much more flexible and can be developed -- and modified -- much more rapidly than purpose--built ``hardware" correlators. Software correlators are thus a natural choice for testing new RFI detection and mitigation techniques for interferometers. The ease with which software correlators can be adapted to test RFI detection algorithms is demonstrated by the addition of kurtosis detection and plotting to the widely used DiFX software correlator, which highlights previously unknown short--duration RFI at the Hancock VLBA station.} \FullConference{RFI mitigation workshop - RFI2010,\\ March 29-31, 2010\\ Groningen, the Netherlands} \begin{document}
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\label{sec:conclusions} A simple kurtosis estimator has been successfully implemented in the DiFX software correlator. This initial foray into high time resolution in--correlator RFI detection was made with minimal effort, requiring only an afternoon of code development and testing, and provided useful information on previously unknown RFI at the Hancock VLBA station. The present simplistic implementation should soon be updated to the unbiased SK estimator, and a thorough analysis of the effects of the coarse front--end quantisation should also be undertaken. Ultimately, it may be possible to generalise the SK concept to the crosscorrelation outputs of the correlator. With or without these improvements, however, the calculated kurtosis could be used to flag data on the fly (on timescales shorter than a full integration) or to produce a flag table which could be inspected and applied later in offline analysis. The latter case, whilst allowing the flexibility to deselect some flags, loses the ability to save data which has been affected for only a short subset of a given integration. In either case, this approach is still an excision procedure which ultimately results in data loss. "Kurtosis blanking" and "kurtosis flagging" will be added as DiFX features in the near feature, and the usability and presentation of the saved kurtosis results will be improved. Development of other forms of in--correlator RFI mitigation are planned for the DiFX correlator. One example is the rejection of bright sources outside the desired array field of view, which can be achieved through field of view shaping using tapered time integrations (see \cite{lonsdale04a}). This will allow a determination of the relative trade-off between accuracy of the tapering function and the performance overhead incurred. The success of the kurtosis development detailed in this work indicates that these approaches can be rapidly and effectively implemented.
| 10 | 12 |
1012.0325
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1012
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1012.0439_arXiv.txt
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The nature of the broad-band spectra of supermassive accreting black holes in active galactic nuclei (AGNs) is still unknown. The hard X-ray spectra of Seyferts as well as of Galactic stellar-mass black holes (GBHs) are well represented by thermal Comptonization, but the origin of the seed photons is less certain. The MeV tails observed in GBHs provide evidence in favour of non-thermal electron tails and it is possible that such electrons are also present in the X-ray emitting regions of AGNs. Using simulations with the kinetic code that self-consistently models electron and photon distributions, we find that the power-law-like X-ray spectra in AGNs can be explained in terms of the synchrotron self-Compton radiation of hybrid thermal/non-thermal electrons, similarly to the hard/low state of GBHs. Under a very broad range of parameters the model predicts a rather narrow distribution of photon spectral slopes consistent with that observed from LINERs and Seyferts at luminosities less than 3 per cent of the Eddington luminosity. The entire infrared to X-ray spectrum of these objects can be described in terms of our model, suggesting a tight correlation between the two energy bands. We show that the recently found correlation between slope and the Eddington ratio at higher luminosities can be described by the increasing fraction of disc photons in the emitting region, which may be associated with the decreasing inner radius of the optically thick accretion disc. The increasing flux of soft photons is also responsible for the transformation of the electron distribution from nearly thermal to almost completely non-thermal. The softer X-ray spectra observed in narrow-line Seyfert galaxies may correspond to non-thermal Comptonization of the disc photons, predicting that no cutoff should be observed up to MeV energies in these sources, similarly to the soft-state GBHs.
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Accreting radio-quiet supermassive black holes (SMBHs) residing in the centres of quasars, Seyfert galaxies, narrow-line Seyfert 1 galaxies (NLSy1) and some fraction of low-ionization nuclear emission-line regions (LINERs) in many respects are analogous to the Galactic stellar-mass black holes (GBHs) in X-ray binaries. In the X-ray/soft $\gamma$-ray band, the spectra of Seyferts can be represented by a sum of a power-law-like continuum which cuts off at a few hundred keV and a reflection component with the iron fluorescent $\rm K{\rm\alpha}$ line believed to be produced by reprocessing of the intrinsic power-law by cold opaque matter, probably the accretion disc \citep{NanPoun94}. The 2--10 keV intrinsic spectral energy slope (defined as $F_E \propto E^{-\alpha}$) of the power-law $\alpha \sim$ 0.9--1.0, which is ubiquitously found in Seyferts, is somewhat larger than what is measured in GBHs in their hard state $\alpha \sim$ 0.6--0.8 \citep*[e.g.][]{ZLS99}. NLSy1 having softer X-ray spectra than Seyferts, probably represent a state with higher accretion rate (in Eddington units) similar to the soft/very high state of GBHs \citep*{PDO95}. The seemingly similar X-ray slopes of Seyferts triggered the efforts to find a physical model which would explain such spectral stability. The non-thermal models became popular \citep[see][ for a review]{Sve94}. These invoke injection of high-energy leptons into the system with subsequent Compton cooling by accretion disc photons and photon-photon pair productions initiating pair cascades. Saturated cascades produce intrinsic spectra with $\alpha \approx0.9$ \citep{ZL85,Sve87}, which becomes consistent with the X-ray spectra of Seyferts after accounting for the hardening due to Compton reflection \citep{GF91}. The pair-cascade models, on the other hand, predict a strong tail above 300 keV and an annihilation line which have never been observed. The \textit{CGRO}/OSSE observations of the brightest Seyfert 1 galaxy NGC 4151 constrain the fraction of energy going to non-thermal injection to be less than 50\%, while the rest of the power going to thermal heating \citep*{ZLM93,ZJM96}. The existing upper limits on the average flux of Seyferts above 100 keV are compatible with the presence of weak non-thermal tails \citep[e.g.][]{Gon96, Johns97}. The more recent \textit{INTEGRAL} observations \citep{LZ10} have not improved those constraints. Non-detectable high-energy tail in Seyfert galaxies can also be interpreted in terms of pure thermal models, where the power is equally shared among all the thermal particles. The spectra indeed can be well described by Comptonization on \textit{thermal} 50--150 keV electrons \citep[see e.g.][]{ZJP97,P98,Zdz99}. The stability of the spectral slopes can be interpreted as the evidence of the radiative feedback between the hot X-ray emitting plasma and the cool accretion disc, which is assumed to be a sole source of the seed photons for Comptonization. The slab-corona model, where the hot plasma sandwiches the accretion disc \citep[see e.g.][]{HM91, HM93}, predicts too soft spectra due to large flux of the reprocessed UV photons \citep{SPS95}. \citet*{MDM05} showed that the hard spectrum can be achieved if one assumes large ionization parameter of the disc, but in this case the model fails on the predicted reflection properties. Localized active regions atop of a cool disc \citep*{HMG94,SPS95,PS96,Sve96} are more photon starved and can produce much harder spectra in agreement with observations. However, in this case it is difficult to understand why there is a preferential slope, as it is a strong function of the separation of the active region from the disc. The detection of the MeV tails in the spectra of the GBH Cyg X-1 in both hard and soft states \citep{McConnell02} imply the presence of a significant non-thermal component in the electron distribution. The power-law looking spectra extending to the MeV energies in other soft-state GBHs \citep{Grove98,ZGP01} are consistent with being produced by non-thermal Comptonization \citep{PC98,P98}. By analogy, the electrons in Seyferts also might have a non-thermal population. More probably in both types of sources, the electron distribution is hybrid in both hard and soft states \citep{PC98,P98,coppi99}, with the non-thermal fraction increasing for softer spectra. However, non-thermal particles could be spatially separated from the thermal ones. The presence of the non-thermal particles, even if they are not energetically dominant, has a strong impact on the emitted spectrum, because the synchrotron emission can increase by orders of magnitude even if only 1 per cent of electron energy is in the power-law tail \citep{WZ01}. As most of the synchrotron emission is self-absorbed, this process plays an important role in shaping the electron distribution by thermalizing particles via emission and absorption of synchrotron photons, the so called synchrotron boiler \citep*{GGS88,GHS98}. It was recently shown that under the conditions of GBHs, Coulomb collisions and the synchrotron boiler efficiently thermalize electrons producing thermal population at low energies even if the electrons were originally non-thermal \citep{VP08,PV09, MB09}. The hard spectrum of GBHs can be fully accounted for by the synchrotron self-Compton (SSC) mechanism in the resulting hybrid electrons. In this picture the disc photons are not required to interact with the hot X-ray emitting plasma at all. This mechanism also gives a rather stable X-ray slopes for a large range of parameters, relieving the need for the feedback between the disc and the active region. As many of the radiative processes depend only on the compactness of the source, but not on its luminosity or size separately, one may think that the spectra of Seyferts can also be described in terms of hybrid Comptonization models. However, not all processes can be scaled away. The aim of this paper is to study in details the spectral formation in relativistic hybrid plasmas in the vicinity of a SMBH. One of our goals is to learn how the results scale with the mass of the compact object. We also test the role of bremsstrahlung in SSC models and, finally, we compare the model with the data on Seyferts and NLSy1s.
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We have studied the spectral formation in hot non-thermal plasmas under the conditions relevant to the vicinities of SMBHs in AGNs, using a self-consistent SSC model. We show that the SSC model can reproduce the exponential cut-off at energies above $\sim100$ keV and the power-law X-ray continuum with energy spectral indices $\alpha \sim 0.8-1.0$ observed in low-luminosity Seyferts and LINERs. The entire infrared/optical to X-ray spectrum in this objects can be produced by the SSC mechanism, suggesting a strong correlation between the two energy bands. No specific conditions are required to stabilize the slope for a wide range of parameters (magnetization, optical depth, injection slope, Eddington ratio and black hole mass). We have demonstrated the difference in scaling of the SSC model and the two-phase disc-corona models. While in the latter case one would expect systematically softer spectra from objects with higher masses, in the case of SSC, there exists a limiting slope (when the mass of the central object tends to infinity) $\alpha = 5/(s+2)$, which depends only on the slope of the tail of the electron distribution. Comptonization spectra can get systematically softer for objects of higher masses, if the electron injection function is sufficiently hard. We have tested the role of bremsstrahlung emission in SSC models. We found that even though the process is not energetically significant, in certain cases, it leads to the hardening of the power-law above a few keV and makes the overall X-ray spectrum concave. We have also studied the role of additional soft blackbody photons from the accretion disc. With the increasing amount of soft photons, the equilibrium Maxwellian temperature drops and the X-ray spectrum softens. The resulting spectra look similar to the GBHs in their soft state, which in the case of SMBHs are likely to be represented by NLSy1s and quasars. We have found that it is possible to explain the spectral slope -- flux correlation, widely discussed in the literature, by parametrizing the fraction of the disc photons in the medium as a power-law function of the Eddington ratio. This implies, that the inner radius of the truncated accretion disc is a strong function of the accretion rate.
| 10 | 12 |
1012.0439
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1012
|
1012.2116_arXiv.txt
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We discuss the properties of an object in the Subaru Deep Field (SDF) classified as a galaxy in on-line data bases and revealed on the Subaru images as a genuine polar-ring galaxy (PRG) candidate. We analyse available photometric data and conclude that this object consists of a $\gtrsim5$ Gyr old early-type central body surrounded by a faint, narrow inner ring tilted at a $\sim25^\circ$ angle relative to the polar axis of the host galaxy. The halo surrounding the main stellar body exhibits a diversity of spatially extended stellar features of low surface brightness, including a faint asymmetric stellar cloud and two prominent loops. These faint features, together with the unperturbed morphology of the central host, are clear signs of a recent coalescence of two highly unequal mass galaxies, most likely a pre-existing early-type galaxy and a close-by gas-rich dwarf galaxy. The presumed stellar remnants observed near the edges of the ring, including possibly the surviving captured companion itself, indicate that the merger is still taking place.
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Mutual interactions between galaxies are probably one of the main processes leading to the formation of presently observed galaxies. The discovery of shells and ripples around elliptical galaxies (Kormendy \& Djorgovski 1989; Tal et al.\ 2009) and of extended stellar tidal features around normal disc galaxies (Mart\'{i}nez-Delgado et al.\ 2010) are among the most compelling evidence for this scenario. In more extreme and rare cases the remnants of the interacting galaxies are not mixed in one smooth object, but some of the material settles quasi-statically into the equatorial plane of the main body, forming a polar-ring galaxy (PRG; Whitmore et al.\ 1990). Alternatively, simulations show that the polar structure can form as the result of a ``secondary event'' during the evolution of the galaxy, involving the accretion of gas from a close, gas-rich dwarf companion (Reshetnikov \& Sotnikova 1997). However, both formation scenarios are difficult to validate observationally since they predict the formation of a polar ring on short time scales of less than 1-2 Gyr (Bekki 1998; Bournaud \& Combes 2003). Since only about 0.5 per cent of all early-type galaxies have observable polar rings (Whitmore et al.\ 1990), on-going galaxy interactions that form a PRG are probably extremely rare events. A galaxy tidally disrupted while still in orbit around a much more luminous and massive companion should leave distinctive tidal stellar features before merging with the central galaxy or mixing into a seemingly smooth component. These stellar debris can extend well beyond the central galaxy and mostly occur at a very low surface brightness level, making it difficult to detect in sky survey images. In this paper we present an interpretation of existing observational data of the object SDSS J132533.22+272246.7 (hereafter referred to as RG) identified as a ring galaxy during a supernova survey in the Subaru Deep Field (SDF; Graur et al., in preparation). The object is centred at $\alpha=13^{h}25^{m}33.2^{s}, \delta=+27^{\circ}22\arcmin46.7\arcsec$ (J2000) at a redshift of $z=0.061$, which puts it at a distance of about 250 Mpc, assuming $H_0=73$ km sec$^{-1}$ Mpc$^{-1}$. This paper is organized as follows: Section \ref{obs} gives a description of the observations and data reduction; in Section \ref{results} we present the observational results along with our interpretation of the object and the main conclusions are summarized in Section \ref{conc}. \begin{figure*} \caption{Panel a: SDSS $g$-band image of RG. Panels b-d: Combined Subaru image displayed in different contrast levels to enhance structural features. Stellar debris and loops are marked by arrows. \label{RG}} \vspace{2mm} \begin{center} \begin{tabular}{cc} \includegraphics[trim=0.5cm 0cm 0.5cm 0.5cm,width=6cm]{acontrast.eps} & \includegraphics[trim=0.5cm 0cm 0.5cm 0.5cm,width=6cm]{acontrast1.eps} \\ \vspace{5cm} \includegraphics[trim=0.5cm 0cm 0.5cm 0.5cm,width=6cm]{acontrast2.eps} & \includegraphics[trim=0.5cm 0cm 0.5cm 0.5cm,width=6cm]{acontrast3.eps} \vspace{-5cm} \end{tabular} \end{center} \end{figure*}
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\label{conc} We present a photometric study of a candidate PRG in the SDF. The object has a complex structure consisting of:\\ 1. A central host which we identify as a $\gtrsim5$ Gyr old early-type galaxy with an exponential bulge and a nearly edge-on disc;\\ 2. A faint off-centre narrow ring at a $\sim25^\circ$ angle to the polar axis of the host; \\ 3. A low surface brightness stellar halo extending asymmetrically on both sides of the ring;\\ 4. At least two disrupted stellar debris in intermediate position angles between the ring and the polar axis of the host;\\ 5. Two outer loops on both sides in the direction of the stellar debris.\\ Our morphological study suggests that this galaxy belongs to the group of bulge-dominated PRGs and that the ring was probably formed around the pre-existing host during a merger with a gas-rich dwarf companion. This event probably took place less than 1-2 Gyr ago, since it appears that the captured material has not yet fully dispersed into the halo or settled into the polar structure.
| 10 | 12 |
1012.2116
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1012
|
1012.2266_arXiv.txt
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{For decades now, scientific data volumes have experienced relentless, exponential growth. As a result, legacy astronomical data formats are straining under a burden not conceived when these formats were first introduced. With future astronomical projects ensuring this trend, ASTRON and the LOFAR project are exploring the use of the Hierarchical Data Format, version 5 (HDF5), for LOFAR radio data encapsulation. Most of LOFAR's standard data products will be stored using the HDF5 format. In addition, HDF5 analogues for traditional radio data structures such as visibility data and spectral image cubes are also being developed. The HDF5 libraries allow for the construction of distributed, entirely unbounded files. The nature of the HDF5 format further provides the ability to custom design a data encapsulation format, specifying hierarchies, content and attributes. The LOFAR project has designed several data formats that will accommodate and house all LOFAR data products, the primary styles and kinds of which are presented in this paper. With proper development and support, it is hoped that these data formats will be adopted by other astronomical projects as they, too, attempt to grapple with a future filled with mountains of data.} \FullConference{ISKAF2010 Science Meeting - ISKAF2010\\ June 10-14, 2010\\ Assen, the Netherlands} \ShortTitle{LOFAR and HDF5} \FullConference{ISKAF2010 Science Meeting - ISKAF2010\\ June 10-14, 2010\\ Assen, the Netherlands} \begin{document}
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The promising advent of the LOFAR telescope's operational epoch holds forth both great scientifiic potential and challenges to current and legacy information technologies: volume and complexity of the data will continue to push the envelope of commonly used data protocols. Recognizing that this envelope is already strained, the LOFAR project has embarked on an ambitious project to design and define a set of radio data standard formats that are capable of encapsulating the full spectrum of, not just LOFAR data products, but astronomical radio data in general. It is with this ambition in mind that the LOFAR data formats group have been developing these format specifications and associated software infrastructure, an effort now ongoing for over two years. It was determined that HDF5 would be a robust, viable data framework that can handle the size, scope, diversity, distributed nature and parallel I/O processing requirements of LOFAR data. This work also has potential use beyond the radio community. New large scale optical telescopes such as the LSST are also investigating the viability of using HDF5. Furthermore, the 20 year history of HDF and its continuing use by NASA's earth orbiting/observing missions ensure broad, ongoing use and support. In addition to the format descriptions themselves, the LOFAR project is currently developing a set of software tools for creating and working with these formats. The Data Access Library (DAL) in C++, along with an associated Python interface (pyDAL), are designed to allow for the easy construction and manipulation of these data formats. There are also a number of tools already available to read and visualize HDF5 files, such as HDFView, ViSiT, PyTables, h5py and IDL.
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In order that adoption of HDF5 in astronomy prove useful in the real world, the LOFAR project is committing resources to help develop the next generation of astronomical tools for, not only LOFAR data, but other SKA pathfinder projects and, more broadly, astronomical radio data in general. Though our immediate goal is, of course, to meet LOFAR's many and varied scientific requirements, we hope these HDF5-based formats will be more broadly useful to our colleagues in the radio community for their data products as well. The large effort is the development of the Data Access Library (DAL), which will ultimately provide interfaces through FITS, the Casa/AIPS++ Measurement Set and LOFAR HDF5 formats. A python interface to the DAL, pyDAL, is pending development of the DAL. All LOFAR products will be accessible through DAL and pyDAL tools. Future work involves developing an interface for DS9 and HDF5 LOFAR data files --- this will allow users to open and examine LOFAR data with the de facto standard astronomical image viewer. Ultimately we would like to see these formats grow into a true set of standards for radio data that can meet the demands of the next generation of radio observatories. Such standards are something sorely lacking in the radio community at present and something we will certainly need as we move into the SKA era. The large effort by LOFAR to design an HDF5 radio data standard is driven in great part by consideration that there are no effective standards for astronomical radio data. And, as indicated earlier, the expected data volumes produced by LOFAR will, in many cases, swamp currently employed file technologies. It is this reality that has led LOFAR on this work, and it is hoped that other institutes and telescope projects will join this effort toward building a radio data standard, one essential to a cooperative future for radio astronomy.
| 10 | 12 |
1012.2266
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1012
|
1012.0780_arXiv.txt
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We show that, in a model of modified gravity based on the spectral action functional, there is a nontrivial coupling between cosmic topology and inflation, in the sense that the shape of the possible slow-roll inflation potentials obtained in the model from the nonperturbative form of the spectral action are sensitive not only to the geometry (flat or positively curved) of the universe, but also to the different possible non-simply connected topologies. We show this by explicitly computing the nonperturbative spectral action for some candidate flat cosmic topologies given by Bieberbach manifolds and showing that the resulting inflation potential differs from that of the flat torus by a multiplicative factor, similarly to what happens in the case of the spectral action of the spherical forms in relation to the case of the $3$-sphere. We then show that, while the slow-roll parameters differ between the spherical and flat manifolds but do not distinguish different topologies within each class, the power spectra detect the different scalings of the slow-roll potential and therefore distinguish between the various topologies, both in the spherical and in the flat case.
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Noncommutative cosmology is a new and rapidly developing area of research, which aims at building cosmological models based on a ``modified gravity" action functional which arises naturally in the context of noncommutative geometry, the {\em spectral action} functional of \cite{CC}. As we discuss more in detail in \S \ref{NCcosmSec} below, this functional recovers the usual Einstein--Hilbert action, with additional terms, such as a conformal gravity, Weyl curvature term. It also has the advantage of allowing for interesting couplings of gravity to matter, when extended from manifolds to ``almost commutative geometries" as in \cite{CoSM} and later models \cite{CCM}, \cite{BroSu}. Thus, this approach makes it possible to recover from the same spectral action functional, in addition to the gravitational terms, the full Lagrangian of various particle physics models, ranging from the Minimal Standard Model of \cite{CoSM}, to the extension with right handed neutrinos and Majorana mass terms of \cite{CCM}, and to supersymmetric QCD as in \cite{BroSu}. The study of cosmological models derived from the spectral action gave rise to early universe models as in \cite{MaPie} and \cite{KoMa}, which present various possible inflation scenarios, as well as effects on primordial black holes evaporation and gravitational wave propagation. Effects on gravitational waves, as well as inflation scenarios coming from the spectral action functional, were also recently studied in \cite{NeOSa1}, \cite{NeOSa2}, \cite{NeSa}. Our previous work \cite{MaPieTeh} showed that, when one considers the nonperturbative form of the spectral action, as in \cite{CC2}, one obtains a slow-roll potential for inflation. We compared some of the more likely candidates for cosmic topologies (the quaternionic and dodecahedral cosmology, and the flat tori) and we showed that, in the spherical cases (quaternionic and dodecahedral), the nonperturbative spectral action is just a multiple of the spectral action of the sphere $S^3$, and consequently the inflation potential only differs from the one of the sphere case by a constant scaling factor, which cancels out in the computation of the slow-roll parameters, which are therefore the same as in the case of a simply connected topology and do not distinguish the different cosmic topologies with the same spherical geometry. This result for spherical space forms was further confirmed and extended in \cite{Teh}, where the nonperturbative spectral action is computed explicitly for {\em all} the spherical space forms and it is shown to be always a multiple of the spectral action of $S^3$, with a proportionality factor that depends explicitly on the 3-manifold. Thus, different candidate cosmic topologies with the same positively curved geometry yield the same values of the slow-roll parameters and of the power-law indices and tensor-to-scalar ratio, which are computed from these parameters. In \cite{MaPieTeh}, however, we showed that the inflation potential obtained from the nonperturbative spectral action is different in the case of the flat tori, and not just by a scalar dilation factor. Thus, we know already that the possible inflation scenarios in noncommutative cosmology depend on the underlying geometry (flat or positively curved) of the universe, and the slow--roll parameters are different for these two classes. The slow-roll parameters alone only distinguish, in our model, between the flat and spherical geometries but not between different topologies within each class. However, in the present paper we show that, when one looks at the amplitudes for the power spectra for density perturbations and gravitational waves (scalar and tensor perturbations), these detect the different scaling factors in the slow-roll potentials we obtain for the different spherical and flat topologies, hence we obtain genuinely different inflation scenarios for different cosmic topologies. We achieve this result by relying on the computations of the nonperturbative spectral action, which in the spherical cases are obtained in \cite{MaPieTeh} and \cite{Teh}, and by deriving in this paper the analogous explicit computation of the nonperturbative spectral action for the flat Bieberbach manifolds. A similar computation of the spectral action for Bieberbach manifolds was simultaneously independently obtained by Piotr Olczykowski and Andrzej Sitarz in \cite{OlSi}. Thus, the main conclusion of this paper is that {\em a modified gravity model based on the spectral action functional predicts a coupling between cosmic topology and inflation potential, with different scalings in the power spectra that distinguish between different topologies, and slow-roll parameters that distinguish between the spherical and flat cases.} The paper is organized as follows. We first describe in \S \ref{InflGeomSec} the broader context in which the problem we consider here falls, namely the cosmological results relating inflation, the geometry of the universe, and the background radiation, and the problem of cosmic topology. We then review briefly in \S \ref{NCcosmSec} the use of the spectral action as a modified gravity functional and the important distinction between its asymptotic expansion at large energies and the nonperturbative form given in terms of Dirac spectra. In \S \ref{BiebSec} we present the main mathematical result of this paper, which gives an explicit calculation of the nonperturbative spectral action for certain Bieberbach manifolds, using the Dirac spectra of \cite{Pfa} and a Poisson summation technique similar to that introduced in \cite{CC2}, and used in \cite{MaPieTeh} and \cite{Teh}. Finally, in \S \ref{InflPotSec} we compare the resulting slow-roll inflation potentials, power spectra for density perturbations and slow-roll parameters, for all the different possible cosmic topologies. \subsection{Inflation, geometry, and topology}\label{InflGeomSec} It is well known that the mechanism of cosmic inflation, first proposed by Alan Guth and Andrei Linde, naturally leads to a flat or almost flat geometry of the universe (see for instance \S 1.7 of \cite{Linde}). It was then shown in \cite{KaSpeSu} that the geometry of the universe can be read in the cosmic microwave background radiation (CMB), by showing that the anisotropies of the CMB depend primarily upon the geometry of the universe (flat, positively or negatively curved) and that this information can be detected through the fact that the location of the first Doppler peak changes for different values of the curvature and is largely unaffected by other parameters. This theoretical result made it possible to devise an observational test that could confirm the inflationary theory and its prediction for a flat or nearly flat geometry. The experimental confirmation of the nearly flat geometry of the universe came in \cite{dBL} through the Boomerang experiment. Thus, the geometry of the universe leaves a measurable trace in the CMB, and measurements confirmed the flat geometry predicted by inflationary models. The cosmic topology problem instead concentrates not on the question about the curvature and the geometry of the universe, but on the possible existence, for a given geometry, of a non-simply connected topology, that is, of whether the spatial sections of spacetime can be compact 3-manifolds which are either quotients of the 3-sphere (spherical space forms) in the positively curved case, quotients of 3-dimensional Euclidean space (flat tori or Bieberbach manifolds) in the flat case, or quotients of the 3-dimensional hyperbolic space (hyperbolic 3-manifolds) in the negatively curved space. A general introduction to the problem of cosmic topology is given in \cite{LaLu}. Since the cosmological observations prefer a flat or nearly flat positively curved geometry to a nearly flat negatively curved geometry (see \cite{dBL}, \cite{WDP}), most of the work in trying to identify the most likely candidates for a non-trivial cosmic topology concentrate on the flat spaces and the spherical space forms. Various methods have been devised to try to detect signatures of cosmic topology in the CMB, in particular through a detailed analysis of simulated CMB skies for various candidate cosmic topologies (see \cite{RWULL} for the flat cases). It is believed that perhaps some puzzling features of the CMB such as the very low quadrupole, the very planar octupole, and the quadrupole--octupole alignment may find an explanation in the possible presence of a non-simply connected topology, but no conclusive results to that effect have yet been obtained. The recent results of \cite{MaPieTeh} show that a modified gravity model based on the spectral action functional imposes constraints on the form of the possible inflation slow-roll potentials, which depend on the geometry and topology of the universe, as shown in \cite{MaPieTeh}. While the resulting slow-roll parameters and spectral index and tensor-to-scalar ratio distinguish the even very slightly positively curved case from the flat case, these parameters alone do not distinguish between the different spherical topologies, as shown in \cite{Teh}. As we show in this paper, the situation is similar for the flat manifolds: these same parameters alone do not distinguish between the various Bieberbach manifolds (quotients of the flat torus), but they do distinguish these from the spherical quotients. However, if one considers, in addition to the slow-roll parameters, also the power spectra for the density fluctuations, one can see that, in our model based on the spectral action as a modified gravity functional, the resulting slow-roll potentials give different power spectra that distinguish between all the different topologies. \subsection{Slow-roll potential and power spectra of fluctuations} We first need to recall here some well known facts about slow-roll inflation potentials, slow-roll parameters, and the power spectra for density perturbations and gravitational waves. We refer the reader to \cite{SKamCoo} and to \cite{Lidsey}, \cite{StLy}, as well as to the survey of inflationary cosmology \cite{Baumann}. Consider an expanding universe, which is topologically a cylinder $Y\times \R$, for a $3$-manifold $Y$, with a Lorentzian metric of the usual Friedmann form \begin{equation}\label{Friedmann} ds^2 = - dt^2 + a(t)^2 ds_Y^2 \end{equation} where $ds_Y^2= g_{ij} dx^i dx^j$ is the Riemannian metric on the $3$-manifold $Y$. In models of inflation based on a single scalar field slow-roll potential $V(\phi)$, the dynamics of the scale factor $a(t)$ in the Friedmann metric \eqref{Friedmann} is related to the scalar field dynamics through the acceleration equation \begin{equation}\label{accel} \frac{\ddot{a}}{a} = H^2 (1-\epsilon), \end{equation} where $H$ is the Hubble parameter, which is related to the scalar field and the inflation potential by \begin{equation}\label{Hubble} H^2 = \frac{1}{3} \left( \frac{1}{2} \dot{\phi}^2 + V(\phi)\right), \ \ \text{ and } \ \ \ddot{\phi} + 3 H \dot{\phi} + V^\prime(\phi) =0, \end{equation} and $\epsilon$ is the slow-roll parameter, which depends on the potential $V$ as described in \eqref{slowrollparam} below, see \cite{Baumann} for more details. It is customary to decompose perturbations of the metric $ds^2$ of \eqref{Friedmann} into scalar and tensor perturbations, which correspond, respectively, to density fluctuations and gravitational waves. One typically neglects the remaining vector components of the perturbation, assuming that these are not generated by inflation and decay with the expansion of the universe, see \S 9.2 of \cite{Baumann}. Thus, one writes scalar and tensor perturbations in the form \begin{equation}\label{scaltensper} ds^2 = - (1+2\Phi) dt^2 + 2 a(t) \, dB \, dt + a(t)^2 ((1-2\Psi) g_{ij}+2\Delta E + h_{ij}) dx^i dx^j \end{equation} with $dB=\partial_i B dx^i$ and $\Delta E= \partial_i\partial_j E$, and where the $h_{ij}$ give the tensor part of the perturbation, satisfying $\partial^i h_{ij}=0$ and $h^i_i=0$. The tensor perturbations $h_{ij}$ have two polarization modes, which correspond to the two polarizations of the gravitational waves. One considers then the intrinsic curvature perturbation \begin{equation}\label{intrcurv} \cR = \Psi - \frac{H}{\dot{\phi}}\delta \phi, \end{equation} which measures the spatial curvature of a comoving hypersurface, that is, a hypersurface with constant $\phi$. After expanding $\cR$ in Fourier modes in the form \begin{equation}\label{Rfourier} \cR = \int \frac{d^3k}{(2\pi)^{3/2}} \cR_k\, e^{ikx} , \end{equation} one obtains the power spectrum $\cP_s(k)$ for the density fluctuations (scalar perturbations of the metric) from the two-point correlation function, \begin{equation}\label{Ps2point} \langle \cR_k \cR_{k'} \rangle = (2\pi^2)^3\, \cP_s(k)\, \delta^3(k+k') . \end{equation} In the case of a Gaussian distribution, the power spectrum describes the complete statistical information on the perturbations, while the higher order correlations functions contain the information on the possible presence of non-Gaussianity phenomena. The power spectrum $\cP_t(k)$ for the tensor perturbations is similarly obtained by expanding the tensor fluctiations in Fourier modes $h_k$ and computing the two-point correlation function \begin{equation}\label{tens2point} \langle h_k h_{k'} \rangle = (2\pi^2)^3\, \cP_t(k)\, \delta^3(k+k'). \end{equation} See \cite{Baumann}, \S 9.3, and \cite{StLy} for more details. In slow-roll inflation models, the power spectra $\cP_s(k)$ and $\cP_t(k)$ are related to the slow-roll potential $V(\phi)$ through the leading order expression (see \cite{SKamCoo}) \begin{equation}\label{PstV} \cP_s(k) \sim \frac{1}{M_{Pl}^6} \frac{V^3}{(V^\prime)^2} \ \ \text{ and } \ \ \cP_t(k) \sim \frac{V}{M_{Pl}^4}, \end{equation} up to a constant proportionality factor, and with $M_{Pl}$ the Planck mass. Here the potential $V(\phi)$ and its derivative $V^\prime(\phi)$ are to be evaluated at $k= a H$ where the corresponding scale leaves the horizon during inflation. These can be expressed a power law as (\cite{SKamCoo}) \begin{equation}\label{powerlawP} \begin{array}{rl} \cP_s(k) \sim & \cP_s(k_0) \displaystyle{\left(\frac{k}{k_0} \right)^{1 - n_s + \frac{\alpha_s}{2} \log(k/k_0)}} \\[3mm] \cP_t(k) \sim & \cP_t(k_0) \displaystyle{\left(\frac{k}{k_0} \right)^{n_t + \frac{\alpha_t}{2} \log(k/k_0)}} , \end{array} \end{equation} where the spectral parameters $n_s$, $n_t$, $\alpha_s$, and $\alpha_t$ depend on the slow--roll potential in the following way. In the slow-roll approximation, the slow-roll parameters are given by the expressions \begin{equation}\label{slowrollparam} \begin{array}{rl} \epsilon = & \displaystyle{\frac{M_{Pl}^2}{16\pi} \left( \frac{V^\prime}{V} \right)^2} \\[3mm] \eta = & \displaystyle{\frac{M_{Pl}^2}{8\pi} \frac{V^{\prime\prime}}{V}} \\[3mm] \xi = & \displaystyle{\frac{M_{Pl}^4}{64\pi^2} \frac{V^\prime V^{\prime\prime\prime}}{V^2}} \end{array} \end{equation} Notice that we follow here a different convention with respect to the one we used in \cite{MaPieTeh} on the form of the slow-roll parameters. The spectral parameters are then obtained from these as \begin{equation}\label{spectralparam} \begin{array}{rl} n_s \simeq & 1 - 6 \epsilon + 2 \eta \\[2mm] n_t \simeq & - 2 \epsilon \\[2mm] \alpha_s \simeq & 16 \epsilon \eta - 24 \epsilon^2 - 2 \xi \\[2mm] \alpha_t \simeq & 4 \epsilon\eta - 8 \epsilon^2 \end{array} \end{equation} while the tensor-to-scalar ratio is given by \begin{equation}\label{tensorscalar} r = \frac{P_t}{P_s} = 16 \epsilon. \end{equation} {}From the point of view of our model, the following observation will be useful when we compare the slow-roll potentials that we obtain for different cosmic topologies and how they affect the power spectra. \begin{lem}\label{scaleVandP} Given a slow-roll potential $V(\phi)$ and the corresponding power spectra $\cP_s(k)$ and $\cP_t(k)$ as in \eqref{PstV} and \eqref{powerlawP}. If the potential $V(\phi)$ is rescaled by a constant factor $V(\phi) \mapsto \lambda V(\phi)$, then the power spectra $\cP_s(k)$ and $\cP_t(k)$ are also rescaled by the same factor $\lambda >0$, while in the power law \eqref{powerlawP} the exponents are unchanged. \end{lem} \proof This is an immediate consequence of \eqref{PstV}, \eqref{slowrollparam}, \eqref{spectralparam}, and \eqref{powerlawP}. In fact, from \eqref{PstV}, we see that $V(\phi) \mapsto \lambda V(\phi)$ maps $\cP_t \mapsto \lambda \cP_t$, and also $\cP_s \mapsto \lambda \cP_s$, since it transforms $V^3 (V^\prime)^{-2} \mapsto \lambda V^3 (V^\prime)^{-2}$. On the other hand, the expressions $(V^\prime /V)^2$ and $V^{\prime\prime}/V$ and $V^\prime V^{\prime\prime\prime}/V^2$ in the slow-roll parameters \eqref{slowrollparam} are left unchanged by $V \mapsto \lambda V$, so that the slow-roll parameters and all the resulting spectral parameters of \eqref{spectralparam} are unchanged. Thus, the power law \eqref{powerlawP} only changes by a multiplicative factor $\cP_s(k_0)\mapsto \lambda \cP_s(k_0)$ and $\cP_t(k_0)\mapsto \lambda \cP_t(k_0)$, with unchanged exponents. \endproof
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\label{InflPotSec} We have seen in this paper that, in a modified gravity model based on the non-perturbative spectral action functional, different cosmic topologies, either given by spherical space forms or by flat Bieberbach manifolds, leave a signature that can distinguish between the different topologies in the form of the slow roll inflation potential that is obtained from the variation of the spectral action functional. The amplitude of the potential, and therefore the amplitude of the corresponding power spectra for density perturbations and gravitational waves (scalar and tensor perturbations), differs by a factor that depends on the topology, while the slow-roll parameters only detect a difference between the spherical and flat cases. As one knows from \cite{Lidsey}, \cite{SKamCoo}, \cite{StLy}, both the slow-roll parameters and the amplitude of the power spectra are constrained by cosmological information, so in this kind of modified gravity model, one in principle obtains a way to constrain the topology of the universe based on the slow-roll inflation potential, on the slow-roll parameters and on the power spectra for density perturbations and gravitational waves. The factors $\lambda_Y$ that correct the amplitudes depending on the topology are given by the following table. \begin{center} \begin{tabular}{| c | c || c | c |} \hline $Y$ spherical & $\lambda_Y$ & $Y$ flat & $\lambda_Y$ \\ \hline & & & \\ sphere & $1$ & flat torus & $1$ \\ & & & \\ \hline & & & \\ lens $N$ & $\frac{1}{N}$ & $G2(a)(b)(c)(d)$ & $\frac{HSL}{2}$ \\ & & & \\ \hline & & & \\ binary dihedral $4N$ & $\frac{1}{4N}$ & $G3(a)(b)$ & $\frac{HL^2}{2\sqrt{3}}$ \\ & & & \\ \hline & & & \\ binary tetrahedral & $\frac{1}{24}$ & $G4(a)(b)$ & $\frac{HL^2}{4}$ \\ & & & \\ \hline & & & \\ binary octahedral & $\frac{1}{48}$ & $G5$ & ? \\ & & & \\ \hline & & & \\ binary icosahedral & $\frac{1}{120}$ & $G6$ & $\frac{HLS}{4}$ \\ & & & \\ \hline \end{tabular} \end{center} \medskip Notice that some ambiguities remain: the form of the potential and the value of the scale factor $\lambda$ alone do not distinguish, for instance, between a lens space with $N=24$, a binary dihedral quotient with $N=6$ and the binary tetrahedral quotient, or between the Poincar\'e dodecahedral space (the binary icosahedral quotient), a lens space of order $N=120$ and a binary dihedral quotient with $N=30$. At this point we do not know whether more refined information can be extracted from the spectral action that can further distinguish between these cases, but we expect that, when taking into account a more sophisticated version of the spectral action model, where gravity is coupled to matter by the presence of additional (non-commutative) small extra-dimensions (as in \cite{CCM}, \cite{CoSM}), one may be able to distinguish further. In fact, instead of a trivial product $X\times F$, one can include the non-commutative space $F$ using a topologically non-trivial fibration over the 4-dimensional spacetime $X$ and this allows for a more refines range of proportionality factors $\lambda_Y$. We will discuss this in another paper.
| 10 | 12 |
1012.0780
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1012
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1012.0749_arXiv.txt
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If thick disks are ubiquitous and a natural product of disk galaxy formation and/or evolution processes, all undisturbed galaxies which have evolved during a significant fraction of a Hubble time should have a thick disk. The late-type spiral galaxy NGC~4244 has been reported as the only nearby edge-on galaxy without a confirmed thick disk. Using data from the {\it Spitzer Survey of Stellar Structure in Galaxies} (S$^{4}$G) we have identified signs of two disk components in this galaxy. The asymmetries between the light profiles on both sides of the mid-plane of NGC~4244 can be explained by a combination of the galaxy not being perfectly edge-on and a certain degree of opacity of the thin disk. We argue that the subtlety of the thick disk is a consequence of either a limited secular evolution in NGC~4244, a small fraction of stellar material in the fragments which built the galaxy, or a high amount of gaseous accretion after the formation of the galaxy.
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Thick disks are detected in edge-on and very inclined galaxies as an excess of flux at some distance from the mid-plane, typically at a few thin disk scale-heights. They usually have exponential profiles with a larger scale-height than the `canonical' or thin disks. Thick disks were first detected by Tsikoudi (1979) and were defined by Burstein (1979). A thick disk has subsequently been found in the Milky Way (Gilmore \& Reid 1983) and in basically all kinds of disk galaxies (see the introduction of Yoachim \& Dalcanton 2008 and references therein). Several theories of thick disk formation have been proposed. It has been suggested that thick disks are formed through mechanisms of internal evolution, namely disk heating by gravitational instabilities (Bournaud et al.~2007; 2009) and/or stellar scattering by encounters with features such as dark matter subhaloes, spiral arms or giant molecular clouds (Villumsen 1985; H\"anninen \& Flynn 2002; Benson et al.~2004; Hayashi \& Chiba 2006; Haywood 2008; Kazantzidis et al.~2008; Sch\"onrich \& Binney 2009a; Sch\"onrich \& Binney 2009b ; Bournaud et al.~2009). In this picture the thick disk formation may take less than one gigayear in the case of fast dynamical heating by a very clumpy initial disk (Bournaud et al.~2009) or may be much slower (secular evolution; Villumsen 1985; H\"anninen \& Flynn 2002; Benson et al.~2004; Kazantzidis et al.~2008). The specific type of secular evolution which causes a slow heating of the stellar orbits has been termed `blurring' by Sch\"onrich \& Binney (2009a; 2009b). Another possibility is that the thick disk is formed through accretion of stellar material from merging satellites at the time of galaxy formation (Abadi et al.~2003; Yoachim \& Dalcanton 2006). A third possibility is that the thick disk is a consequence of {\it in situ} star formation (Brook et al.~2004, Elmegreen \& Elmegreen 2006) or of star formation with a high initial velocity dispersion in very massive star clusters (Kroupa 2002). Whatever the thick disk formation mechanism is, it has to act over most, if not all, disk galaxies as thick disks have been found to be ubiquitous (Dalcanton \& Bernstein 2002; Yoachim \& Dalcanton 2006). Therefore, if any disk galaxy was found not to possess a thick disk, it would be likely to have suffered a peculiar origin and/or evolution. A search in the literature yields a candidate, NGC~4244, which has been claimed either to be a doubtful case (Yoachim \& Dalcanton 2006), or not to have a thick disk at all (Fry et al.~1999). NGC~4244 is a highly inclined apparently undisturbed Scd galaxy located 4.4\,Mpc away (Seth et al.~2005a). At that distance, one arcsecond is equivalent to roughly 21\,pc in linear scale. NGC~4244 is nearly bulgeless, except for a compact and bright nuclear star cluster (see Seth et al.~2008 for a detailed study on the nuclear cluster; the bulge-to-disk ratio is 0.03 according to Fry et al.~1999). NGC~4244 exhibits systematic displacements of the mid-plane of a few tens of parsecs with a wave-length of a couple of kiloparsecs (this phenomenon has been termed `corrugation' and in the case of NGC~4244 it has been studied by Florido et al.~1991). Kodaira \& Yamashita (1996) found NGC~4244 to have less warm dust and star formation than galaxies of the same Hubble type and \hi\ content, thus labeling it as an `anemic' galaxy. Finally, de Jong et al.~(2007) found a sharp disk cut-off which they suggest to have a dynamical origin. NGC~4244 could be weakly interacting with NGC~4214, which has a similar recession velocity and is found at a projected distance of 120\,kpc from NGC~4244. Several studies have attempted to identify the structural components of NGC~4244. Bergstrom et al.~(1992) found no evidence for a massive stellar halo. Fry et al.~(1999) detected only a thin disk. Hoopes et al.~(1999) found that, unlike some other edge-on galaxies, NGC~4244 has its ionized gas concentrated in the thin disk. Strickland et al.~(2004) found no evidence for a radio, UV, or X-ray emitting halo. Seth et al.~(2005b) found that the scale-heights of the older stellar components are larger than those of the younger components, which could be an indicator of the existence of a thick disk; however, the difference in scale-heights between the main sequence population and the red giant branch (RGB) population is smaller than in all the other galaxies studied in their paper. They also detected what could be interpreted as a halo. Tikhonov \& Galazutdinova (2005) claimed to have discovered a tenuous extended halo with a scale-length of several kiloparsecs in NGC~4244. Seth et al.~(2007) found a faint halo component 2.5\,kpc above the mid plane at the position of the minor axis. To sum up, previous studies on NGC~4244 show little to no emission at any wavelength outside its well-known thin disk. As nearly all very inclined disk galaxies have been shown to host a thick disk, the absence of it in NGC~4244, if confirmed, would be an indication of a peculiar evolution or formation mechanism of this galaxy. In this Paper we present results of a study of NGC~4244 based on a {\it Spitzer Survey of Stellar Structure in Galaxies} (S$^{4}$G) $3.6\mu{\rm m}$ image (Sheth et al.~2010; Sect.~2), confirming that this galaxy does in fact have signs of two disk components. For this purpose we have fitted the galaxy using {\sc Galfit} (version~3.0; Peng et al.~2002; Peng et al.~2010; Sect.~3) and produced luminosity profiles of NGC~4244 at different galactocentric radii (Sect.~4). We describe the luminosity profiles and compare them with a simple inclined two-disk model with a diffuse dust contribution in the thin disk in Sect.~5. We present a summary of the NGC~4244 main structural components in Sect.~6, and discuss our findings in Sect.~7.
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We find evidence of two disk components in a galaxy which was reported to have only one or as being a doubtful case. Using modeling we find that the fact of the thick disk is not detected in the far side of the galaxy can be explained by the effect of dust in a not perfectly edge-on galaxy. Fry et al.~(1999) did not detect the thick disk in their surface photometry because although the $R$-band image has a limiting surface brightness similar to ours, they assumed that the galaxy was symmetric with respect to the mid-plane and they summed the top and bottom luminosity profiles, which diluted the effect of the subtle thick disk. In addition $R$-band is more sensitive to dust, making such a subtle thick disk hard to detect in this band. Seth et al.~(2005b) deduced the presence of an extended component other than the thin disk from the scale-height of the oldest stars (RGB), which have a slight overdensity above the fitted sech$^{-2}$ profile. However, this component is more extended than the thick disk we found and may be a halo. The sum of all stellar components they have studied (young main sequence stars, Helium-burning stars, AGB stars, and RGB stars) shows some trace of what could be a thick disk in the online color version of their Fig.~3. Tikhonov \& Galazutdinova (2005) found an extended feature by counting RGB stars at the location of the arc-like feature. Seth et al.~(2007) studied the RGB star distribution in a thin strip along the bottom part of the minor axis of NGC~4244 deep enough to find a halo component undetected by us due to its extreme faintness, but they did not notice any trace of the thick disk. This is probably due to the fact that the evidence for a thick disk is more subtle at $r=0\arcsec$ than at other galactocentric radii (see Fig.~\ref{profiles4244}). The scale-height for the disk they fit ($\sim300$\,pc or $\sim14\arcsec$) is exactly the same as the one we find for the thin disk at $r=0\arcsec$. One possible thick disk formation mechanism is dynamical heating of a thin disk. According to recent modeling (Bournaud et al.~2007; 2009), the thick disk should, at least in part, form due to gravitational instabilities. These are also responsible for the creation of giant star-forming clumps which will eventually merge to form the galaxy bulge during the first gigayear after the galaxy formation. As NGC~4244 has virtually no bulge and a thick disk which is not very differentiated from the thin disk, it is conceivable that for some reason the disk dynamical instabilities have been suppressed, which would lead to a lack of thin/thick disk differentiation. The other kind of dynamical heating is a consequence of the encounters of stars with disk substructure and is considered to be a secular evolutionary process due to thin disk irregularities such as spiral arms and giant molecular clouds, or due to dark matter subhaloes. The lack of a very prominent thick disk would indicate that NGC~4244 has got a very smooth structure during most of its life. This smooth structure would have led to a modest `blurring' in the disk of the galaxy. Alternatively, NGC~4244 could have been formed much more recently than other disk galaxies studied so far. Part of the thick disk could also have formed through accretion of stellar material from merging satellites at the time of the galaxy build-up (Yoachim \& Dalcanton 2006). According to this model the accreted gas builds the thin disk. In this formation scenario, the lack of an extended thick disk would imply that the merging proto-galaxies did not have a substantial stellar component at the time of the formation of NGC 4244. Last, but not least, the thick disk may have formed {\it in situ}. According to Elmegreen \& Elmegreen (2006) edge-on galaxies with redshifts $z=0.5-4.5$ (corresponding to look-back times of $4.8-12.2$\,Gyr) have disk morphologies whose ratio between long and short axes are compatible with those of local thick disks. However, these disks should shrink when the galaxy accretes gaseous material, which makes the potential well deeper. In the context of this thick disk origin theory NGC~4244 may be a case of a galaxy having accreted more gaseous material than average, thus making the thick disk thinner than for most disk galaxies. This material accretion must have been through gentle gas inflows which have settled in the thin disk. To sum up, we have found evidence for two disk components in NGC~4244. The reason for the low degree of differentiation between the thin and thick disk remains unknown and may be related to a lack of strong secular evolution, to a lack of stellar material in the fragments which built up the galaxy, or to a high inflow rate of gaseous material after the galaxy formation.
| 10 | 12 |
1012.0749
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1012
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1012.2607.txt
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\ifdraft \vspace{-58ex} \hfill \framebox{\Blb{\LARGE\bf Draft of \isodate}} \vspace{54ex} \fi We present a catalogue of \note{accurate} positions and correlated flux densities \note{for} 410 flat-spectrum, compact extragalactic radio sources previously detected in the AT20G survey. The catalogue spans the declination range $[-90\degr, -40\degr]$ and was constructed from four 24-hour VLBI observing sessions with the Australian Long Baseline Array at 8.3~GHz. The \note{VLBI} detection rate in these experiments is 97\%, the median uncertainty of \note{the} source positions is 2.6~mas, \note{and} the median correlated flux density \note{on projected baselines longer than 1000\,km} is 0.14~Jy. The \note{goals of this work are} 1)~to provide a pool of \note{southern} sources with positions \note{accurate to a few milliarcsec, which can be used} for phase referencing observations, geodetic VLBI and space navigation; 2)~to extend the complete flux-limited sample of compact extragalactic sources to the southern hemisphere; and 3)~to investigate \note{the} parsec-scale properties of high-frequency selected sources from the AT20G survey. As a result of \note{this VLBI campaign}, the number of compact radio sources \note{south of declination $-40\degr$ which have measured VLBI correlated flux densities and positions known to milliarcsec accuracy has} increased by a factor of 3.5. The catalogue and supporting material is available at {\tt http://astrogeo.org/lcs1}.
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Catalogues of positions of compact extragalactic radio sources with the highest accuracy are important for many applications. \Note{These include} imaging faint radio sources in the phase referencing mode, differential astrometry, space geodesy, and space navigation. The method of VLBI first proposed by \citet{r:mat65} allows us to derive the position of sources with nanoradian precision (1 nrad $\approx$ 0.2~mas). The first catalogue of source coordinates determined with VLBI contained 35~objects \citep{r:first-cat}. Since then hundreds of sources have been observed under geodesy and astrometry VLBI observing programs at 8.6 and 2.3~GHz (X and S bands) using the Mark3 recording system at the International VLBI Service for Geodesy and Astrometry (IVS) network. Analysis of these observations resulted in the ICRF catalogue of 608~sources \citep{r:icrf98}. \Note {The Very Long Baseline Array} (VLBA) \Note {was later used to measure the positions of $~\!4000$ compact radio sources} in the VLBA Calibrator Survey (VCS) \citep{r:vcs1,r:vcs2,r:vcs3,r:vcs4,r:vcs5,r:vcs6} and the geodetic program RDV \citep{r:rdv}. All sources with declinations above $-45\degr$ detected using Mark3/Mark4 under IVS programs were re-observed with the VLBA \Note{in the} VCS and RDV programs, which significantly improved the accuracy of their positions. As a result of these efforts, the probability of finding a calibrator with a VLBI-determined position greatly increased. In the declination range $\delta > -30\degr$ the probability of finding a calibrator within a $3\degr$ radius of a given position reached 97\% by 2008. \Note{Since} the VLBA is located in the northern hemisphere, observations in the declination zone $[-50\degr, -30\degr]$ are difficult and the array \Note{cannot observe} sources with $\delta < -52\degr$. In 2008, the probability of finding a calibrator within a radius of $3^\circ$ was 75\% in the declination zone $[-40^\circ, -30^\circ]$ and 42\% for declinations \Note{south of }$-40^\circ$. The VLBI calibrator list\footnote{Available at {\tt http://astrogeo.org/rfc}} in 2008 had 524~sources in the zone $[+52^\circ, +90^\circ]$, \Note{but} only 98~objects in the zone $[-52^\circ, -90^\circ]$ \Note{which cannot be reached by} the VLBA. These southern sources were observed during geodetic experiments and during two dedicated southern hemisphere astrometry campaigns \citep{r:fey_south1,r:fey_south2}. The reason for this disparity is the scarcity of VLBI antennas in the southern hemisphere, particularly stations with dual frequency S/X receivers and geodetic recording systems. Also until recently there has been a lack of good all-sky \Note{catalogues suitable for finding candidate VLBI calibrators}. The Australian Long Baseline Array (LBA) consists of six antennas located in Australia with the South Africa station {\sc hartrao} often joining in. \Note{This VLBI network operates 3--4 observing sessions a year, each about one week long}. Although the hardware used by the LBA was not designed for geodesy and absolute astrometry observations, it was demonstrated by \citet{r:geo_lba} that despite significant technical challenges, absolute astrometry VLBI observations with the LBA network is feasible. In a pilot experiment in June 2007, \Note{the} positions of participating stations were determined with accuracies 3--30~mm, and positions of \Note{five} new sources were determined with accuracies 2--5~mas. Inspired by these astrometric results, we launched the X-band LBA Calibrator Survey observing campaign (LCS), with the aim of determining milliarcsecond positions and correlated flux densities \Note{for} one thousand compact extragalactic radio sources \Note{at declinations south of $-30^\circ$}. The overall objective of this campaign is to match the density of calibrators in the northern hemisphere and \Note{so} eliminate the disparity. We \Note{have three long-term goals in} this campaign. \Note{Firstly, setting up} a dense grid of calibrators with precisely known positions within several degrees of any target \Note{will make} make phase-referencing observations of weak targets \Note{feasible}. According to \citet{r:wrobel_pr}, \Note{63\% of VLBA observations in 2003--2008} were made in the phase referencing mode. A dense grid of calibrators also makes {\note it possible to do} differential astrometry of Galactic objects such as pulsars and masers, and allows direct determination of the parallax at distances up to several kiloparsec \citep{r:del09}. These sources form the pools of targets for observations under \Note{the geodesy programs} and for space navigation. \Note{Our} second goal is to extend the complete flux-limited sample of compact extragalactic radio sources (with emission from milliarcsecond-size regions) to the entire sky. According to Kovalev (private communication, 2010), analysis of the $\log N$--$\log S$ diagram of the VLBI calibrator list suggests the sample of radio--loud Active \Note{Galactic} Nuclei (AGN) is complete at the level of correlated flux density 200~mJy at X-band at spatial frequencies 25~$M\lambda$ at $\delta > -30\degr$. Extending \Note{this} complete sample to the entire sky will make it possible to generalize the properties of the sample, such as distribution of compactness, distribution of brightness temperature, bulk motion, viewing angle, irregularities of the spatial distribution, to the entire population of radio loud AGN. The third goal \Note{is} to investigate \Note{the} properties of high-frequency selected radio sources from \Note{the} Australia Telescope 20~GHz (AT20G) survey \citep{r:at20g,r:mass10}. Obtaining observations of a subsample of AT20G sources with milliarcsecond resolution will allow us to investigate \Note{the} properties, such as spectral index, polarization fraction and variability, of a population of extremely compact sources. In this paper, we present the results from the first four 24-hour experiments observed in 2008--2009. \Note{The selection of candidate sources from} the AT20G catalogue is discussed in section~\ref{s:selection}. The station setups during the observing sessions is described in section~\ref{s:observations}. \Note{The} correlation and post-correlation analysis, which is rather different from ordinary VLBI experiments, is discussed in sections~\ref{s:correlation} and \ref{s:analysis}. \Note{An error analysis} of single-band observations, including evaluation of ionosphere-driven systematic errors, is given in subsection~\ref{s:errors}. The catalogue of source positions and correlated flux densities is presented in section~\ref{s:catalog}, \Note{and} the results are summarized in section~\ref{s:summary}.
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\label{s:summary} The absolute astrometry LBA observations turned out highly successful. The overall detection rate was 97\% --- the highest rate ever achieved in a VLBI survey. If we exclude extended sources, non-AT20G sources and the \Note{six} planetary nebulae \Note{included} in the candidate list due \Note{to an} oversight, the detection rate is 99.8\%! We attribute this high detection rate to two factors. Firstly, the AT20G catalogue is highly reliable and is biased \Note{towards very} compact objects. Selecting candidates based on \Note{the} simultaneous AT20G spectral index proves to be a good methodology. Secondly, the LBA has very high sensitivity. The baseline detection limit over 2~minute of integration time varied from 7~mJy to 30~mJy, with 7--12~mJy at baselines with \parkes. We have successfully resolved group delay ambiguities with spacing 3.9~ns for all observations. This became possible using the innovative algorithm exploiting relatively low level of instrumental group delay errors. We have determined positions of 410 target sources never before observed using VLBI, with \Note{a} median uncertainty \Note{of} 2.6~mas. Error analysis showed a moderate contribution of the mismodeled ionosphere path delay to the overall error budget. Both random and systematic errors are accounted for in the \Note{uncertainties} ascribed to source positions exploiting the overlap of 111 additional sources observed in LCS1 experiments with their positions known from prior observations. The positional accuracy of the LCS1 catalogue is a factor of 350 greater than the positional accuracy of the AT20G catalogue, \Note{corresponding} to the ratio of the maximum baseline \Note{lengths} of the LBA and the ATCA. The new catalogue has increased the number of sources with declinations $< -40\degr$ from 158 to 568, i.e. by a factor of 3.5. We determined correlated flux densities for 410 target and 111 calibrator sources, and presented their median values in three ranges of baseline projection lengths. The correlated flux density of \Note{the} target sources varied from 0.02 to 2.5~Jy, \Note{and was in the range 80--300~mJy for 70\% of the} sources. \Note{The uncertainties} of the correlated flux densities \Note{are estimated to be typically 5--8\%}. This observing program is continuing. By November 2010, four additional twenty four hour experiments \Note{had been} observed with several more observing sessions planned.
| 10 | 12 |
1012.2607
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1012
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1012.2099_arXiv.txt
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} The Calar Alto Observatory is located at 2168m height above the sea level, in the Sierra de los Filabres (Almeria-Spain) at ~45 km from the Mediterranean Sea. It is the second largest European astronomical site in the northern hemisphere just behind the Observatorio del Roque de los Muchachos (located in the island of La Palma), and the most important in the continental Europe (with excellent communications, making logistics easy, unexpensive and reliable). Currently there are six telescopes located in the complex, three of them directly operated by the Centro Astron\'omico Hispano Alem\'an A.I.E., a partnership between the Spanish National Research Council (CSIC\footnote{www.csic.es}) and the German Max-Plank Society (MPG\footnote{www.mpg.de}). These telescopes include the Zeiss 3.5m, the largest telescope in the continental Western Europe. The observatory is under operations since 1975, when its 1.23 m Zeiss reflector saw first light. The observatory operates a very large array of optical and near-infrared astronomical instrumentation, including imagers and spectrographs with different field-of-view and resolutions. There has been different attempts to characterize some of the main astronomical properties during its 35 years of operations: (i) Leinert et al. (1995) determined the sky brightness corresponding to the year 1990; (ii) Hopp \& Fernandez (2002) studied the extinction curve corresponding to the years 1986-2000; (iii) Ziad et al. (2005) estimated the median seeing in the observatory from a single campaign in May 2002; and more recently (iv) S\'anchez et al. (2007) where the optical night sky spectrum was presented, including an analysis of the light pollution, together with a more accurate estimation of the night-sky extinction, the typical seeing, the night-sky brightness and the fraction of useful time; and (v) S\'anchez et al. (2008), where the night sky brightness in the near-infrared and the fraction of useful time was presented. Several of these features are discussed below. The comprehensive database for the weather is public\footnote{www.caha.es/WDXI/wdxi.php} and it can also be obtained upon request. \begin{figure} \center \includegraphics[width=6.8cm]{barradodF1a.jpg} ~ \includegraphics[width=7.5cm]{barradodF1b.jpg} ~ \caption{\label{fig1} {\bf Left.-} Extinction due to dust at Calar Alto. The values are comparatively smaller than representative values of observatories closer to the Equator, both in the Northern and the Southern Hemispheres. The behaviour is highly seasonal. {\bf Right.-} Seeing distribution (June 2001 -- Sept. 2005) at the Calar Alto Observatory. The median value is about 0.9 arcsec (S\'anchez et al. 2007). } \end{figure}
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The Calar Alto observatory has two main properties: a very well-characterized and excellent astronomical site (specially for spectroscopy) and outstanding logistics (excellent communications and a location in continental Europe). The current and future suite of instruments for the optical and near-IR telescopes, and the support of two of the main research institutions in Europe (the Spanish National Research Council and the German Max-Plank Society, channeled through the Instituto de Astrof\'{\i}sica de Andaluc\'{\i}a and the Max-Plank Institut f\"ur Astronomie), provide the necessary conditions for its continuation as relevant astronomical facility. \small %
| 10 | 12 |
1012.2099
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1012
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1012.0235_arXiv.txt
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We stack 4.6\,Ms of high spectral resolution \emph{XMM-Newton} Reflection Grating Spectrometer spectra from galaxy clusters, groups of galaxies and elliptical galaxies. For those objects with a central temperature of less than 1~keV, we detect O\,\textsc{vii} for the first time, with a probability of false detection of $2.5\times10^{-4}$. The flux ratio of the O\,\textsc{vii} to Fe\,\textsc{xvii} lines is $1/4$ to $1/8$ of the emission expected for isobaric radiative cooling in the absence of heating. There is either a process preventing cooling below 0.5~keV, anomalous O/Fe abundance ratios, absorbing material around the coolest X-ray emitting gas or non-radiative cooling taking place. The mean N\,\textsc{vii} emission line is strong in the sub-keV sample. As the ratio of the hydrogenic N and O lines is largely independent of temperature, we measure a mean N/O ratio of $4.0 \pm 0.6 \Zsun$. Although the continuum around the C\,\textsc{vi} lines is difficult to measure we can similarly estimate that the C/O ratio is $0.9 \pm 0.3 \Zsun$.
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The intracluster medium (ICM) is the few $\times10^{6}$ to $\sim 10^8$~K X-ray emitting plasma which fills the potential wells of objects ranging in scale from galaxy clusters down to massive elliptical galaxies. The X-ray surface brightness profile of these objects is often steeply peaked \citep[e.g.][]{Stewart84}. The implied mean radiative cooling times are less than 1~Gyr. In the absence of heating 10s to 100s of solar masses per year should be cooling below $\sim10^{6}$K to form a cooling flow \citep{Fabian94}. The relative strength of X-ray emission lines is a good probe of the temperature distribution of the ICM. The major spectral lines in the $5-38${\AA} spectral range can be observed in bright objects using the Reflection Grating Spectrometer (RGS) instruments on \emph{XMM-Newton}. The weak or missing Fe\,\textsc{xvii} emission lines, indicating material around $\sim 0.5$~keV, in RGS spectra showed that there is much less cool gas present in clusters than expected from cooling in the absence of heating \citep[e.g.][]{Peterson01,Peterson03}. However, these lines have been detected in several objects using deep observations \citep{SandersRGS08,SandersA220409}. There can be a wide range of $\times10-15$ in temperature, but less than 10 per cent of the emission expected for radiative cooling is seen. Active galactic nuclei (AGN) in the cores of clusters are believed to be the mechanism by which energy is supplied to the ICM to combat a large fraction of the radiative cooling \citep[for reviews see][]{PetersonFabian06,McNamaraNulsen07}. \begin{figure} \includegraphics[width=\columnwidth]{line_strengths.pdf} \caption{Emissivity of lines in the RGS waveband from various strongly X-ray emitting ions as a function of temperature at Solar metallicity.} \label{fig:linestrengths} \end{figure} Fig.~\ref{fig:linestrengths} shows the emissivity of emission lines from certain strong X-ray emitting ions as a function of temperature. To create this plot we summed the emissivity of strong lines in the RGS spectral range calculated using \textsc{spex} 2.00.11 \citep{Kaastra00}. The emissivity in Fe\,\textsc{xvii} emission peaks at a temperature of around 0.5~keV. To observe gas at lower temperatures requires the detection of O\,\textsc{vii} lines. No significant detections of O\,\textsc{vii} have been reported in clusters, groups or elliptical galaxies, implying there is little material below 0.5~keV temperature. Very deep \emph{XMM} observations of bright high-metallicity cool objects may detect the presence of O\,\textsc{vii} emission lines. The calorimeters planned for \emph{ASTRO-H} and \emph{IXO} should also be able to detect these lines. In the absence of any such observations at this time we conducted a stacking analysis of the best observations in the \emph{XMM-Newton} archive. In this paper we assume the relative Solar abundances of \cite{AndersGrevesse89}.
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We examine the mean RGS spectrum of objects from the \emph{XMM-Newton} archive in different temperature bins. For those objects with a RGS-measured temperature of less than 1~keV we detect for the first time O\,\textsc{vii} emission. This is indicative of material below 0.5~keV. If the gas in the objects is cooling without any heating, there is $1/4$ to $1/8$ of the expected O\,\textsc{vii} emission compared to the 16.8 and 17.1{\AA} Fe\,\textsc{xvii} flux. This may imply that a small fraction of the gas is cooling in these objects, the abundance ratios of O and Fe are anomalous, or that there is absorbing material around the coolest gas. An alternative explanation is non-radiative cooling of the coolest gas. The ratio of the hydrogenic lines in the spectrum are independent of temperature, allowing estimates of the mean N/O and C/O ratios. In the sample of objects below 1~keV the N/O ratio is $4.0 \pm 0.6 \Zsun$. The C/O ratio is $0.9 \pm 0.3 \Zsun$.
| 10 | 12 |
1012.0235
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1012
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1012.2006_arXiv.txt
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Cosmological Inflation~\cite{LindeBook,MukhanovBook} naturally generates a spectrum of density fluctuations responsible for large scale structure formation which is consistent with the observed CMB anisotropies.\cite{Komatsu2010} It also generates a spectrum of gravitational waves, whose amplitude is directly related to the energy scale during inflation and which induces a distinct B-mode polarization pattern in the CMB.\cite{DurrerBook} Moreover, Inflation typically ends in a violent process at preheating,\cite{preheating} where large density waves collide at relativistic speeds generating a stochastic background of GW~\cite{GWpreh} with a non-thermal spectrum characterized by a prominent peak at GHz frequencies for GUT-scale models of inflation (or at mHz-kHz for low scale models of inflation), and an amplitude proportional to the square of the mass scale driving/ending inflation. Such a background could be detected with future GW observatories like Adv-LIGO~\cite{ligo}, LISA~\cite{lisa}, BBO~\cite{bbo}, DECIGO~\cite{decigo}, etc. In case the end of inflation involves the presence of gauge fields associated with the breaking of some symmetry, like in most scenarios of hybrid inflation, then the PGW spectrum presents an extra peak associated with the mass scale of the corresponding gauge fields~\cite{DFG2010}. This would constitute a very clear signature of the physics of reheating, and its detection would open a new window into the early universe. Furthermore, if inflation ended with a global O(N) phase transition, like in certain scenarios of hybrid inflation, then there is also a GWB due to the continuous self-ordering of the Goldstone modes at the scale of the horizon,\cite{Krauss} which is also scale-invariant on subhorizon scales,\cite{FFDGB} with an amplitude proportional to the quartic power of the symmetry breaking scale, that could be detectable with laser interferometers as well as indirectly with the B-mode polarization of the CMB.\cite{GBDFFK} \begin{figure} \vspace{-0.5cm} \centerline{ \includegraphics[width=13.0 cm,height=10 cm]{energyGW.pdf} } \vspace{-0.5cm} \caption{The time evolution of the GW energy density during the initial stages of preheating after hybrid inflation, from Ref.[6]. Note the three stages of tachyonic growth, bubble collisions and turbulence.} \label{fig:GWpreh} \end{figure} Gravitational waves produced during inflation arise exclusively due to the quasi-exponential expansion of the Universe~\cite{MukhanovBook}, and are not sourced by the inflaton fluctuations, to first order in perturbation theory. They have an approximately scale invariant and Gaussian spectrum whose amplitude is proportional to the energy density during inflation. GUT scale inflation has good chances to be discovered (or ruled out) by the next generation of CMB anisotropies probes, Planck~\cite{Planck} COrE~\cite{Bpol} and CMBpol~\cite{CMBpol}.% At the end of inflation, reheating typically takes place in several stages. There is first a rapid (explosive) conversion of energy from the inflaton condensate to the fields that couple to it. This epoch is known as preheating~\cite{preheating} and occurs in most models of inflation. It can be particularly violent in the context of hybrid inflation, where the end of inflation is associated with a symmetry breaking scenario, with a huge range of possibilities, from GUT scale physics down to the Electroweak scale. Gravitational waves are copiously produced at preheating from the violent collisions of high density waves moving and colliding at relativistic speeds~\cite{GWpreh}, see Fig.~1. The GW spectrum is highly peaked at the mass scale corresponding to the symmetry breaking field, which could be very different from the Hubble scale. In low-scale models of hybrid inflation it is possible to attain a significant GWB at the range of frequencies and sensitivities of LIGO or BBO. The origin of these GW is very different from that of inflation. Here the space-time is essentially static, but there are very large inhomogeneities in the symmetry breaking (Higgs) field due to the random spinodal growth during preheating. Although the transition is not first order, ``bubbles" form due to the oscillations of the Higgs field around its minimum. The subsequent collisions of the quasi-bubble walls produce a rapid growth of the GW amplitude due to large field gradients, which source the anisotropic stress-tensor~\cite{GWpreh}. The relevant degrees of freedom are those of the Higgs field, for which there are exact analytical solutions in the spinodal growth stage, which later can be input into lattice numerical simulations in order to follow the highly non-linear and out-of-equilibirum stage of bubble collisions and turbulence. However, the process of GW production at preheating lasts only a short period of time around symmetry breaking. Soon the amplitude of GW saturates during the turbulent stage and then can be directly extrapolated to the present with the usual cosmic redshift scaling. Such a GWB spectrum from preheating would have a characteristic bump, worth searching for with GW observatories based on laser interferometry, although the scales would be too small for leaving any indirect signature in the CMB polarization anisotropies. Moreover, the mechanism generating GW at preheating is also active in models where the SB scenario is a local one, with gauge fields present in the plasma, and possibly related to the production of magnetic field flux tubes~\cite{PMFpreh}. In such a case, one could try to correlate the GWB amplitude and the magnitude and correlation length of the primordial magnetic field seed.
| 10 | 12 |
1012.2006
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1012
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1012.0003_arXiv.txt
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The Coma cluster of galaxies hosts the brightest radio halo known and has therefore been the target of numerous searches for associated inverse Compton (IC) emission, particularly at hard X-ray energies where the IC signal must eventually dominate over thermal emission. The most recent search with the \suzakus Hard X-ray Detector (HXD) failed to confirm previous IC detections with \rxtes and \sax, instead setting an upper limit 2.5 times below their nonthermal flux. However, this discrepancy can be resolved if the IC emission is very extended, beyond the scale of the cluster radio halo. Using reconstructed sky images from the 58--month \swifts BAT all sky survey, the feasibility of such a solution is investigated. Building on Renaud et al., we test and implement a method for extracting the fluxes of extended sources, assuming specified spatial distributions. BAT spectra are jointly fit with an \xmms EPIC-pn spectrum derived from mosaic observations. We find no evidence for large-scale IC emission at the level expected from the previously detected nonthermal fluxes. For all nonthermal spatial distributions considered, which span the gamut of physically reasonable IC models, we determine upper limits for which the largest (most conservative) limit is $\la 4.2\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (20--80 keV), which corresponds to a lower limit on the magnetic field $B > 0.2$ $\mu$G. A nominal flux upper limit of $< 2.7\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$, with corresponding $B > 0.25$ $ \mu$G, is derived for the most probable IC distribution given the size of the radio halo and likely magnetic field radial profile.
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\label{sec:comabat:intro} The X-ray emission from clusters of galaxies is primarily thermal in origin and is produced by a diffuse population of intergalactic electrons in the ionized intracluster medium (ICM). These electrons coexist with a nonthermal, relativistic electron population in at least some clusters -- inferred from observations in the radio regime -- which should also radiate at X-ray energies. While thermal emission clearly dominates in the kilo-electron volt (keV) energy range, it declines rapidly outside this range, allowing the detection of a nonthermal spectral signature as soft or hard excess emission. This possibility is especially promising at hard ($>$10 keV) energies, where the exponential decline of the thermal bremsstrahlung continuum is distinctly steeper than the expected nonthermal spectrum. Measurements of nonthermal X-ray emission are critical to the determination of the total amount of relativistic energy in the ICM, which is currently poorly constrained. While no more than $\sim$10\% of this energy is tied up in nonthermal components, amounts at or near this level will affect the dynamics and structure of the thermal gas \citep[e.g.,][]{VBG09}. Specifically, studies that attempt to infer the total masses of clusters from the hydrostatic state of the thermal gas will produce biased mass estimates if the pressure support of relativistic particles and fields is not accurately included. The mass functions built from these estimates can be used to constrain cosmological parameters; these studies are already underway using observables derived in both the X-ray \citep[e.g.,][]{MAE+08, Vik+09} and microwave \citep[e.g.,][through the Sunyaev-Zel'dovich effect]{Van+10} regimes. A measurement of the total energy in relativistic ICM components is possible when X-ray and radio nonthermal fluxes are combined. Diffuse, cluster-wide synchrotron radio emission, called radio halos or relics depending on their morphology, imply that both magnetic fields and relativistic electron populations are present on large scales. The total luminosity of a synchrotron-emitting electron is given by \begin{equation} \label{eq:lsync} L_R = \frac{4}{3}\sigma_Tc\gamma^2\epsilon_{B} \, , \end{equation} where $\sigma_T$ is the Thomson cross-section, $c$ is the speed of light, $\gamma$ is the Lorentz factor of the electron, and $\epsilon_B=B^2/8\pi$ is the energy density of the magnetic field. For a collection of relativistic electrons, the value of $L_R$ depends both on the number of electrons and on $B$ and cannot independently determine either. However, these same electrons will up-scatter cosmic microwave background (CMB) photons through inverse Compton (IC) interactions, which have a luminosity $L_X$ equivalent in form to equation~(\ref{eq:lsync}) but with $\epsilon_B$ replaced by the energy density of the CMB. Since both luminosities are proportional to the number of electrons, their ratio gives the volume-averaged magnetic field, \begin{equation} \label{eq:comabat:synicratio} \frac{L_R}{L_X} = \frac{B^2/8\pi}{aT_{CMB}^4} \, , \end{equation} where $a$ is the radiation constant and $T_{CMB}$ is the temperature of the CMB. The IC radiation should be observable at hard X-ray energies \citep{Rep77}. Thus far, IC emission has only been detected at low significance \citep{NOB+04, MA09} or, in one case at higher significance, in the Ophiuchus cluster (\citealt{EPP+08}; but see also \citealt{Aje+09} and \citealt{Fuj+08}), although the diffuse radio emission in this cluster is restricted to a smaller scale mini-halo \citep{Mur+10}. The measurement of an IC flux from a synchrotron source directly leads to a simultaneous determination of the average value of $B$ and the relativistic electron density \citep{HR74, Sar88}. Therefore searches for IC emission coincident with a radio halo or relic are an excellent way to constrain the contribution of relativistic materials in clusters. The first, and brightest, radio halo was discovered by \citet{Wil70} in the Coma cluster, and its radio properties have perhaps been the best studied \citep[e.g.][]{GFV+93, DRL+97, TKW03}. Coma has been observed by all the major observatories with hard X-ray capabilities \citep{HM86, Baz+90, HBS+93, RUG94}, and more recently non-thermal detections have been claimed by \citet{RG02} with {\it RXTE} and by \citet{FDF+99, FOB+04} with {\it BeppoSAX}, though the latter detection is controversial \citep{RM04, FLO07}. Due to the large field of view (FOV) of these non-imaging instruments and the simple characterization of the thermal gas, the source of this emission remains uncertain. Even more recently, long ($\sim 1$ Msec) observations with {\it INTEGRAL} have imaged extended diffuse hard X-ray emission from Coma, though it was found to be completely consistent with thermal emission \citep{RBP+06, ENC+07, LVC+08}. Most recently, \citet{WSF+09} performed a joint analysis of spectra from the \xmms EPIC-pn and \suzakus HXD-PIN instruments -- the most sensitive instruments at soft and hard energies to date -- of the Coma cluster and were unable to detect IC emission. Instead, they found an upper limit 2.5 times below the detections of \citet{RG02} and \citet{FOB+04}. However, the narrower FOV of the HXD relative to the collimators of the \rxtes PCA/HEXTE and \saxs PDS leaves open the possibility that the spatial distribution of IC photons is highly extended, and therefore much of the flux was missed by the HXD. The IC would have to be much broader than the size of the radio halo. A uniform IC surface brightness of at least 30\arcmin\ in radius from the cluster center is sufficient to reconcile these results. Therefore, an imaging analysis at hard X-rays is required to confirm this picture; unfortunately, no focussing hard ($>$10 keV) X-ray telescope has yet been deployed. In the meantime, it is possible to perform a crude imaging analysis with coded mask instruments, as previously discussed by \citet{RGL+06}. In this work, we report on the spatial and spectral hard X-ray emission from the Coma cluster using the 58--month accumulation of the \swifts Burst Alert Telescope (BAT) all-sky survey. After the first 9 months of the survey, Coma was seen to be clearly extended \citep{Aje+09}, so an accurate measurement of its flux must account for its resolved nature; the standard method of extracting fluxes from coded mask instruments assumes the underlying source to be point-like. Using models for the spatial distribution of thermal and potential nonthermal emission, we measure the total, extended flux in the 8 energy bands that make up the survey. These fluxes are then converted into spectra, which we jointly fit with an \xmms EPIC-pn spectrum from a {\it spatially identical region}. In this way, despite poor spatial resolution ($\sim20$\arcmin), we are sensitive to any large-scale, extended emission above the detection threshold for the survey. While the sensitivity of the BAT detector is lower than instruments such as the \suzakus HXD-PIN, the survey's large exposure time -- thanks to a FOV that sees $1/8^{\rm th}$ of the sky in a single pointing -- gives it a comparable, if not superior, overall sensitivity. In Section~\ref{sec:comabat:obs}, we describe the \swifts BAT survey in general and the \xmms EPIC-pn and BAT observations of the Coma cluster specifically. The extraction of spatially extended fluxes from models, along with the specific models themselves, is discussed in Section~\ref{sec:comabat:spatial}. Spectra constructed from these spatial fits are presented in Section~\ref{sec:comabat:specfits}, along with the results of joint fits with the \xmms spectrum. In Section~\ref{sec:comabat:uls}, we provide upper limits on spatially extended, nonthermal emission, and in Section~\ref{sec:comabat:disc} we discuss the implications of our non-detection for the relativistic phase of the ICM of the Coma cluster. In the appendices, we describe the calibration of the survey such that joint fits with \xmms are straightforward, and we demonstrate that the BAT instrument intrinsically detects extended emission on the scales of interest here, though with higher uncertainty than for a point source. We assume a flat cosmology with $\Omega_M = 0.23$ and $H_0 = 72$ km/s/Mpc and a luminosity distance to Coma of 98.4 Mpc. Unless otherwise stated, all uncertainties are given at the 90\% confidence level.
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\label{sec:comabat:disc} By taking advantage of the crude imaging capabilities of the \swifts BAT instrument and the impressive sensitivity of the 58--month all sky survey, we are able to constrain the amount of nonthermal, hard X-ray emission -- extended or otherwise -- from the Coma cluster. We find no evidence for an extended, hard excess that could reconcile recent detections from \rxtes \citep{RG02} and \saxs \citep{FOB+04} with the upper limit from \suzakus \citep{WSF+09}; note, however, these detections would still be in conflict with the upper limit of \citet{RM04}. Generic, uniform surface brightness disks, along with a recently proposed IC model \citep{KW10}, were fit to BAT survey images, converted to spectra, and investigated for signs of a nonthermal component. For each spatial model, we compute upper limits on a grid of positions and compare them to previous measurements, being careful to convert detected fluxes into intrinsic source fluxes, given a particular spatial distribution, by accounting for the collimator vignetting functions. These are direct comparisons, in the sense that the instrumental response of all detectors involved have been fully considered, and as such we, like \citet{RM04}, cannot confirm the claimed detections of \citet{RG02} and \citet{FOB+04}. Their observed hard excesses could have had other reasonable sources, if not diffuse IC emission from the nonthermal phase of the ICM. A common difficulty is an accurate determination of both the cosmic and non-X-ray background, the treatment of which is the primary difference between \citet{FOB+04} and \citet{RM04}. Another possibility is the variable nature of nearby point sources, most notably the AGN W Comae, which was once quite bright but has been fading for many years. The concurrence of the \rxtes and \saxs observations could have simply been unlucky and caught W Comae (or another source) in a bright state. A somewhat more subtle, and perhaps more likely, explanation concerns the multi-temperature nature of Coma's ICM. Small amounts of hot gas could dominate the high energy emission, so the extrapolation of an average temperature determined from lower energy data may not be an adequate description of the thermal contribution to the high energy flux. The effect of the multi-temperature gas in Coma is evident in the SW extension at hard energies observed by \integrals \citep{RGL+06, ENC+07} and confirmed here; higher temperatures seen at this location in the temperature map \citep{WSF+09} are sufficient to explain the change in morphology, which points to the increased significance of this gas at higher energies. Even so, single temperature spectral fits do not produce IC detections in this study or in \citet{WSF+09}. The explanation may simply rest in a slight mischaracterization of the hard energy emission weighted temperature; in \citet{FOB+04}, the FOV of the lower energy HPGSPC instrument does not quite match the higher energy PDS, and the temperature of $7.67\pm0.1$ keV found in \citet{RG02} is significantly below that allowed by the \xmms data. Given radio synchrotron emission and an upper limit on the X-ray IC flux, a lower limit on the average ICM magnetic field can be estimated, as described by equation~(13) in \citet{WSF+09} and the accompanying text. A diffuse radio flux of 640 mJy at 1.4 GHz is detected out to a radius of $\sim40\arcmin$ in \citet{DRL+97}. For comparison, we will use the upper limit of $2.7\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$ ($20 < E < 80$ keV) from the $R=40\arcmin$ disk model. These values imply $B>0.25$ $\mu$G, an increase from \citet{WSF+09} but still well below the equipartition value of $B_{\rm eq}=0.5$ $\mu$G for the Coma radio halo \citep{GFV+93}. A slightly lower limit of $B>0.2$ $\mu$G results if a more conservative IC upper limit of $4.2\times10^{-12}$ erg s$^{-1}$ cm$^{-2}$ is used, which considers the limits from all spatial models tested. Regardless, these limits on $B$ fall well below line of sight estimates of several $\mu$G from Faraday rotation measure (RM) observations \citep{FDG+95}, though due to geometric effects these measurements may not represent the average cluster magnetic field \citep{Pet01}. However, the global field may be recovered by combining many RM measurements along different lines-of-sight through the ICM with numerical simulations \citep{Mur+04}. \citet{BFM+10} have applied this method to the Coma cluster, deriving a radial profile where the energy density of the magnetic field falls roughly in proportion with the energy density of thermal gas and with a central field strength of $B_0 \sim 4.7$ $\mu$G. Combining this model of $B(r)$ with an approximate representation of the radial density profile of synchrotron emission, implied by a rough $\beta$-model fit to the point source subtracted image of \citet{DRL+97} ($r_c=18\arcmin$, $\beta=1$, and $I_0=1.23$ mJy arcmin$^{-2}$), directly leads to a prediction of the expected IC surface brightness as a function of radius. Our illustrative -- due to the large uncertainties in all parameters assumed in this exercise -- IC surface brightness distribution is flat out to $\sim 30\arcmin$ with a 20--80 keV flux of $\sim 8\times10^{-17}$ erg s$^{-1}$ cm$^{-2}$ arcmin$^{-2}$, at which point it nearly linearly drops toward zero, though not reaching it, around a radius of 90\arcmin. This surface brightness is about an order of magnitude below that implied by our upper limits, providing a possible explanation for why we are unable to detect an IC signature. An order of magnitude lower hard X-ray flux for Coma is also predicted by \citet[][see their Fig.~5]{BL10}, who have developed the most comprehensive picture yet of radio halo generation by MHD turbulence. On the other hand, larger IC fluxes would be expected if the radio synchrotron emission falls off more gradually than modeled here, since a flatter radial profile would suggest a higher relativistic electron density given the falling magnetic field with cluster radius. More accurate maps of Coma's radio halo, preferably at lower frequencies where the radio electrons correspond more closely to the IC-emitting electrons, will clarify this issue. Ultimately, a true detection of IC emission from Coma will have to wait for upcoming missions with focussing hard X-ray telescopes, namely {\it NuSTAR}\footnote{http://www.nustar.caltech.edu/} and {\it Astro-H} \footnote{http://astro-h.isas.jaxa.jp/}. For {\it NuSTAR} to achieve a sensitivity comparable to our upper limits, a single pointed observation of at least 100 ks will be required \citep{MHK+09}. However, the much finer spatial resolution will remove the uncertainty associated with bright background AGN and allow multiple spatially-resolved joint fits. Assuming the hottest gas, which produces the largest amount of thermal emission at hard energies, is localized, then these regions can be identified and avoided in order to detect a lower surface brightness, but more uniform, IC component. Similarly, if the IC emission is more localized, it will be easier to identify with spatially-resolved joint fits between \xmms and {\it NuSTAR} or {\it Astro-H} spectra, as has been done with {\it Chandra} data alone \citep{MA09}. The unambiguous detection of IC emission associated with radio halos and relics is crucial to determining the energy content in the relativistic phase of the ICM and how significant of an influence this phase has on the dynamics and structure of the thermal gas in clusters.
| 10 | 12 |
1012.0003
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1012
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1012.1143_arXiv.txt
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We have obtained near-IR photometry for the 11 Praesepe white dwarfs, to search for an excess indicative of a dusty debris disk. All the white dwarfs are in the DAZ temperature regime, however we find no indications of a disk around any white dwarf. We have, however determined that the radial velocity variable white dwarf WD0837+185 could have an unresolved T8 dwarf companion that would not be seen as a near-IR excess.
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There are $\approx10$ dusty debris disks known to exist around white dwarfs. The disks are found around relatively cool white dwarfs ($9,000<T<22,000$K) whose atmospheres are polluted with heavy elements (the so-called DAZ white dwarfs). The disks provide the obvious reservoir for the white dwarf to accrete this material, which otherwise would sink from the atmosphere on a timescale of days. DAZs are identified through high resolution, high S/N optical spectra in which lines of Ca, Si and Fe can be detected. \cite{zuckerman03} estimated that 20-25 per cent of single DA white dwarfs show Ca \textsc{ii} K lines, indicating they are DAZ. This statistic has been put into question, however by \cite{koester05} who studied 478 DA white dwarfs with 10 000 K $\geq$ T$_{\rm eff} \geq$30 000 K and found 24 DAZs, 6 of which had been discovered by \cite{zuckerman03}. This put the fraction of DA white dwarfs that are DAZ at 0.5 per cent. It was suggested by \cite{koester05} that this discrepancy has occurred as the \cite{zuckerman03} sample was mainly of objects with T$_{\rm eff}$$<$ 10 000 K, where the Ca absorption lines at lower abundances are easiest to detect, so their sample was biased towards the DAZs. \cite{kilic08} studied a sample of 37 DAZ white dwarfs using IRTF and $Spitzer$, of which 7 had dusty debris disks. They tentatively estimate the fraction of DAZ that harbour disks at $\approx$ 20 per cent, although this is obviously from a small sample. Open star clusters are ideal places within which to search for white dwarfs with disks as all cluster members have a known age. Therefore we can calculate the cooling age of any white dwarf, and hence the mass of the progenitor star. We recently investigated the white dwarf members of the moderately rich nearby Praesepe open cluster, measuring their effective temperatures and gravities \cite{casewell09, dobbie06}. We identified that WD0837+218 has a radial velocity that is inconsistent with it being a cluster member, but have included it in the sample here for completeness. All eleven white dwarfs have temperatures between 14 5000 K and 22 000 K with log g between 8.1 and 8.45, in the DAZ range, although the values for the magnetic white dwarf WD0836+201 should be regarded with some scepticism as the temperature fitting is unreliable. The above statistics suggest we should find a 0-2 DAZs in our sample. At a distance of 177$^{+10.3}_{-9.2}$ pc (as determined from Hipparcos measurements, \cite{mermilliod97}), Praesepe is one of the closest star clusters. It is slightly metal rich with respect to the Sun ([Fe/H] = +0.11, \cite{an07}). Indeed, as both the metallicity and the kinematics of Praesepe are similar to those of the Hyades, the former is often touted as a member of the Hyades moving group and therefore is assumed have an age comparable to the latter, $\tau$=625$\pm$50 Myr (e.g. \cite{claver01}). We note that this age for the Hyades was derived by comparing model isochrones generated from slightly metal enhanced (Z=0.024) stellar models which included moderate convective overshooting to the colors and magnitudes of a sample of cluster members selected using Hipparcos astrometric data \cite{perryman98}.
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\cite{kilic08} suggest that up to 25 per cent of externally polluted DAZ white dwarfs show an infrared excess indicative of a dust disk, and in the absence of prior knowledge of metallicity, \cite{farihi09} expect 1-3 per cent of all single white dwarfs with cooling ages less than $\approx$0.5 Gyr to possess dust disks. Hence, it is not necessarily surprising that we have not detected a dust disk in Praesepe. However, we are able to place constraints on the presence of substellar companions. These white dwarfs have evolved from late B-type main sequence stars (2.9-3.5 M$_{\odot}$ \cite{casewell09}), spectral types that are not observed in radial velocity searches for substellar and planetary companions and we have determined that a T8 brown dwarf cannot be detected in the UKIDSS and UFTI observations of WD0837+185. Such a brown dwarf has a mass of 25 M$_{\rm Jup}$ at 625 Myr. Hence, these limits can be combined with the results of radial velocity searches for substellar companions to place limits on the formation of brown dwarfs as binary companions to late B-stars. \begin{theacknowledgments} The United Kingdom Infrared Telescope is operated by the Joint Astronomy Centre on behalf of the Science and Technology Facilities Council of the U.K. \end{theacknowledgments}
| 10 | 12 |
1012.1143
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1012
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1012.1469_arXiv.txt
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We investigate the serendipitous X-ray source population revealed in {\sl XMM-Newton} observations targeted in the Galactic Plane within the region 315\deg $< l <$ 45\deg~and $|b|<$2.5\deg. Our study focuses on a sample of 2204 X-ray sources at intermediate to faint fluxes, which were detected in a total of 116 {\sl XMM-Newton} fields and are listed in the 2XMMi catalogue. We characterise each source as spectrally soft or hard on the basis of whether the bulk of the recorded counts have energies below or above 2 keV and find that the sample divides roughly equally (56\%:44\%) into these soft and hard categories. The X-ray spectral form underlying the soft sources may be represented as either a power-law continuum with $\Gamma \sim 2.5$ or a thermal spectrum with kT $\sim 0.5$ keV, with $N_H$ ranging from $10^{20-22} \rm~cm^{-2}$. For the hard sources, a significantly harder continuum form is likely, \ie $\Gamma \sim 1$, with $N_H = 10^{22-24} \rm~cm^{-2}$. For $\sim$50\% of the hard sources, the inferred column density is commensurate with the total Galactic line-of-sight value; many of these sources will be located at significant distances across the Galaxy implying a hard band luminosity $L_X > 10^{32} \rm~erg~s^{-1}$, whereas some will be extragalactic interlopers. A high fraction ($^{>}_{\sim}$90\%) of the soft sources have potential near-infrared (2MASS and/or UKIDSS) counterparts inside their error circles, consistent with the dominant soft X-ray source population being relatively nearby coronally-active stars. These stellar counterparts are generally brighter than J=16, a brightness cutoff which corresponds to the saturation of the X-ray coronal emission at $L_X = 10^{-3}~L_{bol}$. In contrast, the success rate in finding likely infrared counterparts to the hard X-ray sample is no more than $\approx15\%$ down to J=16 and $\approx25\%$ down to J=20, set against a rapidly rising chance coincidence rate. The make-up of the hard X-ray source population, in terms of the known classes of accreting and non-accreting systems, remains uncertain.
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\label{sec:intro} The brightest sources discovered in the first all-sky surveys conducted at X-ray wavelengths (\eg {\it Uhuru}, \citealt{Forman78}; {\it Ariel V}, \citealt{Warwick81}), were found to be X-ray luminous close binary systems powered by the accretion of matter onto a compact object. These sources, with intrinsic X-ray luminosities (L$_X$) typically in the range $10^{36-38}$ \ergs~in the 2--10 keV band, can be classed either as low-mass or high-mass X-ray binaries (LMXBs or HMXBs) depending on the nature of the non-degenerate star, the companion objects being either neutron stars or black holes. X-ray catalogues of the time also contained a few examples of other types of Galactic X-ray source, including X-ray bright supernova remnants, cataclysmic variables (CVs) and coronally-active binaries, such as RSCVn systems, albeit with inferred X-ray luminosities typically much less than $10^{36}$ \ergs. Later more sensitive X-ray surveys utilising imaging instruments operating in a somewhat softer spectral regime (\eg {\it Einstein}, \citealt{Hertz84}; {\it ROSAT}, \citealt{Voges99}) showed that the X-ray sky at low Galactic latitude is crowded, with nearby coronally-active stars and binaries dominating the source statistics in the soft ($<$ 2 keV) X-ray band (\citealt{Motch97}). Subsequently, an imaging survey in the Galactic plane carried out by {\it ASCA} (\citealt{Sugizaki01}) provided a first detailed view of the general population of faint Galactic X-ray sources in the 2--10 keV band. Since the impact of X-ray absorption by interstellar gas is greatly diminished above $\sim$ 2 keV, the {\it ASCA} survey was able to detect hard-spectrum sources with X-ray luminosities markedly lower than those of the classical Galactic X-ray binaries across a significant fraction of the inner Galaxy. More recently, the improved sensitivity, spatial resolution and energy range afforded by the {\it Chandra} and {\it XMM-Newton} observatories has provided the opportunity to revisit the issue of the faint Galactic source populations in both the soft and hard X-ray bands. For both missions, the fields-of-view of the on-board X-ray cameras are such that in a typical low-latitude pointing, in addition to the primary target, many tens of serendipitous sources are seen. Both missions have also invested observing time in studying specific regions of the Galactic plane, most notably the Galactic centre, through both deep pencil-beam observations (\citealt{Ebisawa01}; \citealt{Ebisawa05}; \citealt{Muno03}; \citealt{Muno04}; \citealt{Hong09}; \citealt{Rev09}) and mini-surveys in which wider angle coverage is achieved by the mosaicing of multiple (relatively short) observations (\citealt{Wang02}; \citealt{Hands04}; \citealt{Wijnands06}; \citealt{Muno06}; \citealt{Muno09}). The outcome is that 10 years post-launch both the {\it Chandra} and {\it XMM-Newton} archives contain a wealth of data relevant to Galactic X-ray sources at intermediate to faint flux levels, spanning a wide range of intrinsic luminosity. Recent results from {\it Chandra} and {\it XMM-Newton} demonstrate the potential of Galactic X-ray surveys, at readily accessible sensitivity limits, to detect a wide variety of source types. For instance, coronally-active binaries can be detected at distances up to 1 kpc or beyond (\citealt{Herent06}) and young stellar objects (YSOs) can be unveiled in regions of current star-formation. In the case of the latter, evolved protostars and TTauri stars with extreme coronal emission and hard X-ray spectra (kT$\sim$1-4 keV) can be detected in dense molecular clouds despite the large line-of-sight column density (\eg \citealt{Feigelson99}). Isolated Neutron Stars (ISNs), such as those discovered in the ROSAT survey (Haberl 2007), are radio-quiet objects located at distances of no more than a few hundred parsecs that display soft thermal X-ray spectra. CVs constitute a source class that may, potentially, be found in large numbers in sensitive Galactic X-ray surveys. In particular, it has been proposed that intermediate polars may account for a large fraction of the low L$_{X}$ sources which reside in the Galactic Centre region (\citealt{Muno06}) and through their integrated emission may account for a significant fraction of the hard Galactic X-ray Ridge emission (\citealt{Sazonov06}; \citealt{Rev06}). Relatively quiescent X-ray binaries with either low-mass (\eg black hole transients) or high-mass (\eg Be star) secondaries may also make a non-negligible contribution to the source statistics. However, in truth, our knowledge of the makeup of the Galactic X-ray source population at relatively faint levels is quite limited. This is certainly the case in the 2--10 keV band where, in principle, the visible volume extends to the edge of the Galaxy. In order to better define the various populations in terms of their space density, scale height and luminosity function, we need to characterise and, where possible, identify much larger samples of sources than are currently available. More comprehensively characterised source samples might also reveal how X-ray sources map onto structures such as the Galactic spiral arms, the thin and thick disc, the Galactic bulge and the mass concentration within $\sim$ 100 pc of the Galactic Centre. A range of astrophysical issues, for example, relating to the formation and evolution of close binary systems and accretion at low mass-transfer rates, might also be addressed. In the above context, the {\it Chandra} Multiwavelength Plane (ChaMPlane) Survey is aiming at a systematic analysis of low-latitude fields with the objective of measuring or constraining the populations of low-luminosity ( L$_{X}$ $_{\sim}^{>}~10^{31}$ erg s$^{-1}$) accreting white dwarfs, neutron stars, and stellar mass black holes in the Galactic plane and bulge (see \citealt{Grindlay05} for full details). The programme utilises {\it Chandra} X-ray data in combination with follow-up optical and IR photometric and spectroscopic observations, so as to maximise the number of identified sources and explore the populations thereby revealed. Recent publications from the ChaMPlane programme include a study of {\it Chandra} fields in the Galactic anti-centre (\citealt{Hong05}) and the Galactic centre and bulge ({\citealt{Laycock05}; \citealt{Koenig08}; \citealt{Hong09}; \citealt{vandenberg09}). In the case of {\sl XMM-Newton}, the {\it XMM-Newton} Galactic Plane Survey (hereafter XGPS; \citealt{Hands04}) has sampled a flux range which bridges the gap between the relatively shallow {\it ASCA} Galactic Plane Survey (\citealt{Sugizaki01}) and the sensitivity limits reached in deep {\it Chandra} pointings (\eg \citealt{Ebisawa01}). In a paper reporting the XGPS, \citet{Hands04} discussed the results from a programme including 22 pointings which cover a region of approximately 3 square degrees between 19\degree - 22\degree~in Galactic longitude and $\pm$0.6\degree~ in latitude. Subsequent optical follow-up observations of a representative sample of the brightest identified low-latitude hard band sources detected in the XGPS have recently been presented by \citet{Motch10}. In the present paper, we build on the XGPS studies of \citet{Hands04} by carrying out a systematic investigation of the X-ray source population seen at intermediate to faint fluxes in {\it XMM-Newton} pointings targeted at the Galactic plane. More specifically, we consider observations encompassing a narrow strip of the Plane towards the central quadrant of the Galaxy. In the next section of this paper, we give details of the set of {\sl XMM-Newton} observations which provide the basis of the study. In Section~3, we describe the selection criteria we employ to extract a clean ``serendipitous'' X-ray source sample from these fields, using the 2XMMi catalogue (\citealt{Watson09}) as the input database. Section 4 goes on to investigate various properties of the sample with a focus on the available X-ray spectral information. In section 5 we present an investigation of the likely near-infrared counterparts based on a cross-correlation of the X-ray sample with both the 2MASS and UKIDSS surveys. This leads on to a discussion of the nature of the soft- and hard-source populations in section 6, followed by a brief summary of our conclusions of this first paper (Paper I). In subsequent papers (in preparation), we will explore the average X-ray spectral properties of a subset of relatively bright sources drawn from our source sample (Paper II), the log N - log S curves for both the soft and hard source samples (Paper III), and the broad-band colours of likely counterparts to the X-ray sources (Paper IV).
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\label{sec:concl} We have used the 2XMMi source lists pertaining to 116 {\it XMM-Newton} observations targeted at the inner quadrant of the Galactic Plane, to construct a sample of serendipitous Galactic X-ray sources. The main properties of the 2204 sources which comprise the sample are: \begin{enumerate} \item The bulk of the sources have {\it total} count rates in the range 2-100 pn ct ks$^{-1}$ (0.5-12 keV). In the soft (0.5-2 keV) band, 2 pn ct ks$^{-1}$ corresponds to $F_X = 4 \times$10$^{-15}$ erg s$^{-1}$ cm$^{-2}$ (0.5--2 keV), assuming an absorbed power-law spectrum with $\Gamma$ = 2.5 and $\textit{N}_H$ = 10$^{21}$ cm$^{-2}$. The same count rate in the hard (2-12 keV) band equates to $F_X = 2.5 \times$10$^{-14}$ erg s$^{-1}$ cm$^{-2}$ (2--10 keV), in this case assuming $\Gamma$ = 1.0 and $\textit{N}_H$ = 3$\times$10$^{22}$ cm$^{-2}$. \item Using a characterisation based on whether the majority of the counts were recorded below or above 2 keV, the sample splits rather cleanly into 1227 soft sources and 977 hard sources. \item Both the broad-band hardness ratio (HR) distribution and the {\it Band Index} ($BI$) plots reflect a wide spread of underlying source spectra. For the soft sources, the X-ray spectra may be represented as either power-law continua with $\Gamma \sim 2.5$ or as thermal spectra with kT $\sim 0.5$ keV with $N_H$ ranging from $10^{20-22} \rm~cm^{-2}$. For the hard sources, a significantly harder continuum form is likely, \ie $\Gamma \sim 1$, with $N_H = 10^{22-24} \rm~cm^{-2}$. The sources with HR$>$0.8 have column densities commensurate with the total Galactic line-of-sight value. \item A high fraction ($^{>}_{\sim}$90\%) of the soft sources have potential NIR (2MASS and/or UKIDSS) counterparts inside their error circles, consistent with the dominant soft X-ray source population being relatively nearby coronally-active stars. In contrast, the success rate in finding likely NIR counterparts to the hard X-ray sample is no more than $\approx25\%$ down to J=20, set against a much higher chance coincidence rate. The make-up of the hard band population in terms likely contributors such as CVs, active binaries, relatively quiescent LMXBs/HMXB and other classes of object remains uncertain. \end{enumerate} In future papers we will explore the average X-ray spectral properties of a subset of relatively bright sources drawn from our source sample (Paper II), the log N - log S curves for both the soft and hard source samples (Paper III) and the broad-band colours of likely counterparts to the X-ray sources (Paper IV).
| 10 | 12 |
1012.1469
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1012
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1012.1972_arXiv.txt
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From a survey for redshifted \HI\ 21-cm and OH 18-cm absorption in the hosts of a sample of radio galaxies and quasars, we detect \HI\ in three of the ten and OH in none of the fourteen sources for which useful data were obtained. As expected from our recent result, all of the 21-cm detections occur in sources with ultra-violet continuum luminosities of $L_{\rm UV}\leq10^{23}$ \WpHz. At these ``moderate'' luminosities, we also obtain four non-detections, although, as confirmed by the equipartition of detections between the type-1 and type-2 objects, this near-50\% detection rate cannot be attributed to unified schemes of active galactic nuclei (AGN). All of our detections are at redshifts of $z\lapp0.67$, which, in conjunction with our faint source selection, biases against UV luminous objects. The importance of ultra-violet luminosity (over AGN type) in the detection of 21-cm is further supported by the non-detections in the two high redshift ($z\sim3.6-3.8$) radio galaxies, which are both type-2 objects, while having $L_{\rm UV} > 10^{23}$ \WpHz. Our 21-cm detections in combination with those previously published, give a total of eight (associated and intervening) \HI\ absorbing sources searched and undetected in OH. Using the detected 21-cm line strengths to normalise the limits, we find that only two of these eight may have been searched sufficiently deeply in OH, although even these are marginal.
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\label{intro} Although opaque to optical light, the dusty Universe is transparent to radiation at radio wavelengths, thus making the spectroscopic study of the 21-cm spin-flip transition of neutral hydrogen (\HI) a very useful tool in probing the far reaches of the cosmos. The low probability of the transition compounded by the inverse square law, renders \HI\ 21-cm currently undetectable in emission at redshifts of $z\gapp0.1$. However, in the absorption of radio waves emitted from background quasars, the line strength depends only upon the column density of the absorber and the flux of the background source. Therefore by using absorption lines, we can in principal probe \HI\ to redshifts of $z\sim50$ (or when the Universe was 1\% its present age), where the ionosphere begins to affect low frequency radio waves ($\lapp30$ MHz). With such observations we can address several outstanding questions in cosmology and fundamental physics: \begin{enumerate} \item Probe the Epoch of Re-ionisation -- when the first ever stars ignited, re-ionising the gas in the smaller cosmos (e.g. \citealt{cgfo04}). \item Determine the contribution of the neutral gas content to the mass density of the Universe (\citealt{kpec09,cur09a}). \item Measure any putative variations in the values of the fundamental constants of nature at large look-back times, to at least an order of magnitude the sensitivity provided by the best optical data (see \citealt{tmw+06}). This offers one of the few experimental tests of current Grand Unified Theories, thus having profound implications for modern physics. \end{enumerate} This latter point requires the comparison of the redshift of the 21-cm line with other transitions, which may be optical/ultra-violet [from singly-ionised metals, giving $\Delta(\mu\alpha^2 g_{p})/\mu\alpha^2 g_{p}$], where $\alpha$ is the fine structure constant, $\mu$ the electron-to-proton mass ratio and $g_{p}$ the proton g-factor, millimetre-wave [rotational transition of molecules, giving $\Delta(\alpha^2 g_{\rm p})/\alpha^2 g_{\rm p}$] or other decimetre transitions (see \citealt{cdk04} and references therein). Specifically, transitions arising from the hydroxyl radical (OH), which can also be intra-compared \citep{dar03}, thus avoiding possible line-of-sight effects which could mimic a change in the constants. Thus, highly redshifted \HI\ 21-cm and OH 18-cm absorbers are of great interest, although these are currently very rare, with only 73 \HI\ 21-cm absorption systems at $z\geq0.1$ known -- 41 of which occur in galaxies intervening the sight-lines to more distant quasars (see table 1 of \citealt{cur09a}), with the remainder arising in the host galaxies of the quasars themselves (see table 1 of \citealt{cw10}). In the case of OH, the situation is more dire with only five absorbers currently known \citep{cdn99,kc02a,kcdn03,kcl+05}. Four of these were originally found through millimetre-wave molecular absorption, although further surveys have proven fruitless (see \citealt{cmpw03}), which we suggest is due to the traditional optical selection of the sources: The target of choice in many previous surveys have been damped Lyman-$\alpha$ absorption systems (DLAs), since these are known to contain large columns of neutral hydrogen ($N_{\rm HI}\geq2\times10^{20}$ cm$^{-2}$, by definition) at precisely determined redshifts. Although 19 DLAs have been detected in the \MOLH\ Lyman and Werner UV bands (see \citealt{nlps08}, \citealt{jwpc09} and \citealt{sgp+10}), these are at molecular fractions well below the detection thresholds of current microwave instruments \citep{cmpw03}. Furthermore, the molecular abundances appear to be correlated with the colour of the background quasar in that the DLAs have molecular fractions of ${\cal F}\equiv\frac{2N_{\rm H_2}}{2N_{\rm H_2}+N_{\rm HI}}\sim10^{-7} - 0.3$ and $V-K\lapp4$, whereas the millimetre (and OH) absorbers have molecular fractions ${\cal F}\approx0.6 - 1$ and optical--near-infrared colours of $V-K\gapp5$ (see figure 3 of \citealt{cww09}). That is, not only are the radio-band absorbers redder than those of the optical-band, but there may be a correlation between the normalised OH line strength and optical--near infrared colour (\citealt{cwm+06}), although this requires a larger number of detections for confirmation. These points strongly suggest that the quasar light is reddened by the dust in the foreground absorber, which prevents the dissociation of the molecules by the ambient UV field. From this it is apparent that in order to detect redshifted molecular absorption with current radio instruments, targets must be selected on the basis of their optical and near-IR photometry, where we select the reddest objects. However, the obscuration responsible for the quasar reddening could be located anywhere between us and the quasar redshift (the three intervening systems are the strongest absorbers, see Sect. \ref{ohr}) and, although wide-band decimetre scans are more efficient than at millimetre wavelengths \citep{mcw02,cwmk03}, these are very susceptible to radio frequency interference (RFI). Therefore, in addition to our programme of using the wide-band spectrometer on the Green Bank Telescope (GBT) to perform 200 MHz wide frequency scans of the entire redshift space towards very red, radio-loud objects (see \citealt{cdbw07}), we are searching for \HI\ and OH absorption associated with the host galaxy of the quasar. Here we add the results of our recently completed searches for associated absorption and discuss these in the context of our previous search results \citep{cwm+06,cww+08}.
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\subsection{\HI\ results} \label{hir} In Table \ref{res} we summarise the line strengths/limits from these new searches for redshifted \HI\ 21-cm and OH 18-cm absorption. \begin{table*} \centering \begin{minipage}{143mm} \addtocounter{table}{1} % \caption{The \HI\ column densities, $N$, derived from the optical depths given in Table \ref{obs}. $T_{\rm s}$ is the spin temperature of the \HI~21-cm, $T_{\rm x}$ is the excitation temperature of the OH and $f$ the respective covering factor. $z$-range is the redshift range over which the column density limit applies (Table \ref{obs}), followed by the galaxy/quasar classification. The final columns give the AGN type and 1216\AA\ luminosity [\WpHz], determined/calculated as per \citet{cww+08}. } \begin{tabular}{@{}l l l c c c c c c @{}} \hline Source & $z_{\rm host}$ & Line & $N$ [\scm] & $z$-range & Class & Type & Ref & $\log L_{\rm UV}$ \\ \hline 0108+388 & 0.66847 & \HI\ & $2.3\pm0.2\times10^{19}.\,(T_{\rm s}/f)$& 0.66847 & Gal & 2 & L96 & 20.309\\ ... & ... &OH & $<3.6\times10^{13}.\,(T_{\rm x}/f)$ & 0.62987--0.67268 & ... &...& ... &... \\ 0131--001 &0.879 &\HI\ & $<5.7\times10^{16}.\,(T_{\rm s}/f)$ & 0.87488--0.88458 & QSO & --& -- & 20.221\\ ... & ... &OH & $<1.1\times10^{13}.\,(T_{\rm x}/f)$ & 0.87766--0.88232 & ... &...& ...&... \\ 0213--026 & 1.178 & \HI\ & $<2.9\times10^{18}.\,(T_{\rm s}/f)$ & 1.16624--1.18054 & QSO & 2 & D97a & 22.119\\ ... & ... & OH & $<2.4\times10^{13}.\,(T_{\rm x}/f)$ &1.11960--1.24545 &... &...& ...&...\\ 0454+066 & 0.405 & OH & $<3.6\times10^{14}.\,(T_{\rm x}/f)$ &0.40138--0.40873 & QSO & 2 & D97a & 21.567 \\ 0500+019 & 0.58477 & OH &$<1.2\times10^{13}.\,(T_{\rm x}/f)$ &0.58419--0.58736 & Gal & 2 & H03 & 20.367 \\ 0647+415 & 3.79786 & \HI\ &$<9.1\times10^{17}.\,(T_{\rm s}/f)$ & 3.78364--3.81379& Gal & 2 & D97b & 23.258 \\ ... & ... & OH &$<9.5\times10^{14}.\,(T_{\rm x}/f)$ & 3.76988--3.81423& ... &...&...& ...\\ 0847+37& 0.406818 & \HI\ & $<6.8\times10^{17}.\,(T_{\rm s}/f)$ &0.40183--0.41221 & Gal & 2& SDR7& 20.930\\ 1107--187 & 0.497 & \HI\ & $1.5\pm0.5\times10^{18}.\,(T_{\rm s}/f)^*$ & 0.48909 & Gal & 1 & D97a & 19.157\\ % ... & ... & OH & $<1.8\times10^{13}.\,(T_{\rm x}/f)$ & 0.47137--0.51372&... &...&... & ..\\ 1243+036 & 3.5699 & \HI\ & $<7.7\times10^{17}.\,(T_{\rm s}/f)$ & 3.55653--3.58388 & Gal & 2 & R97 & 23.382 \\%lum from SDSS phot. No SDSS spec ... & ... & OH & $<8.8\times10^{13}.\,(T_{\rm x}/f)$ & 3.56273--3.58607 & ... &...&... & ...\\ 1430--155& 1.573 & OH & $<8.8\times10^{13}.\,(T_{\rm x}/f)$ &1.56667--1.58345 & QSO& 1 & D97a& 21.790\\ 1504--166& 0.876 & \HI\ & $<2.2\times10^{17}.\,(T_{\rm s}/f)$ & 0.83385--0.88365 & QSO & 1 & H78& 22.361\\ 1535+004 & 3.497 & OH & $<1.9\times10^{14}.\,(T_{\rm x}/f)$ &3.48456--3.51125 & QSO & ---& -- & --- \\ 1654--020 & 1.99 & OH & $<3.3\times10^{13}.\,(T_{\rm x}/f)$ & 1.97924--1.99858 & Gal & 1 & D97a & 22.151 \\ 1706+006 & 0.449 & \HI\ & $<2.7\times10^{17}.\,(T_{\rm s}/f)$ & 0.43436--0.47902$^{\dagger}$ & Gal & 2 & D97a &19.838\\ ...& ...& OH & $<2.5\times10^{13}.\,(T_{\rm x}\/f)$ & 0.42388--0.46221 &...& ...&...&...\\ 2252--090 & 0.6064 & \HI\ & $2.0\pm0.4\times10^{19}.\,(T_{\rm s}/f)$ & $0.60748$ & Gal & 2 & D97a & 20.802\\ ...& ...& OH & $<1.0\times10^{13}.\,(T_{\rm x}/f)$ &0.60246--0.6162 &...& ...&... & ...\\%Sept 2352+495 & 0.23783 & OH & $<1.0\times10^{13}.\,(T_{\rm x}/f)$&0.20788--0.26921& Gal & 2 & L96 & 19.030 \\ \hline \end{tabular} {Notes: $^*$$N_{\rm HI}<1.1\times10^{17}.\,(T_{\rm s}/f)$ \scm\ at $z=0.48381-0.49898$ over the absorption free region if a false detection, $^{\dagger}$with RFI at $z = 0.453-0.459$.\\References: H78 -- \citet{hms78}, L96 -- \citet{lzr+96}, D97a -- \citet{dwf+97}, D97b -- \citet{dvva97}, R97 -- \citet{rvm+97}, H03 -- \citet{hsj+03}, SDR7 -- SDSS DR7.} \label{res} \end{minipage} \end{table*} For the 21-cm searches, we have obtained one, possibly two new detections, as well as confirming and improving upon a previous detection. From the top panel of Fig. \ref{2-N-L}, we see that all of the sources have been searched as deeply as in previous surveys and from the bottom panel, we see a range of 1216 \AA\ luminosities (given in Table~\ref{res}): \citet{cww+08} found a critical luminosity ($L_{\rm UV}\sim10^{23}$ \WpHz) at this wavelength, above which 21-cm has never been detected. All but two of the sample lie below this threshold, these being 0647+415 and 1243+036 (Sect. \ref{sect:od}), which, as our previous $z\sim3 - 4$ searches \citep{cww+08}, are above the critical luminosity due to their high redshifts % causing the selection of the brightest sources, despite their relatively faint magnitudes (Table \ref{obs}, cf. figure 5 of \citealt{cww09}). \begin{figure*} \centering \includegraphics[angle=270,scale=0.70]{2-N-L.eps} \caption{The scaled velocity integrated optical depth of the \HI\ line ($1.823\times10^{18}.\int \tau dv$) [top] and the ultra-violet ($\lambda\approx1216$ \AA) luminosity [bottom] versus the host redshift for the $z\geq0.1$ radio galaxies and quasars searched in associated 21-cm absorption. The filled symbols/hatched histogram represent the 21-cm detections and the unfilled symbols/unfilled histogram the non-detections, with the large coloured symbols designating the new results presented here (1107--187 is treated as a detection, Sect. \ref{det1107}). The shapes represent the AGN classifications, with triangles representing type-1 objects and squares type-2s ({\bf +} and {\sf x} designate an undetermined AGN type for a detection and non-detection, respectively). The legend shows the number of each AGN type according to the $L_{\rm UV}=10^{23}$ \WpHz\ partition.} \label{2-N-L} \end{figure*} These two new high redshift sources differ from the all of the other $L_{\rm UV}\gapp10^{23}$ \WpHz\ targets searched in 21-cm in that they are type-2 objects. Their inclusion increases the significance of the ultra-violet luminosity effect, with the binomial probability of 0 out of 19 detections occuring by chance being just $1.9\times10^{-6}$, if a 21-cm detection and non-detection are equally probable. Assuming Gaussian statistics, this corresponds to a significance of $4.76\sigma$. As mentioned in Sect. \ref{ss}, the distribution of 21-cm detections in radio galaxies and quasars is usually attributed to unified schemes of active galactic nuclei \citep{ant93,up95}, where, due to the edge-on torus of dense circumnuclear material, type-2 objects (usually galaxies) present a thick column of intervening gas along our sight-line and thus absorb in 21-cm (\citealt{jm94,cb95}), whereas type-1 objects (usually quasars) do not. Of the new detections, one is type-2 and one is type-1 (assuming that 1107--187 is a detection, Sect. \ref{det1107}), and contributing to a type-2 detection rate of 46\% and a type-1 rate of 48\%, these confirm our previous finding that unified schemes of AGN cannot be used to explain the incidence of 21-cm absorption in these objects (cf. \citealt{mot+01,pcv03,gss+06,gs06a}). That is, these results further support our suggestion the bulk of the cool gas is located in the main galactic disk, which is randomly oriented with respect to the torus of obscuring material invoked by unified schemes of AGN \citep{cw10}. \subsection{OH results} \label{ohr} Although OH was not detected in any of the sample, from the \HI\ detections\footnote{0108+388, 0500+019, 1107--187, 2352+495 (this paper), 0902+343 \citep{cb03}, J1124+1919, J1347+1217 and J2316+0404 \citep{gss+06}.} we can obtain normalised OH line strength limits in order to compare with the five detected OH absorbers. % \begin{figure*} \centering \includegraphics[angle=270,scale=0.70]{3-colour.eps} \caption{The normalised OH $^{2}\Pi_{3/2} J = 3/2$ (1667 MHz) line strength ($2.38\times10^{14}\int \tau_{_{\rm OH}}\, dv/1.82\times10^{18}\int\tau_{_{\rm HI}}\, dv$) versus the blue--near-IR (left), optical--near-IR (middle) and red--near-IR (right) colour. The filled symbols show the five known OH absorbers (circles -- associated, stars -- intervening absorbers, see \citealt{cwm+06} and references therein) and the unfilled symbols the \HI\ 21-cm detections with OH upper limits, colour coded by reference (green -- \citealt{cb03}, blue -- \citealt{gss+06} and red -- this paper). The probability of each distribution occuring by chance is shown, along with the associated significance (see main text). The line shows the least-squares fit to those also detected in millimetre-band molecular transitions for the optical--near-IR colours \citep{cwm+06}.} \label{OHoverH} \end{figure*} In Fig. \ref{OHoverH} we add the new results to the molecular line strength/optical--near-IR colour correlation found by \citet{cwm+06}, showing also the corresponding distributions using the blue and red magnitudes. Since the full-width half maxima (FWHM) of the OH lines are expected to be close to those of the 21-cm profiles \citep{cdbw07}, as per \citet{cww+08}, we have rescaled the OH column densities limits by $\sqrt{{\rm FWHM_{HI}}/\Delta v}$, in order to give the limit of a single channel ``smoothed'' to ${\rm FWHM_{OH} \approx FWHM_{HI}}$. This should give a more accurate estimate of the upper limit than quoting this per $\Delta v$ channel.\footnote{$\Delta v$ is the original resolution of the observations or the 10 \kms\ used in Table \ref{obs} and ${\rm FWHM_{HI}}$ is obtained from \citet{mir89,cmr+98,cb03,vpt+03,gss+06}, as well as this paper (Sect. \ref{hidet}).} From this, we see that, even if the reddening of the quasar light does occur within its host galaxy and not at some intervening redshift, according the to the five known OH absorbers, only 0500+019 and (from the $R-K$ plot) 0108+388 may have been searched sufficiently deeply. Although the limit is close to the expected detection threshold, 0500+019 is also undetected in HCO$^+$, to limits which are significantly stronger according to the $N_{\rm HCO+}$---$V-K$ correlation \citep{cwc+10}, and so perhaps the reddening of this source does not occur in the host galaxy but is the cause of some intervening absorber.\footnote{$<30$\% of the redshift space towards 0500+019 has been scanned for HCO$^+$ (\citealt{mcw02}), although the detection of OH does not ensure the detection of a millimetre transition \citep{kcl+05}.} Note finally, that the addition of these limits through the {\sc asurv} survival analysis package \citep{ifn86}, increases the significance of the $V-K$ correlation over that for the OH detections only ($S(\tau) = 1.96\sigma$, \citealt{cwm+06}). For the $B-K$ and $R-K$ correlations the significance is somewhat lower, although, due to the limited availability of these magnitudes for the known absorbers, this is not surprising being based upon only two or three detections.
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1012.1972
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1012
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1012.1696_arXiv.txt
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The IRAS source, 19312+1950, exhibits SiO maser emission, which is predominantly detected in evolved stars enshrouded by a cold molecular envelope. In fact, the mojority of the observational properties of IRAS~19312+1950 is consistent with the nature of an asymptotic giant branch (AGB) star or post-AGB star. Interestingly, however, some of the observational properties cannot be readily explained within the standard scheme of stellar evolution, and those are rather reminiscent of young stellar objects. In the present research we considered the evolutionary status of IRAS~19312+1950 as revealed by the VLBI and MERLIN observations in SiO, H$_2$O and OH maser lines. The double-peaked profile of the 22~GHz H$_2$O maser line is clearly detected, with the emission regions of its red and blue-shifted components separately located, leaving a space of about 10.9~mas between them. The kinematic properties of H$_2$O maser emission region appear to be more consistent with a bipolar flow rather than other interpretations such as the Keplerian rotation of a disk. The red-shifted component of the SiO maser emission, which exhibits a double-peak profile in previous single-dish observations, is clearly detected in the present interferometry, while the 1612~MHz OH maser line exhibits a complicated line profile consisting of a single strong peak and many weak, high-velocity spikes. The structure of OH maser emission region is partially resolved, and the kinematic properties of the OH maser emission region are reminiscent observations of a spherically expanding shell, even though the evidence is scant. Collectively, the maser observations described here provide additional support for the evolved star hypothesis for IRAS~19312+1950.
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The maser source, IRAS 19312+1950, exhibits various peculiarities. Its extended infrared nebulosity was first noticed in the 2MASS image by \citet{nak00}. The angular size of the near-infrared structure is roughly 30$''$. Soon after the initial identification of extended structure in the 2MASS image, a near-infrared (NIR) image with high angular resolution was obtained by the SIRIUS camera mounted on the University of Hawaii 2.2~m telescope \citep[see, Figure~1 in][]{deg04}. This image showed a point-symmetric structure with complicated filaments in its nebulosity. Since SiO maser emission has been detected toward IRAS~19312+1950 \citep{nak00}, the largely extended NIR structure is truly remarkable. This is because SiO maser sources, which are usually identified with an oxygen-rich (O-rich) asymptotic giant branch (AGB) star, rarely show extended structure that can be resolved by ground-based telescopes. \citet{nak04b} discuss the spectral energy distribution (SED) of IRAS~19312+1950, which exhibits a flat-topped (or moderately double-peaked) profile (in a range from $\lambda=1$~$\mu$m to 100~$\mu$m) with a weak absorption feature of silicate at 9.7~$\mu$m. They estimated a distance to the object by matching it to the absolute flux for a typical AGB star as 2.5--3.9 kpc [for reference, more recently, a trigonometric parallax distance to IRAS~19312+1950 was measured with VERA to be 3.8$^{+0.8}_{-0.6}$~kpc \citep{ima11}; this is roughly consistent with the distance given by Nakashima et al.~2004; in this paper we adopt 3.5~kpc as a representative distance if not otherwise specified]. This luminosity distance is not inconsistent with the kinematic distance estimated by assuming circular rotation for the Milky Way. However, a flat-topped SED is not usual in AGB stars, even though it is occasionally found in post-AGB stars \citep[see, e.g.,][]{uet00} and young stellar objects (YSOs). \citet{mur07} analyzed an $H$-band image taken with the CIAO camera mounted on the 8~m Subaru telescope. They constructed a radiative transfer model for the dust envelope by assuming a spherical geometry, to derive some of the fundamental stellar parameters for the central star; these are broadly consistent with an evolved stellar status. Further, \citet{nak04b} have estimated its mass-loss rate at $\sim10^{-4}$~M$_{\odot}$~year$^{-1}$ from their CO observations; this large mass-loss rate suggests that if the central star is an evolved star, it must be at a late-AGB (or early post-AGB) phase, where a dense superwind can take place \citep[in fact, such a large mass-loss rate has been identified in relatively young post-AGB stars; see, e.g.,][]{hri05}. IRAS~19312+1950 exhibits a pretty complicated chemistry. So far, 22 molecular species have been detected \citep{deg04,nak04b}. The profile of the molecular radio lines consists of two different kinematic components: broad-weak ($\Delta v \sim 30$ km~s$^{-1}$) and narrow-strong ($\Delta v < 5$ km~s$^{-1}$) components. The peak velocities of these two components are roughly same. \citet{nak04b} found that the broad-weak component is emitted from a relatively small region with a angular size less than $\sim15''$ and the structure of the narrow-strong component is clearly resolved even by a single-dish telescope beam. \citet{nak05} made interferometric observations using BIMA in the CO ($J=1$--0 and 2--1), $^{13}$CO ($J=1$--0 and 2--1), C$^{18}$O ($J=1$--0), CS ($J=2$--1), SO ($J_{K}=3_{2}$--$2_{1}$), and HCO$^{+}$ ($J=3$--2) lines, and the structure of both the broad-weak and narrow-strong components were spatially resolved. The BIMA observation revealed that the morphological and kinematic properties of the broad-weak component are explained by an expanding sphere, while those of the narrow-strong component are best interpreted as a bipolar flow. The angular size of the narrow-strong component is 5-6 times larger than that of the broad-weak component. This larger size of the narrow component further complicates the situation, because a system consisting of a smaller spherical wind and larger bipolar flow is not naturally explained within the standard scheme of the stellar evolution unless the bipolar flow had a faster velocity than the spherical component at some point in the past. One possible interpretation of IRAS~19312+1950 might be an O-rich AGB or a very young O-rich post-AGB star contingently embedded in an isolated small dark cloud \citep{deg04}. However, at the current moment, we have to say that we cannot be definitive about the true nature of this star. In this paper, we report the result of the first VLBI observations of IRAS~19312+1950 in the SiO and H$_2$O maser lines using the Very Long Baseline Array\footnote{The VLBA(HSA)/National Radio Astronomy Observatory is a facility of the National Science Foundation, operated under a cooperative agreement by Associated Universities, Inc.} (VLBA), High Sensitivity Array (HSA) and Japanese VLBI Network (JVN)\footnote{NRO and VERA/Mizusawa VLBI observatory are branches of the National Astronomical Observatory, an interuniversity research institute operated by the Ministry of Education, Culture, Sports, Science and Technology.}. We also report the result of radio interferometric observation of the OH maser line using the Multi-Element Radio Linked Interferometer Network (MERLIN). The purpose of this research is to aid our understanding of the true nature of IRAS 19312+1950 through the maser properties revealed by the interferometric observations. The outline of this paper is as follows. In section 2, we give details of the observation and data reduction. In section 3, we present the results of observations. In section 4, we discuss the kinematic properties of maser emission regions, and consider the evolutionary status of IRAS~19312+1950 on the basis of the present observations. Finally, the results are summarized in Section 5.
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The present observations in maser lines have revealed the kinematics of maser emission regions in the molecular envelope of IRAS~19312+1950 for the first time. As mentioned in Section~1, one of the most important scientific problems on this peculiar IRAS object is the identification of its evolutionary status. In this section, we first discuss the kinematic properties of the H$_2$O and SiO maser emission region, and consider how the OH maser properties can be understood within in the context of the stellar evolution. Then, we revisit the discussion about the evolutionary status of IRAS~19312+1950 on the basis of the present maser observations and previous observational results. \subsection{Properties of the H$_2$O and SiO Masers and its Kinematic Interpretation} The most notable results in the present observations are visible in Figure~2. The features seen in the top and middle panels of Figure~2 are clearly different from the maser-spot distribution of the H$_2$O maser of typical AGB stars. Since the usual H$_2$O maser emission of AGB stars predominantly comes from a spherically expanding envelope, the distribution of maser spots are usually much more extended or scattered, and often exhibits a ring (or patchy ring) structure \citep[see, e.g.,][]{bow93,bow94}. In the case of IRAS~19312+1950, the detected maser spots are lying roughly on a straight line as seen in the top panel of Figure~2. This maser spot distribution is quite different from that of typical AGB stars. Even though the existence of aspherical motions are occasionally suggested in AGB stars, based on VLBI observations in the H$_2$O and SiO maser lines \citep[see, e.g.,][]{dia03,nkg08}, even in such cases a spherical component is still dominant, and therefore the maser spot distribution is much more scattered compared with the case of IRAS~19312+1950. Then, how can we interpret the kinematics of the maser emission region? One possibility might be a Keplerian rotating disk showing a double-peak line profile, which originates in the two edges of the disk. If we assume that the distance to the object is 2.5--3.9~kpc \citep{nak04b} and Keplerian rotation, the central mass is estimated to be 4.4--6.7 M$_{\odot}$. This is not inconsistent with AGB mass (i.e., 0.8--8.0 M$_{\odot}$), and the existence of a Keplerian rotation disk is sometimes suggested in AGB/post-AGB stars and PNe \citep[see., e.g.,][]{ber00,cot04,buj05,nak05,bab06,der07}. However, Keplerian rotation disks in AGB stars present significant challenges in angular momentum transport, absent a binary companion. For IRAS~19312+1950, firstly the feature seen in the $p$-$v$ diagram (see, lower panel of Figure~1) does not show a Keplerian-rotation-like curve, which should show a point-symmetric pair of two arches. Moreover, if we assume a rotating disk, we should see emission at velocities close to the systemic velocity, because the material lying in the near- and far sides of the disk has no relative velocity to us, but no such emission is seen in the spectra. (The spectra given in Figure~1 have a gap at velocities close to the systemic velocity, but we have confirmed that there is no emission in this velocity range through previous single-dish observations.) Additionally, we also should see emission at the spatially-intermediate location between the blue and red-shifted components if we assume a rotating disk model \citep[see, e.g.,][]{bab06}, but no such emission is confirmed in Figure~2. For these reasons, it seems that a Keplerian rotation disk is most likely excluded. Another possibility might be a binary system consisting of two maser sources. The tentative detection of the SiO maser (mentioned in Section 3.2) suggests this idea, because the location of the tentative SiO emission (i.e., blue-shifted component), which coincides with the H$_2$O maser emission, is somewhat unusual for a single AGB star. Maser spots of the SiO, H$_2$O and OH maser lines usually exhibit a stratified distribution in an AGB envelope in order of decreasing excitation temperature (i.e., SiO masers locate in the innermost region, H$_2$O masers are surrounding the SiO maser emission, and OH maser locate in the outermost region). Therefore if the tentative detection of the SiO maser emission is real we need to consider the reason for the spatial coincidence, although we do note however that the relative location of different maser species can be affected by projection effects if there is a highly asymmetric distribution of circumstellar material. One possible explanation for the spatial coincidence of different masers and the kinematic properties of the H$_2$O emission region might therefore be a binary system including two maser stars (AGB stars or red supergiants, for example) emitting both H$_2$O and SiO masers. This is most likely not the case however, because the maser spot distribution again presents a problem: the straight-line distribution of maser spots seen in Figure~2 is not consistent with the maser spot distribution of AGB stars and red supergiants, which show more scattered distribution originated in a spherical envelope. Even if we assume that the straight-line distribution can be explained, for example, by tidal interaction, the total luminosity of the binary system including two maser stars raises further problems. Two AGB stars increase the absolute flux of the binary system up to at least 16,000 -- 20,000 L$_{\odot}$ (if we assume that red supergiants are included in the system, the total luminosity must be more than that; see, Section 4.3), and this large absolute flux causes an unreasonably large distance, which does not match up with the other physical parameter such as the kinematic and trigonometric parallax distances and infrared flux. Therefore the explanation of a binary system of two maser stars is ruled out here. If we try to understand the spatial coincidence of the SiO and H$_2$O masers (if this is real), we should consider the other possibilities: for example, a peculiar pumping mechanism different from that working in usual circumstellar envelopes of AGB stars. However, before going to such an astrophysical interpretation of SiO maser emission, we should confirm the detection of the blue-shifted component of the SiO maser line by doing further VLBI observations, because the detection is just tentative at this time. For the moment, a reasonable explanation of the kinematics of the maser emission region would be a bipolar outflow. Even though we have not very clearly confirmed a proper motion of maser spots in a 197-days period, this is presumably because the observing period is too short and/or the effects of the inclination angle. In Section 3.1, we mentioned the velocity gradient ($\sim$1.0 mas/km~s$^{-1}$) found in the blue-shifted component of the H$_2$O line. If we assume a Hubble-type flow (i.e., the motion in which the velocity increases in proportional to the distance from the central star), we can estimate the distance (and dynamical timescale) between the driving source of the jet and the blue-shifted component. However, the estimated distance, 18~mas, does not match up with the geometrical center of the maser spot distribution as seen in Figure~2. Presumably, the kinematics of the jet is not a simple Hubble-type flow, and the acceleration seems to change intricately with the time and distance. On the other hand, if we assume a ballistic motion of the jet and find a proper motion of maser spots, we can estimate the dynamical timescale. With the same method applied in our previous studies (see, e.g., Imai et al.~2002, Yung et al.~2010), as mentioned in Sect 3.1 we found a proper motion of only one maser-spot, which exhibits a proper motion of 0.64 mas along with the jet axis (see, the arrow in Figure 2). The dynamical time scale of the jet based on this proper motion (assuming a ballistic motion) is about 5.5 years. However, we should note that this value of a proper motion is not very reliable, because we seen a possible acceleration in Figure~3, and therefore because the identification of the maser spots over different epochs is not very reliable. To reveal the details of the jet kinematics, we need to monitor the proper motion of maser spots in a much longer time scale. \subsection{Properties of the OH Maser} According to a previous single-dish observation (B. M. Lewis 2000, private communication), the 1612~MHz line is clearly stronger than the 1665 and 1667~MHz lines, and therefore IRAS~19312+1950 should be classified into the type~II OH maser source \citep[][see, Figure~7: Note that in Figure~7 the Arecibo profiles of the 1665 and 1667~MHz lines are magnified in intensity]{hab96}. OH maser emission is often detected in YSOs, but the 1665 and 1667~MHz lines are always stronger than the 1612~MHz line in YSOs \citep[see., e.g.,][]{edr07}, and therefore the relative intensity ratio of the OH maser lines supports the evolved star status of IRAS~19312+1950. As well as the H$_2$O and SiO masers, however, the line profile of the 1612~MHz OH maser is also clearly different from the typical case of AGB, which usually show a double-peak profile \citep{dia85,siv90,dav93} caused by the approaching and receding sides of an expanding spherical envelope. Note that AGB and post-AGB stars occasionally show a single-peak profile, but such a single-peak profile usually can be interpreted as a part of a double-peak profile originated from a spherically expanding shell. In the case of IRAS~19312+1950, emission peaks of the 1612~MHz OH line are seen only in the blue-shifted side of the systemic velocity ($\sim$34--36~km~s$^{-1}$), and the line profile is very complicated, consisting of a single strong peak and multiple weak spikes as seen in Figure~7. Since \citet{nak05} detected a compact spherical flow (with a size of $~$3--5$''$) in their CO line mapping, the distribution of the OH maser spots might trace the inner part of the compact spherical flow. In fact, if we assume that we are looking at a part of a spherically expanding shell with the systemic velocity of $\sim$34--36~km~s$^{-1}$, the behavior of OH maser spots seen in the $p$--$v$ diagram (in PA$=$90$^{\circ}$) might be explained by a spherically expanding shell with an expanding velocity of 15--25~km~s$^{-1}$ \citep[in fact, the compact spherical flow detected in the CO lines has an expanding velocity of roughly 20~km~s$^{-1}$,][]{nak05}. Here we note that the size of an spherically expanding shell should be growing up as the velocity comes closer to the systemic velocity. However, a problem is that such a spherically-expanding-shell-like feature cannot be found at PA$=$0$^{\circ}$, suggesting the prospective spherical flow somehow has deficiency of a part of its structure. This asymmetric nature might be explained by the effect of a bipolar flow lying at the central part of the envelope, otherwise it might be explained by a torus-like structure (having its axis roughly in the direction of PA$=$90$^{\circ}$), which blanks out the spherical flow. Unfortunately, in the present case, it is quite difficult to check the sphericity of the emission region of the 1612 MHz line, because the angular resolution of the MERLIN observation is insufficient. Interestingly, the Arecibo profiles of the 1665 and 1667~MHz lines exhibit an intensity peak shifted from that of the 1612 ~MHz line. Since the OH main lines (i.e., 1665 and 1667~MHz lines) are excited closer to the central star than the 1612~MHz line in general \citep[see, e.g.,][]{cha94}, this velocity shift might suggest the systematic difference of the gas motions between the outer and inner regions. \citet{dea04} investigated the relation between mid-infrared colors and the line profile of the OH maser lines, and found that there is a likely evolutionary trend in the shape of the OH maser profile with some sources evolving from double-peaked to irregular (caused by a bipolar jet) profiles. The red mid-infrared color of IRAS~19312+1950 [$\log(F_{25}/F_{12})=0.5$; here, $F_{12}$ and $F_{25}$ are the IRAS flux densities at $\lambda=25$ and 12~$\mu$m, respectively] suggests that this object is lying at the very late stage of AGB or already at early post-AGB \citep{nak00}, and therefore it is not very strange even if an irregular profile and asymmetric nature caused by the interaction between a spherical AGB wind and a bipolar flow is seen in the OH maser properties, because well-developed bipolar jets are often found in post-AGB stars and planetary nebulae. \subsection{Considerations on the Evolutionary Status} The detection of the SiO maser emission seems to support the evolved star status of IRAS~19312+1905, because SiO maser sources are usually identified to evolved stars except for three, anomalous YSOs in star forming regions \citep[i.e., Ori IRc 2, W51 IRs 2, and Sgr B2 MD5; see, e.g.,][]{has86}. On the other hand, the largely extended nebulosity, rich-chemistry and large mass of IRAS~19312+1950 \citep{nak04b,deg04} have caused a problem in the identification of its evolutionary status. Even though we have repeatedly discussed this problem in our previous papers \citep{nak00,nak04a,nak04b,deg04,nak05,mur07}, the current data provide a fresh insight into this question. (For readers' convenience, in Table 2 we briefly summarize the discussions about the evolutional status in previous papers.) As we discussed in Sections 4.1 and 4.2, the kinematics of the maser emission region most likely can be explained as a bipolar outflow, but molecular bipolar flows are seen in a wide variety of objects: YSOs, AGB and post-AGB stars, proto-planetary nebulae (PPNe; note that some groups apply the word ``pre-planetary nebulae'' to PPNe to avoid confusion with proto-planetary disks in star forming regions) and planetary nebulae (PNe), and therefore the bipolarity of the molecular outflow, by itself, cannot be a definitive probe of the evolutionary status. However, the other properties of the maser emission seem to support the evolved star hypothesis, as we mentioned above: for example, the intensity ratios of the OH maser lines (i.e., classification into the type II OH maser sources) are consistent with the evolved star status, and the kinematical nature of the OH maser is reminiscent of an spherically expanding flow (although some asymmetric nature is found). Additionally, we note that there are no known OH maser sources toward low-mass YSOs \citep[see, e.g.,][]{gar99,edr07,sah07}, and therefore if IRAS~19312+1950 is a YSO, it must be related to a high-mass star forming region (SFR), suggesting the presence of an ultracompact H$_{\rm II}$ region with detectable continuum emission. However, the NRAO VLA Sky Survey at $\lambda=$20~cm \citep{con98} shows no continuum sources toward IRAS~19312+1950. In addition, Br$\gamma$ emission was also negative in a several-minute integration using the 2.3~m telescope at the Australian National University (P. R. Wood 2002, private communication). Even though there are three exceptional SiO maser sources in star SFRs \citep{nak00,nak05} as mentioned above, all the three sources exhibit a strong continuum emission of an ionized region, because those are lying in high-mass SFRs. Thus, we can say that the nature of IRAS 19312+1950 is clearly different from that of the SiO maser sources previously known in SFRs. If we assume IRAS 19312+1950 is an evolved star enshrouded by a cold envelope, what kind of evolved stars are appropriate for the true identity of this object according to its observational properties? Although red supergiants (RSGs) often emit SiO masers, we can exclude RSGs from the candidates, because its large luminosity (e.g., $>2\times10^{4}$~L$_{\odot}$, Wood et al.~1983; $>1.3\times10^{5}$~L$_{\odot}$, Groenewegen et al.~2009) causes an unreasonably large distance, which does not match up with other physical parameters such as the kinematic distance \citep{nak04b} and trigonometric parallax distance \citep{ima11}. Therefore, we should consider only intermediate- or low-mass evolved stars that have a main-sequence mass less than about 8.0~$M_{\odot}$. Although intermediate-mass stars finally reach to the PPN and PN phases via the AGB and post-AGB phases, PNe and PPNe harboring an ionized region also can be excluded from the possibilities, because no continuum emission is observed indicating an ionized region. Practically, possibilities are limited to AGB and post-AGB stars, which do not harbor a significant ionized region but can emit SiO maser emission. In our previous papers \citep{nak00,deg03,deg04,nak04a,nak04b,nak05}, we have mentioned the resemblance between IRAS~19312+1950 and OH~231.8+4.2 (Rotten Egg Nebula), which is a PPN with SiO maser emission. In fact, OH~231.8+4.2 exhibits a rich set of molecular species \citep{mor87}, SiO and H$_2$O maser emission \citep[see, e.g.,][]{des07} and an extended bipolar nebulosity \citep{mea03}. The infrared colors of OH~231.8+4.2 is also similar to that of IRAS~19312+1950 \citep[see, e.g.,][]{nak00}. The H$_2$O maser emission coming from a tip of a bipolar jet was mapped with VLBA \citep{des07}, showing a similar morphological and kinematic properties to IRAS~19312+1950. However, the nature of the SiO maser emission seems to be different in the two objects. In the case of OH~231.8+4.2, a M9-10 III Mira variable is lying at the center of the nebulosity, and the circumstellar rotating ring (or disk) around this Mira variable is considered to be the source of the SiO maser emission \citep{des07}. On the contrary, so far we have no observational evidence of a Mira variable in IRAS~19312+1950; for example, a monitoring observation in near-infrared band for roughly one-month period did not show any evidence of a time variation in intensity (Fujii et al. 2002, private communication) and the SiO maser properties does not show a rotation or spherical expansion, which is a prospective characteristic of circumstellar SiO maser emission of a Mira-type star. At the current moment, we believe that the nature of the SiO maser emission of IRAS~19312+1950 is different from that of OH~231.8+4.2. Here, we point out notable resemblance between IRAS 19312+1950 and a particular class of young post-AGB (or late-AGB) stars, based on the maser properties: namely the ``water fountains (WFs)'' -- that is young post-AGB (or late-AGB) stars exhibiting a velocity range of the H$_2$O maser emission larger than that of the OH maser emission \citep[see, e.g.,][]{ima02}. Since water fountains are harboring a high velocity, tiny molecular jet, which is usually detected in the H$_2$O maser line, the H$_2$O line shows a large velocity range. Tiny molecular jets found in WFs have been considered to be the onset of the formation of asymmetric PN morphology. Therefore, the kinematics of the tiny molecular jets in WFs has been frequently investigated using the VLBI technique \citep[see, e.g.,][]{ima02,vle06,yun10}. The primary characteristics of WF jets are its high expanding velocity (often larger that 100 km~s$^{-1}$) and highly collimated morphology, and in addition the jets occasionally exhibit a precessing/spiral pattern \citep{ima02,yun10}. OH maser emission is detectable from the mojority of know WFs \citep[a dozen of WFs are known so far; see, Table 1 in][]{ima07} except for IRAS 19134+2131 (for this object, no OH emission is detected; the WF status is confirmed only with the high velocity range of the H$_2$O maser emission). Since the nature of OH maser emission of WFs are basically same with that of AGB stars \citep[a typical expansion velocity of the OH line is 10--25 km~s$^{-1}$; see, e.g.,][]{tel89}, the emission seems to come from an spherical envelope, which is not yet affected by the bipolar jet \citep[a spherical distribution of OH maser emission is confirmed with VLBI observations; see, e.g.,][]{ima02}. In the case of IRAS~19312+1950, according to Figure~7, the expansion velocity of the OH component is roughly 15--25~km~s$^{-1}$ if we assume the systemic velocity of the object is 35~km~s$^{-1}$ (this expansion velocity is consistent with that of other known WFs). On the other hand, if we take into account the 0~km~s$^{-1}$ component (detected in previous observations; see, Section 3.1), the expansion velocity of the H$_2$O component must be larger than 25~km~s$^{-1}$. This situation is quite reminiscent of WFs, even though in the case of IRAS~19312+1950 only the blue-shifted side of the OH emission is detected. In fact, infrared colors of IRAS~19312+1950 are very similar to those of W43A, which is a proto-typical example of a water fountain \citep[see Figure~6 in][]{nak03}. There, as well as in IRAS~19312+1950, SiO maser emission, which exhibits a double peak profile, is detected, and the velocity separation of the two peaks is roughly 30~km~s$^{-1}$ \citep{nak03,ima05}; this SiO maser properties of W43A is similar to those of IRAS~19312+1950. Additionally, the large mass-loss rate estimated from CO observations \citep[$\sim10^{-4}$~M$_{\odot}$~yr$^{-1}$;][]{nak04b} is also consistent with the WF status, because the mass-loss rate of intermediate-mass AGB stars can often reach such a high mass-loss rate at late-AGB or early post-AGB due to the super wind \citep[see, e.g.,][]{hri05}. What is more, we see the difference of position angles of bipolar structure between different observational methods, as mentioned in Section 3.1. This fact might be due to a precessing motion, which is often seen in water fountains \citep[see, e.g.,][]{ima02}. The origin of the rich chemistry and large ambient mass of IRAS 19312+1950 is, however, still left unsolved. In order to consider the relation and resemblance between IRAS 19312+1950 and WFs, further investigations of WFs are required (for example, CO mapping, molecular line surveys and searching for new WFs), because the number of well-studied WFs are still very limited. Additionally, checking the sphericity of the emission region of the OH maser is also important; for this purpose, more sensitive (to observe 1665 and 1667 MHz lines) and higher-resolution (to resolve each maser emission source) VLBI observations of the OH maser lines are required. \subsection{Short Remarks on Another Possibility} As discussed in Section 4.3 and previous papers, in the standard scheme of stellar evolution, some observational properties of IRAS~19312+1950 cannot be explained unless we assume a rare situation (for example, a intermediate-mass evolved star embedded in a dark cloud). An interesting object related to this problem is the peculiar nova, V838~Mon. This object exhibits detectable SiO maser emission \citep{deg05}, and does not fall into any conventional categories of SiO maser sources. V838 Mon erupted in the beginning of 2002 January, exhibiting the spectrum of A--F-type supergiant around optical maximum, and then the spectral type has changed to that of an M-type supergiant \citep[therefore, bright in infrared bands,][]{mun02,cra03}. At the current moment, the merger of two stars, on the main sequence or evolving toward the main sequence, is touted as the most promising model to explain the observational properties of V838 Mon \citep{sok07,mas10}. This merging hypothesis might be an advantage when we explain the large ambient mass of IRAS~19312+1950, if the merged star exhibits a luminosity similar to a single AGB star (as a large luminosity causes a problem as discussed in Sections 4.1 and 4.3). In fact, \citet{kam08} observed the CO gas of V838 Mon, and the derived mass of the ambient material is roughly 25~M$_{\odot}$; this value is similar to the mass of IRAS~19312+1950 \citep[$\sim10$~M$_{\odot}$,][]{nak04b}. However, to trigger the merging event producing the V838 Mon-type object \citep[recently, the word of ``red novae'' has come into use for meaning this type of star, see, e.g.,][]{mas10}, the star should be lying in a star forming region or open cluster, in which main-sequence or young stars are tightly-packed. As we discussed in the previous papers, so far there is no clear evidence of star-forming activities around IRAS~19312+1950, but the census of the nearby stars might be useful to consider the possibility of the red nova origin of IRAS~19312+1950.
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1012.1696
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1012
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1012.3370_arXiv.txt
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We present the current status of two new fully automated reduction and analysis pipelines, built for the Euler telescope and the CORALIE spectrograph. Both pipelines have been designed and built independently at the Universidad de Chile and Universidad Catolica by the two authors. Each pipeline has also been written on two different platforms, IDL and Python, and both can run fully automatically through full reduction and analysis of CORALIE datasets. The reduction goes through all standard steps from bias subtraction, flat-fielding, scattered light removal, optimal extraction and full wavelength calibration of the data using well exposed ThAr arc lamps. The reduced data are then cross-correlated with a binary template matched to the spectral type of each star and the cross-correlation functions are fit with a Gaussian to extract precision radial-velocities. For error analysis we are currently testing bootstrap, jackknifing and cross validation methods to properly determine uncertainties directly from the data. Our pipelines currently show long term stability at the 12-15m/s level, measured by observations of two known radial-velocity standard stars. In the near future we plan to get the stability down to the 5-6m/s level and also transfer these pipelines to other instruments like HARPS.
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The CORALIE spectrograph is a proven instrument for the detection of extrasolar planets (aka exoplanets) around stars like the Sun via the radial-velocity (RV) method, with a number of exciting discoveries already having been made using this instrument (e.g. \citealp{queloz00}; \citealp{eggenberger06}; \citealp{segransan10}). Hunting for exoplanets with CORALIE requires knowledge and application of the so called \emph{simultaneous Thorium} technique, which does not employ any absorption cell in the light beam but instead is based around the concept of point spread function stablisation throughout the optical train of the instrument. Two fibres are used; fibre A monitors the star under observation and feeds the light to the spectrograph, whereas fibre B simultaneously feeds light from a Thorium-Argon (ThAr) gas lamp to the spectrograph, such that the ThAr lines can serve as a reference to monitor the drift of the spectrograph throughout the observation \citep[see][]{baranne96}. CORALIE has been used as the test bed for the successful ESO-HARPS spectrograph \citep{pepe00} but in contrast to HARPS it is mounted on a telescope with a primary mirror diameter of only 1.2m, it operates at around half the resolving power of HARPS and it does not actively stabilise environmental pressure changes that will cause the RV zero point to drift throughout an observing night due to changes in the refractive index of air. In spite of this, and as highlighted above, CORALIE is more than sufficient to discover and fully characterise planetary mass bodies with various orbital characteristics orbiting bright solar-type stars, particularly since the June 2007 hardware upgrade of the instrument that gives rise to a gain in magnitude of 1.5. We have began a number of RV projects that can be accomplished using CORALIE. One such project is part of the Calan-Hertfordshire Extrasolar Planet Search (CHEPS) which is a search for new short period gas giant planets that have a high probability of transiting their host star (see \citealp{jenkins08}; \citealp{jenkins09}). Another CORALIE project shall be the follow-up of transiting systems detected by the new HAT-South project \citep{bakos09}.
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We have given a status update on ongoing work to develop two new Doppler pipelines for the CORALIE spectrograph located on the ESO la Silla site in Chile. We have explained the pipeline steps as they currently stand and have highlighted a few steps we are currently working on to increase the stability and precision of our codes. Finally we have shown that the pipelines have thus far reached a precision of around 14m/s, sufficient to discover and characterise planets around stars like the Sun.
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1012.3370
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1012
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1012.0765_arXiv.txt
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In this short review, we summarize our present understanding (and non-understanding) of exoplanet formation, structure and evolution, in the light of the most recent discoveries. Recent observations of transiting massive brown dwarfs seem to remarkably confirm the predicted theoretical mass-radius relationship in this domain. This mass-radius relationship provides, in some cases, a powerful diagnostic to distinguish planets from brown dwarfs of same mass, as for instance for Hat-P-20b. If confirmed, this latter observation shows that planet formation takes place up to at least 8 Jupiter masses. Conversely, observations of brown dwarfs down to a few Jupiter masses in young, low-extinction clusters strongly suggests an overlapping mass domain between (massive) planets and (low-mass) brown dwarfs, i.e. no mass edge between these two distinct (in terms of formation mechanism) populations. At last, the large fraction of heavy material inferred for many of the transiting planets confirms the core-accretion scenario as been the dominant one for planet formation.
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\subsection{General overview} The realm of extrasolar planet discoveries now extends from gaseous giants of several Jupiter masses down to objects of a few Earth masses. Detailed models of planet structure and evolution have been computed by different groups (Fortney et al. 2007, Baraffe et al. 2008, Burrows et al. 2007, Leconte et al. 2009; see Baraffe et al. 2010 for a recent review). These calculations include various internal compositions, based on presently available high-pressure equations of state (EOS) for materials typical of planetary interiors. A detailed discussion and a comparison of these models can be found in Baraffe et al. (2008)\footnote{Models are available at http://perso.ens-lyon.fr/isabelle.baraffe/PLANET08/}. This latter paper also explores the effect of the location of the heavy element material in the planet, either all gathered at depth as a central core or distributed throughout the gaseous H/He envelope, on the planet's radius evolution. These different possible distributions of heavy elements can in some cases have an important impact on the planet's contraction. This paper also shows that the presence of even a modest gaseous (H/He) atmosphere hampers an accurate determination of the planet's internal composition, as the highly compressible gas contains most of the entropy of the planet and thus governs its cooling and contraction rate. In such cases, only the average internal composition of the planet can be inferred from a comparison of the models with the observed mass and radius determinations, for transiting objects. Objects below about 10 Earth-masses, globally denominated Super-Earth or Earth-like planets, on the other hand, are not massive enough to retain a significant gaseous atmosphere by gravitational instability (Mizuno 1980, Stevenson 1982, Rafikov 2006). For these objects, the lack of a substantial gaseous atmosphere allows a more detailed exploration of the planet's internal composition than for the gaseous planets (Valencia et al. 2007, Seager et al. 2007, Sotin et al. 2007). It should be kept in mind, however, that, even for these Super-Earth or Earth-like planets, present uncertainties in the EOS of the various heavy elements (e.g. H$_2$O, Fe) under relevant P-T interior conditions prevent an accurate determination of their internal composition (see e.g. Fig. 2 of Baraffe et al. 2008). The melting lines of water or Iron are not even known under such conditions, so the exact thermodynamics state of these elements is unknown. The situation should improve in the coming years with the advent of high-pressure experiments conducted with the high-power laser facilities developed in the US and in France. \begin{figure}[htbp] % \centering \resizebox{.7\hsize}{!}{\includegraphics{fig1.ps}} \caption{Mass-radius relationship from the stellar to the planetary regime, from one solar mass to one Saturn mass. The four curves display four isochrones, namely, from top to bottom, $10^8$ (dot), $5\times 10^8$ (short-dash), $10^9$ (long-dash) and $5\times 10^9$ (solid) yr. Some objects are identified on the figure, including the recent field M-dwarf/BD system NLTT 41135a,b (Irwin et al. 2010). The group of top 4 objects at $\sim 10\mjup$ includes Wasp\,14b, Hat-P-2b, Wasp\,18b and XO\,3b. } \label{fig:MR} \end{figure} Figure \ref{fig:MR} displays the overall mass-radius relationship in the stellar and substellar domains, from a solar mass down to a Saturn mass. The lines denote the low-mass star, brown dwarf and planet models of the Lyon group for 4 isochrones. The vertical dash-line corresponds to the mass limit to reach thermal equilibrium, i.e. balance between nuclear H-burning energy and gravitational contraction energy, $M_{\mathrm{HBMM}}=0.075\,\msol$ (Chabrier \& Baraffe 1997). This defines the limit between the stellar and brown dwarf domains. The general behaviour of this m-R relationship is discussed in detail in Chabrier \& Baraffe (2000) and Chabrier et al. (2009) and will not be repeated here, where we will focus on the brown dwarf and planetary part of the domain. It is important however, to point out the recent observations of the $61\,\mjup$ and $63\,\mjup$ transiting brown dwarfs WASP-30b (Anderson et al. 2010) and LHS 6343 C (Johnson et al. 2010), respectively, which remarkably confirm the predicted theoretical mass-radius relation in the brown dwarf domain (Chabrier \& Baraffe 2000). As seen in the figure and explored in detail in Baraffe et al. (2008) and Leconte et al. (2009), for several of these transiting planets, the observed mass-radius relation can be adequately explained by the planet "standard" evolution models mentioned in \S1.1. A typical case, for instance, is CoRoT-Exo-4b, a 0.72$\mjup$ planet whose 1.17 $\rjup$ radius is reproduced at the 1$\sigma$ level by a model including a 10$\mearth$ water core surrounded by a gaseous H/He envelope, i.e. a global $\sim 5\%$ mass fraction of heavy material, more than twice the solar value (Leconte et al. 2009)\footnote{Models are available at:\\ http://perso.ens-lyon.fr/jeremy.leconte/JLSite/JLsite/Exoplanets\_Simulations.html}. Choosing rock or a mixture of water and rock instead of water as the main component of the core only slightly changes this value. Several other transiting planet mass-radius signatures are well explained by standard models with moderate to high (up to $\sim 95\%$ for Neptune-mass planets) heavy element enrichment, as expected from the standard "core accretion" scenario for planet formation (see \S\ref{sec:formation}). \begin{figure}[htbp] % \centering \resizebox{.7\hsize}{!}{\includegraphics{fig2.ps}} \caption{Planetary radii at 4.5 Gyr as a function of mass, from 0.1 $\mearth$ to 20 $\mjup$. Models with solar metallicity (Z = 2\%) and with different amounts of heavy material (water, "rock" (i.e. olivine or dunite), or iron) are shown (Baraffe et al. 2008, Fortney et al. 2007). Solid curves are for non-irradiated models while dash-dotted curves correspond to irradiated models at 0.045 AU from a Sun. The positions of Mars, the Earth, Uranus, Neptune, Saturn and Jupiter are indicated by solid points, while the most recent transiting Earth-like (Corot-Exo-7b, GJ1214b) and Neptune-mass (Hat-P-26b, GJ436b, Kepler-4b, Hat-P-11b) planets are indicated by solid triangles. } \label{fig:MRth} \end{figure} Figure \ref{fig:MRth} focuses on the lowest mass part of the planetary domain, from 20 Jupiter masses down to Mars, i.e. going from gaseous giants to nearly incompressible matter. The figure displays the behaviour of the mass-radius relationship in this domain for various internal compositions, and highlights also the impact of stellar irradiation for a typical HD209458b-like system on the radius of a gas-dominated planet. Also indicated on the figure are the locations of the Solar System planets and of the recently discovered Earth-like\footnote{Note that Fig. 2 displays the revised 1$\sigma$ mass determination of Corot-7b (Pont et al. 2010).}, Super Earth or Neptune-like transiting objects. \subsection{The planet radius anomaly} \label{sec:anom} On the other hand, as seen in Figure \ref{fig:MR}, a large number of transiting planets exhibit a radius significantly larger than predicted by the theory, even when including irradiation effects from the parent star. Denoting $R_{\mathrm{irrad}}$ such a theoretical radius, the \textit{radius anomaly} of "Hot Jupiters" is thus defined as $(R_{\mathrm{obs}}-R_{\mathrm{irrad}})/R_{\mathrm{irrad}}$ (see e.g. Leconte et al. 2009, 2010c or Fig. 10 of Baraffe et al. 2010). Several physical mechanisms have been suggested to explain this radius anomaly. The most promising ones are discussed in details in Baraffe et al. (2010) and are quickly summarized below: - tidal heating due to circularization of the orbit, as originally suggested by Bodenheimer et al. (2001). This suggestion has been revisited recently by Leconte et al. (2010a) using orbital equations which are valid at any order in eccentricity (Hut 1981, see also Eggleton et al. 1998). Indeed, all the previous calculations addressing this issue were based on a tidal model valid only for nearly circular orbits, as developed initially for our Solar System planets (Goldreich \& Soter 1966). As rigorously demonstrated in Leconte et al. (2010a) and Wisdom (2008), such a model severely underestimates the tidal dissipation rate as soon as the (present or initial) eccentricity is larger than about 0.2-0.3. Using tidal equations valid at any order in eccentricity shows that tidal dissipation, although providing a substantial source of energy and - for moderately bloated planets - leading to the appropriate radius, cannot explain the very bloated objects such as HD 209458b (Leconte et al. 2010a, Hansen 2010). It should be stressed that the aforementioned limitation of the so-called constant-$Q$ model does not have anything to do with the description of the dissipation mechanism in the star or the planet, as often misunderstood, but stems from the truncated expansion of the orbital equations. A dedicated discussion of these tidal effects is given in this volume by Leconte et al. - downward transport of kinetic energy originating from strong winds generated at the planet's surface by a small amount ($\sim 1\%$) of absorbed incident stellar radiation (Showman \& Guillot 2002). Although appealing, such a mechanism still needs to be correctly understood. Simulations par Burkert et al. 2005, for instance, do not produce such a dissipation (see Shownan et al. 2008 for a recent review). The identification of a robust mechanism for transporting this energy deep enough is still lacking and an accurate (so far missing) description of the (small-scale) dissipative processes in such natural heat engines is mandatory to assess the validity and the importance of this mechanism for hot-Jupiters (see e.g. Goodman 2009). - ohmic dissipation in the ionized atmosphere of hot-Jupiters (Batygin \& Stevenson 2010). This scenario has received some support from recent 3D resistive MHD atmospheric circulation simulations of HD 209458b's weakly ionized atmosphere (Perna et al. 2010). According to these simulations, for magnetic field strengths $B\gtrsim 10$ G, enough ohmic dissipation occurs at deep enough levels (from a few bars to several tens of bars) to affect the internal adiabat and to slow down enough the planet's contraction to yield a significantly inflated radius. These results have to be confirmed by further studies, as quantifying the impact of non-ideal MHD terms and induced currents in numerical simulations is a challenging task. - enhanced opacities ($\sim 10\times$ the solar mixture) in hot-Jupiter atmospheres, stalling the planet's cooling and contraction (Burrows et al. 2007). It should be stressed, however, that if the planet H/He envelope's global {\it metallicity} is enhanced at this level, the increased molecular weight will cancel or even dominate the opacity effect and will lead to a similar or smaller radius than the one obtained with solar metallicity (Guillot 2008). This scenario is thus so far an ad-hoc procedure and enhanced sources of opacities for a global solar-like metallicity must be identified, both theoretically and observationally. - inefficient (layered or oscillatory) convection in the planet's interior, due to a gradient of heavy elements either inherited from the formation stages or due to core erosion during the planet's evolution (Chabrier \& Baraffe 2007). Although layered convection is observed in many situations in Earth lakes or oceans, due to the presence of salt concentrations (the so-called thermohaline convection), it remains unclear, however, whether this process can occur and persist under giant planet interior conditions. In contrast to the first three scenarios, the last two ones (i) do not invoke an extra source of heating in the planet but rather an hampered output flux during the evolution, leading to a slower contraction rate, (ii) do not necessarily apply to short-period, irradiated planets only but could possibly also occur in planets at large orbital distances. None of these mechanisms has either been confirmed or ruled out so far. Note that they are not exclusive from each other and it might be possible that they all contribute, at some level, to the puzzling anomalously large radius problem.
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In this review, we have examined our present understanding and non-understanding of exoplanet formation, structure and evolution. The results can be summarized as follows: \begin{itemize} \item the theoretical mass-radius relationship in the brown dwarf and planetary regime seems to be confirmed by recent radius determinations of transiting massive brown dwarfs. When the object's radius is smaller than the one predicted for a gaseous body with solar composition, this m-R relationship enables us to distinguish planets from brown dwarfs in their overlapping mass domain and thus provides a key diagnostic to identify these two distinct populations. In other cases, the diagnostic remains ambiguous and the very nature of the transiting object can not be determined. \item Present models of planet interior structure and evolution stand on relatively robust grounds and enable us to infer with reasonable confidence the gross internal structure and composition of these objects. Uncertainties in the EOS of various elements under the relevant conditions, however, prevent a detailed determination of this composition. \item a large fraction of {\it gas dominated} transiting planets still exhibit a radius significantly larger than predicted by the models. Several physical mechanisms have been proposed to solve this "radius anomaly" problem but, so far, no firm conclusion about which one, if any, of these mechanisms is the correct one has been reached. \item tidal energy dissipation due to circularization of the orbit, in the planet's interior, although providing a significant extra source of energy to the planet, has been shown not to be sufficient to explain the radius of the most bloated planets, including HD 209458-b. Indeed, when properly calculating the orbital evolution equations in case of a finite present or initial eccentricity, tidal dissipation is shown to occur too quickly during the planet's evolution to explain its present radius. Although a proper treatment of the contribution of dynamical tides is presently lacking, the equilibrium tide contribution calculated with the complete tidal equations still provides a lower limit for tidal dissipation and must be correctly calculated. Interestingly enough, recent observations of spin-orbit misalignment for planets orbiting F stars seem to point to a tidal dissipation in the star, and thus to a dynamical evolution of the system, which depends on the stellar mass, more precisely on the size of the stellar outer convection zone (Winn et al. 2010). \item an update of the presently discovered transiting systems confirms the previous analysis of Levrard et al. (2009). Only a handful of these systems have enough total angular momentum to reach an orbital equilibrium state. For the vast majority of these systems, the planet experiences ongoing orbital decay and will eventually merge with the star, with the dynamical evolution timescale for the orbit semimajor axis and the stellar spin and obliquity being essentially the lifetime of the system itself (Levrard et al. 2009, Matsumura et al. 2010). \item departure of both the parent star and the transiting planet from sphericity, because of either rotational or tidal forces, affects both the depth of the transit light curve and the planet's radius determination, and leads to an {\it underestimate} of this latter. This bias must be corrected to get a proper determination of the planet's genuine equilibrium radius, the one calculated with 1D structure models. \item observations of Hat-P-20-b, if confirmed, show that planets form up to at leat about 8 $\mjup$ and thus the brown dwarf and planet mass regimes very likely overlap. Therefore, there is no common mass limit between these two populations of astrophysical bodies, stressing again the inadequacy of the definition put forward by the IAU. \item the large number of transiting planets whose radius implies a substantial fraction of heavy material strongly supports the core accretion scenario formation for planets. In this scenario, the planet embryo originates from accretion of solids onto a core in the protoplanetary disk, leading eventually to dynamical accretion of gas dominated material above about 10 $\mearth$. Conversely, this same large metal enrichment excludes the gravitational instability scenario as the dominant formation mechanism for planets. \end{itemize}
| 10 | 12 |
1012.0765
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1012
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1012.0881_arXiv.txt
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Understanding cosmic acceleration mechanisms, such as jet formation in black holes, star collapses or binary mergers, and the propagation of accelerated particles in the universe is still a `work in progress' and requires a multi-messenger approach, exploiting the complementarities across all possible probes: ultra-high energy cosmic rays (UHECR), gamma-rays and neutrinos. In this report I will summarize some of the IceCube results concerning searches for astrophysical neutrino point sources and diffuse fluxes from populations of sources widely distributed in the sky or from the interactions of protons on the cosmic microwave background producing the GZK cut-off in the cosmic ray spectrum. I will compare the results to other neutrino telescopes and to astrophysical models of neutrino production in sources. Another unresolved question concerns the nature of dark matter. Indirect searches have the opportunity to observe where it is located in the universe through the observation of secondary photons, neutrinos or antiparticles such as positrons and antiprotons. The potential for the search of neutrinos from the annihilation of WIMPs in IceCube is greatly enhanced by the addition of more compact strings, the DeepCore. \vspace{1pc}
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\label{sec1} The IceCube Neutrino Observatory will be composed of a deep array of 86 strings holding 5,160 digital optical modules (DOMs) deployed between 1.45 and 2.45 km below the surface of the South Pole ice and a surface array, IceTop, for extensive air shower measurements on the composition and spectrum of cosmic rays (CRs). The strings are typically separated by about 125 m with DOMs vertically apart by about 17 m along each string. A DOM is a spherical, pressure-resistant glass vessel, housing a Hamamatsu R7081-02, 252~mm-diameter phototube (PMT) \cite{PMT}. IceCube construction started with a first string installed in the 2005--6 season and will be completed in the austral summer of 2010--11. The running configuration consists of 79 strings. Six of these strings hold DOMs with higher quantum efficiency (HQE-DOMs). Laboratory measurements of the HQE-DOM quantum efficiency relative to a standard reference 2-inch PMT indicate values around $\sim 30\%$ compared to the standard DOMs with $\sim 20\%$ at about 390 nm \cite{deepcorepaper}. The HQE-DOMs are located at smaller spacing of about 70 m horizontally and 7 m vertically in the layout shown in Fig.~\ref{fig:fig1}. Together with two additional strings that will be deployed during the last construction season and seven standard strings of IceCube they make up DeepCore, designed to enhance the physics performance of IceCube below 1 TeV such as dark matter searches. Fig.~\ref{fig:fig2} shows the layout of the 86 strings of the full detector and of the 40-strings configuration, that was used for most of the analysis results I will summarize here. Fig.~\ref{fig:fig3} shows the occupancy (the fraction of events with a DOM that detected at least 1 photon on a string where at least seven other DOMs have also detected photons) versus depth \cite{deepcorepaper}. In a uniform medium, such a plot would exhibit a decreasing trend due to muon energy losses in the ice, but in IceCube it reflects the convolution between energy losses and ice depth dependent optical properties. As a matter of fact, the effect of a big dust layer, most probably with a component of volcanic nature, is visible between -2,150 and -1,950 m. Moreover, the deep ice is so transparent that light propagates up to $\sim 600$~m and the hit rate is higher than it would be in a uniform medium. The plot explains why no DeepCore DOM is installed in the dust region. The plot also shows the increase in efficiency of the HQE-DOMs at the expense of some increase of the dark noise rate (from about 500 Hz for normal DOMs to about 700 Hz in HQE-DOMs). \begin{figure}[t] \includegraphics[width=3.in,height=3.5in]{deepcore.eps} \caption{A schematic layout of IceCube DeepCore. The upper diagram shows a top view of the string positions in relation to current and future IceCube strings. It includes two additional strings, situated close to the central DeepCore string. The lower diagram shows the instrumented DeepCore region with the surrounding IceCube strings. \label{fig:fig1}} \end{figure} The charge and time of Cherenkov photons induced by relativistic charged particles passing through the ice sheet is detected and the PMT signal is digitized with dedicated electronics included in the DOMs \cite{WF}. The DAQ hardware is made of two waveform digitization systems using Analog Transient Waveform Digitizers (ATWD) that record the details of the first 400 ns of the waveform at 300 MSPS and fADCs (fast ADC) using a 40 MSPS commercial ADC chip covering up to 6.4 $\mu$s. A digitization cycle is initiated by a discriminator trigger with threshold set at a voltage corresponding to about 1/4 photoelectron. When this happens, the FPGA starts ATWD and fADC digitization on the next clock edge and data are sent to the surface when a coincidence of at least one other hit in the nearest or next-to-nearest neighboring DOMs within $\pm 1$~$\mu$s is satisfied (Hard Local Coincidence - HLC). After the 40-string run, Soft Local Coincidences have been added to the acquisition and then to reconstructions. In this case even isolated hits are stored saving only fADC information. \begin{figure}[t] \includegraphics[width=3.in,height=3.in]{ic40layout} \caption{\label{fig:fig2} Overhead view of the 40-string configuration, along with the additional strings that will make up the complete IceCube detector. } \end{figure} Various triggers are used in IceCube. Most of the results shown here are based on a simple multiplicity trigger requiring that the sum of all HLC hits in a rolling time window of $5\,\mu$s is above 8 (SMT8). The duration of the trigger is the amount of time that this counter stays at or above 8 as the time window keeps moving. Once the trigger condition is met, all local coincidence hits are read in a readout window of $\pm 10\,\mu$s for the 40-string run and of $^{+6}_{-4} $ $\mu$s (to reduce the noise rate) in the 79-string run. IceCube triggers primarily on down-going muons at a rate of about 950~Hz in the 40-string configuration and about 1.8 kHz in the 79-string one. Variations in the trigger rate by about $\pm10\%$ are due to seasonal changes. During the austral summer the atmosphere above the South Pole gets thinner and the probability of pions generated in the CR induced cascades to decay rather than interact increases - and hence the muon rate \cite{Tilav:2010hj}. \begin{figure}[ht] \includegraphics[width=3.in,height=2.in]{occupancy} \caption{\label{fig:fig3} In situ measurement of the occupancy versus depth. The lower black stars are for DOMs with standard PMTs, the upper red crosses are for HQE-DOMs.} \end{figure} Direction of events can be reconstructed using the time of hit PMTs and the amplitude as well, and the energy can be inferred exploiting the stochastic energy loss properties of muons and the charge measurement. The muon energy resolution on an event by event basis is limited (about an uncertainty of a factor of 2 on the muon energy at the core of the detector between $10^4 -- 10^8$~GeV) but the capability of reconstructing spectra (as shown in Sec.~\ref{sec2}) is helped by the wide lever arm of many orders of magnitude in energy to which IceCube is sensitive. The point spread function is shown in Fig.~\ref{fig:fig4}. It is determined by the kinematic angle between the neutrino and the reconstructed secondary muon and by the intrinsic angular resolution of the detector (limited mainly by scattering of photons in the ice and by the PMT transit time spread). Two bins in energy are shown: for higher energies tracks are longer across the array and can be better reconstructed. The high energy muons produced by muon neutrino interactions point back to the neutrino source with degree-level accuracy making IceCube a `neutrino telescope'. Neutrino events would cluster around the point source producing them within an uncertainty of less than $1^o$ for 50\% of the events with $E_\nu \in [10,100]$ TeV and within $0.6^o$ for 50\% of the events with $E_{\nu} \in [1,10]$ PeV for 40 strings. The absolute pointing accuracy has been confirmed studying a deficit of muons due to cosmic rays blocked by the Moon disc. The Moon shadow detection, initially reported in \cite{Moon}, is seen at the level of 6.76$\sigma$ for 14 lunar cycles with the 40-string configuration. New data increase this significance. IceCube is also sensitive to shower events induced by neutral current and electron/tau charged current interactions. These events have no pointing capability but the energy can be reconstructed. \begin{figure}[ht] \includegraphics[width=3.2in,height=2.2in]{PSF_ic40_ic86} \caption{\label{fig:fig4}Point spread function (fraction of events included in the angle shown on the x axis between the neutrino simulated direction and the reconstructed muon one) for the 40 and 86-string configurations and for two energy bins at the final level of cuts for the point source search \protect\cite{jon}.} \end{figure} \begin{figure} [t] \includegraphics[width=3.2in,height=2.5in]{ForPaper_ZenithBig_figconf.eps} \caption{\label{fig:filtering} Distribution of the reconstructed cosine zenith of tracks at trigger level, L1, and final cut level for point source analysis for data and simulation of atmospheric muons (using the composition model in \cite{Hoerandel:2002yg}) and atmospheric neutrinos \cite{Barr:2004br}. The deficit of Monte Carlo (MC) horizontal events compared to data is most likely caused by a deficiency in knowledge of the CR composition in the region around the knee that would require a harder component of protons or of heavier elements \protect\cite{jon}.} \end{figure} In this report, I will focus on results for astrophysical sources of neutrinos and on dark matter searches. Many results concern the 40-string configuration in operation from 2008 April 5 to 2009 May 20. Over the entire period the detector ran with an uptime of 92\%, yielding 375.5 days of total exposure. Deadtime is mainly due to test runs during and after the construction season dedicated to calibrating the additional strings and upgrading data acquisition systems. The data and MC comparison is summarized in Fig.~\ref{fig:filtering} for different levels: trigger level, online filtering (L1) that runs at the South Pole with very basic reconstruction and cuts that reduce for 40-strings the trigger rate of about 1~kHz to 22~Hz and final level of cuts for the point source analysis \cite{jon}. These cuts are based on variables such as the quality and the angular error of the reconstruction, the number of hits due to unscattered Cherenkov light from the track and the total charge detected in one event. Another intriguing subject, covered in another report in these proceedings \cite{simona}, is the observation of anisotropies at large ($\gtrsim 60^o$) and intermediate scale ($\sim 20^o$) in the muon flux produced by cosmic ray (CR) interactions on the atmosphere. \begin{figure}[htb] \includegraphics[width=3.in,height=2.8in]{diffuse.eps} \caption{The flux of atmospheric $\nu_\mu+\bar{\nu}_{\mu}$ measured by AMANDA-II (squares) \protect\cite{julia} and 2 analyses with 40 strings of IceCube, an unfolding (triangles) \protect\cite{Warren} and a forward folding \protect\cite{Sean} (preliminary region between lines showing the range of uncertainty) is compared to atmospheric neutrino calculations (high line \protect\cite{Honda:2006qj}; low line \protect\cite{Barr:2004br}). Horizontal lines are 90\% c.l. upper limit to an $E^{-2}$ muon neutrino flux for AMANDA-II \cite{julia} (807 d), ANTARES (334 d with 9-12 line configurations) \protect\cite{Spurio} and 40 strings of IceCube (375.5 d). The WB limit \protect\cite{Waxman:1998yy} is shown together with some models on GRBs \protect\cite{Waxman:1998yy,Razzaque} and an example of an AGN model \cite{Stecker} rejected at 5$\sigma$ c.l. by the 40-string limit. }\label{fig:diffuse} \end{figure}
| 10 | 12 |
1012.0881
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1012
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1012.0848_arXiv.txt
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{The Taurus-Auriga association is perhaps the most famous prototype of a low-mass star forming region, surveyed at almost all wavelengths. Unfortunately, like several other young clusters/associations, this T association lacks an extensive abundance analysis determination.} {We present a high-resolution spectroscopic study of seven low-mass members of Taurus-Auriga, including both weak-lined and classical T Tauri stars designed to help robustly determine their metallicity. } {After correcting for spectral veiling, we performed equivalent width and spectral synthesis analyses using the GAIA set of model atmospheres and the 2002 version of the code MOOG.} {We find a solar metallicity, obtaining a mean value of [Fe/H]=$-$0.01$\pm$0.05. The $\alpha$-element Si and the Fe-peak one Ni confirm a solar composition. Our work shows that the dispersion among members is well within the observational errors at variance with previous claims. As in other star forming regions, no metal-rich members are found, reinforcing the idea that old planet-host stars form in the inner part of the Galactic disc and subsequently migrate.} {}
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The Taurus-Auriga (Tau-Aur) association is the nearest large star-forming region (SFR), with a close distance of $d$=140 pc (e.g., Kenyon et al. 1994). Given its lack of massive O/B stars, this T association in the past 50 years has become a standard region to study the low-mass star formation processes. Tau-Aur is widely spread across the Northern sky ($\sim$100 square degrees, Scelsi et al. 2007a) and contains hundreds of low-mass members; Rebull et al. (2010) found 148 new candidate members, of which 34 were confirmed by spectroscopic follow-up. The bulk of stars presents an average age of $\sim$1 Myr (e.g., Brice\~{n}o et al. 1999). The whole region has been observed at almost all wavelengths, from infrared to X-rays (e.g., Itoh et al. 1996; Brice\~no et al. 1999; Davis et al. 2008; Luhman et al. 2010). Kenyon \& Hartmann (1995, hereafter KH95) provided a thorough investigation of the stellar population, from class 0 (proto-stars) to class II/III (i.e. classical and weak-lined T Tauri stars, respectively -hereafter CTTs and WTTs). They analysed optical and infrared photometric observations and presented colour-magnitude and colour-colour diagrams, luminosity and mass functions, along with information on the near-infrared excess and accretion properties. Lithium abundance studies of this association have been also performed over the past decade, such as, e.g., Basri et al. (1991), Magazz\'u et al. (1992), Mart\`{i}n et al. (1994), and Sestito et al. (2008), who presented a (re)assessment of Li~{\sc i} abundances in WTTs and CTTs. G\"{u}del et al.~\cite{gud} carried out an extensive X-ray survey (XMM-{\it Newton Survey of the Taurus Molecular Cloud} --XEST) covering an area of $\sim$ 5 square degrees and concentrating mainly on the higher stellar density regions. Whereas all these studies have helped to characterize the stellar population, Tau-Aur shares with other nearby SFRs the lack of an accurate abundance analysis. The determination of chemical composition of SFRs is instead critically important to a variety of astrophysical issues, in both planetary and stellar contexts, as we previously discussed in our pilot project focusing on the Orion association (D'Orazi et al. 2009, hereafter D09). At variance with OB associations, whose abundances can provide an independent test of the so-called {\it triggered} star formation scenario (e.g., Blaauw 1991) and indicate whether there is local enrichment, we would expect to measure a very homogeneous composition for T associations. As is well known, giant planets are preferentially found around metal-rich main sequence stars (Santos et al. 2004, and reference therein). The natural question hence arises as to whether in the first epochs of planetary formation (discs of T Tauri stars are commonly accepted as planet birthplaces) the probability of hosting a giant planet depends on the star's metallicity. On the other hand, and very interestingly, several studies (e.g., Luhman 2004) have found significant differences between the initial mass function (IMF) derived for the Tau-Aur and those derived for denser systems containing massive members, e.g., the Trapezium, with the former containing a too small number of brown dwarfs according to the standard IMF (Brice\~no et al. 2002) and a surplus of late-K and early-M stars (Luhman et al. 2009). These differences can be attributed to different conditions of the environment where the stars were formed. Some studies have claimed that the IMF slope may depend on metallicity, in the sense that metal-rich environments tend to produce more low-mass stars (M $< 0.7 \rm M_{\odot}$) than metal-poor systems (Larson 2005; da Rio et al. 2009 and references therein). The obvious question is: could metallicity play a role in the lower fraction of very low-mass stars and brown dwarfs detected in Taurus with respect to Trapezium (Brice\~no et al. 2002)? More generally, we may ask whether there is a difference between chemical compositions of low-mass and high-mass SFRs or, whether, SFRs represent a chemically homogeneous class of objects? \\ Available information on Tau-Aur abundances comes from the previous works by Padgett (1996) and Santos et al. (2008). First, Padgett (1996) analysed eight WTT members and derived a solar composition, i.e. an average value of [Fe/H]=0.01$\pm$0.13 (rms), with a significant scatter among members ranging from $-$0.11 to +0.22. Santos et al. (2008) presented the iron abundances of three Taurus members: in this case, the spread in metallicity is also large, varying from $-$0.18 to +0.05 with a mean value of [Fe/H]=$-$0.07$\pm$0.12 (rms). Given the quite large star-to-star variation obtained by both studies, a new accurate abundance analysis for this association is warranted, to minimize the impact of the observational uncertainties and ascertain whether the metallicity scatter among members is indeed real or not. In a different framework, Scelsi et al. (2007b) obtained coronal abundances for a sample of X-ray bright members of Taurus, finding that the iron abundance in the corona is significantly (by a factor of $\sim 5$) lower than the solar photospheric value. More generally, their coronal abundances revealed a pattern in agreement with those obtained in previous studies. Specifically, X--ray observations of young and/or active stars have unveiled a so--called FIP (first ionization potential)-dependent abundance trend: low FIP elements (FIP$\leq$10 eV, like iron, silicon, nickel) are under-abundant in the corona with respect to the photospheric values (the so-called inverse-FIP effect, for further details see Brinckman et al. 2001; Audard et al. 2003; Scelsi et al. 2007b). As stressed by Scelsi et al., however, this FIP-related pattern for Tau-Aur, as well as for other young stars, was obtained by comparing the coronal abundances to the solar photospheric values, rather than the stellar photospheric abundances. Different or no FIP effects were indeed found when considering the correct stellar photospheric abundances (e.g., Sanz-Forcada et al. 2004; Maggio et al. 2007). Determining the photospheric abundances of Tau-Aur members with measured coronal abundances thus appears critical to confirm whether a FIP-related effect is present in these stars or not. In this work, we present a high-resolution spectroscopic study based on seven low-mass members of the Tau-Aur association. The sample was selected specifically to determine accurate abundances, by discarding, e.g., fast rotators, strong accretors, and/or binary stars. Most importantly, this is one of the first sample that also includes T Tauri stars that have not yet dissipated their circumstellar discs and reveal moderate accretion from the surrounding material. Details of the observations and data reduction are given in Section~\ref{obstau}; Section~\ref{analysistau} describes the estimate of and correction for spectral veiling, along with the abundance analysis procedure. The results and the scientific implications are illustrated in Sections~\ref{results} and~\ref{discu}, while in Section~\ref{summary} a conclusive summary is given.
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\label{summary} We have presented elemental abundances of seven stars, both CTTs and WTTs, belonging to the Tau-Aur association. We have found a very homogeneous, solar metallicity for this T association, deriving a mean value of [Fe/H]=$-$0.01$\pm$0.05. Both [Si/Fe] and [Ni/Fe] also exhibit solar ratios and agree very well with the observed abundance pattern of thin disc stars at the same metallicity. In contrast to the previous determination of abundances in Tau-Aur, which spanned a wide range in [Fe/H], we conclude that the internal dispersion in metallicity for this association is very similar to values derived in older open clusters. Along with our previous project in Orion and the available estimates in the literature, all the SFRs surveyed to date seem to share a similar chemical composition, suggesting a uniform ISM in the solar surroundings at the present time. In this context, we note that no metal-rich members have been detected in all the analysed young associations: this could validate the idea that metal-rich planet-host stars in the solar circle were formed in the inner disc and subsequently moved to their current location. Their not-so-high metallicity might reflect the rather low frequency of giant planets in the solar neighbourhood, at variance with that of the inner Galactic regions. However, given that only a small number of SFRs have been chemically characterised so far, and a small number of stars per region, further investigations are mandatory to definitely help us address this controversial, but intriguing issue.
| 10 | 12 |
1012.0848
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1012
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1012.2884_arXiv.txt
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Dynamically cold stellar streams are ideal probes of the gravitational field of the Milky Way. This paper re-examines the question of how such streams might be used to test for the presence of ``missing satellites'' --- the many thousands of dark-matter subhalos with masses $10^5-10^7 \rm M_\odot$ which are seen to orbit within Galactic-scale dark-matter halos in simulations of structure formation in $\Lambda$CDM cosmologies. Analytical estimates of the frequency and energy scales of stream encounters indicate that these missing satellites should have a negligible effect on hot debris structures, such as the tails from the Sagittarius dwarf galaxy. However, long cold streams, such as the structure known as GD-1 or those from the globular cluster Palomar 5 (Pal 5) are expected to suffer many tens of direct impacts from missing satellites during their lifetimes. Numerical experiments confirm that these impacts create gaps in the debris' orbital energy distribution, which will evolve into degree- and sub-degree- scale fluctuations in surface density over the age of the debris. Maps of Pal 5's own stream contain surface density fluctuations on these scales. The presence and frequency of these inhomogeneities suggests the existence of a population of missing satellites in numbers predicted in the standard $\Lambda$CDM cosmologies.
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Standard $\Lambda$CDM models of the Universe allow us to explain structure formation on large scales. However, they predict an order of magnitude more dark-matter subhalos within the halos of typical galaxies than the number of known satellite galaxies orbiting the Milky Way \citep{Klypin1999a,Moore1999a,Diemand2007a,Springel2008a}. Recent, large-area stellar surveys have discovered dozens of new satellite galaxies, most notably using the Sloan Digital Sky Survey \citep[SDSS, e.g.][]{Willman2005a,Belokurov2006a,Belokurov2007a, Zucker2006a,Irwin2007a,Koposov2007a,Walsh2007a} but the number discrepancy between simulated dark-matter subhalos and observed satellite populations is still significant. This discrepancy can partially be explained by accounting for the the incomplete sky-coverage of SDSS and the distance-dependent limit on this survey's sensitivity to low-surface brightness objects \citep{Koposov2008a,Tollerud2008a}. Indeed, models which take this into account and consider diffuse, (i.e. undetectable) satellite galaxies can reconcile the number counts for subhalos \citep{Bullock2010b}. However, when they impose the suppression of stellar populations in low mass subhalos (which have masses below $5 \times10^8 \rm \msun$) the number of undetected galaxies significantly declines and the prediction of numerous purely dark-matter subhalos less massive than $5 \times 10^8 \rm \msun$ remains. There could be a genuine absence of ``missing satellites" in the inner halo due to destruction by disk shocks, as illustrated in the calculations \citet{DOnghia2010a}. However, note that these analytic descriptions of disk shocking based on the energy criterion are known to overestimate disruption rates of subhalos significantly\citep{Goerdt2007a}. Once these destructive effects are accurately accounted for, proof of the existence (or lack) of these ``missing satellites'' could provide an important constraint on the nature of dark matter, which sets the minimum scale for the formation of dark-matter subhalos \citep[e.g.][]{Hooper2007a}. Along with the discovery of new satellite galaxies, SDSS has also uncovered a multitude of stellar structures in the Milky Way halo from disrupting globular clusters or satellite galaxies. In many cases, the debris is dynamically cold and distributed narrowly in space \citep{Odenkirchen2001a,Belokurov2006c,Lauchner2006a,Grillmair2006a,Grillmair2006c, Grillmair2006d,Grillmair2009a}. Such cold stellar streams should be sensitive probes of the gravitational potential. On global scales, they can be used to constrain the radial profile, shape and orientation of the Milky Way's triaxial dark-matter halo \citep[e.g.][]{Johnston1999a,Ibata2002a,Johnston2005a,Binney2008a,Eyre2010a,Koposov2010a,Law2010a}. The presence of dark-matter subhalos would add asymmetries to the global potential over a range of smaller scales which will perturb these cold streams or even destroy them. Hence, if the missing satellites do exist they will add random uncertainties to any stellar-dynamical assessment of the global potential. Gravitational lensing has been suggested to be a useful tool to probe the presence of subhalos \citep{Chiba2002a,Metcalf2002a,Chen2003a,Moustakas2003a,Metcalf2004a,Keeton2009a,Riehm2009a,Xu2009a}. These investigations conclude that flux ratio anomalies in lensed images or lensing time delays could be caused by dark-matter subhalos, though the constraints are limited by our knowledge of the spatial distribution of subhalos. However, this method is only applicable to the most massive and most centrally concentrated dark-matter halos, and not to galaxies like the Milky Way more generally. The effect of dark-matter subhalos on stellar streams has been explored in several previous studies. \citet{Ibata2002a} showed that debris from the destruction of a $10^6 \rm M_{\odot}$ globular cluster should be affected by heating due to repeated close encounters of subhalos and concluded that this effect could be detectable with future astrometric surveys. Moreover, \citet{Quinn2008a} found that the inhomogeneities seen in Pal 5's tidal tails could not be accounted for in simulations evolved in a smooth potential. For the streams of larger satellites like the Sagittarius dwarf galaxy (hereafter Sgr), \citet{Johnston2002a} found that although stars in the debris are scattered by encounters with dark matter subhalos, the thickness of the current Sgr stream could be explained as being due to the Large Magellanic Cloud alone. \citet{Siegal-Gaskins2008a} tested the additional influence of different host potentials on debris from satellite galaxies and pointed out that while subhalos can shift the positions of streams and cause clumpy structures, the shape of the halo potential and orbital path can have an overall comparable effect. Most recently, \citet{Carlberg2009a} modeled a simplified stream on a circular orbit and concluded that dynamically old ($>$ 3Gyr) stellar streams cannot survive in the presence of subhalos with the masses and numbers predicted by $\Lambda$CDM. These previous works point to stellar streams as perhaps the most powerful way to find the missing satellites. However, none of these investigations separated the effect on streams of the known (and therefore uninteresting!) satellites with masses $> 10^8 \msun$ \citep[][and references therein]{Bovill2009a,Bullock2010a} from those that are ``missing'' (the pure dark-matter subhalos). In this study we construct a framework for understanding stream inhomogeneities by first isolating and dissecting the characteristics of disturbances caused by dark-matter subhalos alone. In contrast to previous work, which looked at the overall response of cold streams to the complete $\Lambda$CDM subhalo mass spectrum, we look at the expected frequency, influence and characteristic observable signatures of subhalos in each mass decade separately. We also contrast the response of different streams to the same masses, from ribbons such as Pal 5 to the giant stream from Sgr. Our twin aims are: (i) to understand with which streams we are most likely to be able to conclusively prove the existence or absence of missing satellites: and (ii) to learn how signatures of missing satellites that are apparent in streams might be interpreted. It should be noted that the discovery of very cold streams from globular clusters has inspired discussions of how the intrinsic properties of stellar streams themselves could cause inhomogeneities in their density distributions \citep{Kupper2008a,Kupper2010a,Quillen2010a} and these self-induced fluctuations could confuse the conclusive association of observed disturbances with dark-matter subhalo interactions. Our own work is also motivated by these current observations which contain tantalizing suggestions of non-uniformity in some cold stellar streams \citep[e.g. the structure known as GD-1 and those from the globular cluster Pal 5, see][]{Odenkirchen2003a,Grillmair2006b,Grillmair2006c,Koposov2010a}, as well as the prospect of the density of these streams being mapped more extensively (and accurately) in space and velocity with observations in the near future. Such observations could potentially distinguish between the effect of subhalo encounters on different mass scales as well differentiate these signatures from non-uniformities due to intrinsic stream dynamics. We first review our understanding of the properties of dark-matter structures and stellar stream evolution in smooth potentials in \S \ref{back.sec}. We use this understanding to make order-of-magnitude estimates for the frequency and effect of encounters of stellar streams with structures of different masses in \S \ref{analytic.sec}. We then go on to illustrate these expectations with numerical experiments in \S \ref{numerical.sec} as well as discuss the observational signatures of these encounters in \S \ref{discussion.sec}. We summarize our conclusions in \S 6.
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\label{conclusion.sec} The goal of this study was to test if a dynamically cold stellar stream could survive in the presence of dark-matter subhalos and if so, to characterize observable features in the stream produced by subhalo encounters. In particular we were interested in finding signatures of subhalos that would be considered ``missing satellites'' rather than those that we already know exist. We conclude that: \begin{enumerate} \item The mere existence of cold stellar streams does not imply the absence of missing satellites --- dynamically cold streams can survive for many Gyears even when bombarded by subhalos. \item Those streams that are observed should contain the imprint of past direct impacts from subhalos in the form of gaps in surface density and discontinuities in velocities. \item The frequency and scale of the gaps is dependent on the mass spectrum of subhalos and the properties of the stream itself. In the case of Pal 5, there should be observable fluctuations at degree and sub-degree scales due to $\sim 90$ encounters with subhalos in the mass range $10^5-10^7 M_\odot$ (i.e. the missing satellites), while distant encounters with larger subhalos produce less observable effects. Hotter stellar streams, such as Sgr's debris are large enough to hide the signatures of the many encounters it suffers with missing satellites. \item Current observations of Pal 5's stream show that its surface density profile contains fluctuations on comparable scales to those predicted in our simulations, and could be interpreted as direct proof of the existence of missing satellites. \end{enumerate} Our study has built on previous works to look at: effects of individual encounters with different mass scales; the distinct signatures of subhalos in different mass decades; the integrated influences of subhalo encounters; and comparisons to the current observational data. However, while we have succeeded in outlining the expected scales of signatures of different decades of the $\Lambda$CDM mass spectrum in stellar streams, there are still some open questions that need to be answered before we are confident enough to say that we have definitely found the missing satellites. In particular, self-consistent simulations of Pal 5 disruption including stream self-gravity and continuous mass loss are needed to clarify to what extent subhalo signatures could be confused by these effects. With such simulations in hand, it would then make sense to develop a multi-dimensional statistic to compare to all available data in a more robust way. The promise of current and near-future efforts to map Pal 5 and other stellar streams are strong motivators for this work.
| 10 | 12 |
1012.2884
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1012
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1012.0309_arXiv.txt
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It is shown that a tensor-to-scalar ratio close to $r = 0.03$, which can be observed by Planck, is realized in supersymmetric hybrid inflation models with TeV-scale soft supersymmetry breaking terms. This extends our previous analysis, which also found $r \lesssim 0.03$ but employed intermediate scale soft terms. Other cosmological observables such as the scalar spectral index are in good agreement with the WMAP data.
| 10 | 12 |
1012.0309
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1012
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1012.1413_arXiv.txt
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Widefield surveys have always provided a rich hunting ground for the coolest stars and brown dwarfs. The single epoch surveys at the beginning of this century greatly expanded the parameter space for ultracool dwarfs. Here we outline the science possible from new multi-epoch surveys which add extra depth and open the time domain to study.
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Widefield multi-epoch surveys have been a crucial discovery tool for ultracool dwarfs. Photographic plate data was used by \cite{LHS} to identify nearby cool stars by their proper motion. As technology evolved these plates were digitised (\citealt{Hambly2001}, \citealt{Monet2003}) leading to more discoveries of cool nearby stars. The study of brown dwarfs did not begin with widefield digital infrared and red optical surveys, but datasets such as 2MASS (\citealt{Skrutskie2006}) and SDSS (\citealt{York2000}) massively expanded the sample of brown dwarfs. Works such as \cite{Chiu2006}, \cite{Reid2008} and \cite{Burgasser2004} provided large samples and objects from these surveys were used to define the new L and T spectral classes (\citealt{Kirkpatrick1999}, \citealt{Burgasser2006}). Now a new generation of widefield surveys are expanding into relatively unexplored regions, identifying extremely cool objects, finding nearby brown dwarfs by their trigonometric parallax and opening the time domain to the search for variability. Here we outline the leading surveys currently in operation and those about to come online.
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Widefield surveys will continue to play a leading role in the discovery and study of ultracool dwarfs. From the widefield infrared surveys with UKIDSS, VISTA and WISE, to multi-epoch optical surveys with Skymapper, Pan-STARRS, PTF and in the future LSST, such surveys will identify the coolest objects, large samples for population studies and many interesting variable sources. A summary of the surveys discussed in this paper can be found in Figure~2.
| 10 | 12 |
1012.1413
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1012
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1012.3620_arXiv.txt
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{ Relativistic jets are a common feature of radio loud active galactic nuclei. Multifrequency observations are a unique tool to constrain their physics.} { We report on a detailed study of the properties of the jet of the nearby BL Lac object PKS 2201+044, one of the rare cases where the jet is detected from radio to X-rays.} {We use new adaptive optics near-IR observations of the source, obtained with the ESO multi-conjugated adaptive optics demonstrator (MAD) at the Very Large Telescope. These observations acquired in Ground-Layer Adaptive Optics mode are combined with images previously achieved by HST, VLA and Chandra to perform a morphological and photometric study of the jet.} { We find a noticeable similarity in the morphology of the jet at radio, near-IR and optical wavelengths. We construct the spectral shape of the main knot of jet that appears dominated by synchrotron radiation.} { On the basis of the jet morphology and the weak lines spectrum we suggest that PKS 2201+044 belongs to the class of radio sources intermediate between FRIs and FRIIs.}
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Radio loud active galactic nuclei (AGN) are characterized by the presence of relativistic jets mainly detected in radio band. In some cases, the jets can also be observed in the visible and at higher frequencies. In particular, this is so for some nearby BL Lac objects. They represent a subclass of AGN characterized by weakness of lines, luminous, rapidly variable non-thermal continuum, significant polarization, strong compact flat spectrum radio emission, and superluminal motion (e.g. Ulrich et al. 1997, Giroletti et al. 2004). Similar properties are also observed in flat spectrum radio quasars and these two types of AGN are often grouped together into the class of Blazars. Additional evidence of the presence of relativistic jets in BL Lacs is given by their strong gamma emission (Abdo et al. 2010). However, because of the close alignment of the jet with the line of sight, it is very difficult to distinguish it unless the angular resolution is sufficiently high. In fact, the detection of the jet at the different frequencies depends both on the limited angular resolution and on the short lifetime of the high energy electrons producing the non- thermal emission of the jet. The advent of high resolution imaging facilities such as HST and Chandra allowed systematic searches for angular resolved multiwavelength counterparts of radio jets in Blazars (Scarpa et al. 1999a) and also the detection of jets in radio-loud quasars at high redshift (Schwartz et al. 2000, Sambruna et al. 2002, Tavecchio et al 2007). Moreover, the use of adaptive optics assisted imaging on Large Telescope started to permit to investigate from the ground the near-IR jet properties (Hutchings et al. 2004, Falomo et al. 2009). In this paper we present near-IR images acquired with an innovative adaptive optics (AO) camera of PKS 2201+044, the nearest BL Lac object (z= 0.027). Morphological and photometric properties are discussed together with the observations in radio, optical, and X-ray bands (Scarpa et al. 1999b, Sambruna et al. 2007). We adopt the concordance cosmology with $\Omega_m$=0.3, $\Omega_\lambda$=0.7 and H$_{0}$= 70 km s $^{-1}$ Mpc$^{-1}$. In this scenario, at z=0.027 1 arcsec corresponds to 0.54 kpc. We defined $\alpha$ as S $\sim$$\nu^{-\alpha}$ where S is the flux density at frequency $\nu$.
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In this paper, we presented high resolution near-IR images of the jet of the nearly BL Lac object PKS 2201+044 using an innovative adaptive-optics device (MAD) built as a demonstrator for multi-conjugated AO imaging obtained with MAD-VLT instrument. These new data, together with the analysis of data previously obtained in optical (by HST) , radio (by VLA and MERLIN) and X-ray band, have provided us with insight into the jet associated with this nearby extragalactic source. The main results from this study are:\\i) the morphology of the jet is very similar at radio, near-IR and optical frequencies with knot A being the only substructure detected at all wavelengths and resolutions;\\ ii) near-IR MAD observations obtained in a GLAO mode give a good mapping of the diffusive emission intra-knots enhancing a conical shape;\\ iii) the emission from the knot A is dominated by single synchrotron component;\\ iv) we derived for the component A and C an estimate for the synchrotron age, 8.6$\times$10$^{-2}$ Myrs and 6.1$\times$10$^{-2}$ Myrs respectively;\\ v) a detailed study (Fig. \ref{peak}) reveals that knot A emission is not homogeneous, in particular in direction SW-NE. This multiwavelength study of the jet of PKS 2201+044 shows strong similarity with the jets of two other well studied BL Lacs, 3C 371 (Sambruna et al. 2007) and PKS 0521-365 (Falomo et al. 2009). These objects are very peculiar being the only three BL Lacs for which the jet is detected in optical bands. It is worth to note that for all these objects weak broad emission lines are present in their optical spectra (Sbarufatti et al. 2006, Falomo et al. 2009, Sambruna et al. 2007). These suggest that they are intermediate sources between FR I and FR II, having multiwavelength FRIs properties (Ledlow $\&$ Owen 1996) but showing broad optical emission lines that FRI sources are generally inefficient to produce (Baumm et al. 1995). We argue that the presence of broad optical emission lines could be due to intermediate properties of their broad line regions. Future investigations of other similar objects showing jet would clarify this point.
| 10 | 12 |
1012.3620
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1012
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1012.3550_arXiv.txt
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{According to the modern cosmological paradigm, galaxies and galaxy systems form from tiny density perturbations generated during the very early phase of the evolution of the Universe. } {Using numerical simulations, we study the evolution of the density perturbation phases of different scales to understand the formation and evolution of the cosmic web.} {We apply the wavelet analysis to follow the evolution of high-density regions (clusters and superclusters) of the cosmic web.} {We show that the maxima and minima positions of density waves (their spatial phases) almost do not change during the evolution of the structure. Positions of density perturbation extrema of are more stable for large scale perturbations. In the context of the present study we call density waves of scale $\geq 64$~\Mpc\ large, waves of scale $\simeq 32$~\Mpc\ medium, and waves of scale $\simeq 8$~\Mpc\ small, within a factor of 2. } {In the cosmic structure formation of the synchronisation (coupling) of density waves of different scales plays an important role.}
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The basic structural elements of the Universe are filamentary superclusters and voids forming a web-like structure -- the supercluster-void network \citep{{Einasto:1980}, {Zeldovich:1982}, {de-Lapparent:1986}, {Bond:1996}}. The standard cosmological paradigm predicts that a period of accelerated expansion, dubbed inflation \citep{Starobinsky:1980rt,Guth:1981ys,Linde:1982fr,Albrecht:1982}, generated density fluctuations \citep{Mukhanov:1981,Hawking:1982,Starobinsky:1982zr,Guth:1982} as well as primordial gravitational waves \citep{Starobinskii:1979vn} through quantum-gravitational processes. In the simplest form of this scenario, the primordial density field is predicted to form a statistically homogeneous, isotropic and almost-Gaussian random field, after the transition from quantum to approximate classical description of perturbations. If the hypothesis of primordial Gaussianity is correct, then density waves of different scales began with random and uncorrelated spatial phases. As the density waves evolve, they interact with others in a non-linear way. This interaction leads to the generation of non-random and correlated phases which form a spatial pattern of the present cosmic web. Owing to non-linear processes during galaxy formation and the physical biasing problem (almost no galaxies form in low-density regions), the present density field is highly non-Gaussian. There have been a number of attempts to gain quantitative information on the behaviour of phases in gravitational systems; for a review see \citet{Coles:2009} and references therein. Using numerical simulations, \citet{Ryden:1991fk} showed that initial phase information is rapidly lost in short wavelengths during evolution. \citet{Hikage:2005bd} analysed the clustering of SDSS galaxies using the distribution function of the sum of Fourier phases. Fourier phases are statistically independent of the Fourier amplitudes, thus the phase statistics plays a complementary role to the conventional two-point statistics of galaxy clustering. Gaussian fields have a uniform distribution of the Fourier phases over $0\leq \theta_k \leq 2\pi$. Therefore, characterising the correlation of phases is expected to be a direct means to explore non-Gaussian features. From the comparison of observations with mock catalogues constructed from N-body simulations, the authors find that the observed phase correlations for the galaxies agree well with those predicted by the spatially flat $\Lambda$CDM model, evolved from Gaussian initial conditions. The analysis in Fourier space is, however, not sensitive to the location of particular high-density features in real space, such as filaments, clusters, and superclusters. To have a better understanding of the texture of the cosmic web, the web must be studied in the real space. Different statistical measures have been used to describe quantitatively the cosmic texture, for recent reviews see \citet{Martinez:2002ye}, \citet{Saar:2009}, and \citet{van-de-Weygaert:2009bh}. One of these statistics is the wavelet analysis, which analyses properties of waves of various scales in real space \citep[see][and references therein]{Jones:2009dq}. Wavelet analysis has been used to detect voids and filaments in the Center for Astrophysics (CfA) survey first slice \citep{Slezak:1993nx}, to de-noise the galaxy distribution \citep{Martinez:2005kl}, to detect the integrated Sachs-Wolfe effect in the cosmic microwave background (CMB) radiation \citep{Vielva:2006tg}, {to study the discreteness effects in simulations \citep{Romeo:2008}}, and for many other purposes where the spatial position of structural elements is important. The goal of this paper is to investigate the evolution of the texture of the cosmic web. It is well known that on small scales the original phase information is lost during the non-linear stage of the evolution \citep{Ryden:1991fk}. On the other hand, the main large-scale skeleton of the texture of the cosmic web is already determined by the initial gravitational potential field \citep{Kofman:1988}. {If considered from the point of view of the Fourier decomposition, maxima and minima in {\em any} spatial distribution occur in the points where phases of the Fourier modes are synchronised. However, there is a non-trivial problem, which is not completely solved yet. In the classical ``down-top'' model, like the isocurvature model, there are no built-in large-scale features. The structure formation starts from small-scale systems, which grow by random clustering. In the classical ``top-down'' model, like the adiabatic neutrino-dominated hot dark matter model, there is a built-in cut-off scale, which determines the scale of the structure. The presently accepted dark energy dominated $\Lambda$CDM model is essentially a ``down-top'' model, because the structure formation starts from small systems. This model differs from the classical ``down-top'' model in one important detail -- here a broad power spectrum of density perturbations is present. Objects of a smaller scale and mass form earlier. But in the case of a broad power spectrum of perturbations, it becomes a non-trivial question whether extrema of perturbations of a given scale will remain at the same places if perturbations of larger scales are added. This may occur only if some phase synchronisation or coupling between perturbations of {\em different scales} exists. In the case of the broad wave spectrum, extrema of density perturbations should define locations where gravitationally bound objects and voids form first. On the other hand, the gravitational potential defines the location of the skeleton of the cosmic web knots. Consequently, it is not clear at all why extrema of density perturbations coincide with knots defined by the gravitational potential. In other words: Why is the skeleton stable in the ``down-top'' $\Lambda$CDM model? As we see in this paper, just because of this synchronisation between waves of different scales.} Studies of Fourier phases show that the phase coupling in the non-linear regime plays an important role in the formation of the fine texture of the cosmic web \citep{Chiang:2002}. To avoid complications caused by the highly non-linear regime, we concentrate on the evolution of waves at intermediate and large scales using the wavelet decomposition of the evolving density field. The Fourier modes are fully specified by their wavelength, their orientation, and phase. Because the phase determines where the maxima and minima are located along a Fourier mode, we also use the same terminology (somewhat less strict this time) once we speak about the wavelet decomposition of the density field. Here we assume that the (spatial) phase and the locations of the maxima and minima carry the same information, and thus will use these terms interchangeably in the following. Also, quite often the locations of the cells inside the cubical density grid are located as $(i,j,k)$, with $i$, $j$, and $k$ are integers that run from $1$ to $N$, where $N=256$ is most often assumed throughout the work. We shall focus our attention on the two main problems: the evolution of phases (positions of maxima) of density perturbations at medium and large scales, and the phase coupling (synchronisation) of perturbations of different scales. To follow the evolution of perturbations of different size, we use simulations in boxes of various sizes from 100 to 768 ~\Mpc. To find the sensitivity of our results to the resolution, we make simulations with $256^3$ and $512^3$ cells and equal number of particles. For comparison with the real Universe, we shall calculate the density field and its wavelet decompositions for a slice (wedge) of the Sloan Digital Sky Survey (SDSS). Preliminary results of this study have been reported at several conferences \citep{Einasto:2006mz,Einasto:2006zr,Einasto:2009ly}. This paper is a follow-up of a study by \citet{Einasto:2005a} of the environmental effects of clusters in SDSS and simulations. In the next section we describe the numerical models used in this study. We also make a qualitative wavelet analysis of the simulated density field, follow the density evolution in time, and compare the evolution of density waves of various scales. In Section 3 we make a correlation analysis of wavelet-decomposed density fields. In Section 4 we analyse the luminosity density field of the SDSS, and study the role of phase synchronisation in the formation of clusters and superclusters. In Section 5 we discuss our results. The last section gives our main conclusions.
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Main conclusions of the present paper can be formulated as follows. \begin{enumerate} \item The wavelet analysis has demonstrated a good diagnostic value for studying the evolution of galaxy systems of various scales and masses. \item In the formation of cosmic structures the synchronisation (coupling) of density waves of different scales plays an important role. \item Positions of density maxima of waves of large and medium scales practically do not change during the evolution. On smaller scales positions of density maxima change during the evolution, the changes are larger for waves with shorter wavelengths. \item Superclusters are objects where density waves of medium and large scales combine {\em in similar phases to generate high-density regions}. \item Voids are regions in space where density waves of medium and large scales combine {\em in similar under-density phases}. \item Clusters of galaxies are objects where density waves of small scales combine in similar over-density phases. \item The larger is the scale of the highest phase synchronisation, the richer are the clusters and superclusters. \end{enumerate}
| 10 | 12 |
1012.3550
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1012
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1012.3766_arXiv.txt
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We present a spectroscopic catalog of 70,841 visually inspected M dwarfs from the seventh data release (DR7) of the Sloan Digital Sky Survey (SDSS). For each spectrum, we provide measurements of the spectral type, a number of molecular bandheads, and the H$\alpha$, H$\beta$, H$\gamma$, H$\delta$ and Ca II K emission lines. In addition, we calculate the metallicity-sensitive parameter $\zeta$ and 3D space motions for most of the stars in the sample. Our catalog is cross-matched to Two Micron All Sky Survey (2MASS) infrared data, and contains photometric distances for each star. Future studies will use these data to thoroughly examine magnetic activity and kinematics in late-type M dwarfs and examine the chemical and dynamical history of the local Milky Way.
| 10 | 12 |
1012.3766
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1012
|
1012.3080.txt
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% context heading (optional) % {} leave it empty if necessary {\emph{Low ionisation nuclear emission-line region} (LINER) nuclei have been claimed to be different than other \emph{active galactic nuclei} (AGN) due to the presence of complex absorbing structures along the line-of-sight and/or an inefficient mode of accretion onto the supermassive black hole. However, this issue is still open.} % aims heading (mandatory) {To investigate the broad band X-ray spectrum of NGC\,4102, one of the most luminous LINERs in the \emph{Swift}/BAT survey.} % methods heading (mandatory) {We studied a 80 ksec \emph{Suzaku} spectrum of NGC\,4102, together with archival \emph{Chandra} and \emph{Swift}/BAT observations. We also studied the optical (3.5m/TWIN at Calar Alto observatory) and near-infrared (WHT/LIRIS at Observatorio Roque los Muchachos) spectra that were taken contemporaneous to the \emph{Suzaku} data.} % results heading (mandatory) {There is strong evidence that NGC\,4102 is a \emph{Compton-thick} AGN, as suggested by the \emph{Swift}/BAT detected intrinsic continuum and the presence of a strong narrow, neutral FeK$\rm{\alpha}$ emission line. We have also detected ionised Fe$\rm{_{XXV}}$ emission lines in the \emph{Suzaku} spectrum of the source. NGC\,4102 shows a variable soft excess found at a significantly higher flux state by the time of \emph{Suzaku} observations when compared to \emph{Chandra} observations. Finally, a complex structure of absorbers is seen with at least two absorbers apart from the \emph{Compton-thick} one, derived from the X-ray spectral analysis and the optical extinction.} % conclusions heading (optional), leave it empty if necessary {All the signatures described in this paper strongly suggest that NGC\,4102 is a \emph{Compton-thick} Type-2 AGN from the X-ray point of view. The ``soft excess", the electron scattered continuum component, and the ionised iron emission line might arise from \emph{Compton-thin} material photoionised by the AGN. From variability and geometrical arguments, this material should be located somewhere between 0.4 and 2 pc distance from the nuclear source, inside the torus and perpendicular to the disc. The bolometric luminosity ($\rm{L_{bol}=1.4\times 10^{43}erg~s^{-1}}$) and accretion rate ($\rm{\dot{m}_{Edd}=5.4\times10^{-3}}$) are consistent with other low-luminosity AGN. However, the optical and near infrared spectra correspond to that of a LINER source. We suggest that the LINER classification might be due a different spectral energy distribution according to its steeper spectral index.}
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\label{sec:intro} Active galactic nuclei (AGN) emit over the entire electromagnetic spectrum and are widely believed to be powered by the accretion of matter onto a supermassive black hole \citep[SMBH,][]{Rees84}. Several families within the AGN category have been established from the observational point of view. Although their classification is sometimes misleading, it is widely believed that a unified model can explain them all under a single scenario \citep{Antonucci93}. A key ingredient in this scheme is a dusty torus whose inclination with respect to the observer's line of sight is responsible for the dichotomy between optical Type-1 (with broad permitted lines, face-on view) and Type-2 (with narrow permitted lines, edge-on view) AGN. However, this scheme needs to be further refined since there are several sub-classes of objects that cannot be easily fitted into this scenario \citep[for example unobscured Type 2 Seyferts, e.g.][] {Mateos05,Dewangan05,Panessa02}. One of the most intriguing cases are \emph{low ionisation nuclear emission-line regions} \citep[LINERs, ][]{Heckman80}. As suggested by their low X-ray luminosities \citep[$\rm{L(2-10~keV)}$ $\rm{\sim 10^{39-42}erg~s^{-1}}$, see][]{Gonzalez-Martin09A} they could be the link between AGN ($\rm{L(2-10~keV)}$ $\rm{\sim 10^{41-45}erg~s^{-1}}$) and normal galaxies \citep[$\rm{\sim 10^{38-42}erg~s^{-1}}$,][]{Fabbiano89}. Moreover, they are the dominant population of active galaxies in the nearby universe \citep{Ho97}. However, their nature is not yet well understood. Several samples of LINERs have been analysed at X-ray frequencies, a large fraction of them showing AGN signatures \citep{Gonzalez-Martin09A,Gonzalez-Martin06,Dudik05}. In spite of this, it still remains unclear how to fit LINERs into the AGN Unification scenario. A radiatively inefficient accretion flow onto the SMBH \citep{Ho09a} and/or the presence of highly obscuring matter have been proposed to explain the differences between LINERs and more luminous AGN \citep{Goulding09,Dudik09,Gonzalez-Martin09B}. Using the ratio $\rm{log(Fx(2-10~keV)}$ $\rm{/F([O~III])}$) (R$\rm{_{X/[O~III]}}$, hereinafter), \citet{Gonzalez-Martin09B} showed that LINERs have a higher fraction of \emph{Compton-thick} sources than Type-2 Seyfert galaxies. This implies high column densities and significant suppression of the intrinsic continuum emission below 10 keV. Only indirect proof of their Compton-thickness can be obtained with \emph{Chandra} and \emph{XMM-Newton} data. Therefore, the nature of these sources is yet to be confirmed. A more direct evidence comes from the determination of the strength of the neutral iron $\rm{K\alpha}$ emission line and the direct view of the nuclear continuum above 10 keV. NGC\,4102 is a nearby Sb galaxy with a nuclear optical spectrum that was first classified as an HII region by \citet{Ho97} although its UV emission is not compatible with this classification \citep{Kinney93}. \citet{Goncalves99} classified its optical spectrum as composite, concluding that the nucleus is dominated by starburst emission although a weak Type-2 Seyfert component is also present. NGC\,4102 is included in the \citet{Carrillo99} sample of LINERs\footnote{This catalogue included all the nuclei classified as LINERs in the literature to data.} and we have reclassified it as LINER by means of the emission lines given in \citet{Moustakas06}. NGC\,4102 has been observed with the {\it Chandra}/ACIS snapshot survey \citep{Dudik05}. They classified it as an AGN-like source. \citet{Tzanavaris07} pointed out its AGN signatures, and considered it as a good candidate for harbouring a hidden AGN. They claimed the presence of an iron line, although poor statistics did not allow them to accurately constrain its equivalent width. %($\sf{EW(FeK\alpha)\sim 200}$ eV). \citet{Ghosh08} showed that NGC\,4102 has an AGN and strong star formation activity. They also pointed out the existence of a reflection component based on a hint of a strong FeK$\rm{\alpha}$ emission line. According to the $\rm{F_{X}(2-10 keV)/F([O~III])}$ ratio, NGC\,4102 is a good candidate to be a \emph{Compton-thick} source (see Sect. \ref{sec:dis} in this paper). Therefore, NGC\,4102 is an ideal case to study the obscuration in LINERs. Here we present the \emph{Suzaku} spectra of NGC\,4102. We also present %for the first time, optical (TWIN/2.2m in Calar Alto observatory) and near infrared (LIRIS/WHT in El Roque de los Muchachos observatory) spectra which were taken contemporaneously (up to one month apart) with our \emph{Suzaku} observation. \emph{Chandra} archival data are also revisited to study the long term variability of this source. This paper is organised as follows. In Section 2 we describe the X-ray data reduction and observations. In Section 3 we present the X-ray spectral fitting. In Section 4 we review the NGC\,4102 activity classification as seen by optical and near-IR observations. Finally, we discuss the nature of the emission seen in NGC\,4102 in Section 5 and report the main conclusions in Section 6. A distance of 17 Mpc is assumed for NGC\,4102 throughout the analysis \citep{Tully88}. A $\Lambda$CDM cosmology with ($\Omega_{\rm M}$,~$\Omega_{\Lambda}$)~=~(0.3,~0.7) and ${H}_{0}$~=~75~ ${\rm km}~{\rm s}^{-1}~{\rm Mpc}^{-1}$ (i.e. z=0.0042) is also assumed.
| 10 | 12 |
1012.3080
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1012
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1012.3258_arXiv.txt
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{The contribution by massive stars ($M > 9M_{\odot}$) to the weak s-process component of the solar system abundances is primarily due to the $^{22}$Ne neutron source, which is activated near the end of helium-core burning. The residual $^{22}$Ne left over from helium-core burning is then reignited during carbon burning, initiating further s-processing that modifies the isotopic distribution. This modification is sensitive to the stellar structure and the carbon burning reaction rate. Recent work on the $^{12}$C + $^{12}$C reaction suggests that resonances located within the Gamow peak may exist, causing a strong increase in the astrophysical S-factor and consequently the reaction rate. To investigate the effect of an increased rate, $25 M_{\odot}$ stellar models with three different carbon burning rates, at solar metallicity, were generated using the Geneva Stellar Evolution Code (GENEC) with nucleosynthesis post-processing calculated using the NuGrid Multi-zone Post-Processing Network code (MPPNP). The strongest rate caused carbon burning to occur in a large convective core rather than a radiative one. The presence of this large convective core leads to an overlap with the subsequent convective carbon-shell, significantly altering the initial composition of the carbon-shell. In addition, an enhanced rate causes carbon-shell burning episodes to ignite earlier in the evolution of the star, igniting the $^{22}$Ne source at lower temperatures and reducing the neutron density.} \FullConference{11th Symposium on Nuclei in the Cosmos, NIC XI\\ July 19-23, 2010\\ Heidelberg, Germany} \begin{document}
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The s-process components identified to contribute to the solar abundance distribution are the weak component, that is produced in massive stars (M > 9$M_{\odot}$), and the main and strong components, that are produced in AGB stars. In particular, the weak s-process component is responsible for most of the isotopes in the mass range $60 < A < 90$. During helium-core burning in massive stars, $^{22}$Ne is formed from $^{14}$N synthesized by the CNO cycle, via the reaction chain $^{14}$N$(\alpha,\gamma)^{18}$F$(\beta^+)^{18}$O$(\alpha,\gamma)^{22}$Ne. At the end of helium burning, when the temperature reaches $0.25$ GK ($1$ GK = $10^9$ K), the $^{22}$Ne$(\alpha,$n$)^{25}$Mg reaction becomes efficient, resulting in an s process characterised by an average neutron density $n_n \sim 10^{6}$ n cm$^{-3}$ and a neutron exposure (for a 25$M_{\odot}$ star) $\tau_n \sim 0.2$ mbarn$^{-1}$. During the advanced stages, convective carbon-shell burning reignites the remaining $^{22}$Ne with a much higher neutron density but lower neutron exposure ($n_n \simeq 10^{11}$ n cm$^{-3}$ and $\tau_n \simeq 0.06$ mbarn$^{-1}$ \cite{1991ApJ...371..665R}). The s process also occurs during (radiative) carbon-core burning via the $^{13}$C($\alpha$, n)$^{16}$O neutron source \cite{1998ApJ...502..737C}. However, in standard $25 M_{\odot}$ stars heavy elements synthesized in the core are further processed and are not ejected during the supernova explosion. Thus they do not contribute to the final yields. Changes to the $^{12}$C + $^{12}$C reaction have an effect on the s process in massive stars \cite{2010JPhCS.202a2023B}, but a detailed analysis has so far been limited to the effect of a reduced rate due to fusion hindrance \cite{2007PhRvC..76c5802G}, although the consequences of an increased rate have been considered in superbursts on accreting neutron stars in X-ray binaries \cite{2009ApJ...702..660C}.
| 10 | 12 |
1012.3258
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1012
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1012.3772_arXiv.txt
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\noindent Energetic antiprotons in cosmic rays can serve as an important indirect signature of dark matter. Conventionally, the antiproton flux from dark matter decay or annihilation is calculated by solving the transport equation with a space-independent diffusion coefficient within the diffusion zone of the galaxy. Antiproton sources outside the diffusion zone are ignored under the assumption that they propagate freely and escape. In reality, it is far more likely that the diffusion coefficient increases smoothly with distance from the disk, and the outlying part of the dark matter halo ignored in the conventional approach can be significant, containing as much as 90\% of the galactic dark matter by mass in some models. We extend the conventional approach to address these issues. We obtain analytic approximations and numerical solutions for antiproton flux for a diffusion coefficient that increases exponentially with the distance from the disk, thereby including contributions from dark matter annihilation/decay in essentially the full dark matter halo. We find that the antiproton flux predicted in this model deviates from the conventional calculation for the same dark matter parameters by up to about 25\%.
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While astronomical observations have firmly established that about 80\% of the matter content of the universe exists in the form of non-baryonic dark matter, its microscopic properties remain hitherto unknown. Weakly interacting massive particles (WIMPs) are the best-motivated candidates for dark matter from a theoretical point of view. WIMPs couple weakly to standard model (SM) particles, opening the possibility of dark matter particles decaying or annihilating into SM final states. Energetic cosmic rays produced in such decay or annihilation processes -- in particular antimatter, which is rarely produced in astrophysical processes -- can serve as important indirect signatures of dark matter in the galaxy. Observation of such signals can reveal information on the microscopic properties of dark matter. Our focus in this paper is on antiprotons produced in the annihilation or decay of dark matter in the Milky Way galaxy. This has become a topic of significant interest following recent measurements of the antiproton flux and the ratio of antiproton to proton flux up to 180 GeV by the PAMELA experiment~\cite{pamela, pamela2}. No deviations from the expected astrophysical flux were observed. This data can in principle be used to put bounds on dark matter properties; however, to do this, the effects of propagation of antiprotons between the dark matter decay/annihilation location and the detector must be properly accounted for. Propagation of antiprotons in the galaxy is governed by transport equations, which incorporate interactions with galactic magnetic fields and the interstellar medium (ISM). Conventional calculations solve the transport equation in a two-zone model: a thin disk of $\mathcal{O}$($100$ pc) thickness on the galactic plane where the interactions with the ISM occur, embedded in a larger diffusion region where galactic magnetic fields are appreciable and trap cosmic rays. The diffusion region is taken to be a cylinder of radius $R$ of order 20 kpc and half-thickness $L$ of order 1--10 kpc, and the diffusion coefficient is assumed to be position-independent inside this region. Outside this diffusion region, antiprotons are assumed to propagate freely, leading to vanishing antiproton density on the boundaries of the cylinder in the steady state. With these assumptions, the transport equations can be solved in a straightforward manner. In particular, for antiproton energies of interest for dark matter searches (of order 10 GeV and above), additional approximations can be made that allow an {\it analytic} solution to the antiproton density and flux in the form of a Bessel series~\cite{solutions}. The conventional two-zone model is, however, a rather crude approximation. First, the assumption of a sharp boundary at $L$ between the diffusion and free-propagation zones is unphysical. In reality, magnetic fields are known to decrease gradually with the distance from the disk, with exponential decay providing a reasonable fit to data. The diffusion coefficient likely follows a similar exponential profile~\cite{exponentialk}. Second, a typical dark matter halo is spherically symmetric and extends {\it beyond} the diffusion region and into the free propagation zone, in particular in the vertical direction: for example, for an isothermal dark matter profile in a model with $L=1$ kpc, the diffusion zone contains only 10\% of the dark matter mass of the full halo. The two-zone model completely ignores the antiprotons produced by dark matter decay/annihilation outside the diffusion zone. The aim of this paper is to extend the conventional formalism for antiproton flux calculations to overcome these shortcomings.\footnote{The extension presented here is similar to the formalism developed by us in a previous paper~\cite{positrons} for positrons. However, the absence of energy-loss terms for antiprotons allows us to make significant progress analytically, which was not possible for positrons.} After reviewing the conventional formalism in Sec.~\ref{sec:oldform}, we consider a three-zone model in Sec.~\ref{sec:newform1}. This model still assumes an abrupt change in the diffusion coefficient at $L$, but includes the sources in the free-propagation zone extending to $D\gg L$, so that essentially all of the antiprotons from the dark matter halo are taken into account. We find an analytic solution to this model, and show that in the limit when diffusion in the free-propagation zone is completely absent, the sources in this zone have no effect on the flux measured at Earth. We then consider a model with an exponentially varying diffusion coefficient~\cite{exponentialk} in Sec.~\ref{sec:newform2}, and obtain both a numerical solution and two analytic approximations which converge to it at high energies (above 50 GeV or so, depending on the desired accuracy). In Sec.~\ref{sec:results}, we present numerical results quantifying the effects of this more realistic propagation model on the antiproton fluxes produced by dark matter decay or annihilation, and the corresponding bounds on dark matter properties. We close with concluding remarks in Sec.~\ref{sec:conc}.
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\label{sec:conc} In this paper, we extended the conventional approach to calculation of the antiproton flux from dark matter annihilation/decay to allow for the possibility of a position-dependent diffusion coefficient. We studied two models, the three-zone model where the diffusion coefficient jumps on the boundary between the ``diffusion" and ``free-propagation" zones, and a model with the diffusion coefficient growing smoothly (exponentially) with distance away from the disk. In the first model, we found an analytic solution and showed that in the limit of infinite diffusion coefficient in the free-propagation zone the flux on Earth is not modified by the sources in that zone due to perfect reflection of antiprotons from the zone boundary. This seems to justify the conventional calculation even in situations when there are sources outside the diffusion zone, as is common in dark matter studies. On the other hand, the three-zone model is a rather crude approximation, since magnetic fields, and with them the diffusion coefficient, are expected to vary smoothly with distance from the disk. Such smooth variation was incorporated in our second model, which assumed an exponentially growing diffusion coefficient, which has been shown to produce consistent fits to conventional astrophysical cosmic ray observables in Ref.~\cite{exponentialk}. We found numerical solutions as well as analytic approximations valid at high antiproton energies for this model. We found that the resulting fluxes on Earth differ from the predictions of the conventional calculation with the MED propagation model by at most about 25\%. Such deviations are well within the numerous uncertainties inherent in the calculation, hence the use of the conventional approach is justified, at least at present, in computing the antiproton fluxes from dark matter annihilation/decay and using them to place bounds on dark matter models. Nevertheless, since the model with exponentially growing diffusion coefficient almost certainly captures the correct physics of charged particle propagation in the galaxy better than the conventional one, and since the solutions we obtained in this model -- especially the analytic solutions, which are accurate at high energies -- are essentially as simple as the conventional ones, we advocate the use of this model as a benchmark for dark matter studies. \vskip0.5cm \noindent{\large \bf Acknowledgments} \vskip0.3cm We thank Joakim Edsj\"{o} and Jesse Thaler for useful discussions. B.S. would like to thank the Oskar Klein Center at Stockholm University, where part of this work was done, for their hospitality. This research is supported by the U.S. National Science Foundation through grant PHY-0757868 and CAREER grant No. PHY-0844667.
| 10 | 12 |
1012.3772
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1012
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1012.4014_arXiv.txt
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We investigate the column density distribution function of neutral hydrogen at redshift $z=3$ using a cosmological simulation of galaxy formation from the OverWhelmingly Large Simulations ({\small OWLS}) project. The base simulation includes gravity, hydrodynamics, star formation, supernovae feedback, stellar winds, chemodynamics, and element-by-element cooling in the presence of a uniform UV background. Self-shielding and formation of molecular hydrogen are treated in post-processing, without introducing any free parameters, using an accurate reverse ray-tracing algorithm and an empirical relation between gas pressure and molecular mass fraction. The simulation reproduces the observed $z=3$ abundance of Ly-$\alpha$ forest, Lyman Limit, and Damped Ly-$\alpha$ $\HI$ absorption systems probed by quasar sight lines over ten orders of magnitude in column density. Self-shielding flattens the column density distribution for $\NHI > 10^{18} \cmsq$, while the transition to fully neutral gas and conversion of $\HI$ to H$_2$ steepen it around column densities of $\NHI = 10^{20.3} \cmsq$ and $\NHI = 10^{21.5} \cmsq$, respectively.
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\noindent Ground-based spectroscopic observations targeting quasars are excellent probes of $z\geq 1.7$ neutral hydrogen \citep[e.g.][]{1998ARA&A..36..267R, 2005ARA&A..43..861W}. The Sloan Digital Sky Survey (SDSS) has produced approximately $1.5 \times 10^4$ moderate resolution quasar spectra \citep{2009ApJS..182..543A}. These spectra provide ample data on $\HI$ absorption lines with column densities $\NHI > 10^{20.3} \cmsq$, so called Damped Ly-$\alpha$ systems (DLAs) \citep{2009ApJ...696.1543P,2009A&A...505.1087N}. Lines with $\NHI < 10^{17.2} \cmsq$, the so called Ly-$\alpha$ forest, are best discovered in high-resolution spectra of bright quasars \cite[e.g.][]{2002MNRAS.335..555K}. Lines with intermediate column densities, Lyman Limit Systems (LLSs), lie on the flat part of the curve of growth, which complicates the determination of their column densities. Traditional methods of measuring $\NHI$ in DLAs can be applied to high-resolution spectra for lines with $\NHI > 10^{19} \cmsq$ when damping wings begin to appear \citep[e.g.][]{2005MNRAS.363..479P, 2007ApJ...656..666O}. Progress on the most difficult lines with $10^{14.5} \cmsq < \NHI < 10^{19} \cmsq $ has recently been made by \cite{2010ApJ...718..392P} by combining independent measurements of the Lyman limit mean free path and integral constraints over the column density distribution. Combining the observations above, one can determine the \HI column density distribution function $\fNHI$, {\em i.e.} the number of lines per unit column density $d\NHI$, per unit absorption distance $dX$, at redshifts $z \approx 3$ from $\NHI = 10^{12} \cmsq$ to $\NHI = 10^{22} \cmsq$. Early determinations of $\fNHI$ at these redshifts were reasonably well described by a single power law, $\fNHI \propto N_{\rm HI}^{-\eta}$, with $\eta = 1.5$ \citep{1987ApJ...321...49T}. As the quality of observations improved, this was no longer the case. \cite{1993MNRAS.262..499P}, showed that a single power law and a double power law with a break at $\NHI = 10^{16} \cmsq$ both failed Kolmogorov-Smirnov tests at the 99\% confidence level. The most recent observations are well fit by a series of six power laws which intersect at $\NHI = $ $\{10^{14.5}$, $10^{17.3}$, $10^{19.0}$, $10^{20.3}$, $10^{21.75} \}$ $\cmsq$ \citep{2010ApJ...718..392P}. Attempts to explain the shape and normalization of $\fNHI$ in a cosmological context have typically focused on sub sets of the full column density range. Analytic \citep[e.g.][]{2001ApJ...559..507S}, semi-analytic \citep[e.g.][]{1997ApJ...479..523B}, and numerical \citep[e.g.][]{ 1998MNRAS.301..478T, 1998MNRAS.297L..49T} models were instrumental in identifying the Ly-$\alpha$ forest lines with the diffuse, photo-ionized, intergalactic medium. Numerical work has also played a large role in determining properties of higher column density systems \citep[e.g.][]{1996ApJ...457L..57K, 1997ApJ...484...31G, 1998ApJ...495..647H, 2003ApJ...598..741C, 2004MNRAS.348..421N, 2006ApJ...645...55R, 2007ApJ...655..685K, 2008MNRAS.390.1349P, 2009MNRAS.397..411T, 2010arXiv1008.4242H, 2010arXiv1010.5014C, 2010ApJ...725L.219N, 2011arXiv1101.1964M}. Although self-shielding is crucial for modelling optically thick absorbers, only \cite{2006ApJ...645...55R}, \cite{2007ApJ...655..685K}, \cite{2008MNRAS.390.1349P}, and \cite{2011arXiv1101.1964M} have used 3-D radiative transfer to calculate the attenuation of the UV background. Additionally, conversion of $\HI$ to H$_2$ is thought to determine the high end cut off in $\fNHI$ \citep{2001ApJ...562L..95S,2009ApJ...701L..12K}, yet only \cite{2010arXiv1010.5014C} included this process when modelling $\HI$ absorption. We present a cosmological simulation of structure formation, to which we have applied a radiative transfer self-shielding calculation and a prescription for the conversion of $\HI$ to H$_2$ without introducing any free parameters. We show that this simulation reproduces observational determinations of $\fNHI$ around $z=3$ over the entire range in column density. In addition, we determine the typical neutral fractions and total hydrogen number densities for $\HI$ absorbers as a function of column density $\NHI$. \begin{figure*} \center \epsscale{1.0} \plotone{figure1} \caption{ $\HI$ column density distribution function, $\fNHI$, at $z\sim 3$; simulation results are shown as curves and observational data as symbols. The low $\NHI$ curve is obtained using mock spectra fitted with {\sc VPFIT}. Self-shielding and H$_2$ are unimportant in this range. The high $\NHI$ curve is obtained by projecting the simulation box onto a plane and includes self-shielding and H$_2$. The gap around $\NHI\sim 10^{17}$~cm$^{-2}$ separates low and high $\NHI$. Poisson errors on the simulation curves are always smaller than their thickness. We also show high-resolution observations of the Ly-$\alpha$ forest (Kim et al. 2002, ``Kim02''), LLSs (P{\'e}roux et al. 2005, ``Per05''; O'Meara et al. 2007, ``Ome07''), analysis of SDSS DLA data (Noterdaeme et al. 2009, ``NPLS09''), and power law constraints (Prochaska et al. 2010, ``POW10'', open circles are spaced arbitrarily along power law segments and do not represent $\NHI$ bins or errors). } \label{cddf_lrg_fig} \end{figure*}
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We have used a hydrodynamic simulation of galaxy formation together with an accurate ray-tracing treatment of self-shielding from the UV background and an empirical prescription for H$_2$ formation, to compute the $z \approx 3$ $\HI$ column density distribution function. We find agreement between the reference \OWLS model and the entire column density range probed by observations ($10^{12} \cmsq < \NHI < 10^{22} \cmsq$). We have shown that $\fNHI$ flattens above $\NHI = 10^{18} \cmsq$ due to self-shielding, and steepens around $\NHI = 10^{20.3} \cmsq$ and $\NHI = 10^{21.5} \cmsq $ due to the absorbing gas becoming fully neutral, and the transition from atomic to molecular hydrogen, respectively. In future work, we will examine the systems causing this absorption in greater detail and repeat these analyses on a large sample of \OWLS models.
| 10 | 12 |
1012.4014
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1012
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1012.1407_arXiv.txt
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Despite 100 years of effort, we still know very little about the origin of ultra-high energy cosmic rays. The observation of neutrinos produced when cosmic-ray protons with energies above $4\times 10^{19}$ eV interact with the cosmic microwave background radiation, or in the neutrino sources, would tell us much about the origin and composition of these particles. Over the past decade, many experiments have searched for radio waves emitted from the charged particle showers produced when EHE neutrinos interact with Antarctic or Greenland ice or the moon. These experiments have not yet observed a neutrino signal. Two groups are now proposing to instrument 100 km$^3$ of Antarctic ice with radio antennas, producing a detector large enough to observe a clear EHE neutrino signal in a few years of operation.
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\label{Introxy} Despite 100 years of effort, we do not know the source of high-energy cosmic rays. For extremely high-energy particles (EHE, particles with energies above about $10^{17}$ eV), the mystery is even deeper. We do not know of any likely sources within our galaxy. However, at the highest energies, above $4\times10^{19}$ eV (40 EeV), the range of cosmic-rays is very limited. More energetic protons will interact with the $3^0$K cosmic microwave background radiation, and be excited to a $\Delta^+$ resonance, which will decay to $p\pi^0$ or $n\pi^+$, followed by $n\rightarrow pe^-\overline{\nu_e}$, as was first described by Greisen, Zatsepin and Kuzmin (GZK) \cite{GZK}. The end result is a somewhat lower energy proton; this reaction recurs until the proton energy drops below 40 EeV. For nuclei, a similar limitation is present, because the nuclei are photodissociated by the microwave photons. These process limit the range of more energetic cosmic-rays to about 100~megaparsecs (Mpc). Any sources must be relatively close. If they are protons, then the highest energy cosmic-rays should travel relatively straight lines bending a few degrees in 100~Mpc in the inter-stellar magnetic fields. The Auger collaboration studied the arrival directions of EHE cosmic-rays. They compared the arrival direction of cosmic-rays with energies above 60 EeV, with a list of active galactic nuclei within 75~Mpc of earth \cite{AugerAGN,Auger2}, and found a statistically significant correlation. These correlations are only expected if cosmic-rays are mostly protons, because heavier nuclei will be bent in interstellar magnetic fields. However, another Auger analysis indicates that, at energies above 10 EeV, cosmic-rays are mostly heavier nuclei \cite{Auger3}. Clearly, an alternate probe of EHE cosmic-rays is needed. Neutrinos are that probe. They are produced at any accelerator, by 'beam gas' interactions, when cosmic-rays undergoing acceleration collide with remnant gas and/or photons \cite{PT}. Beyond that, there is a guaranteed source of EHE neutrinos. The $\pi^+$ produced by the GZK process decay, producing $\mu^+\nu_\mu$. Subsequently, the $\mu^+$ decays to $e^+ \overline\nu_e \nu_\mu$. Over cosmic distances, these neutrinos oscillate, and arrive at Earth with a flavor mixture of roughly $\nu_e:\nu_\mu:\nu_\tau$ = $1:1:1$ (we neglect the difference between $\nu$ and $\overline\nu$). Because these neutrinos interact weakly, and can travel cosmic-distances, Earthly EHE neutrino detectors can observe neutrinos out to redshifts of 3-4 \cite{ESS}. The down side of the small cross-sections is that a very large detector is required to observe GZK neutrinos. Optical detectors, like the 1 km$^3$ IceCube observatory \cite{RSI} are too small. Analyses using the partially completed IceCube have set limits \cite{IC3EHE}, but even the completed detector is expected to observe less than one event/year. A new approach is needed.
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The observation of GZK neutrinos would finally give a definitive answer about the composition of EHE cosmic rays, and, at the same time, give us directional information about their probable sources. However, because the EHE neutrino flux and cross-sections are small, they have not yet been observed. Two new experiments have been proposed to search for these neutrinos. ARIANNA and ARA will have active volumes of order 100 km$^3$. If EHE cosmic rays are mostly protons, this is big enough to observe of order 100 neutrinos in 3-5 years of operation. This work was funded in part by the U.S. National Science Foundation under grant numbers 0653266 and 0969661 and the Department of Energy under contract number DE-AC-76-00098.
| 10 | 12 |
1012.1407
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1012
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1012.2297_arXiv.txt
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It has been suggested that coronal mass ejections (CMEs) remove the magnetic helicity of their coronal source region from the Sun. Such removal is often regarded to be necessary due to the hemispheric sign preference of the helicity, which inhibits a simple annihilation by reconnection between volumes of opposite chirality. Here we monitor the relative magnetic helicity contained in the coronal volume of a simulated flux rope CME, as well as the upward flux of relative helicity through horizontal planes in the simulation box. The unstable and erupting flux rope carries away only a minor part of the initial relative helicity; the major part remains in the volume. This is a consequence of the requirement that the current through an expanding loop must decrease if the magnetic energy of the configuration is to decrease as the loop rises, to provide the kinetic energy of the CME.
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\label{introduction} The helicity of the solar magnetic field obeys a hemispheric preference which is invariant with respect to the sign reversal of the global magnetic field with the activity cycle (\cite[Hale 1925]{Hale1925}; \cite[Seehafer 1990]{Seehafer1990}). This has led to the suggestion that coronal mass ejections (CMEs) must remove magnetic helicity from the Sun to prevent indefinite accumulation of the helicity in each hemisphere (\cite[Rust 1994]{Rust1994}; \cite[Low 1996]{Low1996}). It was also found that the accumulation and removal of helicity can control the rate of mean-field dynamo action, so that the evolution of the activity cycle may be related to the flow of helicity through the Sun (\cite[Blackman \& Field 2001]{Blackman&Field2001}; \cite[Brandenburg \& Subramanian 2005]{Brandenburg&Subramanian2005}). Careful studies of the long-term helicity budget of two solar active regions (\cite[D{\'e}moulin et al. 2002]{Demoulin&al2002}; \cite[Green et al. 2002]{Green&al2002}) appear to confirm the conjecture of efficient helicity shedding by CMEs. However, both investigations used the linear force-free field approximation to estimate the helicity in the active region atmosphere. The accuracy of this estimate is not known. Similarly, the estimates of the helicity in interplanetary CMEs are still subject to considerable uncertainty (\cite[Demoulin 2007]{Demoulin2007}). In this paper, numerical simulation is used to quantify the transport of magnetic helicity from the source volume of CMEs.
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\label{conclusions} The simulated CME ejects only a minor part of the initial relative magnetic helicity from its source volume. Although this result requires substantiation through the study of its parametric dependence and of other equilibria, the necessary decrease of the current through an expanding unstable flux loop leads us to expect that it holds generally. The number of CMEs per active region varies within very wide limits. Between 30 and 65 CMEs have been estimated to occur throughout the lifetime of the two very CME-prolific active regions studied in \cite[D{\'e}moulin et al.\ (2002)]{Demoulin&al2002} and \cite[Green et al.\ (2002)]{Green&al2002}. On the other hand, the majority of active regions produces no CME at all, or only one CME in their lifetime. Hence, the shedding of helicity by CMEs may be of lower importance than originally conjectured. It appears natural to assume that, perhaps generally, much of an active region's helicity submerges when the region disperses and the major part of its flux submerges below the photosphere. The helicity may then follow the slow journey of magnetic flux in the course of the solar cycle. Annihilation of helicity in the interior of the Sun, following the transport of the helicity-carrying flux to the equatorial plane by the meridional flow, is one possibility to prevent the helicity in each hemisphere from accumulating indefinitely. Another possibility, opposite to a common conjecture, is that the helicity in the solar interior, like magnetic energy, undergoes a normal (or direct) turbulent cascade towards small spatial scales, where it is dissipated. Also, the cascade directions may be different for small-scale and large-scale fields (\cite[Alexakis et al.\ 2006]{Alexakis&al2006}), with the helicity of active-region magnetic fields, considered to be small-scale fields, subject to a direct cascade.
| 10 | 12 |
1012.2297
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1012
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1012.4772_arXiv.txt
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A spherical hydrodynamical expansion flow can be described as the gradient of a potential. In that case no vorticity should be produced, but several additional mechanisms can drive its production. Here we analyze the effects of baroclinicity, rotation and shear in the case of a viscous fluid. Those flows resemble what happens in the interstellar medium. In fact in this astrophysical environment supernovae explosion are the dominant flows and, in a first approximation, they can be seen as spherical. One of the main difference is that in our numerical study we examine only weakly supersonic flows, while supernovae explosions are strongly supersonic.
| 10 | 12 |
1012.4772
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1012
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1012.1227_arXiv.txt
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An increasing number of both photometric and spectroscopic observations over the last years have shown the existence of distinct sub-populations in many Galactic globular clusters and shattered the paradigm of globulars hosting single, simple stellar populations. These multiple populations manifest themselves in a split of different evolutionary sequences in the cluster color-magnitude diagrams and in star-to-star abundance variations. In this paper we will summarize the observational scenario.
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\label{intro} In recent years an increasing amount of photometric and spectroscopic observational evidence have shattered the paradigm of globulars as the prototype of single, simple stellar populations (see Piotto\ 2009 for a recent review). Spectroscopic studies have demonstrated that most globular clusters (GC) have no detectable spread in their iron content and also $s$-process elements do not exhibit large star-to-star variations in the majority of globulars (e.g. Carretta et al.\ 2009a and references therein). On the contrary, every time we have at our disposal a large sample of stars for a given GC, star-to-star variations in the light elements C, N, O, Na, and Al have been clearly detected (e.g. Carretta et al.\ 2009b, Pancino et al.\ 2010 and references therein). These variations are related to correlations and anticorrelations, which indicate the occurrence of high temperature hydrogen-burning processes (including CNO, NeNa, MgAl cycles) and cannot occur in presently observed low mass GC stars. Today it is widely accepted that the observed light-elements variations provide strong support to the presence of multiple stellar populations in GCs with the second generations formed from the material polluted by a first generation of stars. On the contrary the debate on the nature of possible polluters is still open (e.g. D'Antona et al.\ 2004, Decressin et al.\ 2007). While abundance variations are well known since the early sixties, it was only the recent spectacular discovery of multiple sequences in the color-magnitude diagram (CMD) of several GCs that provides an un-controversial prove of the presence of multiple stellar populations in GCs and brought new interest and excitement in GCs research (e.g. Piotto et al.\ 2007). Photometric clues, often easy to detect simply by the inspection of high-accuracy CMDs, arise in form of multiple main sequences (MS, Bedin et al.\ 2004, Piotto et al.\ 2007, Milone et al.\ 2010), split sub-giant branch (SGB, Milone et al.\ 2008, Anderson et al.\ 2009, Piotto \ 2009), and multiple red-giant branch (RGB, Marino et al.\ 2008, Yong et al.\ 2008, Lee et al.\ 2009). Many population properties, like the chemical composition, the spatial distribution, the fraction of stars in each population and their location in the CMD apparently differ from cluster to cluster. Multiple stellar populations have been detected for the first time in the Milky Way satellite $\omega$ Centauri in form of either multiple MSs (e.g. Anderson\ 1997, Bedin et al.\ 2004, Bellini et al.\ 2010), multiple SGBs (e.g. Sollima et al.\ 2005), multiple RGBs (Lee et al.\ 1999, Pancino et al.\ 2000) and large star-to-star variation in iron and $s$-elements (e.g. Johnson et al.\ 2010, Marino et al.\ 2010). Due to its large mass this GCs have been always considered as a peculiar stellar system and often associated to the remnant of a dwarf galaxy. The `extreme' case of $\omega$ Centauri is not analyzed in this work where we focus on `normal' GCs. The following sections are an attempt to define some groups of `normal' clusters that share similar properties.
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For several decades GCs have been considered as the best approximation of simple stellar populations consisting of coeval and chemically homogeneous stars. This picture have been mainly challenged by two observational facts. Since the seventies we know that GCs exhibit a peculiar pattern in their chemical abundances with large star-to-star variations in the abundances of C, N, Na, O, Mg, and Al. These variations are primordial since they are observed in stars at all the evolutionary phases and are peculiar to GC stars. Field stars only changes in C and N abundance expected from typical evolution of low-mass stars. In addition, since the early sixties we know that the HB of some GCs are quite peculiar. The distribution of stars along the RGB can be multimodal with the presence of one or more gaps and in same cases the HB can be extended toward very high temperatures. It is well known metallicity is the first parameter governing the HB morphology there are some GCs with almost the same iron content but different HB morphology demonstrating the metallicity alone is not enough to reproduce the observational scenario. This problem, known as the {\it second-parameter} problem, still lacks of a comprehensive understanding. The recent discoveries of multiple stellar populations in GCs have shattered once and for all the long-held paradigm of GCs as simple stellar populations and brought new interest on these stellar systems. It is very tempting to relate the second parameter HB problem to the complex abundance pattern of GCs as well as to the multiple sequences observed in the CMD of some clusters. As already mentioned the observed variations of light elements indicate the presence of material processed through hot H-burning processes and should be also He-enriched. While small variations in helium content should have a small impact on colors and magnitudes for MS stars a large impact is expected on the colors of HB stars since He-rich stars should be also less massive. In summary the discovery of multiple stellar populations started a new era on globular cluster research. While the observational scenario is still puzzling and there is a rather incoherent picture of the multipopulation phenomenon, for the first time we might have the key to solve a number of problems, like the abundance anomalies and possibly the second parameter problem, as well as the newly discovered multiple sequences in the CMD.
| 10 | 12 |
1012.1227
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1012
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1012.2154_arXiv.txt
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Three-body model fits to Arecibo and Goldstone radar data reveal the nature of two near-Earth asteroid triples. Triple-asteroid system 2001 SN263 is characterized by a primary of $\sim$10$^{13}$ kg, an inner satellite $\sim$1$\%$ as massive orbiting at $\sim$3 primary radii in $\sim$0.7 days, and an outer satellite $\sim$2.5$\%$ as massive orbiting at $\sim$13 primary radii in $\sim$6.2 days. 1994 CC is a smaller system with a primary of mass $\sim$2.6 $\times 10^{11}$ kg and two satellites $\sim$2$\%$ and $\lesssim$$1\%$ as massive orbiting at distances of $\sim$5.5 and $\sim$19.5 primary radii. Their orbital periods are $\sim$1.2 and $\sim$8.4 days. Examination of resonant arguments shows that the satellites are not currently in a mean-motion resonance. Precession of the apses and nodes are detected in both systems (2001 SN263 inner body: $d\varpi/dt \sim$1.1 deg/day, 1994 CC inner body: $d\varpi/dt \sim$ -0.2 deg/day), which is in agreement with analytical predictions of the secular evolution due to mutually interacting orbits and primary oblateness. Nonzero mutual inclinations between the orbital planes of the satellites provide the best fits to the data in both systems (2001 SN263: $\sim$14 degrees, 1994 CC: $\sim$16 degrees). Our best-fit orbits are consistent with nearly circular motion, except for 1994 CC's outer satellite which has an eccentric orbit of $e \sim$ 0.19. We examine several processes that can generate the observed eccentricity and inclinations, including the Kozai and evection resonances, past mean-motion resonance crossings, and close encounters with terrestrial planets. In particular, we find that close planetary encounters can easily excite the eccentricities and mutual inclinations of the satellites' orbits to the currently observed values.
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The existence and prevalence ($\sim$16\%) of binary asteroids in the near-Earth population (Margot et al. 2002; Pravec et al. 2006) naturally lead to the search and study of multiple-asteroid systems (Merline et al. 2002; Noll et al. 2008). Triple systems are known to exist in the outer Solar System, the main belt, and the near-Earth population. Among the trans-neptunian objects (TNOs), there are currently two well-established triples, 1999 TC36 (Margot et al. 2005; Benecchi et al. 2010) and Haumea (Brown et al. 2006), and one known quadruple, Pluto/Charon (Weaver et al. 2006). In the main belt population, four triples are known to exist: 87 Sylvia, 45 Eugenia, 216 Kleopatra, and 3749 Balam (Marchis et al. 2005; Marchis et al. 2007; Marchis et al. 2008a; Marchis et al. 2008b). There are currently only two well-established asteroid triples in the near-Earth population, 2001 SN263 (Nolan et al. 2008a) and 1994 CC (Brozovic et al. 2009), both of which are the focus of this study. There is also another possible triple, near-Earth asteroid 2002 CE26, that may have a tertiary component but the limited observational span of this object prevented an undisputable detection (Shepard et al. 2006). (153591) 2001 SN263 has been unambiguously identified as a triple-asteroid system, and its orbit around the Sun is eccentric at 0.48 with a semi-major axis of 1.99 AU and inclined 6.7 degrees with respect to the ecliptic. It is an Amor asteroid with a pericenter distance of 1.04 AU. The system is composed of three components: a central body (equivalent radius $\sim$ 1.3 km) and two orbiting satellites. In this paper, we use the following terminology for triple systems: the central body (most massive component) is called Alpha, the second most massive body is termed Beta, and the least massive body is named Gamma (Figure \ref{diagrams}). In the case of 2001 SN263, Beta is the outer satellite and Gamma is the inner satellite. The other near-Earth triple is (136617) 1994 CC with a primary of equivalent radius R $\sim$ 315 m, where the opposite is true; Beta is the inner body and Gamma is the outer body. Its heliocentric orbit has a semi-major axis of 1.64 AU and is also eccentric (0.42) and inclined (4.7 degrees) with respect to the ecliptic. 1994 CC is an Apollo asteroid with a pericenter distance of 0.95 AU. In this work, we present dynamical solutions for both triple systems, 2001 SN263 and 1994 CC, where we derived the orbits, masses, and Alpha's $J_2$ gravitational harmonic using N-body integrations. We utilize range and Doppler data from Arecibo and Goldstone, and these observations as well as our methods are described in Section 2. In Section 3, we present our best orbital solutions and their uncertainties. We also include discussion regarding the satellite masses and the primary's oblateness (described by $J_2$), and the observed precession of the apses and nodes are compared to analytical predictions. Section 4 describes the origin and evolution of the orbital configurations, including Kozai and evection resonant interactions, as well as the effects of planetary encounters. Previous studies of multiple TNO systems include the analytic theory of Lee \& Peale (2006) for Pluto, where they treated Nix and Hydra as test particles. Tholen et al. (2008) used four-body orbit solutions to constrain the masses of Nix and Hydra; they did not find evidence of mean-motion resonances in the system. Ragozzine \& Brown (2009) determined the orbits and masses of Haumea's satellites using astrometry from Hubble Space Telescope and the W. M. Keck Telescope. They used a three-body model and found that their data was not sufficient to constrain the oblateness, described by $J_2$, of the non-spherical central body. Their orbital solutions yielded a large eccentricity ($\sim$0.249) of the inner, fainter satellite, Namaka, and a mutual inclination with the outer satellite, Hi'iaka, of $\sim$13.41 degrees. They postulated that the excited state of the system could be conceptually explained by the satellites' tidal evolution through mean-motion resonances. In the main belt, Winter et al. (2009) studied the orbital stability of the satellites in the Sylvia triple system and Marchis et al. (2010) presented a dynamical solution of Eugenia and its two satellites.
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In this work, we found dynamical solutions for two triple systems, 2001 SN263 and 1994 CC, where we have derived the orbits, masses, and Alpha's $J_2$ gravitational harmonic using full N-body integrations. We used range and Doppler data from Arecibo and Goldstone to solve this non-linear least-squares problem with a dynamically interacting three-body model that provided an excellent match to our radar observations. Given the three-body nature of these systems, we also measured the precession rates of the apses and nodes, and compared them to our corresponding analytical expressions from $J_2$ and secular contributions. No resonant arguments were found to be librating in either triple system. For both systems, we detected significant mutual inclinations (2001 SN263: $\sim$14 deg, 1994 CC: $\sim$16 deg) between the orbital planes of Beta and Gamma. We also found a nonzero orbital eccentricity ($\sim$0.2) for 1994 CC's outer body, Gamma. The eccentricity and inclination damping timescales are long, suggesting that both systems are in excited states. We investigated excitation mechanisms that could explain the observed orbital configurations, including Kozai and evection resonances, mean-motion resonance crossings, and close encounters with the terrestrial planets. Close encounters that occur on million-year timescales can reproduce the observed mutual inclinations in both systems and 1994 CC Gamma's eccentricity. \newpage
| 10 | 12 |
1012.2154
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1012
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1012.5859_arXiv.txt
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We investigate the origin of the relations between stellar mass and optical circular velocity for early-type (ETG) and late-type (LTG) galaxies --- the Faber-Jackson (FJ) and Tully-Fisher (TF) relations. We combine measurements of dark halo masses (from satellite kinematics and weak lensing), and the distribution of baryons in galaxies (from a new compilation of galaxy scaling relations), with constraints on dark halo structure from cosmological simulations. The principle unknowns are the halo response to galaxy formation and the stellar initial mass function (IMF). The slopes of the TF and FJ relations are naturally reproduced for a wide range of halo response and IMFs. However, models with a universal IMF and universal halo response cannot {\it simultaneously} reproduce the zero points of both the TF and FJ relations. For a model with a universal Chabrier IMF, LTGs require halo expansion, while ETGs require halo contraction. A Salpeter IMF is permitted for high mass ($\sigma \gta 180 \kms$) ETGs, but is inconsistent for intermediate masses, unless $V_{\rm circ}(R_{\rm e})/\sigma_{\rm e} \gta 1.6$. If the IMF is universal and close to Chabrier, we speculate that the presence of a major merger may be responsible for the contraction in ETGs while clumpy accreting streams and/or feedback leads to expansion in LTGs. Alternatively, a recently proposed variation in the IMF disfavors halo contraction in both types of galaxies. Finally we show that our models naturally reproduce flat and featureless circular velocity profiles within the optical regions of galaxies without fine-tuning.
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\label{sec:intro} Galaxy properties obey several fundamental relations, which have long been thought to hold important clues about the physical processes that influenced their formation and evolution. The relations between rotation velocity and luminosity (for late-types) and velocity dispersion and luminosity (for early-types) are particularly interesting as they connect luminous mass with dynamical mass (which includes not only baryons but also dark matter). These relations are also known as the Tully-Fisher (Tully \& Fisher 1977) and Faber-Jackson (Faber \& Jackson 1976) relations. The scatter in these relations is small, 0.07 dex in velocity for FJ (e.g., Gallazzi \etal 2006) and 0.05 dex in velocity for TF (e.g., Courteau \etal 2007b; Pizagno \etal 2007), which has enabled these relations to be used as secondary distance indicators. The smallness and source of the scatter is also interesting from a galaxy formation point of view. For early-type galaxies the scatter in velocity dispersion correlates with galaxy size, or surface brightness. This leads to the so-called fundamental plane of early-type galaxies, a correlation between size, surface brightness and velocity dispersion (Dressler \etal 1987; Djorgovski \& Davis 1987). By contrast the scatter in the TF relation is independent of size or surface brightness (Zwaan \etal 1995; Courteau \& Rix 1999; Courteau \etal 2007b; Pizagno \etal 2007), suggesting that the TF relation is the edge on projection of the fundamental plane for late-type galaxies. The origin of these scaling relations in \LCDM cosmologies is typically thought to be the relation between halo virial velocity and virial mass, which scale as $\Vvir \propto \Mvir^{1/3}$. Accounting for the higher halo concentrations in lower mass dark matter haloes results in a shallower slope for the relation between the maximum circular velocity of dark matter haloes and the halo virial mass: $V_{\rm max,h} \propto \Mvir^{0.29}$ (Bullock \etal 2001). This scaling is similar to the observed stellar mass TF and FJ relations: $\sigma_{\rm e} \propto \Mstar^{0.29}$ (Gallazzi \etal 2006) and $V_{2.2}\propto \Mstar^{0.28}$ (Dutton \etal 2010b), where $\sigma_{\rm e}$ is the velocity dispersion within the half-light radius, and $V_{2.2}$ is the rotation velocity at 2.2 disk scale lengths. In what follows we define $\Vopt=V_{2.2}$ for late-types, and $\Vopt \propto \sigma_{\rm e}$ for early-types. However, for the $V_{\rm max,h}-\Mvir$ relation to be the direct origin of the TF and FJ relations requires that $\Vopt/V_{\rm max,h}$, and $\Mstar/\Mvir$ are constants. In Dutton \etal (2010b) we showed that this is at best only approximately the case. The relation between $\Vopt$ and $\Vvir$ depends on three factors: 1) The contribution of baryons to $\Vopt$; 2) The structure of the ``pristine'' dark matter halo (i.e., without the influence of baryons); and 3) The response of the dark matter halo to galaxy formation. For low mass star-forming galaxies gas dominates their baryonic budget, but for high mass star-forming galaxies and most non star-forming galaxies, stars dominate the baryon budget. A key uncertainty in measuring stellar masses is the stellar initial mass function (IMF). There is a factor of $\sim 2$ difference between the masses derived assuming the traditional Salpeter (1955) IMF compared with those derived assuming a Chabrier (2003) or Kroupa (2001) IMF. These latter IMFs are based on more recent measurements in the solar neighbourhood. The structure of ``pristine'' dark matter haloes, has been extensively studied using cosmological N-body simulations (e.g., Navarro, Frenk, \& White 1996a; Navarro, Frenk, \& White 1997; Bullock \etal 2001; Eke, Navarro \& Steinmetz 2001; Zhao \etal 2003; Navarro \etal 2004; Diemand \etal 2005, 2007; Macci\`o \etal 2007; Neto \etal 2007; Macci\`o \etal 2008; Zhao \etal 2009; Navarro \etal 2010; Klypin \etal 2010; Mu{\~n}oz-Cuartas \etal 2011). While the nature of the density profile in the inner 0.1\% of the virial radius is still uncertain, and the three parameter Einasto profile provides better fits than the broken power law of Navarro, Frenk, \& White (1997, hereafter NFW), (Merritt \etal 2006), on the scales relevant for modeling galaxy kinematics \LCDM dark matter haloes are well described by the two parameter NFW function. These two parameters are correlated with small scatter, such that the structure of dark matter haloes is almost completely determined by their mass. The response of the halo to galaxy formation has traditionally been modeled assuming galaxy formation is adiabatic, i.e., changes in potential are slow compared to the dynamical time. Under this assumption the halo is expected to contract, resulting in so-called adiabatic contraction (Blumenthal \etal 1986). The standard model assumes dark halo particles are on circular orbits. Using more realistic orbits results in weaker halo contraction (Wilson 2003; Gnedin \etal 2004; Sellwood \& McGaugh 2005). Recent hydrodynamical simulations of galaxy formation yield even weaker halo contraction (Abadi \etal 2010; Duffy \etal 2010; Pedrosa \etal 2010; Tissera \etal 2010). If galaxy formation is adiabatic, then the halo response should depend only on the final distribution of the baryons, and not on the assembly history. In these simulations the assembly history does matter, and thus this calls into question the basic assumption that halo response to galaxy formation is adiabatic. It is also possible for haloes to expand in response to galaxy formation, via a number of processes: rapid mass loss from the galaxy, e.g., driven by supernovae (Navarro, Eke, \& Frenk 1996b; Gnedin \& Zhao 2002; Read \& Gilmore 2005; Governato \etal 2010); dynamical friction operating on baryonic clumps (El-Zant, Shlosman \& Hoffman 2001; El-Zant \etal 2004; Elmegreen \etal 2008; Jardel \& Sellwood 2009) or galactic bars (Weinberg \& Katz 2002; Holley-Bockelmann \etal 2005; Sellwood 2008). Thus while the underlying structure of dark matter haloes is well understood, in order to understand the origin of the TF and FJ relations, it is necessary to know the stellar IMF, and how dark matter haloes respond to galaxy formation. Upper limits to stellar mass-to-light ratios, and hence the IMF, can be obtained from galaxy dynamics. For spiral galaxies, a Salpeter IMF is ruled out from maximal disk fits to resolved rotation curves (Bell \& de Jong 2001). An IMF with stellar masses 0.15 dex lower than a Salpeter, the so-called diet-Salpeter IMF, is the upper limit. For elliptical galaxies a Salpeter IMF is also ruled out for some galaxies (Cappellari \etal 2006). But massive elliptical galaxies ($\sigma>200 \kms$) are consistent with a Salpeter IMF (Bernardi \etal 2010), which may even be favored over lighter IMFs (Treu \etal 2010; Auger \etal 2010a). An alternative constraint on the stellar IMF comes from choosing the stellar mass-to-light ratio normalization that minimizes the scatter in the baryonic\footnote{The baryonic TF relation is the relation between rotation velocity and baryonic mass (i.e., stars plus cold gas).} TF relation (Stark, McGaugh, Swaters 2009). This favours a diet-Salpeter IMF with an uncertainty of $\pm 0.1$ dex in the stellar mass-to-light ratio. For a Chabrier IMF, elliptical galaxies require halo contraction to explain the mass discrepancy between the observed dynamical masses and those predicted when galaxies are embedded in NFW haloes (Schulz \etal 2010; Tollerud \etal 2011). However, for late-type galaxies, models with halo contraction over-predict the rotation velocities at fixed luminosity or stellar mass (Dutton \etal 2007; Dutton \& van den Bosch 2009; Trujillo-Gomez \etal 2010). In this paper we construct bulge-disk-halo models of early-type and late-type galaxies. These models are constrained to reproduce the distribution of stars and gas in galaxies, the relation between stellar mass and halo mass, and the structure of dark matter haloes in cosmological simulations. The key unknowns are the stellar IMF and the halo response to galaxy formation. We use the observed TF and FJ relations to place constraints on these two unknowns. As a by-product of this exercise, we measure the average dark matter fractions within the optical regions of early-type and late-type galaxies as a function of stellar mass. This paper is organized as follows. In \S 2 we describe the mass models. In \S 3 we discuss the observational constraints. In \S 4 we present model TF and FJ relations and compare to the observations. A discussion is given in \S5 and a summary in \S 6.
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We have established that either the stellar initial mass function, or the halo response to galaxy formation cannot be the same for early-type galaxies (ETGs) and late-type galaxies (LTGs). We now discuss physical mechanisms that could cause halo response to be different from the standard adiabatic contraction hypothesis, and others which could cause the IMF to vary. We discuss how these mechanisms might vary with galaxy type, and hence if they can explain, qualitatively our results. Two defining differences between ETGs and LTGs that may have consequences of the IMF and halo response, are the concentration of the stars and the star formation histories. ETGs have higher concentrations (associated with a spheroidal component) and they formed the bulk of stars at earlier times than LTGs. \subsection{What could cause halo contraction to vary with galaxy type?} There are a number of processes that occur during galaxy formation which can alter the structure of the dark matter halo. Some of these can also alter the morphological type of a galaxy. We outline these processes below. \begin{itemize} \item Smooth accretion: Smooth accretion of gas onto the central galaxy is the quintessential process that is expected to result in (adiabatic) halo contraction (Blumenthal \etal 1986). This process is expected to result in disk growth and star formation, and thus LTGs, but it could also occur during the history of ETGs. \item Dissipative major mergers: Major mergers which involve significant amounts of gas a.k.a. ``wet mergers'' are expected to result in halo contraction due to the mass that accumulates at the center of the galaxy. This process turns LTGs into ETGs. \item Non-dissipative major mergers: Major mergers which do not involve gas a.k.a. ``dry mergers'' are expected to result in effective halo contraction due to mixing of stars and dark matter through violent relaxation. This contraction effect partially undoes the segregation of baryons and dark matter produced by dissipation. Note that this is a different physical process than the standard adiabatic halo contraction. This process only occurs for ETGs. \item Clumpy accretion/minor mergers: Clumpy accretion of stars or gas in the form of minor mergers can cause halo expansion due to dynamical friction (e.g., El-Zant \etal 2001, Elmegreen \etal 2008; Romano-Diaz \etal 2008; Johansson \etal 2009). The clumps need to be baryon dominated, or else dark matter will be brought in to replace the dark matter that is removed by dynamical friction. The clumps also need to be dense, or else they will be tidally disrupted before they can alter the center of the halo. Clumpy cold accretion is more likely to occur at high redshifts (Dekel \etal 2009). This process can occur in both ETGs and LTGs. \item Feedback: The energy and momentum feedback from supernova, stellar winds and AGN can cause halo expansion under the following conditions: (1) By removing large amounts of baryons on a timescale much faster than in which they were accumulated (e.g., Navarro \etal 1996b; Gnedin \& Zhao 2002; Read \& Gilmore 2005; Governato \etal 2010); or (2) By inducing large scale bulk motions (Mashchenko \etal 2006; 2008). These processes are expected to be more effective in galaxies with the following properties: (1) lower mass and lower bulge fraction (i.e., LTGs), due to the shallower potential wells; (2) higher gas fractions, because a large fraction of the baryons need to be expelled for this process to be effective, and stellar mass cannot be expelled by feedback; (3) higher redshifts, due the order of magnitude higher specific star formation rates. \item Bars: Galactic bars can cause halo expansion due to dynamical friction on the bar from the dark matter halo (Weinberg \& Katz 2002). Since gas can get driven to the center of the galaxy, bar formation could also result in halo contraction. Bars require disks in which to form, and since bulges help to stabilize disks, bars are more frequent in late-type galaxies. However, bars are known to exist in S0 galaxies, which are considered to be ETGs, and thus the effects of bars on dark matter haloes are not necessarily restricted to LTGs. \end{itemize} Major mergers have long been thought to be the key process that determines the morphological type of a galaxy (Toomre \& Toomre 1972), with major mergers destroying disks, and producing spheroidal galaxies. Galaxy types are observed to vary strongly with stellar mass: low mass galaxies are predominately late-types (disk dominated, gas rich, star forming), while high mass galaxies are predominantly early-types (bulge dominated, gas poor, non star forming). This basic trend can be understood as a consequence of the mass dependence of the frequency of major mergers and the mass dependence of gas fractions of the progenitor galaxies (Maller 2008; Hopkins \etal 2009a,b). The mass dependence of the effects of dissipative major mergers can also qualitatively explain the tilt of the Fundamental Plane (Dekel \& Cox 2006). Thus the key physical process that occurs during early-type galaxy formation, which does not occur during late-type galaxy formation, is a major merger. Whether the merger is dissipative or non-dissipative we expect the haloes to contract, but for different reasons. This expectation needs to be tested, and quantified, with numerical simulations of galaxy/halo mergers. For late-type galaxies, a combination of clumpy cold accretion, and feedback during the early phases of galaxy formation could plausibly result in net halo expansion (e.g., Mo \& Mao 2004). Under this scenario, galaxies with a higher fraction of their stars in a spheroid (i.e., a classical bulge) should have experienced more halo contraction. This is qualitatively consistent with our result that for a fixed IMF early-type galaxies have more halo contraction than late-type galaxies (e.g., Fig.~\ref{fig:zp_ac}). To test this further for early-types and late-types separately, would require the velocity - stellar mass, halo mass - stellar mass, and structural scaling relations to be measured for galaxies with different bulge fractions. In the meantime we note that earlier type spirals have higher rotation velocities at fixed K-band luminosity and stellar mass than spiral galaxies in general (Noordermeer \& Verheijen 2007; Williams \etal 2010). Such a trend would be expected if earlier type spirals experienced more halo contraction than later type spirals. But there could be other explanations, such as more compact baryons, so it is too early to say if this supports our simple scenario. \subsection{What could cause the IMF to vary with galaxy type?} Observations in the Galactic disk suggest that the IMF has a power-law shape at masses above $1 \Msun$, and that it turns over at lower masses (Kroupa 2001; Chabrier 2003). This turnover can be modeled by a log-normal distribution with a characteristic turnover mass $m_{\rm c}$ (Chabrier 2003). The value of $m_{\rm c}$ is $\sim 0.1 \Msun$ in the disk of the Milky Way. Larson (1998, 2005) has argued that the characteristic turnover mass may largely be determined by the thermal Jeans mass, which strongly depends on the temperature of the ISM. An increased ISM temperature at higher redshifts is robustly expected based on the temperature of the cosmic microwave background, and also plausibly from higher star formation rates (SFR) which result in more supernova heat input, and lower metallicities, which result in less efficient cooling. An evolving IMF, in which the characteristic mass increases with increasing redshift, provides an explanation for a number of discrepancies: The difference in evolution of dynamical $M/L$ ratios and colors of early-type galaxies (van Dokkum 2008); The difference in evolution of the galaxy SFR - stellar mass relation between \LCDM galaxy formation models and observations (Dav\'e 2008); The difference between the observed stellar mass density of the universe and the implied stellar mass density from integrating the cosmic star formation history (Larson 2005; Hopkins \& Beacom 2006; Fardal \etal 2007). As shown by van Dokkum (2008), for IMFs with $m_{\rm c} < 0.08$ (i.e., more bottom heavy than a Chabrier IMF), the stellar $M/L$ ratios increase for decreasing $m_{\rm c}$. For $m_{\rm c} > 0.08$ (i.e., more bottom light than a Chabrier IMF) the stellar M/L ratio decreases. However, the relation between stellar $M/L$ and $m_{\rm c}$ is not monotonic. As $m_{\rm c}$ increases beyond $\sim 0.3 \Msun$ the stellar $M/L$ can actually {\it increase}. This is because the mass function becomes dominated by stellar remnants. For old enough stellar populations (Age $\sim 5-10$ Gyr) with $m_{\rm c} \sim 1$ the stellar $M/L$ can equal or even exceed that of a Salpeter IMF. Since ETGs form their stars at higher redshifts than LTGs, the evolving IMF as proposed by Van Dokkum (2008) and Dav\'e (2008) would cause ETGs to have higher present day stellar $M/L$ ratios than LTGs. The normalizations of the $M/L$ ratios are expected to be Salpeter like for ETGs, and not lower than 0.1 dex below Chabrier for LTGs (van Dokkum 2008). In the context of our results, stellar $M/L$ ratios close to Salpeter for ETGs are inconsistent with halo contraction. If the IMF for LTGs is close to Chabrier, then this would also favor models with halo expansion or no halo contraction. Thus this evolving IMF requires that the haloes of both ETGs and LTGs do not contract in response to galaxy formation. Recently van Dokkum \& Conroy (2010) derived constraints on the stellar IMF in the cores of massive elliptical galaxies using stellar absorption lines in the near-IR. They find strong evidence for an IMF with a steeper low mass slope than a Salpeter IMF, i.e., a bottom-heavy IMF. If the IMF is bottom-heavy throughout massive elliptical galaxies, and not just in their centers, then this is the opposite result to what is expected from the evolving IMF models of van Dokkum (2008) and Dav\'e (2008). This bottom heavy IMF results in stellar $M/L$ ratios a factor of $\sim 1.4$ higher than a regular Salpeter IMF. As shown in Fig.~\ref{fig:fdm} the dark matter fractions within the effective radii for the most massive early-type galaxies are $\sim 0.3$ for a Salpeter IMF. Thus, given the current uncertainties in $V_{\rm c}/\sigma_{\rm e}$, the IMF from van Dokkum \& Conroy (2010) is permitted to apply globally in massive early-type galaxies, and not just in their centers. However, this IMF would strongly over-predict the total masses within the effective radii of intermediate mass ($\Mstar \sim 10^{10}\Msun$ ) early-type galaxies. Thus based on our mass models we do not expect the bottom-heavy IMF of van Dokkum \& Conroy (2010) to be universal across early-type galaxies of different masses. \subsection{Comparison with previous studies} There are several recent studies that have addressed dark halo contraction and the stellar initial mass function of galaxies (Treu \etal 2010; Schulz \etal 2010; Trujillo-Gomez \etal 2010; Auger \etal 2010a; Napolitano \etal 2010). Most of these have focused on massive early-type galaxies. Although the individual conclusions vary, all of them are consistent with the following: ETGs with a Chabrier IMF plus un-contracted NFW haloes with standard halo concentrations do not have enough mass within the effective radius. Schulz \etal (2010) and Trujillo-Gomez \etal (2010) advocate models with halo contraction (Gnedin \etal 2004, and Blumenthal \etal 1986, respectively) to provide this additional mass, whereas Treu \etal (2010) and Auger \etal (2010a) advocate a Salpeter IMF. Schulz \etal (2010) argue against a Salpeter IMF based on the results of Cappellari \etal (2006). However, there are some caveats to this line of reasoning. Firstly the dynamical masses from Cappellari \etal (2006) are consistent with a Salpeter IMF for the most massive galaxies. Secondly the dynamical masses from Cappellari \etal (2006) assume mass follows light, which is expected to result in an underestimate of the dynamical masses if the dark matter fractions within the effective radii are significant. In \S \ref{sec:vsigma} we showed that the Cappellari \etal (2006) dynamical masses imply that, on average, $V_{\rm c}/\sigma_{\rm e} = 1.44 \pm 0.01$. In Fig.~\ref{fig:fdm} we showed that if $V_{\rm c}/\sigma_{\rm e} \gta 1.6$ then a Salpeter IMF is consistent for ETGs of all masses. Thus a Salpeter IMF is allowed for the most massive ETGs, and it is not yet robustly ruled out for intermediate mass ETGs (where the dark matter fractions are expected to be the lowest). Auger \etal (2010a) combined strong lensing, weak lensing and stellar dynamics for a sample of 53 massive elliptical galaxies to place constraints on the stellar IMF. They conclude that, given their model assumptions, the data strongly prefer a Salpeter like IMF over a lighter IMF such as Chabrier of Kroupa. While we agree that a model with a Salpeter like IMF can reproduce the observations, we find that our data cannot distinguish between models with Salpeter and Chabrier IMFs. Below we discuss two areas that could contribute to these differences: anisotropy and halo masses. The use of strong lensing and stellar dynamics can constrain the slope of the total mass profile within an effective radius (e.g., Koopmans \etal 2006). As we show in \S~\ref{sec:alpha} this information can help distinguish between models with different IMFs. However, a major source of systematic uncertainty is the anisotropy of the stellar orbits (Koopmans \etal 2009). Auger \etal (2010a) assumed isotropic orbits, $\beta=0$, and thus may have inadvertently favoured a particular IMF. In order to constrain the halo contraction model the halo mass needs to be accurately determined. This is for 2 reasons. Firstly, the halo mass is needed to predict the typical pristine halo concentration using cosmological N-body simulations. Secondly, to provide an accurate normalization of the halo mass profile. For a fixed IMF, the relation between stellar mass and halo mass has been determined using halo abundance matching, weak lensing and satellite kinematics. These techniques yield consistent results, and in particular for massive early-type galaxies (Dutton \etal 2010b; More \etal 2011), and thus provide a consistency check on the models of Auger \etal (2010a). Taking a model with a Chabrier IMF and Gnedin \etal (2004) halo contraction, Auger \etal (2010a) find halo masses a factor of $\sim 0.5$ dex lower than obtained by Moster \etal (2010). The relation from Moster \etal (2010) is in good agreement with the relations we use in this paper for massive early-type galaxies (see Fig. 1 in Dutton \etal 2010b). Thus Auger \etal (2010a) is inferring abnormally low halo masses at fixed stellar mass, which may be biasing their results (as previously discussed by Tortora \etal 2010). Schulz \etal (2010) argue that the halo contraction models of Abadi \etal (2010) and Blumenthal \etal (1986) are inconsistent with the data. We disagree with this conclusion, as we show that for a Chabrier IMF, all three halo contraction models are consistent with the data, given reasonable systematic uncertainties. Furthermore, distinguishing between the Gnedin \etal (2004) and Abadi \etal (2010) models requires stellar masses to be measured to an accuracy of 0.1 dex. Such accuracy may in principle be achievable for early-type galaxies (Gallazzi \& Bell 2009), but the current limiting factors are uncertainties in stellar population synthesis models (Conroy, Gunn, \& White 2009). The degeneracy between halo contraction and stellar IMF for early-types was also discussed by Napolitano \etal (2010), with similar qualitative conclusions as we find here. However, a limitation of this study was the treatment of the stellar mass to halo mass ratio as a free parameter. In our analysis the halo masses are constrained through results from weak lensing and satellite kinematics, which enables us to make more quantitative conclusions regarding the nature of dark halo response for a given IMF. Finally, we note that the conclusion of Napolitano \etal (2010) that the relation between the central dark matter density and effective radius provides evidence for cuspy dark matter haloes is in fact degenerate with the IMF. Our conclusions for late-type galaxies are in agreement with those of Dutton \etal (2007), namely that for a Chabrier IMF, halo expansion is required to match the zero point of the TF relation (as well as galaxy sizes). Trujillo-Gomez \etal (2010) claim their results are incompatible with the conclusions of Dutton \etal (2007). However, their figures show that a model with Blumenthal \etal (1986) adiabatic halo contraction is consistent with the velocity-luminosity (VL) relation of early-types, but it does not match the VL relation of late-types. Their model without halo contraction provides a better fit to the VL relation of late-types. Their results are thus consistent with our findings. \subsection{Future prospects} There are a number of techniques that are capable of constraining the IMF and/or dark matter fractions in galaxies. Here we give a brief outline of these. Upper limits to stellar $M/L$ ratios are obtainable from both strong lensing and dynamical models. The strongest constraints are expected for intermediate mass early-type galaxies, as these are expected to have the highest baryon fractions within the effective radius (Fig.~\ref{fig:fdm}). The ATLAS3D project (Cappellari \etal 2011) contains $\sim 10$ times more galaxies than studied by Cappellari \etal (2006), and thus promises to provide stronger constraints on $V_{\rm circ}(R_{\rm e })/\sigma_e$ over a wider range of galaxy masses than previous dynamical studies. Strong lensing has the potential to provide stronger constraints than dynamical models, but it is currently limited by the sparsity of known strong lenses with $\Mstar\sim 10^{10}\Msun$. Furthermore, low mass early-type galaxies tend to be satellites, which adds an extra complication to inferring total masses from strong lensing. Strong lensing will be able measure the total projected mass accurately, but the problem will be disentangling the mass of the satellite from that of its host. There is a well known disk-halo degeneracy that plagues the decomposition of galaxy rotation curves into baryonic and dark matter components (e.g., van Albada \& Sancisi 1986; van den Bosch \& Swaters 2001; Dutton \etal 2005). Strong gravitational lensing of high inclination disk-dominated galaxies can provide a complementary information to that obtainable from kinematics. Specifically, strong lensing measures projected mass, and ellipticity of projected mass. Both of which depend on the disk mass fraction, and thus a combined strong lensing and dynamics analysis can place constraints on the stellar $M/L$ ratio (Dutton \etal 2011, in prep). This technique has not been fully exploited due to the lack of known disk dominated strong lenses. However, searches for spiral galaxy strong lenses are underway (F{\'e}ron \etal 2009; Sygnet \etal 2010; Treu \etal 2011, in prep), and thus the primary limitation of this method will be soon overcome. An absolute constraint on disk masses can be obtained by using the fact that the disk surface mass density is a function of the vertical velocity dispersion and the disk scale height (Bottema 1993). This method is being applied by the Disk Mass Survey (Verheijen \etal 2007; Bershady \etal 2010). They are measuring the disk mass density profile from vertical velocity dispersions and a statistical measurement of disk scale heights. By subtracting off the observed gas mass density this gives the stellar mass density profile. This method is limited to regions of galaxies where the disk dominates the baryons, i.e., it does not apply to elliptical galaxies or the bulges of spiral galaxies. It is also a statistical method, since it requires knowledge of two parameters that cannot be measured simultaneously. Constraints on dark matter fractions can be obtained from the scatter in the velocity-mass (VM) and size-mass (RM) relations. The basic idea is that the strength of correlation between residuals of the VM and RM relations depends on the dark matter fraction. This method has been applied to late-type galaxies (Courteau \& Rix 1999; Dutton \etal 2007; Gnedin \etal 2007), but the interpretation in terms of dark matter fractions are not always unique. We plan to apply this method to early-type galaxies in a future paper (Dutton \etal 2011, in prep.) The low-mass end of the IMF can be constrained with stellar absorption lines (van Dokkum \& Conroy 2010). The lines are weak and at $\sim 900$ nm, so that this method is only applicable to non-star forming galaxies. The galaxy redshifts are currently limited to be very low by detector technology. This method has only currently been applied to the centers of massive early-type galaxies, with evidence for an IMF more bottom heavy than Salpeter. It would be very interesting to see this method applied radially and in lower mass early-types.
| 10 | 12 |
1012.5859
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1012
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1012.0198_arXiv.txt
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Phase-resolved observations of the solar-type star $\xi$ Bootis A were obtained using the NARVAL spectropolarimeter at the Telescope Bernard Lyot (Pic du Midi, France) during years 2007, 2008, 2009 and 2010. The data sets enable us to study both the rotational modulation and the long-term evolution of various magnetic and activity tracers. Here, we focus on the large-scale photospheric magnetic field (reconstructed by Zeeman-Doppler Imaging), the Zeeman broadening of the FeI 846.84 nm magnetic line, and the chromospheric CaII H and H$\alpha$ emission.
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$\xi$ Boo A is the primary component of a visual binary system which follows a 151-yr orbital period \citep{hershey77}. Its mass is $0.86 \pm 0.07 M_{\odot}$ \citep{fernandes98} and its mean effective temperature $T_{eff}$ is known to be 5,550 K \citep{gray94}. \citet*{toner88} determined a rotational period of 6.43 days which we used to determine the rotational phases of our observations. The equatorial velocity projected on the line of sight is equal to $2.9 \pm 0.4$ km/s \citep{gray84}. $\xi$ Boo A is a very active star, with irregular fluctuations of chromospheric emission \citep{baliunas95}. Its photospheric magnetic field was first detected by \citet*{robinson80}. A first attempt at modelling the magnetic geometry of the star was proposed by \citet{petit05}, who reconstructed a surface distribution of the magnetic field dominated by a large-scale toroidal magnetic component.
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An obvious decrease of the global magnetic field is visible between 2007 and 2010 in the magnetic maps. The same kind of evolution occurs for the CaII H and H$\alpha$ index during the same period, as well as for the width of FeI 846.84 nm magnetic line. During the same time interval, the fraction of large-scale magnetic energy stored in the poloidal field component is multiplied by a factor of two, while the opposite trend is observed in the unsigned magnetic flux. During the timespan of our monitoring, the large-scale magnetic geometry of $\xi$ Bootis A does not experience dramatical changes like the global magnetic polarity reversals recently reported for HD190771 \citep{petit09} or $\tau$ Bootis \citep{fares09}. Future monitoring of the star will tell us wether its magnetic evolution is associated to some kind of cyclicity, and if its global magnetic field can undergo polarity switches, as only observed on more massive solar-type stars up to now.
| 10 | 12 |
1012.0198
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1012
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1012.2362_arXiv.txt
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We present the results of an ongoing survey of cool, late-type supergiants - the descendants of massive O- and B-type stars - that has systematically detected magnetic fields in these stars using spectropolarimetric observations obtained with ESPaDOnS at the Canada-France-Hawaii Telescope. Our observations reveal detectable, often complex, Stokes $V$ Zeeman signatures in Least-Squares Deconvolved mean line profiles in a significant fraction of the observed sample of $\sim$30 stars.
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Supergiants are the descendants of massive O and B-type main sequence stars. Unlike their main sequence progenitors, cool supergiants are characterized by a helium-burning core and a deep convective envelope. Due to their extended radii, low-atmospheric densities, slow rotation and long convective turnover times, supergiants provide an opportunity to study stellar magnetism at the extremes of parameter space. In fact, observations of late-type supergiants show characteristics consistent with magnetic activity, such as luminous X-ray emission and flaring, and emission in chromospheric UV lines - phenomena suggesting the presence of dynamo-driven magnetic fields. Motivated by the activity-related puzzles of late-type supergiants, the near complete lack of direct constraints on their magnetic fields, and recent success of measuring fields of red and yellow giants (e.g. Auri\`{e}re et al. 2008), we have initiated a program to search for direct evidence of magnetic fields in these massive, evolved stars. Here we summarize the recent results of Grunhut et al. (2010).
| 10 | 12 |
1012.2362
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1012
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1012.3086_arXiv.txt
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Galactic Cosmic Rays (GCRs) are mainly protons confined in the galactic magnetic field to form an isotropic flux inside the galaxy. Before reaching the Earth orbit they enter the Heliosphere and undergo diffusion, convection, magnetic drift and adiabatic energy loss. The result is a reduction of particles flux at low energy (below 10 GeV), called solar modulation. We realized a quasi time-dependent 2D Stochastic Simulation of Solar Modulation that is able to reproduce CR spectra once known the Local Interstellar Spectrum (LIS). We were able to estimate the different behaviors associated to the polarity dependence of the Heliospheric modulation for particles as well as for antiparticles. We show a good agreement with the antiproton/proton ratio measured by AMS-01, Pamela, BESS, Heat and Caprice and we performed a prediction for the AMS-02 Experiment.
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The effect of the heliospheric structure on GCRs propagation can be reproduced by a two dimensional (radius and helio-colatitude) Stochastic model solving numerically the Parkers's equation \cite{prot_art}\!. If we do not take into account the effects of the Earth magnetosphere \cite{mi_jgr}\! modulated fluxes depends not only on the level of solar activity but also on particles charge sign and solar magnetic field polarity \cite{art_midrift}\!. The study of the modulation of $\bar{p}/p$ ratios is particularly important, because it includes explicitly the combination of charge sign and polarity dependence. The Local Interstellar Spectra (LIS) used as input of the code, both for protons and antiprotons, are taken by the Galprop model\cite{galprop}\!.
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\label{sec:conc} We built a 2D stochastic Monte Carlo code for particles propagation across the heliosphere. Present model takes into account drift effects and shows quantitatively a good agreement with measured values, both for positive and negative periods and for different particles and charge sign. This is relevant because particles with opposite charge sign undergo a different solar modulation\cite{art_midrift}\!. We compared our simulations with antiproton/proton ratios measured by BESS and PAMELA. We used {\it dynamic} parameters values ($K_{0}$, $\alpha$ and $V_{sw}$) for the related periods, in order to reproduce the propagation of incoming GCR through magnetic disturbances carried by the outgoing solar wind. The {\it dynamic} description of the heliosphere and the forward approach seem to reproduce better the real physical propagation of GCR in the solar cavity. In order to have a more sophysticated model we need to introduce a dependence on the particle time spent in the heliosphere and a larger statistics of measured data during negative solar field periods, as AMS-02 will provide in the next years. Recent measurements\cite{pamela}\! have pointed out the needs to reach a high level of accuracy in the modulation of the fluxes, in relation to the charge sign of the particles and the solar field polarity\cite{mi_ap}\!. This aspect will be even more crucial in the next generation of experiments\cite{zuccon,casaus}\!.
| 10 | 12 |
1012.3086
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1012
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1012.4459_arXiv.txt
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We model the cosmological substratum by a viscous fluid that is supposed to provide a unified description of the dark sector and pressureless baryonic matter. In the homogeneous and isotropic background the \textit{total} energy density of this mixture behaves as a generalized Chaplygin gas. The perturbations of this energy density are intrinsically non-adiabatic and source relative entropy perturbations. The resulting baryonic matter power spectrum is shown to be compatible with the 2dFGRS and SDSS (DR7) data. A joint statistical analysis, using also Hubble-function and supernovae Ia data, shows that, different from other studies, there exists a maximum in the probability distribution for a negative present value $q_{0} \approx - 0.53$ of the deceleration parameter. Moreover, different from other approaches, the unified model presented here favors a matter content that is of the order of the baryonic matter abundance suggested by big-bang nucleosynthesis.
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According to the prevailing interpretation, our Universe is dynamically dominated by a cosmological constant $\Lambda$ (or a dynamical equivalent, called dark energy (DE)) which contributes more than 70\% to the total cosmic energy budget. More than 20\% are contributed by cold dark matter (CDM) and only about 5\% are in the form of conventional, baryonic matter. Because of the cosmological constant problem in its different facets, including the coincidence problem, a great deal of work was devoted to alternative approaches in which a similar dynamics as that of the $\Lambda$CDM model is reproduced with a time varying cosmological term, i.e., the cosmological constant is replaced by a dynamical quantity. Both dark matter (DM) and DE manifest themselves so far only through their gravitational interaction. This provides a motivation for approaches in which DM and DE appear as different manifestation of one single dark-sector component. The Chaplygin-gas model and its different generalizations realize this idea. Unified models of the dark sector of this type are attractive since one and the same component behaves as pressureless matter at high redshifts and as a cosmological constant in the long time limit. While the homogeneous and isotropic background dynamics for the (generalized) Chaplygin gas (GCG) is well compatible with the data, the study of the perturbation dynamics resulted in problems which apparently ruled out all Chaplygin-gas type models except those that are observationally almost indistinguishable from the $\Lambda$CDM model. To circumvent this problem, nonadiabatic perturbations were postulated and designed in a way to make the effective sound speed vanish. But this amounts to an ad hoc procedure which leaves open the physical origin of nonadiabatic perturbations. There exists, however, a different type of unified models of the dark sector, namely viscous cosmological models, which are intrinsically nonadiabatic \cite{VDF}. In the homogeneous and isotropic background a one-component viscous fluid shares the same dynamics as a GCG. Now, what is observed in the redshift surveys is not the spectrum of the dark-matter distribution but the baryonic matter spectrum. Including a baryon component into the perturbation dynamics for a universe with a Chaplygin-gas dark sector, there appears the new problem that the unified Chaplygin-gas scenario itself is disfavored by the data. It is only if the unified scenario with a fixed pressureless (supposedly) baryonic matter fraction of about $0.043$ (according to the results from WMAP and primordial nucleosynthesis) is \textit{imposed} on the dynamics, that consistency with the data is obtained \cite{chaprel}. If the pressureless matter fraction is left free, its best-fit value is much larger than the baryonic fraction. In fact it becomes even close to unity, leaving only a small percentage for the Chaplygin gas, thus invalidating the entire scenario. In other words, a Chaplygin-gas-based unified model of the dark sector is difficult to reconcile with observations. One may ask now, whether the status of unified models can again be remedied by replacing the Chaplygin gas by a viscous fluid. It is exactly this question that we are going to investigate in the present paper. It is our purpose to study cosmological perturbations for a two-component model of baryons and a viscous fluid, where the latter represents a one-component description of the dark sector \cite{BaVDF}. We shall show that such type of unified model is not only consistent for a fixed fraction of the baryons but also for the case that the matter fraction is left free. Our analysis demonstrates that the statistically preferred value for the abundance of pressureless matter is compatible with the mentioned baryon fraction $0.043$ that follows from the synthesis of light elements.
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The approach presented here is based on the fact that the two-component system of a bulk viscous fluid and a separately conserved baryon component behaves in the background as a generalized Chaplygin gas with $\alpha = -\frac{1}{2}$. While the baryon component may be considered dynamically negligible in the background, the situation is different on the perturbative level, since the observed matter agglomerations are related to baryonic density fluctuations. These fluctuations are obtained from a combination of the nonadiabatic total energy density perturbations and relative entropy perturbations in the two-component system where the former source the latter. The probability distribution for the deceleration parameter has a maximum at $q_{0} \approx -0.53$ which partially removes the degeneracy of previous studies which, taken at face value, were incompatible with an accelerated expansion and thus in obvious tension with results for the background. Perhaps still more important is the test of the unified model itself. Many investigations on approaches with a unified dark sector fix the pressureless matter component to be that of the favored (by the WMAP data) baryon fraction and then check whether or not the resulting dynamics can reproduce the observations. But this does not say anything on how probable the division of the total cosmic substratum into roughly 96\% of a dark substance and roughly 4\% of pressureless matter is. To decide this question, one has to consider the pressureless matter fraction as a free parameter and to find out which abundance is actually favored by the data. Our analysis revealed that the matter fraction probability is indeed highest for values smaller than roughly 8\%. This is a result in favor of the unified viscous model. We recall that a corresponding analysis for a Chaplygin gas results in values close to unity \cite{chaprel} which seems to rule out such type of approaches. The present viscous model, on the other hand, remains an option for a unified description of the dark sector, at least as far as the matter power spectrum is concerned. \ack Support by CNPq and FAPES is gratefully acknowledged.
| 10 | 12 |
1012.4459
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1012
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1012.1560_arXiv.txt
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The number of known transiting exoplanets is rapidly increasing, which has recently inspired significant interest as to whether they can host a detectable moon. Although there has been no such example where the presence of a satellite was proven, several methods have already been investigated for such a detection in the future. All these methods utilize post-processing of the measured light curves, and the presence of the moon is decided by the distribution of a timing parameter. Here we propose a method for the detection of the moon {\it directly in the raw transit light curves.} When the moon is in transit, it puts its own fingerprint on the intensity variation. In realistic cases, this distortion is too little to be detected in the individual light curves, and must be amplified. Averaging the folded light curve of several transits helps decrease the scatter, but it is not the best approach because it also reduces the signal. The relative position of the moon varies from transit to transit, the moon's wing will appear in different positions on different sides of the planet's transit. Here we show that a careful analysis of the scatter curve of the folded light curves enhances the chance of detecting the exomoons directly.
| 10 | 12 |
1012.1560
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1012
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1012.3753_arXiv.txt
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{ Though Arp\,220 is the closest and by far the most studied ULIRG, a discussion is still ongoing on the main power source driving its huge infrared luminosity. } { To study the molecular composition of Arp\,220 in order to find chemical fingerprints associated with the main heating mechanisms within its nuclear region. } { We present the first aperture synthesis unbiased spectral line survey toward an extragalactic object. The survey covered the 40~GHz frequency range between 202 and 242~GHz of the 1.3~mm atmospheric window. } { We find that 80\% of the observed band shows molecular emission, with 73 features identified from 15 molecular species and 6 isotopologues. The $^{13}$C isotopic substitutions of HC$_3$N and transitions from H$_2^{18}$O, $^{29}$SiO, and CH$_2$CO are detected for the first time outside the Galaxy. No hydrogen recombination lines have been detected in the 40~GHz window covered. The emission feature at the transition frequency of H31$\alpha$ line is identified to be an HC$_3$N molecular line, challenging the previous detections reported at this frequency. Within the broad observed band, we estimate that 28\% of the total measured flux is due to the molecular line contribution, with CO only contributing 9\% to the overall flux. We present maps of the CO emission at a resolution of $2.9''\times1.9''$ which, though not enough to resolve the two nuclei, recover all the single-dish flux. The 40~GHz spectral scan has been modelled assuming LTE conditions and abundances are derived for all identified species. } { The chemical composition of Arp\,220 shows no clear evidence of an AGN impact on the molecular emission but seems indicative of a purely starburst-heated ISM. The overabundance of H$_2$S and the low isotopic ratios observed suggest a chemically enriched environment by consecutive bursts of star formation, with an ongoing burst at an early evolutionary stage. The large abundance of water ($\sim10^{-5}$), derived from the isotopologue H$^{18}_2$O, as well as the vibrationally excited emission from HC$_3$N and CH$_3$CN are claimed to be evidence of massive star forming regions within Arp\,220. Moreover, the observations put strong constraints on the compactness of the starburst event in Arp\,220. We estimate that such emission would require $\sim2-8\times10^6$ hot cores, similar to those found in the Sgr~B2 region in the Galactic center, concentrated within the central 700~pc of Arp\,220. }
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At a redshift of $z=0.018$, Arp\,220 is the closest ultraluminous infrared galaxy (ULIRG). This galaxy is an advanced merger system as evidenced by the large tidal tails observed in the optical \citep{Joseph1985,Kim2002,Koda2009} and the double nuclei (separation $\sim0.98''$) observed in radio \citep{Norris1988}, mm \citep{Scoville1997,Downes1998,Sakamoto1999}, sub-mm \citep{Sakamoto2008} and near-IR wavelenghts \citep{Graham1990,Scoville1998}. These nuclei are surrounded by two counterrotating gas disks as well as a larger outer disk encompassing both \citep{Sakamoto1999,Mundell2001}. The main nuclear powering sources in ULIRGs are thought to be starbursts events, active galactic nuclei (AGN), or a combination of both. Mid-IR surveys of ULIRGs with ISO \citep{Genzel1998} suggested that $70\%-80\%$ of the population are powered by star formation while only $20\%-30\%$ are AGN powered. This result is consistent with that from near-IR searches for hidden broad-line regions (BLRs) towards ULIRGs that lack BLR signatures in the optical \citep{Veilleux1999}. The identification of the main power source becomes very elusive in extremely obscured nuclei like those of Arp\,220. The nuclei in Arp\,220 are affected by a severe obscuration at $2.2\mu$m \citep{Scoville1998}. Even at 1~mm, the dust towards the more massive western nucleus is found to be significantly optically thick \citep[$\tau\sim1$,][]{Downes2007}. Thus, a number of arguments have been proposed in favor of both powering scenarios to explain the large observed IR luminosity of Arp\,220, though none of them have been conclusive. Among the evidence favoring the AGN-powered scenario are studies showing hard X-ray emission strongly concentrated towards the nuclei, a hard spectrum X-ray point source close to the position of the Western nucleus, and a softer point source towards the Eastern nucleus \citep{Clements2002}. Additionally, the luminosity ratio $L_{X(2-10\,\rm keV)}/L_{FIR}$ is unusually low relative to what is observed in starburst galaxies \citep{Iwasawa2001,Iwasawa2005}, and the 1.98\,keV equivalent width of the 6.7\,keV line of Fe is too large to be purely starburst-driven \citep{Teng2009}. The high column densities and therefore obscuration could hide a Compton thick AGN \citep{Sakamoto1999,Sakamoto2008,Downes2007}. However, the starburst-driven scenario also appears to be supported by a number of observations. The detection of dozens of radio supernovae (RSNe) in both nuclei of Arp\,220 is a clear indication of the starburst events in this galaxy \citep{Lonsdale2006,Parra2007}. The supernova rate of $4\pm2\,\rm yr^{-1}$ agrees with the derived star formation rate based on its FIR luminosity \citep[$\sim300~M_\odot\,{\rm yr}^{-1}$,][]{Dopita2005}, which implies that the starburst traced by the detected RSNe is able to produce the observed FIR luminosity \citep{Lonsdale2006}. However this conclusion is based on assumptions about the truncation of the Initial Mass Function (IMF) and that all supernova events result in RSNe, which requires a dense and compact starburst environment \citep{Smith1998}. Towards the Eastern nucleus, the agreement between the supernovae and the dust surface density suggests a starburst heated dust \citep{Sakamoto2008}. A comparison of the OH megamaser emission \citep{Lonsdale1998} with the sub-mm continuum peak suggests that even if the masers might be associated to an AGN, it would not be the main contributor to the dust heating \citep{Sakamoto2008}. Bright water vapor megamaser emission at 183\,GHz suggests the presence of $\sim10^6$ Sgr\,B2-like hot cores within the central kiloparsec of Arp\,220 \citep{Cernicharo2006}. Extended faint soft X-ray emission is detected beyond the optical boundaries of the galaxy with bright plumes extending 11\,kpc, claimed to be the result of the starburst driven superwinds \citep{McDowell2003}. The dense gas phase of the interstellar medium (ISM) in Arp220 has been targeted by numerous observations of molecular tracers emission in the mm and sub-mm wavelenghts. These observations provided additional constraints on the nature of the power source from a ISM chemical composition point of view. HNC emission has been observed to be overluminous with respect to HCN in Arp\,220 \citep{Huettemeister1995,Cernicharo2006,Aalto2007}. A similar relative enhancement of HNC is only observed in Mrk\,231, with a dominant AGN, and NGC\,4418, with a putative buried AGN, which supports the high HNC/HCN ratio to be a chemical indicator of an X-ray dissociation region \citep[XDR,][]{Aalto2007}. The two nuclei of Arp\,220 are revealed to show different physical conditions as observed from the high resolution maps of HNC \citep{Aalto2009}. Model calculations attribute the observed H$_3$O$^+$ emission to enhanced X-ray irradiation \citep{vanDerTak2008}. In the FIR, however, the lower level absorption lines in the spectrum of Arp\,220 match those observed in the diffuse clouds in the envelope of the Sgr~B2 molecular cloud complex hosting a massive star forming event \citep{Gonz'alez-Alfonso2004}. Although the molecular line observations in Arp\,220 are still limited to the brighter species, the compilation of molecular transitions by \citet{Greve2009} as well as the detection of complex organic species \citep{Salter2008} show that this galaxy is one of the brightest molecular emitters outside the Galaxy together with the starbursts galaxies NGC\,253 and M\,82. This fact turns Arp\,220 into a well suited candidate for molecular line surveys. It is clear that no matter which powering source drives the large IR luminosity in Arp\,220, it will certainly have an imprint on the physical properties and chemical composition of the ISM in this galaxy. In this paper we present the first unbiased chemical study in the 1.3~mm spectral band of Arp\,220 with the aim of finding additional physicochemical clues on the nature of its hidden power source.
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Up to now, the only available extragalactic molecular studies available as comparison templates are the unbiased study toward NGC\,253 \citep{Mart'in2006} and M\,82 (Aladro et al., in prep), and the targeted observations towards NGC\,4945 \citep{Wang2004}. These studies have characterized the chemical complexity of starburst galaxies. The molecular abundances in Arp\,220, resulting from the detected transitions in the 1.3~mm atmospheric window, do resemble those found towards the starburst galaxy NGC\,253. The only species showing an outstanding abundance towards Arp\,220 is H$_2$S. Although the abundance of H$_2$S has only been measured towards NGC\,253, this overabundance is likely to be the result of grain disruption and subsequent H$_2$S injection into gas phase in the early stages of star formation \citep{Hatchell1998}. This sets the burst of star formation in Arp\,220 in a very early stage of evolution and thus, on average, younger than the starburst in NGC\,253. This comparison suggests that the chemistry in Arp\,220 is mostly driven by bursts of star formation where differences in the observed abundances can be accounted for by a difference in the state of evolution of the starbursts \citep{Mart'in2006}. Unfortunately, no complete chemical templates of the ISM in the vicinity of AGNs are available and therefore the best tracers in highly X-ray irradiated regions are based on chemical models. The results from such models \citep{Meijerink2005,Meijerink2006,Meijerink2007,Loenen2008} have motivated a number of observational studies focused on the species CO, HCN, HNC, HCO$^+$, and CN \citep{P'erez-Beaupuits2007,Krips2008,Baan2008,P'erez-Beaupuits2009}. The HCN/HCO$^+$ ratio was used by \citep{Aalto2007a} to evaluate the presence of an AGN on the LIRG NGC\,4418 which, like Arp\,220, has a heavily obscured nucleus. The HCN and HCO$^+$ comparative study by \citet{Krips2008} shows how the HCN/HCO$^+$ ratios in Arp\,220 are similar to those measured in NGC\,1068, the prototypical AGN galaxy. However the excitation of HCN differs significantly in both galaxies \citep{Krips2008}. On the other hand, though overluminous HNC emission in Arp\,220 could be attributed to either pumping through mid-IR photons or the presence of an XDR region \citep{Aalto2007}, the observed HNC/HCN ratio suggest a PDR rather than a XDR origin \citep{Baan2008,Baan2010}. We estimate a column density ratio CN/HCN$\sim0.5$, where we have used the single-dish HCN observations by \citet{Krips2008} and our survey measurement of the CN emission. To estimate the HCN column density we assumed a source size of $2''$ similar to the one used in this paper. This ratio is lower than the range CN/HCN$\sim1-4$ that \citep{P'erez-Beaupuits2009} claims as indicative of XDR regions. Therefore, the CN emission does not appear to be enhanced as observed in the circumnuclear disk of NGC\,1068 \citep{Garc'ia-Burillo2010}. Moreover, \citet{Garc'ia-Burillo2010} found an enhancement of SiO in this region, which does not seem to be particularly prominent in Arp\,220 as compared to NGC\,253. Therefore, we do not find any molecular emission that could be attributed to X-ray driven chemistry in the nuclear region of Arp\,220. Measured isotopologues ratios appear to point to both large opacities affecting the main isotopologues of most observed species and an enriched molecular material in the nuclear region of Arp\,220. ISM enrichment towards Arp\,220 is similar or even enhanced with respect to that found in other starbursts, such as NGC\,253 and M\,82, as a consequence of a series of consecutive short and intense starburst events \citep{Parra2007}. Among the isotopologues, the observed H$^{18}_2$O luminosity can be accounted for by the emission of a few $10^6$ hot molecular cores associated with the massive star formation within the central 700 pc of Arp\,220. Such a concentration of cores could be responsible for the whole bolometric luminosity of this galaxy, rendering unnecessary a significant contribution to the luminosity by a deeply embedded AGN. Far from the H$_2$O atmospheric absorption, the optically thin emission of H$^{18}_2$O is proven to be one of the best tracers of massive star forming hot cores in highly obscured nuclei of galaxies. The importance of the star formation within the nuclear region of Arp\,220 is further supported by the detection of vibrationally excited emission of HC$_3$N and CH$_3$CN, with vibrational temperatures $>300$\,K. Such emission is also tracing the molecular component associated with hot cores as reported by Mart\'in-Pintado (in prep.). However, vibrationally excited HC$_3$N emission with a temperature of 500~K towards NGC\,4418 could be understood as a compact and deeply embedded AGN, heating up the surrounding material \citep{Costagliola2010}. It is remarkable that such vibrationally excited emission has never been reported towards nearby starbursts. The detection towards the ULIRG Arp\,220 and the LIRG NGC\,4418, might be due to the significantly larger contribution of hot core emission as a consequence of the higher star formation rates in these galaxies. Large opacities affect the pure rotational transitions of HC$_3$N (Sect.~\ref{sect.isotRat}). Due to the large continuum opacity at 1mm \citep{Downes2007}, the molecular emission we observed cannot arise from the vicinity of an AGN. Indeed, the simple models presented \citet{Schleicher2010} show the heating sphere of influence of a supermassive AGN to be limited to the central $\sim100$\,pc while the starburst heating dominates outside this volume. Therefore, the hot gas where the vibrationally excited emission of CH$_3$CN and HC$_3$N cannot be tracing the regions around the Compton thick AGN but must arise from other regions likely unaffected by the AGN radiation and purely heated by the starbursting events in Arp\,220.
| 10 | 12 |
1012.3753
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1012
|
1012.4804_arXiv.txt
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We present measurements of dust reddening using the colors of stars with spectra in the Sloan Digital Sky Survey. We measure reddening as the difference between the measured and predicted colors of a star, as derived from stellar parameters from the SEGUE Stellar Parameter Pipeline \citep{Lee:2008a}. We achieve uncertainties of 56, 34, 25, and 29 mmag in the colors $u-g$, $g-r$, $r-i$, and $i-z$, per star, though the uncertainty varies depending on the stellar type and the magnitude of the star. The spectrum-based reddening measurements confirm our earlier ``blue tip'' reddening measurements \citep[S10]{Schlafly:2010}, finding reddening coefficients different by $-3$\%, 1\%, 1\%, and 2\% in $u-g$, $g-r$, $r-i$, and $i-z$ from those found by the blue tip method, after removing a 4\% normalization difference. These results prefer an $R_V=3.1$ \citet[F99]{Fitzpatrick:1999} reddening law to \citet{O'Donnell:1994} or \citet{Cardelli:1989} reddening laws. We provide a table of conversion coefficients from the \citet[SFD]{Schlegel:1998} maps of $E(B-V)$ to extinction in 88 bandpasses for 4 values of $R_V$, using this reddening law and the 14\% recalibration of SFD first reported by S10 and confirmed in this work. \emph{Subject headings: } dust, extinction --- Galaxy: stellar content --- ISM: clouds
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Dust is composed of heavy elements produced by the nuclear burning of stars. These heavy elements are blown out of the stars in winds and explosions, and are reprocessed in the interstellar medium to eventually form dust grains \citep{Draine:2003}. The dust scatters and absorbs light, especially in the ultraviolet through infrared, according to the dust reddening law. The dust also emits photons thermally in the far-infrared. Accordingly, mapping the dust is a central problem in astronomy. In previous work with the blue tip of the stellar locus \citep[S10]{Schlafly:2010}, we examined the reddening law and the accuracy of the \citet[SFD]{Schlegel:1998} dust map using photometry from the Sloan Digital Sky Survey \citep[SDSS]{York:2000} and the uniformity of the color of the blue tip of the stellar locus over the sky. This blue tip work recommended a 14\% recalibration of the SFD dust map in the sense that $E(B-V) = 0.86\cdot\EBVSFD$, and a preference for a \citet[F99]{Fitzpatrick:1999} reddening law over other reddening laws. In this work we set out to test that result using an independent set of data. The SDSS stellar spectra provide an independent test of reddening. Stellar spectra sensitively test reddening because the broadband photometry of a star is almost completely determined by three parameters: temperature, metallicity, and gravity. These paramaters can be determined using only line information in the spectra, allowing the intrinsic broadband colors of the star to be predicted independently from the observed colors of the star. Dust intervening between us and the star, however, will shift the observed colors relative to the intrinsic colors. The difference between the predicted intrinsic colors and measured colors constitutes a measurement of the reddening to that star. This method is broadly similar to that of \citet{Peek:2010} and \citet{Jones:2011}, which use SDSS galaxy and M-dwarf spectra, respectively, to similar effect. The eighth data release of the SDSS has spectra for about 500,000 stars \citep{Aihara:2011}. The SEGUE Stellar Parameter Pipeline \citep[SSPP]{Lee:2008a} has uniformly processed these spectra to measure the temperature, metallicity, and gravity of each of these stars using a variety of methods, including ones that are independent of the observed colors of the star. We use one such method to predict the intrinsic colors of each star. The difference between the observed colors and intrinsic colors is used as a measurement of reddening to each star. These reddening measurements are then used to test the calibration of SFD and the reddening law. We test the calibration of SFD and the reddening law by comparing our reddening measurements with the predictions from SFD and a reddening law. A reddening law predicts, in each color $a-b$, the reddening $E(a-b)$ relative to the reddening in some reference color, usually $B-V$. The SFD dust map predicts $E(B-V)$, effectively giving the normalization of the reddening law in any direction on the sky. We use our reddening measurements to find $R_{a-b} = E(a-b)/\EBVSFD$ for each of the SDSS colors, testing both the reddening law and the SFD normalization. We also find the ratios of the $R_{a-b}$, which test the reddening law without using SFD as a reference. This work additionally lays the foundation for future investigations of the three-dimensional distribution of dust. Because we have stellar parameters for each of the stars, we can obtain accurate absolute magnitudes and distances. The SDSS targets both blue-horizontal-branch (BHB) stars and M-dwarf stars, in principle permitting the dust to be studied over a wide range of distances. In this work, however, we have focused on the two-dimensional distribution of the dust. The paper is organized as follows: in \textsection \ref{sec:data}, we present the data sets used in this work. In \textsection \ref{sec:methods}, we describe our method for transforming SSPP stellar parameters and SDSS photometry into reddening measurements. In \textsection \ref{sec:results}, the reddening measurements are presented, calibrated, and mapped. In \textsection \ref{sec:discussion} and \textsection \ref{sec:conclusion}, we discuss the implications of these results, especially for the reddening law, and conclude.
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\label{sec:conclusion} The SDSS provides a wealth of information for testing reddening. Using the SSPP stellar parameters, our spectrum-based technique achieves reddening measurements with empirical uncertainties of 56, 34, 25, and 29 mmag in the colors $u-g$, $g-r$, $r-i$, and $i-z$, comparable with the expected uncertainties from the photometry and stellar parameters. Spectra of a few blue stars on a sight line are sufficient to tightly constrain reddening provided that well-calibrated photometry is available and that enough spectra are available to calibrate the predicted colors to the measured colors. The use of individual stars with spectra additionally allows extinction to be studied as a function of distance. Stellar spectra permit good estimates of the intrinsic luminosity of a star in addition to its reddening, making three-dimensional studies of the dust feasible. The wide range of SEGUE target types---M and K dwarfs through BHB stars---will likely make such analyses especially fruitful, at least at low Galactic latitudes where all of the stars are not behind the entire dust column. At high Galactic latitudes, we could in principle use large numbers of stars to test for reddening through intermediate and high-velocity clouds. We defer to later work the attempt to extend this analysis into the third dimension, having established the feasibility and accuracy of the method in two dimensions. The combination of large photometric and spectroscopic surveys like Pan-STARRS \citep{Kaiser:2002} and LAMOST \citep{Su:1998} seems particularly promising. These tests should also provide useful feedback to models of synthetic spectra and to stellar parameter estimates. We look forward to incorporating the DR8 version of the SSPP into this analysis, and potentially pairing it with alternative synthetic spectral grids. However, we note that while this analysis provides an effective test of the colors predicted by synthetic spectral grids, nevertheless it does not depend on the accuracy of these colors. The wide range of spectral types targeted by the SDSS permits the predicted colors to be well calibrated to the observed colors, rendering the final results independent of the color accuracy of the synthetic spectra. The spectrum-based reddening measurements give best-fit $R_{a-b}$ that closely agree with the S10 values. Accordingly, we have gained confidence in the F99 reddening law with $R_V = 3.1$ and normalization $N = 0.78$ proposed in S10. The variation in the best-fit normalization of the reddening law seen in that analysis remains a problem. The spectrum-based reddening tests permit the variation in best-fit normalization with extinction to be seen more clearly; when $\EBVSFD \lesssim 0.2$, the best-fit normalization is about 15\% higher than when $E(B-V) \gtrsim 0.2$, though this conclusion relies on the relatively small fraction of the sky where $E(B-V) \gtrsim 0.3$ and SDSS data is available. Nevertheless, the agreement of the S10 and spectrum-based reddening measurements demands that the F99 reddening law be used to predict reddening over \citet{Cardelli:1989} or \citet{O'Donnell:1994} reddening laws. Appendix~\ref{app:ext} gives the extinction per unit $\EBVSFD$ predicted by these sets of measurements in 88 bandpasses for 4 values of $R_V$. We have focused on high Galactic latitudes where we observe $R_V=3.1$; other values of $R_V$ are provided only for convenience. Because we have seen that the best fit reddening law normalization varies over the sky and as a function of $\EBVSFD$, it is possible that outside the SDSS footprint a different normalization might be preferable. However, the shape of the reddening law seems constant over the SDSS footprint, and the normalization we suggest is unambiguously the best choice over the large area of sky covered by the SDSS. Therefore we propose that this reddening law and normalization become the default choice to be used in the absence of other information. Multiple sensitive, mutually consistent measurements of reddening over the SDSS footprint are now available. Extension of this work to larger areas of sky seems readily possible as Pan-STARRS and LAMOST data become available. These measurements will permit a next generation dust map to be constructed and tested. We defer to future work the construction of a new dust map that incorporates the insight gained from these measurements. David Schlegel suggested using SDSS stellar spectra to measure reddening several years ago, and his encouragement and advice were invaluable at the inception of this project. D.F. and E.S. acknowledge support of NASA grant NNX10AD69G for this research. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy. The SDSS-III web site is http://www.sdss3.org/. SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, University of Cambridge, University of Florida, the French Participation Group, the German Participation Group, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University. \ \appendix \label{app:ext}
| 10 | 12 |
1012.4804
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1012
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1012.1610_arXiv.txt
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We present the results of a ~6.1 square degree survey for clusters of galaxies via their Sunyaev-Zel'dovich (SZ) effect at 31~GHz. From late 2005 to mid 2007 the Sunyaev-Zel'dovich Array (SZA) observed four fields of roughly 1.5 square degrees each. One of the fields shows evidence for significant diffuse Galactic emission, and we therefore restrict our analysis to the remaining 4.4 square degrees. We estimate the cluster detectability for the survey using mock observations of simulations of clusters of galaxies; and determine that, at intermediate redshifts ($z\sim$0.8), the survey is 50\% complete to a limiting mass (${\rm M_{200 \overline{\rho}}}$) of ${\rm \sim 6.0\times 10^{14}} M_{\odot}$, with the mass limit decreasing at higher redshifts. We detect no clusters at a significance greater than 5 times the \rms\ noise level in the maps, and place an upper limit on \sigE, the amplitude of mass density fluctuations on a scale of 8$h^{-1}$ Mpc, of ${\rm 0.84 + 0.07}$ at 95\% confidence, where the uncertainty reflects calibration and systematic effects. This result is consistent with estimates from other cluster surveys and CMB anisotropy experiments.
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\label{sec:intro} The number density of massive clusters of galaxies depends strongly on cosmology, in particular through the normalization of the matter power spectrum, and through the dependence of the volume element on the geometry of the universe. As the low-energy photons in the CMB traverse the hot (${\rm \sim 10^8~K}$) gas of a massive cluster, about 1\% of the photons are inverse-Compton scattered. The result is a distortion in the CMB spectrum, the magnitude of which is proportional to the integrated pressure of the intra-cluster medium (ICM), i.e., the density of electrons along the line of the sight, weighted by the electron temperature (\cite{sunyaev72,sunyaev80}; see also \cite{birkinshaw1999}). The SZ flux of a cluster is therefore a measure of its total thermal energy. The change in the observed brightness of the CMB due to the SZ effect is given by \begin{equation} \label{eq:y} \frac{\Delta T_{\rm CMB}}{T_{\rm CMB}}= f(x)\int \sigma_{\rm T}n_e\frac{k_BT_e}{m_ec^2} dl \equiv f(x)y \end{equation} where $T_{\rm CMB}$ is the cosmic microwave background temperature, $\sigma_{\rm T}$ is the Thomson scattering cross section, $k_B$ is Boltzmann's constant, $c$ is the speed of light, $m_e$, $n_e$, $T_e$ are the electron mass, number density and temperature, respectively, and $f(x)$ contains the frequency dependence of the SZ effect (where ${\rm x \equiv \frac{h\nu}{k_B T_{CMB}}}$). Equation \ref{eq:y} defines the Compton $y$-parameter. The SZ effect appears as a temperature decrement at frequencies below $\approx 218~{\rm GHz}$, and as an increment at higher frequencies. As seen in Equation~\ref{eq:y}, the ratio of $\Delta T/T$ is independent of the distance to the cluster. This means that the SZ effect is redshift-independent in both brightness and frequency, offering enormous potential for finding high-redshift clusters. A cluster catalog resulting from an SZ survey of uniform sensitivity has a cluster mass threshold that is only weakly dependent on redshift for $z\gsim 0.7$ (via the angular diameter distance). As a result, SZ cluster surveys are approximately mass-limited and therefore potentially powerful probes of cosmology \cite[e.g.,][]{carlstrom2002}. Experiments such as the South Pole Telescope \citep{ruhl04} and the Atacama Cosmology Telescope \citep{fowler2004} are surveying hundreds of square degrees of sky searching for galaxy clusters through their SZ effect. A precursor survey to those being performed by SPT and ACT was performed with the Sunyaev-Zel'dovich Array (SZA), an 8-telescope interferometer designed specifically for detecting the SZ effect towards clusters of galaxies. Over the span of two years, the SZA surveyed a small region of sky at 31~GHz. This survey has been valuable in characterizing both compact and diffuse cm-wave CMB foregrounds. It has resulted in a measurement of the power spectrum of the CMB at small scales \citep{sharp2010}, a characterization of extragalactic compact source populations \citep{muchovej2010}, and further evidence for large-scale dust-correlated microwave emission (Leitch et al. 2010, in prep). In this paper, we present results of the survey as they pertain to cosmological parameter estimation. The paper is organized as follows: in \S 2 we describe the SZA observations, including a brief description of the instrument, data reduction and calibration, data quality tests, and foreground source extraction. We present the results of the survey in \S 3, followed by the calculation of the expected number of clusters in \S 4. In \S 5 we determine a constraint on $\sigma_8$ and systematic uncertainties associated with this analysis are presented in \S 6. We discuss our results in \S 7.
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From 2005 to 2007, the SZA performed a ~6.1 square degree survey for clusters of galaxies via their SZ effect at 31~GHz. In one of the fields there is evidence for large-scale dust-correlated emission; this field was excluded from the analysis presented here. Of the remaining 4.4 square degrees of the survey suitable for cosmological analysis, we estimate that the survey is $50\%$ complete to a mass of ${\rm M_{200 \overline{\rho}} \sim 6\times 10^{14}~M_{\odot}}$, averaged over redshift. By comparison with simulations, we place an upper limit on the value of \sigE\ of ${\rm 0.84~(+ 0.07)}$ at 95\% confidence, where the uncertainty reflects calibration and systematic uncertainties discussed in \S \ref{sec:errors}, excluding the error associated with the simulated cluster gas physics. Although this last uncertainty is potentially the dominant one, to properly quantify it requires a calculation of the completeness for a wide range of simulations with realistic gas models, which is beyond the scope of this paper. Our limit on \sigE\ is consistent with recent results from SZ surveys performed over larger areas of sky, such as with the South Pole Telescope \citep{vanderlinde2010}, and with determinations of \sigE\ from gravitational lensing and X-ray cluster surveys \citep{smith2003, allen2003}. In addition, it is consistent with determinations of \sigE\ from CMB anisotropy measurements, namely those of WMAP \citep{dunkley2009}, the South Pole Telescope \citep{lueker2010}, and the Atacama Cosmology Telescope \citep{act2010}. Our constraint is also in agreement with that of \cite{sharp2010}, based on CMB anisotropy measurements with the SZA itself. Although the data were collected with the same instrument, we note that both the data sets and the analyses of this paper and of \cite{sharp2010} are completely independent.
| 10 | 12 |
1012.1610
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1012
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1012.5802_arXiv.txt
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According to the hierarchical model, small galaxies form first and merge together to form bigger objects. In parallel, galaxies assemble their mass through accretion from cosmic filaments. Recently, the increased spatial resolution of the cosmological simulations have emphasised that a large fraction of cold gas can be accreted by galaxies. In order to compare the role of both phenomena and the corresponding star formation history, one has to detect the structures in the numerical simulations and to follow them in time, by building a merger tree.
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Recent simulations (\citet{kere_do_2005} for example) have emphasised the role of smooth cold accretion on galaxy formation. We aim at comparing the roles of mergers and gas accretion on galaxy growth by studying numerical simulations.
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The study of these simulations shows that baryonic mass assembly of galaxies seems to be dominated by smooth accretion, although we still have to perform further consistency tests. The next step is to perform statistical studies to confirm the preliminary results, then further physical exploitation can be made such as the role of the environment on the SFR.
| 10 | 12 |
1012.5802
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1012
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1012.5037.txt
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In order to precisely determine temperature and density of molecular gas in the Large Magellanic Cloud, we made observations of optically thin $^{13}$CO($J=3-2$) transition by using the ASTE 10m telescope toward 9 peaks \tma{of }{where} $^{12}$CO($J=3-2$) clumps \tma{}{were} previously detected with the same telescope. The molecular clumps include those in giant molecular cloud (GMC) Types I (with no signs of massive star formation), II (with HII regions only), and III (with HII regions and young star clusters). We detected $^{13}$CO($J=3-2$) emission toward all the peaks and found that their intensities are 3 -- 12 times lower than those of $^{12}$CO($J=3-2$). We \tma{derived }{determined} the intensity ratios of $^{12}$CO($J=3-2$) to $^{13}$CO($J=3-2$), $R^{12/13}_{3-2}$, and $^{13}$CO($J=3-2$) to $^{13}$CO($J=1-0$), $R^{13}_{3-2/1-0}$, at 45$\arcsec$ resolution\tma{,}{.} \tma{and these }{These} ratios \tma{are }{were} used for radiative transfer calculations in order to estimate temperature and density of the clumps. \tma{The clumps range from 15K to 200K in kinetic temperature, and from 8$\times 10^2$ to 7$\times 10^3$ cm$^{-3}$ in density.}{The parameters of these clumps range kinetic temperature $T\mathrm{_{kin}}$ = 15 -- 200 K, and molecular hydrogen gas density $n(\mathrm{H_2})$ = 8$\times 10^2$ -- 7$\times 10^3$ cm$^{-3}$.} We confirmed that the higher density clumps show higher kinetic temperature and that the lower density clumps lower kinetic temperature at a better accuracy than in the previous work. The kinetic temperature and density increase generally from a Type I GMC to a Type III GMC. We interpret that this difference reflects an evolutionary trend of star formation in molecular clumps. The $R^{13}_{3-2/1-0}$ and kinetic temperature of the clumps are well correlated with H$\alpha$ flux, suggesting that the heating of molecular gas \tma{whose density is }{$n(\mathrm{H_2})$ =} $10^3$ -- $10^4$ cm$^{-3}$ can be explained by stellar FUV photons.
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Star formation is of fundamental importance in understanding the evolution of galaxies. Stars are formed in dense clumps of giant molecular clouds (GMCs)\tma{}{,} and kinetic energy and heavy elements are ejected from stars back into the interstellar medium (ISM) through stellar winds and supernova explosions. This cycle enriches metal abundance in the ISM and drives the evolution of galaxies both physically and chemically. It is therefore important to better understand the evolution of GMCs, the principle sites of the star formation in galaxies. The Large Magellanic Cloud (LMC) is the most suitable galaxy to study star formation and natal GMCs because of its ideal location\tma{ to us}{}. The LMC offers a unique opportunity to achieve the highest \tma{resolutions }{resolution} due to its proximity, 50kpc \citep{Keller2006, Feast1999}, and \tma{its}{} nearly face-on position with an inclination angle of $\sim$ 35\degr \citep{vanderMarel2001}\tma{}{. This also} provides an advantage of less contamination along the line of sight. The LMC shows active on-going star formation and many massive young clusters are being formed \citep{Hodge1961}. In the LMC, the metallicity is a factor of $\sim$ 3 -- 4 lower \citep{Dufour1984, Rolleston2002} and the gas-to-dust ratio is a factor of $\sim$ 4 higher \citep{Koornneef1984} than \tma{those of }{values for} the solar neighborhood. The far-ultraviolet (FUV) radiation field is more intense in the LMC than in the Milky Way \citep{Israel1986} and the visual extinction is a few times lower than in the Milky Way. These \tma{characterize }{influence} the physical properties of GMCs, and may affect the initial condition of star and cluster formation. The first spatially resolved surveys of GMCs in the whole LMC were made by the NANTEN 4m telescope in the $^{12}$CO($J=1-0$) transition \citep{Fukui1999, Fukui2001, Fukui2008, MizunoN2001}. These surveys revealed the distribution of the GMCs within a single galaxy at a $\sim$ 40 pc resolution. \citet{Fukui2008} derived the physical properties, such as size, line width, and virial mass, of the GMCs, and found that the $^{12}$CO($J=1-0$) luminosity and virial mass of the clouds show a good correlation. Assuming \tma{}{that} the clouds are in virial equilibrium, \citet{Fukui2008} derived the $X_{\mathrm{co}}$ factor, conversion factor of the $^{12}$CO($J=1-0$) intensity to total molecular column density. The derived $X_{\mathrm{co}}$ factor was similar to that of other Local Group galaxies, such as Small Magellanic Cloud, M31, M33, IC10, and Milky Way, and the GMC mass distribution, dN/dM, was also similar to that of M31, M33, and IC10, suggesting that GMCs in the Local Group galaxies have similar properties \citep{Fukui2010, Blitz2007}. Comparisons between the GMCs and signs of star formation such as HII regions and young stellar clusters were used to classify GMCs into three types in terms of star formation activities \citep{Fukui1999, Kawamura2009}; Type I shows no signs of massive star formation, Type II associated with only small HII regions, and Type III associated with both HII regions and young stellar clusters, and these Types are interpreted as an evolutionary sequence \citep{Kawamura2009}. Comparative studies of \tma{}{the} $^{12}$CO($J=1-0$) to HI \tma{}{ratio} \tma{is }{are} also a key to understand the evolution of GMCs and the GMC formation via conversion of HI into H$_2$. \citet{Wong2009} made \tma{a}{} two dimensional, i.e., spatial, comparison\tma{}{s} of the second NANTEN $^{12}$CO($J=1-0$) survey \citep{Fukui2008} with \tma{ATCA+Parkes HI data sets }{HI data set combined ATCA and Parkes Telescope surveys }\citep{Kim2003}. They found that significant HI column \tma{density }{densities} ($>$ 10$^{21}$ cm$^{-2}$) and peak brightness \tma{temperature }{temperatures} ($>$ 20K) are necessary but not sufficient conditions for CO detection. \citet{Fukui2009} compared \tma{the same data sets in three dimensions }{the three dimensional data cubes}, including a velocity axis in addition to the two spatial axes. They found that GMCs are associated with HI envelopes on scales of $\sim$ 50 -- 100pc, and the HI envelopes may be in dynamical equilibrium or may be accreting onto GMCs to increase the mass via HI -- H$_2$ conversion. \tma{A subsequent }{Subsequently a} higher resolution (45$\arcsec$ $\sim$ 10 pc at 50kpc) survey of $^{12}$CO($J=1-0$) molecular cloud with the Mopra 22m telescope, the Magellanic Mopra Assessment (MAGMA), revealed \tma{}{a} more detailed distribution of the $^{12}$CO($J=1-0$) emission in the individual GMCs \citep{Hughes2010}. \citet{Hughes2010} presented that the physical properties of star-forming GMCs are very similar to the properties of GMCs without signs of massive star formation. The physical properties such as kinetic temperature and density of the molecular gas in the LMC have been investigated in the higher-$J$ transitions ($J=2-1$, $J=3-2$, $J=4-3$, $J=7-6$) of CO \tma{spectra}{} \citep[e.g.,][]{Sorai2001, Johansson1998, Heikkila1999, Israel2003, Bolatto2005, Kim2004, Kim2006, Nikolic2007, Pineda2008, MizunoY2010}. Some authors \citep[e.g.,][]{Bolatto2005, Nikolic2007} suggest two-component model especially for molecular clouds closely associated with HII regions. Most of these studies \tma{target }{targeted} \tma{outstanding }{extraordinary} HII regions \tma{and }{so} the \tma{number of}{} \tma{samples }{sample} is limited. \citet{Minamidani2008} have carried out $^{12}$CO($J=3-2$) observations of 6 GMCs, including one Type I, two Type II, and three Type III GMCs, in the LMC with the ASTE 10m telescope at a spatial resolution of 5 pc, and identified 32 molecular clumps. These data were combined with available $^{12}$CO($J=1-0$) and $^{13}$CO($J=3-2$) data and compared with LVG calculations for 13 clumps. The results show that these clumps range from cool ($\sim$ 10 -- 30K) to warm ($>$ 30 -- 200K) in kinetic temperature, and warm clumps range from less dense ($\sim$ 10$^3$ cm$^{-3}$) to dense ($\sim$ 10$^{3.5}$ -- 10$^5$ cm$^{-3}$) in density, whereas only lower limits in kinetic temperature were obtained in the warm clumps. Most recently, \citet{MizunoY2010} have made $^{12}$CO($J=4-3$) observations of the N159 region with the NANTEN2 4m sub-millimeter telescope. These data were used in LVG analysis combined with $^{12}$CO($J=1-0$), ($J=2-1$), ($J=3-2$), and ($J=7-6$) as well as the isotope transitions of $^{13}$CO($J=1-0$), ($J=2-1$), ($J=3-2$), and ($J=4-3$). The kinetic temperatures and densities were found to be $\sim$ 70 -- 80K and $\sim$ 3$\times 10^3$ cm$^{-3}$ in N159W and N159E, and $\sim$30K and $\sim$ 1.6$\times 10^3$ cm$^{-3}$ in N159S, indicating that an analysis including higher-$J$ transitions of both $^{12}$CO and $^{13}$CO can better constrain kinetic temperature and density. In the present study, we aim \tma{at determining }{to determine} temperature and density of \tma{molecular gas }{molecular hydrogen ,H$_2$, gas} in all three GMC types, Type I, II, and III, with higher accuracy by combining the $^{13}$CO($J=3-2$) data newly obtained by using the ASTE telescope and $^{12}$CO($J=3-2$) and $^{13}$CO($J=1-0$) data obtained with the ASTE and the SEST telescopes, respectively. \tma{These }{All our} data sets \tma{are all }{were} convolved to the same resolution of 45$\arcsec$, corresponding to $\sim$ 10pc at a distance of the LMC, 50kpc \citep{Feast1999}, and the large velocity gradient (LVG) calculations are performed to estimate the line intensities. \tma{This paper is organized as follows:}{}Section 2 describes observations of $^{13}$CO($J=3-2$) transition. Section 3 shows the observational results and describes ancillary data sets. Section 4 shows data analysis, and in Section 5, we discuss the physical properties of clumps and evolutionary sequence of GMCs. In Section 6, we present the summary.
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\subsection{Evolution\tma{}{s} of GMC and clump} The results of our LVG analysis show that the clump kinetic \tma{temperature }{temperatures, $T\mathrm{_{kin}}$,} \tma{distributes }{range} from cool ($\sim$ 15K) to warm ($\sim$ 200K) and the \tma{density }{densities, $n$(H$_2$),} from less dense ($\sim$ 8$\times 10^2$ cm$^{-3}$) to dense ($\sim$ 7$\times 10^3$ cm$^{-3}$). These large variations in the physical properties reflect the different characteristics of the GMCs according to their \tma{types }{Types} I, II, and III\tma{}{, where Type I shows no signs of massive star formation, Type II associated with only small HII regions, and Type III associated with both HII regions and young stellar clusters, and these Types are interpreted as an evolutionary sequence \citep{Kawamura2009}}. A clump in \tma{}{a} Type I GMC, GMC225 No. 1, has the lowest density (\tma{}{$n$(H$_2$)} $\sim$ 0.99$\times 10^2$ cm$^{-3}$ ) and lowest temperature (\tma{}{$T\mathrm{_{kin}}$} $\sim$ 25K). A clump in Type II GMC, N206D No. 1, is a dense (\tma{}{$n$(H$_2$)} $\sim$ 3.2$\times 10^3$ cm$^{-3}$) clump and has intermediate temperature (\tma{}{$T\mathrm{_{kin}}$} $\sim$ 42K). Clumps in Type III GMCs, except for N159 No. 4 (N159S) clump, are dense (\tma{}{$n$(H$_2$)} $\sim$ 4$\times 10^3$ cm$^{-3}$) and warm (\tma{}{$T\mathrm{_{kin}}$} $>$ 50K) clumps. The physical properties are generally correlated with the star formation activity of GMCs. The \tma{kinetic temperature }{$T\mathrm{_{kin}}$} and \tma{density }{$n$(H$_2$)} increase with the evolution of GMCs from Type I to III. This trend becomes clearer with the present determination of \tma{kinetic temperature }{$T\mathrm{_{kin}}$} and \tma{density }{$n$(H$_2$)} than that of the previous study \citep{Minamidani2008}, although the number of clump samples is smaller than that. The N159 No. 4 (N159S) clump is a part of \tma{}{a} Type III GMC, LMC N J0540-7008 \citep{Kawamura2009}. This GMC is quite large and elongated from north to south \citep[275pc $\times$ 53pc, P.A. = 87\degr;][]{Fukui2008}. Young clusters and large HII regions are associated with the northern part of this GMC, and in the southern part, where N159 No. 4 (N159S) is located, young clusters and large HII regions are not associated with. The N159 No. 4 (N159S) clump is less dense (\tma{}{$n$(H$_2$)} $\sim$ 1.5$\times 10^3$ cm$^{-3}$) and has intermediate temperature (\tma{}{$T\mathrm{_{kin}}$} $\sim$ 39K). These properties are quite similar to that of the clumps in Type I or II GMCs, as discussed by \citet{Minamidani2008} and \citet{MizunoY2010}. Comparison of the properties of molecular clumps with signs of star formation activities, such as young star clusters and H$\alpha$ emission, at high spatial resolution will establish the classification of molecular clumps, whose sizes are $\sim$ 7 pc \citep{Minamidani2008}. Because the size of clusters in the LMC are distributed in a range from 0.1 -- 10 pc \citep{Hunter2003}, which is smaller/similar size than the typical size of molecular clumps ($\sim$ 7 pc), it is important to compare molecular clumps with signs of star/clustar formation to understand the evolution of molecular clumps leading to the star/cluster formation. \subsection{Heating of the molecular gas in the LMC} As shown in the previous sections, the $R^{13}_{3-2/1-0}$ is well correlated with H$\alpha$ flux at a 10 pc scale, and the molecular gas kinetic temperatures \tma{}{($T\mathrm{_{kin}}$)} are also well correlated with H$_{\alpha}$ flux. These present findings suggest that the heating of molecular gas whose densities are 10$^3$ -- 10$^4$ cm$^{-3}$ may be dominated by far-ultraviolet (FUV) photons. The intense FUV field controls the physical and chemical processes in the ISM such as formation and destruction of molecules as well as ionization. These regions have been modeled as photo-dissociation regions (PDRs) or photon-dominated regions (PDRs) \citep[e.g.,][and references their in]{Tielens1985, Kaufman1999, Rollig2007}. The FUV flux ($G_0$) is estimated as 3500 in the 30 Doradus region \citep{Bolatto1999, Poglitsch1995, Werner1978, Israel1979} and 300 for the N159 region \citep{Bolatto1999, Israel1996, Israel1979}\tma{,}{.} \tma{and the }{The} gas density is estimated \tma{as }{to be} (1 -- 5)$\times 10^3$ cm$^{-3}$ for these regions\tma{ in the present is work}{}. The PDR surface temperature \tma{can be read as }{is} $\sim$ 400 K for the \tma{30Doradus }{30 Doradus} region and $\sim$ 200 K for the N159 region from Figure 1 of \citet{Kaufman1999}. The PDR gas temperature is relatively constant from the cloud surface to a depth where either heating or cooling changes significantly. The heating is generally dominated by the grain photoelectric heating \citep{Kaufman1999}, and then dust attenuation of FUV flux controls the thermal structure, where the typical size scale is 0.1 -- 2 pc. These suggest that the effect of FUV heating of molecular gas seems to be local and direct phenomena. The warm region can become larger under low-metallicity or high gas-to-dust ratio environments. These temperatures are basically consistent with the kinetic temperatures \tma{}{($T\mathrm{_{kin}}$)} of the warm clumps in the present sample if beam dilution by the present resolution, 10 pc, is taken into account. \subsection{Substructures inferred by the $^{13}$CO($J=3-2$) observations} The $^{13}$CO($J=3-2$) intensities are 3 -- 12 times lower than those of $^{12}$CO($J=3-2$) transition, and the line intensity ratios of $^{12}$CO($J=3-2$) to $^{13}$CO($J=3-2$), $R^{12/13}_{3-2}$\tma{,}{.} \tma{}{These} vary not only from clump to clump but also inside of each clump. In the some clumps, clear two velocity components are detected in the $^{13}$CO($J=3-2$) transition. These suggest some internal structures. \citet{MizunoY2010} showed that the molecular distribution of the N159 No. 1 (N159W) clump in the $^{12}$CO($J=3-2$) transition is similar to that of the $\eta$ Carinae northern cloud in the $^{12}$CO($J=1-0$) emission line smoothed to 5 pc resolution. The original data of $\eta$ Carinae northern cloud has a 2 pc resolution and several substructures are identified \citep{Yonekura2005}. This supports the existence of internal structures inside of $\sim$ 5 pc scale molecular clumps, and indicates that these internal structures can be resolved in the CO transitions with high spatial resolution observations. \tma{The subsequent }{Subsequent} observations of $\eta$ Carinae northern cloud using a high density tracer such as H$^{13}$CO$^{+}$($J=1-0$) at a high spatial resolution, \tma{}{resulted in the detection of} high density molecular cores whose sizes are less than 1 pc \tma{were detected}{}\citep{Yonekura2005}. In the LMC, some \tma{pioneering }{initial} interferometric observations of high density tracers, such as HCO$^{+}$ and NH$_3$, were made \tma{by }{with} ATCA, \tma{whose }{using} resolutions \tma{were }{of} $\sim$ 6$\arcsec$ -- 19$\arcsec$ corresponding to $\sim$ 1.4 -- 5 pc at 50 kpc \citep[e.g.,][]{Wong2006, Ott2008, Ott2010}. \citet{Ott2008} detected two peaks in the HCO$^{+}$($J=1-0$) transition in the N159W region. Further systematic detailed observations using high density tracers at higher resolution, with ALMA for instance, should be important for probing the initial conditions of star/cluster formation.
| 10 | 12 |
1012.5037
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1012
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1012.2425_arXiv.txt
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We study the X-ray and optical properties of 16 Broad Absorption Line (BAL) quasars detected in a $\approx$3 ~deg$^2$ region common to the wide synoptic (W-1) component of the Canada$-$France$-$Hawaii Telescope Legacy Survey (CFHTLS) and the XMM Large Scale Structure survey (XMM-LSS). The BAL fraction is found to be 10\% in full sample, 7\% for the optical colour selected QSOs and as high as 33\% if we consider QSOs selected from their IR colours. The X-ray detected non-BAL and BAL quasars have a mean observed X-ray-to-optical spectral slope ($\alpha_{\rm ox}$) of $-$1.47 $\pm$ 0.13 and $-$1.66 $\pm$ 0.17 respectively. We also find that the BAL QSOs have $\alpha_{\rm ox}$ systematically smaller than what is expected from the relationship between optical luminosity and $\alpha_{\rm ox}$ as derived from our sample. Based on this, we show, as already reported in the literature for quasars with high optical luminosities, our new sample of BAL QSOs have X-ray luminosity a factor of three smaller than what has been found for non-BAL QSOs with similar optical luminosities. Comparison of hardness ratio of the BAL and non-BAL QSOs suggests a possible soft X-ray weakness of BAL QSOs. Combining our sample, of relatively fainter QSOs, with others from the literature we show that larger balnicity index (BI) and maximum velocity ($V_{\rm max}$) of the C~{\sc iv} absorption are correlated with steeper X-ray to optical spectral index. We argue that this is most likely a consequence of the existence of a lower envelope in the distribution of BI (or $V_{\rm max}$) values versus optical luminosity. Our results thus show that the previously known X-ray weakness of BAL QSOs extends to lower optical luminosities as well.
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Broad Absorption Line (BAL) quasars are Active Galactic Nuclei (AGN) characterized by the presence of strong absorption troughs in their UV spectra. They constitute an observed fraction of about 10$-$15\% of optically selected quasars (Reichard et al. 2003; Hewett \& Foltz 2003). Recently, it has been shown that the actual BAL fraction could be higher as optical colour selection of QSOs may be biased against the BAL QSOs (see for example, Dai et al. 2008; Shankar et al. 2008; Urrutia et al. 2009; Allen et al. 2010) The BALs are attributed to material flowing outwards from the nucleus with velocities of 5000 to 50000 km/s (Green et al. 2001). These quasars are classified into three subclasses based on the material producing the BAL troughs. High ionization BAL quasars (HiBALs) have broad absorption from C~{\sc iv}, Si~{\sc iv}, N~{\sc v} and O~{\sc vi}. About 10\% of BAL quasars also show, apart from their HiBAL features, broad absorption lines of lower ionization species such as Mg~{\sc ii} or Al~{\sc iii} and are called low-ionization BAL quasars (LoBALs). Finally, LoBALs with absorptions from excited states of Fe~{\sc ii} or Fe~{\sc iii} are called FeLoBALs (Wampler et al. 1995). BAL quasars in general have higher optical-UV polarization than non-BAL quasars, and the LoBALs tend to have particularly high polarization than HiBALs (Hutsemekers et al. 1998; Schmidt \& Hines 1999; DiPompeo et al. 2010). LoBALs have a more reddened optical continuum compared to non-BAL and HiBAL QSOs(Becker et al. 2000; Sprayberry \& Foltz 1992), thereby suggesting the presence of larger amounts of dust in them. In X-rays too, LoBALs have higher absorbing column densities than HiBALs (Green et al. 2001; Gallagher et al 2002). The dichotomy between BAL and non-BAL quasars is often thought to be a consequence of orientation. The similarity between the optical/UV emission lines and continuum properties of BAL and non-BAL quasars (Weymann et al. 1991; Reichard et al. 2003) supports a scenario where the observed fraction of BAL quasars corresponds to the covering fraction of a wind that could be present in all AGN. Earlier spectropolarimetric observations too support this orientation scheme (Goodrich \& Miller 1995; Hines \& Wills 1995). However, recent spectropolarimetric observations of radio-loud BALs do not favour the orientation dependent scheme for the BAL phenomenon (DiPompeo et al. 2010). On the other hand, Allen et al. (2010) found a strong redshift dependence of the C~{\sc iv} BAL quasar fraction. They conclude that the BAL phenomenon cannot be due to an orientation effect only. Alternative to the orientation scheme it is argued that the observed BAL quasar fraction could correspond to the intrinsic fraction of quasars hosting massive nuclear winds, thereby tracing an evolutionary phase of the AGN lasting $\approx$10$-$15\% of their lives (Hazard et al. 1984; Becker et al. 2000; Giustini et al. 2008). In this scenario, the large amounts of gas and dust surrounding the central source should lead to enhanced far-infrared and sub-millimeter emission in BAL quasars with respect to non-BAL quasars (Giustini et al. 2008). However, sub-millimeter studies found no (Willott et al. 2003) or little (Priddey et al. 2007) differences among the two populations. Also, the mid-IR properties of BAL and non-BAL quasars of comparable luminosities are indistinguishable (Gallagher et al. 2007). BAL quasars are predominantly radio-quiet while few radio-loud BAL quasars are also known (Becker et al. 2000; Brotherton et al. 2005). Although radiative acceleration could be the main driving mechanism in BAL quasars (Arav et al. 1994; Srianand et al. 2002), we still do not have a clear picture of the physics of outflows/winds in BALs. X-ray observations can help constrain the physical mechanisms at play in BAL quasar outflows and the different scenarios proposed (Giustini et al. 2008). Since the ROSAT survey, BAL quasars which are radio-quiet, have been known to have faint soft X-ray to optical luminosity ratio (Green et al. 1995; Green \& Mathur 1996). Their X-ray luminosity is typically 10$-$30 times lower than expected from their UV luminosity, qualifying them as soft X-ray weak objects (Laor et al. 1997). This implies that the soft X-ray continuum of BAL quasars is either (a) strongly absorbed by highly ionized material or (b) intrinsically under luminous. Given the extreme absorption evident in the ultraviolet, this soft X-ray faintness was assumed to result from intrinsic absorption in BAL material of high column density, typically $N_{\rm H} > 10^{22}$~cm$^{-2}$ (Gallagher et al. 2006, hereafter G06; Green et al. 1995; Green et al. 2001; Brotherton et al. 2005; Fan et al. 2009; Gibson et al. 2009). However, it has also been argued that intrinsic X-ray faintness cannot be ruled out as the cause for their observed X-ray weakness (Sabra \& Hamann 2001; Mathur et al. 2000; Gupta et al. 2003, Giustini et al. 2008; Wang et al. 2008; Ghosh \& Punsly 2008). Radio-loud BAL quasars are also found to be X-ray weak when compared with radio-loud non-BAL quasars of similar UV/optical luminosities (Miller et al. 2009). X-ray spectral analysis of BAL quasars considering neutral and ionized absorbers are found to yield low neutral hydrogen ($N_{\rm H} < 10^{21}$~cm$^{-2}$) and high ionized hydrogen ($N_{\rm H}^{i} > 10^{21}$~cm$^{-2}$) column density respectively (Giustini et al. 2008; Streblyanska et al. 2010). It thus seems that the inferred $N_{\rm H}$ values depends on (a) the ionization state of the gas in our line of sight to the BAL quasar and (b) the absorber either fully or partially covering the X-ray source. Here, we investigate the issue of the X-ray weakness of BAL quasars using a new sample of quasars selected in the Canada$-$France$-$Hawaii Telescope Legacy Survey (CFHTLS\footnote{http://www.cfht.hawaai.edu/Science/CFHTLS/}) and overlapping the XMM Large Scale Structure survey (XMM-LSS) and the Spitzer Wide-area InfraRed Extragalactic (SWIRE) Survey with the aims of (i) identifying a homogeneous sample of BAL quasars from a parent quasar sample and (ii) of studying the X-ray nature of those identified BAL quasars. The sample has been selected without a prior knowledge of the BAL nature of the objects. This paper is organized as follows. The data set used and the observations are described in Sect. 2. Identification of BAL quasars is given in Sect. 3. Results of the analysis are presented in Sect. 4 and the conclusions are drawn in Sect. 5. Throughout this paper we adopt a cosmology with $H_{\rm o}$~=~70~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_{\rm m}$~=~0.27 and $\Omega_{\Lambda}$~=~0.73. \begin{figure} \hspace*{-0.5cm}\psfig{file=field1.eps,width=9cm,height=9cm} \caption{Layout of the field used to select quasars in this study. The XMM-LSS and SWIRE regions are delineated. W06, W07, W09 and W10 are the four pointings of CFHTLS and the large circle is the $\sim$3 square degrees region searched for quasars in this work. The quasars are marked as filled circles, the BAL quasars detected in X-ray are shown as crosses and the BAL quasars un-detected in X-rays are shown as open circles. } \end{figure}
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We have presented a new sample of 16 BAL quasars selected from a homogeneous sample of 159 $z_{\rm em} > 1.5$, $g^{\prime} < 22$ mag quasars found in a region of CFHTLS overlapping with the XMM-LSS and SWIRE surveys. \begin{enumerate} \item We find a BAL quasar fraction of $\sim$10\% in the whole sample and of $\sim$8\% among optically selected quasars (120 quasars). This is similar to what is found from other optically selected quasar samples (Weymann et al. 1991; Hewett \& Foltz 2003; Trump et al. 2006; Knigge et al. 2008; Gibson et al. 2009). Also, 7\% of the XMM selected quasars are found to be BAL quasars. If we consider the 12 quasars which were selected from SWIRE colours only (and rejected by the optical colour selection), 4 are found to be BAL quasars which gives a BAL fraction of $\sim$33\%. This agrees with the large 23\% BAL quasar fraction found in 2MASS selected quasars by Dai et al. (2008). We note here that recently Allen et al. (2010), after correcting for selection biases, report an intrinsic C~{\sc iv} BAL quasar fraction of $\sim$41$\pm$5 per cent, though their observed fraction is 8.0$\pm$0.1 per cent. \item The values of $\Delta \alpha_{\rm ox}$, the deviation of the spectral index, $\alpha_{\rm ox}$, from the mean in the overall sample at the same UV luminosity, are distributed symmetrically around zero, whereas the $\Delta \alpha_{\rm ox}$ values of XMM detected BAL quasars are shifted towards lower values with an average of $-$0.169$\pm$0.161. This shows that the X-ray detected BAL quasars are weaker than the X-ray detected non-BAL quasars by a factor of about 3. This X-ray weakness of BAL quasars compared to non-BAL quasars is similar to what was found in earlier studies (G06; Fan et al. 2009; Gibson et al. 2009) however, contrasts with the results of Giustini et al. (2008). The discrepancy we find is however much less than what was reported by G06 who found that optically bright BAL quasars, are X-ray sources weaker by a factor of 30 compared to non-BAL quasars. This might be due to the fact that G06 had used {\it Chandra} observations reaching significantly fainter flux levels than the XMM-Newton observations used here for our sample of BALs. \item We investigated various correlations between the properties of the C~{\sc iv} absorptions and $\alpha_{ox}$ using an extended sample gathered after combining our data with those of Gallagher et al. (2006), Giustini et al. (2008), Gibson et al. (2009) and Fan et al. (2009). For this large HiBAL quasar sample, $\Delta \alpha_{\rm ox}$ was calculated and we find it to be correlated with the balnicity index, BI, and the maximum velocity of the outflow, $V_{\rm max}$. Similarly, $V_{\rm max}$ and BI are correlated with the 2500~\AA~ monochromatic luminosity, $L_{2500\AA}$. This suggests that quasars with high velocity outflows are X-ray weak. While there is a large range in $V_{\rm max}$, for a given $L_{2500\AA}$ at log($L_{2500\AA}$)~$<$~31, there is a lack of objects with $V_{\rm max}$~$<$~10$^4$~km/s at the high luminosity end. Thus there seems to be a lower envelope, the presence of which dominates the above correlation. This probably means that even if radiation pressure is the main driver of the flow, other parameters are important such as the launching radius, the shape of the ionizing spectrum, the mass of the wind, the properties of the confining medium, etc. \item We find the mean X-ray hardness ratio, HR, of XMM detected non-BAL and BAL quasars to be $-$0.61$\pm$0.20 and $-$0.49$\pm$0.20 respectively. While this is consistent with the assumption that the X-ray weakness is due to soft X-ray absorption, the number of X-ray detected BAL quasars is too small to make any statistically significant claim. \item The fit of the X-ray spectra from two BAL quasars shows that they have neutral hydrogen column densities smaller than $10^{22}$~cm$^{-2}$ and close to the Galactic values. This is based on a fully covering neutral absorber model. However, if the absorber is ionized and/or partially covering the X-ray source, the column density derived for these two BAL quasars should be considered lower limits. For these two BAL quasars, the photon spectral index is found to be similar to that of radio-quiet quasars. This agrees with the photon index values found recently from X-ray spectral analysis of a larger sample of BAL quasars (Giustini et al. 2008; Streblyanska et al. 2010). \end{enumerate}
| 10 | 12 |
1012.2425
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