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1609 | 1609.04210_arXiv.txt | We present new spectroscopic observations obtained with GMOS@Gemini-S of a sample of 25 \hii\ regions located in NGC~55, a late-type galaxy in the nearby Sculptor group. We derive physical conditions and chemical composition through the \te-method for 18 \hii\ regions, and strong-line abundances for 22 \hii\ regions. We provide abundances of He, O, N, Ne, S, Ar, finding a substantially homogenous composition in the ionised gas of the disc of NGC~55, with no trace of radial gradients. The oxygen abundances, both derived with \te- and strong-line methods, have similar mean values and similarly small dispersions: 12+$\log$(O/H)=8.13$\pm$0.18~dex with the former and 12+$\log$(O/H)=8.17$\pm$0.13~dex with the latter. The average metallicities and the flat gradients agree with previous studies of smaller samples of \hii\ regions and there is a qualitative agreement with the blue supergiant radial gradient as well. We investigate the origin of such flat gradients comparing NGC~55 with NGC~300, its companion galaxy, which is also twin of NGC~55 in terms of mass and luminosity. We suggest that the differences in the metal distributions in the two galaxies might be related to the differences in their K-band surface density profile. The flatter profile of NGC~55 probably causes in this galaxy higher infall/outflow rates than in similar galaxies. This likely provokes a strong mixing of gas and a re-distribution of metals. | NGC~55 is the nearest edge-on galaxy at a distance of 2.34 Mpc \citep{k16} and it is member of the nearby Sculptor group consisting of approximately 30 galaxies \citep{cote97, jerjen00} and being dominated by the spirals NGC~300 and NGC~253. The main properties of NGC~55 are listed in Table~\ref{tabNGC55}. The nature of the NGC~55 galaxy has been debated for long time: its high inclination \citep[79$^{\circ}$;][]{puche91} has allowed different interpretation of its morphology. It has been sometimes defined as a late-type spiral galaxy \citep{ST87}, while in other works it has been considered as a dwarf irregular galaxy, similar to the Large Magellanic Cloud (LMC), as, e.g., \citet{dV61}. Following \citet{dV61} the main light concentration at visible wavelengths, which is offset from the geometric center of the galaxy, is a bar seen end-on. The structure of the disc of NGC~55 shows asymmetric extra-planar morphology \citep{ferguson96}. The extra-planar asymptotic giant branch (AGB) population is essentially old with ages of about 10~Gyr \citep{Davidge05}. \citet{tanaka11} studied the structure and stellar populations of the Northern outer part of the stellar halo: from the stellar density maps they detected an asymmetrically disturbed, thick disc structure and possible remnants of merger events. Its interstellar medium (ISM) has been studied in several aspects: the neutral component \citep[e.g.,][]{hummel86, puche91, westmeier13}, the molecular component \citep[e.g.,][]{DH89,HD90}, and the ionised component \citep[e.g.,][]{WS83, hoopes96, ferguson96,tullmann03}. The star formation activity is located throughout the disc planar region of NGC~55, but there are also large quantities of gas off of the disc plane still forming stars. \citet{OD99} found shell structures and chimneys outside the planar regions that are signatures of supernovae explosions and stellar winds. The composition of its exceptionally active population of \hii\ regions was studied in the past by \citet{WS83}, and more recently by \citet{tullmann03} and \citet{TR04} who studied some regions, inside and outside the disc, respectively. The radial metallicity gradient of NGC~55 was first outlined by the study of \hii\ regions of \citet{WS83} who found a substantially flat metallicity radial profile. A similar result was obtained by \citet{Pilyugin14} in their re-analysis of the radial metallicity profiles of several late-type spiral galaxies including NGC~55. \citet{tikhonov05} analysed the spatial distribution of the AGB and red giant branch (RGB) stars along the galactocentric radius of NGC~55, revealing again a very small metallicity gradient also for the older stellar populations. The absence of metallicity gradients might suggest a coherent formation of all the disc or very efficient mixing processes since its formation. From a sample of 12 B-type supergiant stars, \citet{castro08} found a mean metallicity of $-$0.40~dex, a value quite close to the LMC metallicity \citep[see, e.g.,][]{hunter07} and a flat gradient. The metallicities of the two extra planar \hii\ regions were found to be slightly lower than those of the central \hii\ regions, suggesting that they might have formed from material that did not originate in the thin disc \citep{tullmann03}. From a dynamical point of view, there are a number of works that confirm tidal interactions among the three pairs (NGC~55 and 300, 247 and 253, and 45 and 7793) of major galaxies in Sculptor Group \citep[see, e.g.,][]{dV68, whiting99, westmeier13}. On top of the dynamical effects of the nearby galaxies on NGC~55, the kinematics of its central regions within the bar shows a gradient in radial velocities towards the galactic centre, which is due to flow of material along the bar \citep{CA88,westmeier13}. \begin{table*} \begin{minipage}{75mm} \caption{Properties of NGC~55} \label{tab:properties_NGC55} \begin{tabular}{lll} % \hline Parameter & Value & Reference \\ \hline Type & SB(s)m & \citet{dV61} \\ Centre &00h14m53s.6 -39$^{\circ}$11'47" & NED (J2000.0) \\ Distance & 2.34$\pm$0.2 Mpc& \citet{k16} \\%TOBEUPDATED!!!!LAURA Total mass & 2.0$\pm$10$^{10}$M$_{\odot}$ &\citet{westmeier13} \\ % Inclination & 79$\pm$4$^{\circ}$ & \citet{puche91} \\ Position angle & 109$^{\circ}$ & HyperLeda\\ \hline \end{tabular} \label{tabNGC55} \end{minipage} \end{table*} The present work is part of our study of the structure and evolution of Local Group and nearby galaxies through the spectroscopy of their emission-line populations \citep[see, e.g., ][]{magrini05,goncalves07, MG09, magrini09, goncalves12, goncalves14, stanghellini15}. In this framework, we have carried on a deep spectroscopic campaign with the multi-object spectrograph GMOS@Gemini-South telescope of the strong-line emitters of NGC~55, as illustrated in Figure~\ref{fig_allsources}. The aim of the present work is to explore the distribution of abundances in this galaxy, studying \hii\ regions located in the disc --from the inner disc to the outskirts-- as well as extra planar regions. The paper is structured as follows: in Section~\ref{sec_obs} we describe the observations --imaging and spectroscopy-- and the data reduction process. In Section~\ref{sec_ana} we present the spectroscopic analysis, whereas in Section~\ref{sec_grad} we describe the spatial distributions of the abundances. In Section~\ref{sec_dis} we discuss our results and compare them with those of its companion galaxy NGC~300. In Section~\ref{sec_conclu} we give our conclusions. \section[]{GMOS@Gemini-S: imaging and spectroscopy} \label{sec_obs} \begin{figure*} \centering \includegraphics[width=16.5truecm]{image.eps} \caption{ {\it Top:} HSDSS IIIaJ4680\AA\ image of NGC~55. The entire image is 30$\times$15~arcmin$^2$, and it is centred at RA=00:15:12.27 and DEC=-39:12:57.89. Superposed to it the three GMOS-S FoV we observed (masks M1, M2 and M3), of 5.5$\times$5.5~arcmin$^2$ each, are highlighted. The green symbols (plus, diamond and circle) are our \hii\ regions located within these three FoV, respectively. Other symbols are as follows: cyan plus, the X-rays sources up to D$_{25}$ from \citet{stobbart06}; and the two blue crosses indicate the extra-planar \hii\ regions studied by \citet{tullmann03}. {\it Bottom:} Our GMOS-S continuum subtracted images (H$\alpha$-H$\alpha$C) of the three FoV. The \hii\ regions we discuss in this paper are identified, following the IDs of Table~\ref{tab_objid}. The orientation, North to the top and East to the left is the same in all the panels. } \label{fig_allsources} \end{figure*} \begin{table} \centering \begin{minipage}{75mm} {\scriptsize \caption{GMOS-S mask identification, classification and coordinates of the \ha\ line-emitters selected from the GMOS-S pre-imaging. M1, M2 and M3 stands for the masks ID. The object classification shown is based on the follow-up spectroscopic analysis of the present study.} \begin{tabular}{@{}llllll@{}} \hline MaskId & Field-ID & Class & RA & Dec \\ & & & J2000.0 & J2000.0 \\ \hline M1S1 & H-1 & HIIr & 00:16:14.93 & -39:15:51.62 \\ M1S4 &H-2 & HIIr & 00:16:10.12 & -39:16:11.49 \\ M1S5 &NGC~55 StSy-1 & SySt & 00:16:07.25 & -39:16:31.91 \\ M1S6 &H-3 & HIIr & 00:16:07.99 & -39:15:45.94 \\ M1S7 &H-4 & HIIr & 00:16:05.73 & -39:16:42.96 \\ M1S8 &H-5 & HIIr & 00:16:02.59 & -39:15:52.35 \\ M1S10 &H-6 & HIIr & 00:16:03.96 & -39:16:26.61 \\ M1S12 &H-7 & HIIr & 00:16:01.04 & -39:15:42.52 \\ M1S13 &H-8 & HIIr & 00:15:56.26 & -39:16:25.79 \\ M1S15 &H-9 & HIIr & 00:15:53.42 & -39:15:41.12 \\ M1S16 &H-10 & HIIr & 00:15:59.58 & -39:16:37.63 \\ M1S17 &H-11 & HIIr & 00:15:54.64 & -39:16:25.90 \\ M1S18 &H-12 & HIIr & 00:15:49.66 & -39:16:23.08 \\ \\ M2S1 & H-13 & HIIr & 00:15:36.94 & -39:14:21.08 \\ M2S2 & NGC~55 SySt-2 & SySt & 00:15:39.02 & -39:14:40.60 \\ M2S3 & H-14& HIIr & 00:15:29.36 & -39:15:18.69 \\ M2S5 & H-15 & HIIr & 00:15:32.49 & -39:14:50.50 \\ M2S6 & H-16& HIIr & 00:15:30.75 & -39:14:32.54 \\ M2S14 & H-17 & HIIr & 00:15:24.71 & -39:13:51.88 \\ M2S15 & H-18 & HIIr & 00:15:29.41 & -39:12:25.09 \\ \\ M3S1 & H-19& HIIr & 00:14:46.02 & -39:11:01.79 \\ M3S2 & H-20& HIIr & 00:14:47.32 & -39:11:32.85 \\ M3S3 & H-21& HIIr & 00:14:49.52 & -39:10:59.99 \\ M3S4 & H-22& HIIr & 00:14:52.18 & -39:11:26.70 \\ M3S5 & H-23& HIIr & 00:14:53.18 & -39:11:53.30 \\ M3S7 & H-24& HIIr & 00:14:56.52 & -39:11:58.05 \\ M3S8 & NGC~55 StSy-3 & SySt & 00:14:58.61 & -39:11:59.14 \\ M3S10 & H-25& HIIr & 00:14:59.94 & -39:12:14.97 \\ \hline \end{tabular} } \end{minipage} \label{tab_objid} \end{table} The data analysed in the present paper were obtained with the Gemini Multi-Object Spectrographs (GMOS) at Gemini South telescope in 2012 and 2013. The two programs through which the data were taken are GS-2012B-Q-10 and GS-2013B-Q-12, with D. R. Gon\c calves as Principal Investigator (PI) in both cases. In total we observed three fields of view of GMOS-S, each of 5.5\arcmin$\times$5.5\arcmin. In the following, we refer to the three fields as M1, M2 and M3 (see Figure~ \ref{fig_allsources}). \subsection*{Pre-Imaging} We obtained the pre-imaging of NGC~55 with the GMOS-S camera in queue mode on the 28$^{th}$ (M1) and 27$^{th}$ (M2, M3) of August 2012. We used the on- and off-band imaging technique to identify the strongest H$\alpha$ line emitters. Their location is shown in Figure~\ref{fig_allsources}. For the three fields of view (FoV) the on-band H$\alpha$ images were sub-divided in 3 exposures of 60~s each, while the three off-band (the continuum of H$\alpha$) H$\alpha$C sub-exposures were of 120~s each. These narrow-band filters have central $\lambda$ ($\lambda$ interval) of 656nm (654-661nm) and 662nm (659-665nm) for H$\alpha$ and H$\alpha$C, respectively. The location of the three fields (5.5\arcmin$\times$5.5~\arcmin) is shown in Figure~\ref{fig_allsources}, and the central coordinates of each field are: for M1 R.A. 00:16:02.62 and Dec. -39:14:41.77; for M2 R.A. 00:15:26.59 and Dec. -39:12:49.08; and for M3 R.A. 00:14:59.06 and Dec. -39:11:30.97. During the pre-imaging observations, the seeing varied from 0.9'' to 1\farcs0. In Table~\ref{tab_objid} we give the coordinates of the observed emission-line objects (25 \hii\ regions, HIIr, and 3 candidate symbiotic systems, SySt). We based the classification of the nebulae on the analysis of their spectra, which will be introduced in the following sections. We discriminate SySts from \hii\ regions on the bases of the presence of absorption features and continuum of late-type M giants, of the strong nebular emission lines of Balmer \hi, \heii, the simultaneous presence of forbidden lines of low- and high-ionization, like \oii, \neiii, \oiii\, \citep{bel00} and of the Raman scattered line at $\lambda$6825\AA, a signature almost exclusively seen in symbiotic stars \citep{schmid89, bel00}. In Figure~\ref{fig_allsources} the positions of our \hii\ regions are shown in contrast with objects investigated in previous works: X-rays sources up to D$_{25}$\footnote{D$_{25}$ is the diameter that corresponds to a surface brightness of 25 mag/arcsec$^2$} from \citet{stobbart06}; and the extra-planar \hii\ regions from \citet{tullmann03}. \subsection*{Spectroscopy} The spectroscopic observations were obtained in queue mode with two gratings, R400+G5305 (red) and B600 (blue). For the mask M1, the spectra were taken on the 22$^{nd}$ (blue) and 23$^{th}$ (red) of December 2012. The spectroscopy of M2 and M3, on the other hand, were obtained about one year later. The red spectra of M2 on the 29$^{th}$ of September, while the blue counterpart was observed on the 13$^{th}$ of October, in 2013. As for M3, the red and blue spectra were taken on the nights of 27$^{th}$ and 29$^{th}$ of December of 2013. We obtained three exposures per mask and grating --with 3$\times$1,430s for M1 and 3$\times$1,390s for M2 and M3 both in blue and red. For technical reasons, only one exposure of the blue spectra of M3 and two of the red spectra of M2 were useful for science. In all the cases, we combined the good quality exposures to increase the signal-to-noise ratio (SNR) of the spectra and to remove cosmic rays. For most of the spectra the effective blue plus red spectral coverage range from 3500~\AA\ to 9500~\AA, in several cases allowing a significant overlap of the spectra. Only few lines were measured with wavelength longer than ~7200~\AA, because of the poor (wavelength + flux) calibration at the red end of the spectra (see Figure~\ref{fig_spectra}). We avoided the possibility of important emission-lines to fall in the gap between the three CCDs, by slightly varying the central wavelength of the disperser from one exposure to another. To do this, we centred the red grating R400+G5305 at 750 $\pm 10$~nm and the blue one B600 at 460 $\pm 10$~nm. The masks were built with slit widths of 1\arcsec\ and with varying lengths to include portions of sky in each slit for a proper local sky-subtraction. The spectroscopic observations were spatially$\times$spectrally binned. The final spatial scale and reciprocal dispersions of the spectra were: 0\farcs144 and 0.09~nm per pixel, in blue; and 0\farcs144 and 0.134~nm per pixel, in red. Following the usual procedure with GMOS for wavelength calibration, we obtained CuAr lamp exposures with both grating configurations, either the day before and after the science exposures. The spectrophotometric standard LTT7379 \citep{hamuy92} was observed with the same instrumental setups as for science exposures, on September 2$^{nd}$ 2013, and used for flux calibration of the three masks, since no standards were obtained near the observation of M1, in 2012. In Figure~\ref{fig_spectra} we show a sample with fully reduced and calibrated GMOS spectra, one spectrum per field, on which it is straightforward to see the quality of our data. Data were reduced and calibrated in the standard way by using the Gemini {\sc gmos data reduction script} and {\sc long-slit} tasks, both being part of {\sc IRAF}\footnote{IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under cooperative agreement with the National Science Foundation.} \section[]{Spectral analysis and determination of the physical and chemical properties} \label{sec_ana} \begin{figure} \centering \includegraphics[width=8.5truecm]{spectra_BR-Gimp.eps} \caption{Sample of our GMOS spectra, one per FoV, M1, M2 and M3, for the \hii\ regions H-11, H-17 and H-23, respectively. The CCD gaps in the blue part of H-23's spectrum were masked, since in this case only one spectrum was good enough for science, so the CCD gaps could not be erased by the combination of different frames. Also note the red spectrum of H-17, noisier than in of the other two \hii\ regions due to the fact that 2, instead of 3 spectra were combined to obtain this one.} \label{fig_spectra} \end{figure} We measured the emission-line fluxes and their errors with the package {\sc SPLOT} of {\sc IRAF}. Errors take into account the statistical errors in the measurement of the fluxes and the systematic errors (flux calibrations, background determination and sky subtraction). We corrected the observed line fluxes for the effect of the interstellar extinction using the extinction law of \citet{mathis90} with R$_{\rm V}$=3.1 (which is assumed to be constant as in our Galaxy, but might slightly vary from galaxy to galaxy, see, e.g., \citet{clayton15}) and the individual reddening of each \hii\ region given by the \cbeta, i.e., the logarithmic difference between the observed and theoretical \hb\ fluxes. Since \hd\ and \hg\ lines are fainter than the \ha\ and \hb\ ones, and located in the bluer part of the spectra, they are consequently affected by larger uncertainties. We thus determined \cbeta\ comparing the observed Balmer I(\ha)/I(\hb) ratio with its theoretical value, 2.87 \citep[case B. c.f., ][]{osterbrock06}. We used lines in common between the blue and red part of the spectra to put the two spectral ranges on the same absolute flux scale. The average scaling factor applied to the fluxes of the red spectra is 0.78. In Table~B1 of the Appendix we present the measured and extinction corrected intensities --both normalised to H$\beta$=100-- of the emission lines of the 25 \hii\ regions listed in Table~\ref{tab_objid}. The spectra of the candidate symbiotic systems will be discussed in a forthcoming paper. The method to derive the physical and chemical condition in \hii\ regions is described in the previous papers of this series, as for instance \citet{MG09, goncalves12, goncalves14}: the {\sc temden} and {\sc ionic} tasks of {\sc IRAF}\footnote{The atomic data source is the that of {\sc analysis/nebular -- IRAF}; http://stsdas.stsci.edu/cgi-bin/gethelp.cgi?at\_data.hlp} are used to derive first electron temperature (from the ratio of the \oiii\ lines at 500.7~nm and 436.3~nm), and density (from the lines of the \sii\ doublet at 671.7~nm and 673.1~nm) of the gas, then ionic abundances. Ionic abundances are combined with the ionisation correction factors for \hii\ regions from \citet{izotov06} to obtain the total abundances. Often, literature studies represent \hii\ regions with a two-zone ionisation structure characterised by two different electron temperature, for the \oii\ and \oiii\ emitting regions. If available \teoii\ --or equivalently \tenii\ -- are used for the \oii\ zone, while \teoiii\ is adopted in the \oiii\ emitting zone. When \teoii\ or \tenii\ are not measured, the relation, based on the photoionisation models of \citet{stasinska82}, \teoii=0.7$\times$\teoiii+3000 K (or a similar one) is adopted. If we consider valid without errors the linear relation between \teoii\ and \teoiii, in the temperature range of our \hii\ regions, the use of a single temperature (the measured one, i.e., \teoiii) might imply a maximum underestimation of about $\sim$700 K and up to a maximum overestimation of $\sim$400 K of the temperature adopted for the \oii\ ionic abundance. We have tested the effect of adopting two different temperatures and, for most regions, it has negligible effect on the total oxygen abundance --of the order 0.01-0.03 dex-- since \oiii\ ionic abundance is on average a factor 10 larger than \oii\ ionic abundance, and consequently it is the ionic fraction that contributes the most to the total O/H. Thus we use only the measured \te\oiii\ for the calculation of the abundances of both low- and high-ionisation species. The abundances of \hei\ and \heii\ were computed using the equations of \citet{benjamin99} in two density regimes, that is \ne\ $>$ 1,000 and $\le$ 1,000 cm$^{-3}$. Clegg's collisional populations were taken into account \citep{clegg87}. The results of physical and chemical properties of the \hii\ regions are shown in Table~\ref{tabu2}. The \te\ has been measured in 18 regions and ranges from 8,600 to 12,000~K; \ne\ is available in 10 regions, while an upper limit is measured in other 6 regions. The measurements of electron densities range from 100 to 600 cm$^{-3}$, typical values of \hii\ regions, with a higher density of 1400 cm$^{-3}$ in H-13. Errors on final abundances take into account % the errors on \te\, on \ne\ and on the line fluxes. The average oxygen abundance determined with the \te\ method, excluding lower limit abundances, is 12+$\log$(O/H)=8.13$\pm$0.18. For the other elements we obtain average values: 12+$\log$(N/H)=7.18$\pm$0.28, 12+$\log$(Ne/H)=6.81$\pm$0.14, 12+$\log$(Ar/H)=6.00$\pm$0.25, and 12+$\log$(S/H)=6.04$\pm$0.25. Helium abundance is quite uniform too. Its mean value, linearly expressed, is He/H=0.092$\pm$0.019. The mean O/H is in good agreement with the metallicity measured by \citet{tullmann03} in their \hii\ region located in the disc-halo transition of NGC~55 (8.05$\pm$0.10 with the \te-method). The two extra-planar \hii\ regions of \citet{tullmann03} are instead slightly metal poorer (7.77 and 7.81) than our average O/H, but they are still consistent with the composition of some individual \hii\ regions, as for instance H-19, H-20, H-21. This might be an indication of incomplete mixing in the disc of NGC~55. \begin{table*} \begin{minipage}{1000mm} {\small \caption{Electron temperatures, electron densities, ionic and total abundance of the \hii\ regions.} \begin{tabular}{@{}lrrrrrrrrr@{}} \hline Diagnostic & H-1 & H-2 & H-5 & H-6 & H-7 & H-11 & H-12 & H-13 & H-14 \\ \hline $T_{e}$[O~{\sc iii}](K) & 12000$\pm$200 & 9200$\pm$150 & 9300$\pm$150 & 11500$\pm$200 &12000$\pm$200 & 8700$\pm$150&11000$\pm$200 &10700$\pm$200 &10800$\pm$200 \\ $N_{e}$~[S~{\sc ii}](cm$^{-3}$)& - & $<100$ & 100$\pm$50 & $<100$ &200$\pm$50 & $<$100 & 600$\pm$100 & 1400$\pm$200 &250$\pm$50 \\ He~{\sc i}/H & 0.14 & 0.07 & 0.09 & 0.07 & 0.11 & 0.10 & 0.075 & 0.07 & 0.10 \\ He/H &0.14$\pm$0.03 &0.07$\pm$0.01 & 0.09$\pm$0.01 &0.07$\pm$0.01 &0.11$\pm$0.05 &0.10$\pm$0.05 & 0.075$\pm$0.05 &0.07$\pm$0.02 & 0.10$\pm$0.05 \\ O~{\sc ii}/H & 2.1e-06 & 3.2e-05 & 1.9e-05 & 5.4e-05 & 1.4e-05 & 6.1e-05 & 2.2e-05 & - & 1.6e-05 \\ O~{\sc iii}/H & 1.3e-04 & 1.1e-04 & 1.5e-04 & 1.9e-05 & 7.9e-05 & 2.0e-04 & 1.1e-04 & 8.9e-05 & 9.8e-05 \\ ICF(O) & 1.2 & 1.0 & 1.0 & 1.0 & 1.0 & 1.0 & 1.0 & - & 1.0 \\ O/H & 1.6e-04 & 1.4e-04 & 1.7e-04 & 3.5e-05 & 9.3e-05 & 2.6e-04 & 1.3e-04 & $>$8.9e-05 & 1.1e-04 \\ 12+log(O/H) & 8.21$\pm$0.07 &8.15$\pm$0.06 & 8.23$\pm$0.06 & 7.90$\pm$0.10 & 7.97$\pm$0.12 & 8.42$\pm$0.12& 8.12$\pm$0.08 & $>$7.95 & 8.06$\pm$0.07 \\ N~{\sc ii}/H & 8.9e-07 & 3.5e-06 & 2.8e-06 & 3.9e-06 & 1.5e-06 & 2.9e-06 & 1.4e-06 & 1.17e-06 & 2.2e-06 \\ ICF(N) & 45.0 & 4.5 & 7.6 & 1.6 & 6.2 & 4.4 & 5.5 & - & 6.5 \\ N/H & 3.9e-05 & 1.5e-05 & 2.1e-05 & 6.2e-06 & 9.2e-06 & 1.2e-05 & 7.6e-06 & - & 7.7e-05 \\ 12+log(N/H) & 7.58$\pm$0.15 & 7.18$\pm$0.15 & 7.32$\pm$0.15 & 6.80$\pm$0.15 & 6.96$\pm$0.12 & 7.10$\pm$0.15& 6.88$\pm$0.12 & - & 7.16$\pm$0.12 \\ Ne~{\sc iii}/H & 5.1e-06 & 4.7e-06 & 5.5e-06 & - & 1.0e-05 & 8.0e-06 & 5.2e-06 & - & 8.1e-06 \\ ICF(Ne) & 1.0 & 1.2 & 1.0 & - & 1.1 & 1.2 & 1.1 & - & 1.0 \\ Ne/H & 5.1e-06 & 5.5e-06 & 5.5e-06 & - & 1.1e-05 & 9.5e-06 & 5.7e-06 & - & 8.5e-06 \\ 12+log(Ne/H) & 6.70$\pm$0.14 & 6.74$\pm$0.15 & 6.75$\pm$0.15 & - & 7.03$\pm$0.04 & 6.7$\pm$0.14 & 6.76$\pm$0.14 & - & 6.93$\pm$0.15 \\ Ar~{\sc iii}/H & - & 9.4e-07 & 8.8e-07 & 2.9e-07 & 5.4e-07 & 1.4e-06 & 7.5e-07 & - & - \\ ICF(Ar) & - & 1.1 & 1.2 & 1.2 & 1.1 & 1.1 & 1.1 & - & - \\ Ar/H & - & 1.0e-06 & 1.0e-06 & 3.5e-07 & 6.0e-07 & 1.5e-06 & 8.2e-07 & - & - \\ 12+log(Ar/H) & - & 6.00$\pm$0.20 & 6.00$\pm$0.20 & 5.53$\pm$0.20 & 5.78$\pm$0.22 & 6.20$\pm$0.25& 5.90$\pm$0.23 & - & - \\ S~{\sc ii}/H & - & 1.2e-06 & 8.5e-07 & 7.2e-07 & 3.4e-07 & 8.7e-07 & 4.2e-07 & 2.4e-07 & 6.3e-07 \\ S~{\sc iii}/H & - & - & - & - & 5.1e-07 &- &- &- & - \\ ICF(S) & - & 1.3 & 2.0 & 1.0 & 1.7 & 1.3 & 1.5 & - & 1.7 \\ S/H & - & 1.6e-06 & 1.7e-06 & 7.2e-07 & 1.4e-06 & 1.2e-06 & 6.5e-07 & - & 1.1e-06 \\ 12+log(S/H) & - & 6.19$\pm$0.30 & 6.22$\pm$0.30 & 5.90$\pm$0.30 & 6.16$\pm$0.28 & 6.10$\pm$0.30 & 5.80$\pm$0.30 & - & 6.05$\pm$0.30 \\ \hline \end{tabular} } \end{minipage} \label{tabu1} \end{table*} \begin{table*} \begin{minipage}{1000mm} {\small \contcaption{} \begin{tabular}{@{}lrrrrrrrrr@{}} \hline Diagnostic & H-15 & H-17 & H-18 & H-19 & H-20 & H-21 & H-22 & H-23 & H-25 \\ \hline $T_{e}$[O~{\sc iii}](K) &8600$\pm$150 &10000$\pm$150 &12000$\pm$500 &12000$\pm$200 &12400$\pm$300 &10500$\pm$200 &9900$\pm$150 &9100$\pm$150 &9600$\pm$150 \\ $N_{e}$~[S~{\sc ii}](cm$^{-3}$) & 100$\pm$50 & $<$100 & - & 300$\pm$100 &150$\pm$50 & $<$100 & $<$100 &150$\pm$50 &200$\pm$50 \\ He~{\sc i}/H & 0.075 & 0.07 & 0.10 & 0.08 & 0.10 & 0.09 & 0.095 & 0.094 & 0.115 \\ He/H & 0.075$\pm$0.02 & 0.07$\pm$0.02 & 0.10$\pm$0.05 & 0.075$\pm$0.02 & 0.07$\pm$0.02 &0.09$\pm$0.01 &0.095$\pm$0.01 &0.094$\pm$0.01 & 0.115$\pm$0.03 \\ O~{\sc ii}/H & 2.6e-05 & 4.2e-05 & - & 2.0e-06 & 1.3e-05 & 1.5e-05 & 1.0e-05 & 1.6e-05 & 4.2e-06 \\ O~{\sc iii}/H & 2.1e-04 & 1.1e-04 & 1.7e-04 & 9.1e-05 & 6.0e-05 & 4.2e-05 & 1.3e-04 & 1.6e-04 & 2.4e-04 \\ ICF(O) & 1.0 & 1.0 & - & 1.0 & 1.0 & 1.0 & 1.0 & 1.0 & 1.0 \\ O/H & 2.3e-04 & 1.6e-04 & $>$1.7e-04 & 9.3e-05 & 7.3e-05 & 5.7e-05 & 1.4e-04 & 1.8e-04 & 2.4e-04 \\ 12+log(O/H) & 8.36$\pm$0.12 & 8.20$\pm$0.07 & $>$8.25 & 7.70e$\pm$0.12 & 7.86$\pm$0.13 & 7.80$\pm$0.12 &8.16$\pm$0.08 & 8.25$\pm$0.08 & 8.38$\pm$0.12 \\ N~{\sc ii}/H & - & 3.8e-06 & - &3.9e-07 &1.4e-06 & 3.0e-06 & 2.6e-06 & 1.9e-06 & 1.9e-06 \\ ICF(N) & - & 4.0 & - &34.2 & 5.4 & 4.0 & 11.4 & 9.1 & 41.4 \\ N/H & - & 1.4e-05 & - & 1.5e-05 &7.2e-06 & 1.2e-05 & 2.9e-05 & 1.6e-05 & 7.7e-05 \\ 12+log(N/H) &- & 7.17$\pm$0.20 &- & 7.10$\pm$0.13 & 6.86$\pm$0.10 & 7.10$\pm$0.12 & 7.46$\pm$0.15 & 7.22$\pm$0.10 & 7.88$\pm$0.20 \\ Ne~{\sc iii}/H & 6.9e-06 & 7.9e-06 & - & 4.8e-06 & 3.8e-06 & - & 5.0e-06 & 8.6e-06 & 7.8e-06 \\ ICF(Ne) & 1.2 & 1.2 &- & 1.0 & 1.1 & - & 1.0 & 1.0 & 1.0 \\ Ne/H & 7.1e-06 & 9.8e-06 & - & 4.8e-06 & 4.2e-06 & - & 5.0e-06 & 8.6e-06 & 7.8e-06 \\ 12+log(Ne/H) & 6.85$\pm$0.14 & 7.00$\pm$0.14 & - & 6.66$\pm$0.15 & 6.62$\pm$0.14 & - & 6.70$\pm$0.15 & 6.93$\pm$0.15 & 6.89$\pm$0.15 \\ Ar~{\sc iii}/H & - & 7.2e-07 & - & 4.6e-07 & 6.0e-07 & 8.7e-07 & - & 1.4e-06 & 1.2e-06 \\ ICF(Ar) &- & 1.1 & - &2.8 & 1.1 & 1.1 & - & 1.3 & 3.2 \\ Ar/H & - & 7.6e-07 & - & 1.3e-06 & 6.6e-07 & 9.2e-07 & - & 1.8e-06 & 3.9e-06 \\ 12+log(Ar/H) & - & 5.88$\pm$0.24 & - & 6.10$\pm$0.17 & 5.82$\pm$0.23 & 6.00$\pm$0.24 & & 6.26$\pm$0.17 & 6.59$\pm$0.23 \\ S~{\sc ii}/H & 7.3e-07 & 5.3e-07 & - & 9.8e-08 & 3.0e-07 & 6.2e-07 & 4.6e-07 & 3.0e-07 & 6.1e-07 \\ ICF(S) & 2.0 & 1.2 & & 7.8 & 1.5 & 1.2 & 2.8 & 2.3 & 9.4 \\ S/H & 1.5 e-06 & 6.5e-07 & - & 7.7e-07 & 4.5e-07 & 7.8e-07 & 1.3e-06 & 6.9e-07 & 5.8e-06 \\ 12+log(S/H) & 6.20$\pm$0.30 & 5.80$\pm$0.30 & - & 5.90$\pm$0.30 & 5.65$\pm$0.30 & 5.90$\pm$0.30 & 6.11$\pm$0.27 & 5.84$\pm$0.27 & 6.76$\pm$0.30 \\ \hline \end{tabular} } \end{minipage} \label{tabu2} \end{table*} \subsection{Strong-line metallicities} % We computed oxygen abundances using strong-line methods to increase the number of regions with a determined metallicity. These methods are based on the intensities of lines that are usually easy to measure because they are much stronger than the lines used as \te\ diagnostic \citep[c.f.][for a complete discussion and comparison among the methods]{ac15}. The strong-line ratios can be calibrated in two different ways: using photoionisation models or using abundances of \hii\ regions obtained through the \te-method. Since the empirical calibration works better in the low metallicity regime, we have used the new calibrations based on \te-method abundances by \citet[][hereafter M13]{marino13} of the two well-known indices, N2=$\log$(\nii/\ha) and O3N2=$\log$[(\oiii/\hb)(\nii/\ha)]. The results are shown in Table~\ref{tab_strong} where we present the galactocentric distances, O/H from the \te-method, the metallicities derived with the M31's N2 and O3N2 indices, and an average between the two strong-line calibrators, which is the value adopted in the following figures. Errors on the adopted strong-line O/H take into account the flux uncertainties and the intrinsic errors of the method (0.18 and 0.16 dex, for O3N2 and N2, respectively, as quoted in M13). Comparing the average oxygen abundance derived with the \te-method, 12+$\log$(O/H)=8.13$\pm$0.18, with the average values determined with the strong-line method, we have: for the N2 index 12+$\log$(O/H)=8.14$\pm$0.12, for the O3N2 index 12+$\log$(O/H)=8.20$\pm$0.14, and for the combination of the two indices 12+$\log$(O/H)=8.17$\pm$0.13. They are thus in extremely good agreement. The increment of the number of regions analysed with the strong-line methods provides an even smaller dispersion of the distribution of the abundances in NGC~55 pointing towards a very homogeneous composition for the interstellar medium for this galaxy. \begin{table} \centering \begin{minipage}{75mm} {\scriptsize \caption{Strong-line calibrated oxygen abundances.} \begin{tabular}{@{}llllllll@{}} \hline Field-ID & D &O/H &O/H & O/H &O/H \\ &kpc &T$_e$ &N2& O3N2 & adopted \\ & & &M13 & M13 & \\ \hline H-1 & 12.13$^{+0.45}_{-1.35}$ & 8.21$\pm$0.07 & 8.03 & 7.97 & 8.03$\pm$0.21\\ % H-2 & 10.90$^{+0.14}_{-0.45}$ &8.15$\pm$0.06 & 8.14 & 8.18 & 8.16$\pm$0.21\\ % H-3 & 10.84$^{+0.27}_{-0.84}$ &- & 8.18 & 8.27 & 8.22$\pm$0.30\\ % H-4 & 10.21$^{+0.00}_{-0.00}$ & - & 8.30 & 8.57 & 8.43$\pm$0.21\\ % H-5 & 9.87$^{+0.10}_{-0.32}$ & 8.23$\pm$0.06 & 8.10 & 8.13 & 8.12$\pm$0.21\\ % H-6 & 9.01$^{+0.00}_{-0.01}$ & 7.90$\pm$0.10 & 8.25 & 8.32 & 8.28$\pm$0.22\\ % H-7 & 9.43$^{+0.09}_{-0.29}$ & 7.97$\pm$0.12 & 8.11 &8.13 & 8.12$\pm$0.25\\ % H-8 & 9.01$^{+0.05}_{-0.17}$ & - & 8.26 & 8.37& 8.32$\pm$0.21\\ % H-9 & 8.42$^{+0.00}_{-0.02}$ &- & 8.12 & 8.14 & 8.13$\pm$0.23\\ % H-10 & 9.52$^{+0.03}_{-0.11}$ & - & 8.27 & 8.44 & 8.31$\pm$0.21\\ % H-11 & 8.93$^{+0.07}_{-0.23}$ & 8.42$\pm$0.12 & 8.12 & 8.14 & 8.13$\pm$0.21\\ % H-12 & 8.55$^{+0.19}_{-0.59}$ & 8.12$\pm$0.08 & 8.01 & 8.09 & 8.05$\pm$0.21\\ % \\ H-13 & 6.21$^{+0.06}_{-0.21}$ & $>$7.95 & 8.00 & 8.08 & 8.04$\pm$0.21\\ % H-14 & 6.46$^{+0.50}_{-1.41}$ & 8.06$\pm$0.07 & 8.07 & 8.11 & 8.09$\pm$0.22\\ % H-15 & 5.77$^{+0.08}_{-0.24}$ &8.36$\pm$0.12 &-&- &-\\ % H-16 & 5.40$^{+0.02}_{-0.07}$ &- & 8.22 & 8.21 & 8.21$\pm$0.21\\ % H-17 & 4.42$^{+0.00}_{-0.01}$ & 8.20$\pm$0.07 & 8.12 & 8.18 & 8.15$\pm$0.22\\ % H-18 & 7.63$^{+1.21}_{-3.03}$ &$>$8.25 & -&-&- \\ % \\ H-19 & 1.46$^{+0.18}_{-0.46}$ &7.70$\pm$0.12 & 7.83& 7.99 &7.91$\pm$0.21 \\ % H-20 & 0.95$^{+0.06}_{-0.18}$ & 7.86$\pm$0.13 & 8.10 & 8.16 &8.13$\pm$0.22\\ % H-21 & 2.12$^{+0.52}_{-1.17}$ & 7.80$\pm$0.12 & -&-&- \\ % H-22 & 1.00$^{+0.25}_{-0.56}$& 8.16$\pm$0.08 & 8.13 & 8.13 & 8.13$\pm$0.21\\ % H-23 & 0.28$^{+0.08}_{-0.16}$ &8.25$\pm$0.08 & 8.00 & 8.09 & 8.05$\pm$0.21 \\ % H-24 & 0.43$^{+0.01}_{-0.03}$ & - & 8.27 & 8.25 & 8.26$\pm$0.22\\ % H-25 &0.98$^{+0.00}_{-0.00}$ & 8.38$\pm$0.12 & 8.05 & 8.05 &8.05$\pm$0.21\\ % \hline \end{tabular} \label{tab_strong} } \end{minipage} \end{table} \section[]{Radial abundance gradients in NGC~55} \label{sec_grad} \begin{figure*} \centering \includegraphics[width=18truecm]{grad_all_ngc55.eps} \caption{Radial abundance gradients of elemental abundances in the \hii\ regions in NGC~55. The abundances of metals (O, N, Ne, Ar, S) are expressed in the logarithmic form 12$+\log$(O/H), whereas the abundance of helium is expressed in linear form. The continuous curves are the weighted linear fits to the data, taking into account errors on both distances and abundances. } \label{fig_grad_all} \end{figure*} \begin{table} \caption{Radial abundance gradients in the disc of NGC~55} \begin{tabular}{lll} % \hline El. & Slope & Intercept \\ \hline O/H & +0.0025$\pm$0.0055 & 8.13$\pm$0.04\\ Ne/H & +0.0181$\pm$0.0081 & 6.78$\pm$0.06\\ N/H & -0.0030$\pm$ 0.0082 & 7.14$\pm$0.06\\ Ar/H & -0.0313$\pm$0.0145 & 6.18$\pm$0.10\\ S/H & +0.0074$\pm$0.0196 & 5.99$\pm$0.13\\ He/H & -0.0017$\pm$0.0008 & 0.094$\pm$0.006\\ \hline N/O &-0.0103$\pm$0.0103 & -0.880$\pm$0.074\\ \hline \end{tabular} \label{tab_grad} \end{table} The \hii\ regions for which we can determine plasma conditions, including \te, \ne, ionic and total abundances, give us the opportunity to study the spatial distribution of abundances and abundance ratios of several elements in the thin/thick disc of NGC~55. We have computed the linear galactocentric distances de-projecting and transforming them in linear distances with the inclination, position angle and distance of Table \ref{tabNGC55}. We use the range of 4$^\circ$ in the inclination angle obtained by the disc model of \citet{puche91} to estimate the uncertainties in de-projected galactocentric distance. The new model of \citet{westmeier13} is consistent with the one of Puche in the inner part of the galaxy where our regions are located. The sample of \hii\ regions (see green symbols in Figure~\ref{fig_allsources}) are located in a large galactocentric range of distances (from the centre to about 12~kpc, see Table~\ref{tab_strong}). In Figure~\ref{fig_grad_all} we show the radial abundance gradients of several elements: O, Ne, S, Ar, N and He. In Table~\ref{tab_grad} we report the slopes and the intercepts of the radial abundance gradients computed with the {\sc fitexy} routine that takes into account both the errors on abundances and galactocentric distances. For all the available elements we found null gradients, within the errors. NGC~55 radial gradients of O/H and N/H have been also recently re-analysed from literature data by \citet{Pilyugin14}. They found oxygen and nitrogen gradients essentially flat, and in good agreement with our results (see Table\ref{tab_grad}). Their slopes for oxygen and nitrogen are -0.0059$\pm$0.0104 dex~kpc$^{-1}$ and -0.0042$\pm$ 0.0145 dex~kpc$^{-1}$, respectively. On the other hand, \citet{k16} analysed a sample of 58 blue supergiant stars. For the first time, they detected a non negligible metallicity gradient of -0.22$\pm$0.06~dex/R$_{25}$ (-0.020$\pm$0.005 dex~kpc$^{-1}$ assuming R$_{25}$=11.0~kpc). Their central metallicity relative to the Sun is Z=-0.37$\pm$0.03. The comparison between the \hii\ region and blue supergiant populations is extremely interesting because they both are young populations and should trace the same epoch in the galaxy lifetime. In the upper panel of Figure~\ref{fig_gas_stars} we plotted the metallicity of the supergiants of \citet{k16} and of our \hii\ regions --reported on the Solar scale \citep[12+$\log$(O/H)=8.66;][]{grevesse07}-- versus the galactocentric distance. We have plotted the whole sample of \citet{k16} without removing the possible outliers. The two sets of metallicities and their radial distributions are in surprisingly good agreement considering all the internal uncertainties of the two metallicity derivations, the differences in the metallicities measured from nebular and stellar spectra (oxygen in the former, a global Z metallicity in the latter), and possible dust depletion of oxygen in \hii\ regions \citep[this correction might amount to 0.10~dex for objects with 7.8$<$12+$\log$(O/H)$<$ 8.3); see ][]{pp10}. It is, however, true that the supergiant abundances show a small decreasing gradient, which is not appreciable in the \hii\ region population abundances. \begin{figure} \centering \includegraphics[width=9truecm]{grad_ngc55_ngc300_gas_stars_new.eps} \caption{Upper panel: NGC~55 \hii\ regions (O/H derived with the \te\ method --filled circles-- and from the strong-line method --empty squares. For the strong-line abundances we adopted the average of the two values, as indicated in Table~\ref{tab_strong}) and supergiants from \citet{k16} (empty stars) radial distribution. The continuous line is the fit to the \hii\ regions with \te\ determinations as in Figure \ref{fig_grad_all}, whereas the dotted line is the fit to the whole sample of supergiant abundances. Lower panel: NGC~300 \hii\ regions \citep[filled circles from]{Bresolin09, Stasinska13} and supergiants from \citet{k08} (empty stars) radial distribution. } \label{fig_gas_stars} \end{figure} \subsection{$\alpha$-elements and nitrogen} Giving the similar origin of the four $\alpha$-elements --O, Ne, S and Ar-- we expect them to have similar radial behaviours. We indeed find almost flat gradients within the uncertainties for all of them, with Ar having possibly a small negative slope, still consistent with a null gradient, within the errors. While $\alpha$-elements are mainly synthesised by massive stars (M$>$8~M$_{\odot}$), nitrogen has a more complex nucleo-genesis, having both a \lq\lq primary" and a \lq\lq secondary" origin. The primary origin refers to conversion of the original hydrogen into nitrogen and it happens in stars with 4~M$_{\odot}<$M$<$8~M$_{\odot}$ \citep{RV81} and/or in very massive (M$>$30M$_{\odot}$) low metallicity stars \citep{WW95}, while the secondary channel is related with the production from C and O initially present in the ISM at the formation of the progenitor star. When the primary nitrogen component dominates, the N/O ratio is expected to be independent of the oxygen abundance and this happens at low metallicity. At higher metallicity, the secondary production becomes more important, and N/O increases with O \citep{vZ98}. The comparison of N with any $\alpha$-element can give indication on different star formation history in different radial regions of NGC~55 disc. If a galaxy experiences a dominant global burst of star formation, the ISM oxygen abundance will increase in about 10$^6$~yr, generating a decrease in N/O. Then over a period of several 100$\times$10$^6$~yr N/O will increase at constant O/H. Consequently N/O can be used as a clock that measures the time since the last major burst of star formation: low values of N/O imply a very recent burst of star formation, while high values of N/O imply a long quiescent period \citep[cf.][]{skillman98}. In Figure~\ref{fig_grad_no}, the radial gradient of N/O is shown. A slightly decreasing gradient towards the outskirts is detected. However this gradient is consistent with 2-$\sigma$ with a flat gradient and homogeneous distribution of abundances, indicating no differences in the star formation histories in different parts of the galaxy. We can also compare the typical N/O value in NGC~55 with those of dwarf and spiral star forming galaxies. From Figure~5 of \citet{Pilyugin03} or Figure~7 of \citet{Annibali15} we can infer that at 12$+\log$O/H$\sim$8.0~dex, there is an increasing dispersion in N/O, with values ranging from $-$1.6 to $-1$~dex. N/O in NGC~55 is located towards the upper envelope of this relationship. The average N/O of NGC~55 (N/O$=-$0.93$\pm$0.21) is close to the M33's one % (see, e.g., Figure 12 in Magrini et al. 2007, for a plot of N/O vs O/H) and to that of the LMC \citep[N/O=-0.96;][]{CR15}. The scatter in N/O values at a given O/H seen, e.g., by \citet{Pilyugin03} and \citet{Annibali15}, can be naturally explained by differences in the star formation histories of galaxies \citep{Pilyugin03}. This conclusion was already suggested by \citet{EP78} who explained the observations of the N/O abundance ratio in external galaxies due to the manufacturing of N in low-mass stars of 4-8 M$_{\odot}$. Very recently, \citet{vincenzo16} presented chemical evolution models aimed at reproducing the observed N/O versus O/H abundance pattern of star-forming galaxies in the Local Universe. They found that position of a galaxy in the N/O vs. O/H plane is mostly determined by its star formation efficiency (see their Figure 4). The high N/O ratio in NGC~55 with respect to other late-type dwarf irregular galaxies having the same O/H might be an indication that the bulk of the star formation happened in the recent past, more than 100$\times$10$^6$~yr ago. This is in agreement with the finding of \citet{Davidge05} of a vigorous star formation episode during the past 0.1-0.2 Gyr. The detection of significant numbers of stars evolving on the AGB phase, indicates that there has been vigorous star formation during the past 0.1-0.2 Gyr. In this time lapse, stars with masses M$>$5~M$_{\odot}$ might had time to evolve and to pollute with N the ISM of NGC~55 and thus to increase the N/O ratio. \begin{figure} \centering \includegraphics[width=9truecm]{grad_NO_ngc55.eps} \caption{Radial N/O gradients of \hii\ regions of NGC~55. } \label{fig_grad_no} \end{figure} \subsection{Helium} The measurement of He in low metallicity galaxies has been often used to extrapolate the primordial He abundance. To account for the amount of helium synthesised by stars, \citet{ptp74} suggested for the first time to study a sample of \hii\ regions spanning a wide range of metallicity. With the assumption of a helium enrichment proportional to metallicity, we can obtain the primordial abundance of He extrapolating to zero metallicity: $$Y=Y_{\rm P}+Z(\Delta Y/\Delta Z),$$ where Y is the mass fraction of He, and Z is the metallicity. Using Equation~2 of \citet{IT97} without correction for the neutral He \citep[$<$2\%, cf.][]{IT97}, we computed the mass fraction of He using the mean abundance of NGC~55, obtaining Y=0.27$\pm$0.08. Excluding the \hii\ region with the highest He abundance (H-1) we have a mean He/H=0.088$\pm$0.020 and thus Y=0.26$\pm$0.04. The linear regression of \citet{IT97} of Y versus oxygen, computed at the metallicity of NGC~55 gives a slightly lower value Y=0.250$\pm$0.002, but still consistent within the errors with our determination. The measurement of He abundance in NGC~55 does not give any particular constraint to the primordial He abundance, but it is in line with the determination of Y in several low metallicity galaxies. | \label{sec_conclu} In the present paper we show new spectroscopic observations of a large sample of 25 \hii\ regions in the Sculptor group member galaxy, NGC~55. We derive physical and chemical properties though the \te-method of 18 \hii\ regions, and strong-line abundances for 22 \hii\ regions. We measure also abundances of He, O, N, Ne, S, Ar. We found a homogenous composition of the disc of NGC~55, with average abundances of He/H=0.092$\pm$0.019, 12+$\log$(O/H)=8.13$\pm$0.18, 12+$\log$(N/H)=7.18$\pm$0.28, 12+$\log$(Ne/H)=6.81$\pm$0.14, 12+$\log$(Ar/H)=6.00$\pm$0.25 and 12+$\log$(S/H)=6.04$\pm$0.25. The abundances are uniformly distributed in the radial direction. This agrees with the study of smaller samples of \hii\ regions \citep{WS83, Pilyugin14} and it is in qualitative agreement with the blue supergiant radial gradient \citep{k16}. We investigate the origin of such flat gradient comparing NGC~55 with its companion galaxy, NGC~300, similar in terms of mass and luminosity and located in the same group of galaxies. The most plausible hypothesis is related to the differences in their K-band surface density profile that, as suggested by \citet{Pilyugin15}, which can provide higher mixing of the NGC55 gaseous component than in similar galaxies. | 16 | 9 | 1609.04210 |
1609 | 1609.04817_arXiv.txt | We characterize the column density probability distributions functions (PDFs) of the atomic hydrogen gas, \HI, associated with seven Galactic molecular clouds (MCs). We use 21 cm observations from the Leiden/Argentine/Bonn Galactic \HI~Survey to derive column density maps and PDFs. We find that the peaks of the \HI~PDFs occur at column densities ranging from $\sim 1$--$2\times 10^{21}~\cm$ (equivalently, $\sim 0.5$--1 mag). The PDFs are uniformly narrow, with a mean dispersion of $\sigma_{\rm HI}\approx 10^{20}~\cm$ ($\sim 0.1$ mag). We also investigate the \HI-to-\htwo~transition towards the cloud complexes and estimate \HI~surface densities ranging from 7--16 \sunits~at the transition. We propose that the \HI~PDF is a fitting tool for identifying the \HI-to-\htwo~transition column in Galactic MCs. | Setting the stage for star formation in the local Universe is the transformation of neutral atomic hydrogen, \HI, into molecular hydrogen, since stars are observed to form in the coldest, densest regions within molecular clouds (MCs). In galaxies, the atomic component of the interstellar medium (ISM) is expected to be organized in phases, the warm neutral medium (WNM) and the cold neutral medium (CNM), existing in near pressure equilibrium (e.g., Kulkarni \& Heiles 1987; Wolfire et al. 2003; McKee \& Ostriker 2007). Molecular clouds are expected to form in the coldest gas which, due to its higher number density, can more efficiently shield from dissociating photons than warm gas. Indeed, in the Milky Way and beyond, MCs are commonly observed to be associated with regions of high-column-density \HI~(e.g., Wannier et al. 1983, 1991; Elmegreen \& Elmegreen 1987; Chromey et al. 1989; Engargiola et al. 2003). Both Galactic and extragalactic studies indicate that MCs are frequently surrounded by envelopes of \HI~having surface densities of $\sighi\sim 8-10~\sunits$, (e.g., Engargiola et al. 2003; Imara \& Blitz 2011; Lee et al. 2012). Such observations suggest that \HI~is a necessary ingredient for MC formation. On the other hand, perhaps \HI~envelopes result from the photodissociation of molecular hydrogen, \htwo~(e.g., Allen et al. 2004; Heiner et al. 2011). It is also possible that such \HI~envelopes result from a combination of MC formation and evolution processes. Either way, these interpretations highlight the importance of understanding the function of atomic gas in MCs and stellar evolution. The various roles \HI~may play in MC evolution have been studied from theoretical and observational standpoints. First, some cold atomic clouds provide sufficient shielding against dissociation by the ultraviolet (UV) interstellar radiation field that molecular gas can form in their interiors. The structure of photodissociation regions (PDRs), where the transition from atomic to molecular gas occurs, has been treated theoretically by authors including van Dishoeck \& Black (1986), Hollenbach \& Tielens (1997), Draine \& Bertoldi (1996), Browning et al. (2003), and Krumholz et al. (2008; 2009). The latter investigated the PDR around a spherical cloud of solar metallicity, and estimated that a minimum surface density of $\sighi \sim 10~\sunits$ is required for \htwo~formation. Moreover, Goldbaum et al. (2011) noted that higher initial surface densities of \HI~envelopes may result in higher mass MCs. Second, MCs may continue to accrete atomic gas for an extended period during and their formation (e.g., Chieze \& Pineau Des Forets 1989; Hennebelle \& Inutsuka 2006; V\'azquez-Semadeni et al. 2010; Goldbaum et al. 2011). Authors including Goldbaum et al. (2011) and V\'azquez-Semadeni et al. (2010) demonstrated that while stellar feedback and accretion of atomic gas compete in regulating the mass of MC, accretion can be the deciding factor in the total mass of MCs. Also, accretion of atomic gas may shape the dynamics of MCs by maintaining turbulence (Chieze \& Pineau Des Forets 1989; Hennebelle \& Inutsuka 2006; Goldbaum et al. 2011). A third role of atomic gas in MC evolution is the influence of the CNM and WNM on molecule and star formation, as Stanimirovi\'c et al. (2014) and Lee et al. (2015) demonstrated may be the case in the Perseus molecular cloud. In this paper, we employ a useful tool for studying the \HI~enveloping MCs. We derive the probability distribution functions (PDFs) of the atomic gas associated with a group of Galactic clouds, with the aim of determining what the properties of the \HI~PDFs tell us about the MC evolution as well as the conversion from atomic to molecular gas---the \HI-to-\htwo~transition. In recent years, a great deal of attention has been paid to evaluating the PDFs derived from column density maps of MCs from dust extinction observations (e.g., Lombardi et al. 2006; 2008; 2011; Kainulainen et al. 2009; Alves et al. 2014) and from dust continuum emission observations (e.g., Lombardi et al. 2014; Schneider et al. 2013; 2015a; 2015b; Imara 2015). A number of theoretical studies determine that the widths of the column density PDFs correlate with the level of turbulence in MCs (e.g., Padoan et al. 1997; Federrath et al. 2008; Klessen 2000), while the high-extinction regimes of the PDFs correlate with the amount of high-density, self-gravitating, and potentially star-forming gas in MCs (e.g., Collins et al. 2012; Federrath \& Klessen 2013; Clark \& Glover 2014; Ward et al. 2014). Indeed, a number of observational studies have found that PDFs derived from dust measurements have a ubiquitous log-normal distribution over a narrow range of column densities and exhibit a characteristic power-law ``tail'' at high densities above visual magnitudes of $\av\sim 2$ (e.g., Kainulainen et al. 2009; Froebrich \& Rowles 2010; Lombardi et al. 2015). However, a consensus on how to interpret the widths and \emph{low}-extinction regimes of the dust PDFs has not been reached. This is because the dust extinction or emission observations used to derive the column density map toward a given MC has an inherent lower limit that is defined by the noise level of the observations. Therefore, the width of a PDF---which may or may not be a manifestation of the amount of turbulence in a cloud---depends on the lowest level contour used to determine that PDF (Lombardi et al. 2015). While a number of studies have been dedicated to understanding the column density PDFs of clouds from dust measurements, there are far fewer observational treatments of the PDFs of the atomic gas associated with MCs. Recently, Burkhart et al. (2015) examined the \HI~PDF of the Perseus molecular cloud. They found that the shape of the PDF is log-normal and much narrower than the PDF of the cloud derived from dust extinction measurements. The column density at the peak of the \HI~PDF in Perseus is close to the \HI~column density at the \HI-to-\htwo~transition, $\sim 6$--8 \sunits, measured by Lee et al. (2012). While this study leads to the tempting suggestion that the \HI~PDF may be used to determine the column density at the \HI-to-\htwo~transition, such a proposal would be more persuasive if similar evidence were found for a larger sample of clouds. Here, for the first time, the \HI~column density PDFs for a large group of Galactic MCs are investigated in detail, side by side with the dust extinction PDFs of the clouds. In \S 2, we summarize the observations used to conduct this study. In \S 3, we present our methods and results on \HI~column density PDFs. We discuss implications of our results, including the \HI-to-\htwo~transition and its connection with the \HI~PDF, in \S 4. Concluding remarks are presented in \S 5. | We have presented a detailed investigation of the PDFs of atomic gas associated with seven Galactic molecular clouds having a range of SFRs. Since most of the MCs in our sample lay far enough from the Galactic plane so that blending along the line of sight towards the clouds is mitigated, it was possible to isolate their \HI~envelopes using 21-cm observations and make reliable measurements of the properties. We created column density maps of the \HI~associated with MCs, we analyzed the features of the \HI~PDFs, and we measured the column density of atomic gas at the \HI-to-\htwo~transition. Our most salient results are summarized as follows. \begin{enumerate} \item The \HI~column density PDFs associated with MCs tend to have narrow, log-normal shapes. The peaks of the PDFs occur at column densities ranging from $\sim 1$--$2\times 10^{21}~\cm$. \item The properties of the \HI~PDFs are fairly robust to moderate variations to the spatial criteria used to select the 21-cm emission associated with the cloud complexes. The location of the column density at which an \HI~PDF peaks, \nhi, is sensitive to the chosen velocity range of the 21-cm emission used to create the PDF; however, the dispersion of the \HI~PDF, $\sigma_{\rm HI}$ is fairly insensitive to variations in the kinematic criteria. \item The column density at the peak of the \HI~PDF, \nhi, tends to increase with increasing MC mass. Moreover, higher mass MCs tend to have higher \HI~masses, $M_{\rm HI}$, and slightly higher \HI~mass fractions, $f_{\rm HI}$. % \item Most of the low-extinction material below $\av\lesssim 1$ mag represented in the NICEST column density maps towards a given MC cannot be explained by warm, diffuse \HI~towards the cloud. If this low-extinction material is not dominated by warm atomic gas or dark molecular gas, it is probably due chiefly to foreground or background that is unrelated to the cloud. \item The typical dispersion of the \HI~PDFs is $\sigma_{\rm HI}\approx 0.12\times 10^{21}$ \cm. The sonic Mach numbers we would predict for the \HI~based on measurements of the PDF variance in clouds would be subsonic to transonic. This is in contrast to the typical supersonic CNM sonic Mach number measurements found in both observations and simulations (e.g., Heiles \& Troland 2003; Gazol \& Kim 2013; Burkhart et al 2015). This suggests that there is a substantial WNM component in and around MCs which would lower the average sonic Mach number in the cloud, and/or the \HI~PDF width is truncated in MCs past the \HI-to-\htwo~transition due to the depletion of \HI. \item At the \HI-to-\htwo~transition in cloud complexes, we estimate \HI~surface densities ranging from $\sim 7$--16 \sunits. These values are consistent with $\Sigma_{0, \rm HI}$ obtained from least-squares fitting to the \HI~PDFs. We propose that---assuming an appropriate selection criteria for selecting \HI~emission is adopted---the \HI~PDF is a useful tool for identifying the \HI-to-\htwo~transition column in Galactic MCs. \end{enumerate} | 16 | 9 | 1609.04817 |
1609 | 1609.09104_arXiv.txt | Many models currently exist which attempt to interpret the excess of gamma rays emanating from the Galactic Center in terms of annihilating or decaying dark matter. These models typically exhibit a variety of complicated cascade mechanisms for photon production, leading to a non-trivial kinematics which obscures the physics of the underlying dark sector. In this paper, by contrast, we observe that the spectrum of the gamma-ray excess may actually exhibit an intriguing ``energy-duality'' invariance under $E_\gamma \rightarrow E_\ast^2/E_\gamma$ for some $E_\ast$. As we shall discuss, such an energy duality points back to a remarkably simple alternative kinematics which in turn is realized naturally within the Dynamical Dark Matter framework. Observation of this energy duality could therefore provide considerable information about the properties of the dark sector from which the Galactic-Center gamma-ray excess might arise, and highlights the importance of acquiring more complete data for the Galactic-Center excess in the energy range around 1~GeV. | A robust excess in the flux of gamma-ray photons emanating from the Galactic Center (GC) with energies of $\mathcal{O}(\mathrm{GeV})$ has been observed in Fermi Large Area Telescope (Fermi-LAT) data. This excess was first noted in Ref.~\cite{Goodenough:2009gk} and corroborated by a number of subsequent, independent analyses~\cite{Hooper:2010mq, Hooper:2011ti,Abazajian:2012pn,Hooper:2013rwa,Gordon:2013vta,Huang:2013pda, Abazajian:2014fta,Daylan:2014rsa,Lacroix:2014eea, YizhongGC, Calore:2014xka,Calore:2014nla}, including a dedicated study by the Fermi-LAT collaboration itself~\cite{TheFermi-LAT:2015kwa}. This excess consists not of a spectral line, but rather of a continuum bump which extends over a range of photon energies $0.3\mbox{~GeV}\lesssim E_\gamma \lesssim 50\mbox{~GeV}$ and peaks at approximately $E_\gamma \sim 1$~GeV. A variety of possible explanations have been advanced as to the origin of this gamma-ray excess. Possible astrophysical explanations include emission from a population of millisecond pulsars~\cite{Hooper:2010mq,Hooper:2011ti,Abazajian:2012pn,Gordon:2013vta, Abazajian:2014fta,Abazajian:2010zy} and the decay of neutral pions produced by collisions of cosmic-ray particles with interstellar gas~\cite{Hooper:2010mq,Hooper:2011ti, Abazajian:2012pn,Gordon:2013vta}. However, the spectrum produced by millisecond pulsars is too soft in the sub-GeV region to explain the observed data~\cite{Hooper:2013nhl} and millisecond pulsars born in globular clusters can account for only a few percent or less of the observed excess~\cite{Hooper:2016rap}. In addition, the observed distributions of gas seem to yield a poor fit to the spatial distribution of the signal~\cite{Lacroix:2014eea,Linden:2012iv,Macias:2013vya}. More recently, in Refs.~\cite{Bartels:2015aea,Lee:2015fea}, it has been asserted that the excess can be described by a set of unresolved point sources, and new methods to characterize these sources were devised. An exciting alternative possibility is that the excess is the result of annihilating or decaying dark-matter particles within the galactic halo. Indeed, the spatial morphology of the excess is consistent with dark-matter annihilations from a spherically symmetric density profile, and the excess extends outward more than $10^\circ$ from its center at the dynamical center of the Milky Way~\cite{Daylan:2014rsa}. As a result, a variety of models currently exist in the literature which posit a dark-matter origin for the continuum feature observed in the Fermi-LAT data. In such models, a suitably broad spectrum of gamma rays is realized through a variety of complicated cascade mechanisms. For example, such a gamma-ray spectrum can be generated via the subsequent showering and/or hadronization of Standard-Model (SM) particles initially produced directly from dark-matter annihilation. The observed excess is well reproduced by a dark-matter (DM) particle with a mass $m_\chi \sim (30-50)$~GeV and an annihilation cross-section $\langle\sigma v\rangle \approx (1-3)\times 10^{-26}~\hbox{cm}^3/\hbox{s}$ which annihilates primarily to $b\bar{b}$~\cite{Daylan:2014rsa,Calore:2014nla}. Likewise, a dark-matter particle with a mass $m_\chi \sim 10$~GeV and an annihilation cross-section $\langle\sigma v\rangle \approx (0.5-2)\times 10^{-26} ~\hbox{cm}^3/\hbox{s}$ which annihilates primarily to $\ell^+\ell^-$~\cite{Lacroix:2014eea} also provides a good fit to the Fermi-LAT data, provided that secondary photons produced by inverse Compton scattering and bremsstrahlung processes involving both primary and secondary electrons are taken into account. On the other hand, the recent AMS-02 data on the cosmic-ray antiproton flux~\cite{Aguilar:2015ooa, AMS02} has begun to exclude states in which a $q\bar{q}$ final state dominates~\cite{Giesen:2015ufa}. Concrete models in which the dark-matter candidate annihilates primarily to $b\bar{b}$~\cite{Alvares:2012qv, Okada:2013bna,Modak:2013jya,Alves:2014yha,Ipek:2014gua,Basak:2014sza} and to $\ell^+\ell^-$~\cite{Kyae:2013qna,Kim:2015fpa} have also been identified. Indeed, additional studies on other final states~\cite{Calore:2014nla} and generic model constraints~\cite{Kong:2014haa} have established that there exist further SM channels through which a dark-matter particle can annihilate or decay and reproduce the observed excess. Cascades involving one or more exotic intermediary particles which eventually decay down to SM fermions which in turn subsequently shower or hadronize have also been considered~\cite{Boehm:2014bia,Ko:2014gha,Abdullah:2014lla,SheltonDMCascade,Elor:2015tva}. While dark-matter models of this sort are capable of reproducing the GC excess, the showering and cascade dynamics on which these models rely in order to generate an acceptable gamma-ray spectrum have their disadvantages as well. For example, their complicated dynamics obscures the relationship between the detailed shape of the gamma-ray spectrum and the properties of the underlying dark sector. In this paper, by contrast, we propose a set of models in which the kinematics connecting the gamma-ray spectrum back to the dark sector is more straightforward. As a result, we find that characteristic imprints in the shape of that spectrum can potentially provide direct information about the dark sector. We begin our study by identifying a characteristic feature of the GC gamma-ray excess which points back to a particularly simple photon-production kinematics. In particular, we observe that the spectrum of this excess may potentially exhibit an intriguing ``energy duality'' under which the spectrum remains invariant under the transformation $E_\gamma\rightarrow E_\ast^2/E_\gamma$ for some self-dual energy $E_\ast$. As we shall argue, the presence of such an energy duality is indicative of a particularly simple kinematics in which the signal photons are produced directly via the two-body decays of an intermediary particle. Energy dualities of this sort have been exploited in other contexts involving similar decay kinematics, such as cosmic-ray pion decay~\cite{Stecker} and the decay of heavy (new) particles produced at colliders~\cite{KaustubhEnergyPeak}. At present, due to uncertainties in the astrophysical modeling of the GC region and also due to a paucity of reliable information about the shape of the spectrum at photon energies $\mathcal{O}(10\mbox{~MeV}) \lesssim E_\gamma \lesssim \mathcal{O}(1\mbox{~GeV})$, the information contained in the Fermi-LAT data alone is not sufficient to conclusively determine whether the spectrum of the GC excess in fact displays such an energy duality. Nevertheless, as we shall discuss, if such a duality {\it were}\/ to be confirmed through future experiments, this result would immediately favor a particular class of dark-matter models. Moreover, these observations apply not only for the GC gamma-ray spectrum but also for the spectra from other sources, such as dwarf galaxies, for which backgrounds can be more reliably estimated. While a spectrum with these duality properties can be realized in certain cascade-based models~\cite{DoojinAstroEnergyPeak}, we shall show that a self-dual gamma-ray spectrum also has a natural interpretation within the Dynamical Dark Matter (DDM) framework~\cite{DDM1,DDM2}. Indeed, as we shall show, there exists a simple class of DDM models which yield an energy-dual spectrum that provides an excellent fit to the Fermi-LAT data, with a self-dual energy $E_\ast \sim \mathcal{O}(1\mbox{~GeV})$. These results further highlight the importance of acquiring more complete gamma-ray data in the energy range $10\mbox{~MeV} \lesssim E_\gamma \lesssim 1\mbox{~GeV}$. This paper is organized as follows. In Sect.~\ref{sec:GCExcess}, we examine the energy spectrum of the GC excess and discuss the extent to which it might potentially exhibit an energy-duality invariance under $E_\gamma\rightarrow E_\ast^2/E_\gamma$ with $E_\ast \sim \mathcal{O}(1\mbox{~GeV})$. In Sect.~\ref{sec:Framework}, we then discuss how a gamma-ray spectrum with such an invariance can arise from dark-matter annihilation or decay. In Sect.~\ref{sec:Models}, we introduce a series of simple DDM models which give rise to a gamma-ray spectrum with this invariance and demonstrate that such DDM models provide a successful fit to the Fermi-LAT data. Our conclusions are then presented in Sect.~\ref{sec:conclusion}, where we also discuss the potential implications of energy duality for other astrophysical gamma-ray signals which might be observed in the future. Finally, an Appendix contains a derivation of certain results presented in the text. | } The possibility that the excess in the flux of gamma rays emanating from the vicinity of the GC is the result of annihilating or decaying dark matter is an intriguing one. If dark matter is indeed responsible for this excess, one pressing question is what, if anything, we can learn about the properties of the dark sector from the spectral information associated with that excess. Dark-matter models of the gamma-ray excess typically rely on complicated cascade mechanisms for photon production in order to reproduce the spectrum of the excess --- mechanisms whose non-trivial kinematics obscures the connection between the properties of that spectrum and the properties of the dark-matter candidate. In this paper, by contrast, we have considered an alternative dark-matter interpretation of the gamma-ray excess --- one in which a more direct connection exists between the properties of the underlying dark sector and the spectral shape of the gamma-ray excess to which it gives rise. In particular, we have pointed out that the spectrum of the observed excess in the Fermi-LAT data is potentially invariant with respect to an energy duality transformation of the form $E_\gamma \to E_\ast^2/E_\gamma$ for a self-dual energy $E_\ast \sim \mathcal{O}(1\mbox{~GeV})$. Motivated by this observation, we have presented a broad class of physical scenarios wherein such an energy self-duality is realized. In these scenarios, dark-matter annihilation/decay produces a non-trivial injection spectrum $dN_\phi/dE_\phi$ of intermediary particles $\phi$, each of which subsequently decays into a final state involving one or more photons which are mono-energetic and isotropically distributed in the $\phi$ rest frame. We have also shown that an appropriate injection spectrum of $\phi$ particles for describing the Fermi-LAT data is naturally realized within the context of the DDM framework. It is clear that our scenario relies directly on the existence of a multi-component dark sector, as this is a primary ingredient of the DDM framework. The possibility of non-minimal dark sectors has received increasing attention because many DM models predicated upon such sectors not only have non-trivial cosmological consequences (\eg, ``assisted freeze-out''~\cite{Belanger:2011ww}), but also often interesting phenomenological implications as well (\eg, ``boosted dark matter''~\cite{Agashe:2014yua,Berger:2014sqa,Kong:2014mia} as well as collider, direct-detection, and indirect-detection signatures~\cite{Dienes:2012yz, Dienes:2012cf, Dienes:2013xff, Dienes:2014bka, Dienes:2015bka} that transcend those normally associated with traditional WIMP-like single-component dark-matter scenarios). Indeed, multi-component dark sectors can even give rise to enhanced complementarity relations which can be used to probe and constrain the parameter spaces of such models~\cite{Dienes:2014via}. Thus, our explanation of the GC excess within the context of the DDM framework --- if corroborated by future experiments --- could provide an interesting window into the physics of the dark sector. Indeed, it would be interesting to study the cosmological and phenomenological implications of the particular set of DDM parameters obtained in our fit to the Fermi-LAT data. It is also important to realize that our discussion of the energy duality of the photon spectrum under $E_\gamma \rightarrow E_\ast^2/E_\gamma$ has a broad applicability that extends well beyond its application to the gamma-ray excess observed in the Fermi-LAT data. Indeed, this duality can be used as a tool for deciphering the origins of {\it any}\/ generic continuum excess which might potentially be observed at future X-ray or gamma-ray facilities. As discussed in Sect.~\ref{sec:Framework}, a broad range of spectral shapes can be realized within scenarios of the sort described above. In particular, any bump-like feature in the gamma-ray spectrum can be realized in such a scenario, provided \begin{itemize} \item the spectral feature is self-dual under the transformation $E_\gamma \rightarrow E_\ast^2 / E_\gamma$; and \item the spectral feature has a global maximum at $E_\gamma = E_\ast$ and decreases monotonically as $E_\gamma$ either increases or decreases away from $E_\ast$. \end{itemize} Moreover, we have shown that in scenarios of this sort, the shape of the spectral feature is directly correlated with the behavior of the intermediary injection spectrum at $E_\phi = m_\phi$. In particular, information about the kinematics of $\phi$ production and decay is manifest in the behavior of $dN_\gamma/dE_\gamma$ near its maximum: \begin{itemize} \item If $dN_\phi / dE_\phi $ remains non-zero as $E_\phi \rightarrow m_\phi$, the photon spectrum will exhibit a cuspy peak at $E_\ast$. \item If $dN_\phi / dE_\phi \rightarrow 0$ as $E_\phi \rightarrow m_\phi$, the photon spectrum will be smooth at $E_\ast$. \item If $dN_\phi / dE_\phi$ vanishes below some threshold energy $\overline{E} > m_\phi$, the photon spectrum will exhibit a plateau around $E_\ast$. \end{itemize} Thus, an excess of photons emanating from any astrophysical source which possesses the above features not only lends itself to an interpretation in terms of our annihilating/decaying dark-matter scenario, but can also yield additional information about the properties of the underlying dark sector. In general, an intermediary particle $\phi$ which couples to photon pairs in the manner indicated in Eq.~(\ref{eq:CouplingLagrangian}) will also couple to gluon pairs through an interaction term of the form $\mathcal{L} \ni (c_g/f_{\phi}) \phi G^a_{\mu\nu}\tilde{G}^{\mu\nu a}$, where $G^a_{\mu\nu}$ is the gluon field-strength tensor and $c_g$ is a dimensionless coefficient. Such an interaction term leads to additional, hadronic decay channels for $\phi$. The production of photons in association with these channels (through showering or from the decays of final-state hadrons) gives rise to an additional contribution to the differential photon flux. The production rate for these photons depends on the value of $c_g$, which is highly model-dependent. In this paper, for simplicity, we have assumed that $c_g \ll 1$ and that this showering/hadron-decay contribution to $d\Phi/dE_\gamma$ is therefore negligible. It is nevertheless interesting to consider how our results would be modified in situations in which $c_g \sim \mathcal{O}(1)$ and the showering/hadron-decay contribution is significant. In order to assess the impact of the showering/hadron-decay contribution on the overall photon signal spectrum from dark-matter annihilation in our DDM scenario, we begin by noting that this contribution, like the contribution from $\phi\rightarrow \gamma\gamma$ decay, owes its shape both to the kinematics of photon production in the rest frame of a decaying $\phi$ particle and to the spectrum of boosts imparted to the $\phi$ particles by the DDM ensemble constituents. The spectrum of boosts is governed in large part by the scaling exponent $\xi$, which characterizes how the contribution to the production rate of $\phi$ particles from an individual ensemble constituent $\chi_n$ scales with $m_n$ across the ensemble. As is evident from Fig.~\ref{fig:fit}, the best-fit values for $\xi$ for both of our ROI's are roughly $\xi \approx -2.5$, which implies that this contribution falls off rapidly with $m_n$. As a result, we find that the {\it collective}\/ contribution to the production rate for secondary photons from ensemble constituents $\chi_n$ with $m_n$ above a few GeV is essentially negligible, even when $c_g \sim \mathcal{O}(1)$. Thus, the contribution to the overall signal flux from showering/hadron decay is expected to be significant only for $E_\gamma \lesssim \mathcal{O}(1\mathrm{~GeV})$. The kinematics of photon production from showering/hadron-decay in the rest frame of the decaying intermediary has important ramifications as well. The primary parameter of interest here is $m_\phi$, our best-fit value for which is $m_\phi \approx 1.2$~GeV for both ROI's. Since this is less than twice the proton mass, baryon-number conservation implies that final states consisting primarily of light mesons --- and especially of pions --- should dominate the partial width of $\phi$ to hadrons. The dominant contribution to the secondary-photon spectrum at $E_\gamma \sim \mathcal{O}(1\mathrm{~GeV})$ is therefore likely to be the contribution from on-shell $\pi^0\rightarrow \gamma\gamma$ decay. Photons produced in this way have their own distinctive kinematics. In particular, the energy spectrum associated with these photons manifests an energy duality of its own, with self-dual energy $m_{\pi^0}/2$. The presence of such a duality could be exploited in order to disentangle this contribution from the primary photon spectrum. In principle, one could significantly reduce the contamination of the primary spectrum by subtracting off the contribution to the signal flux which is dual under $E_\gamma \rightarrow \m_{\pi^0}^2/(4E_\gamma)$. Moreover, since the shape of this $\pi^0$-decay contribution to the secondary photon spectrum is correlated with the shape of the primary-photon spectrum, a comparison between these two contributions could provide additional evidence in support of a DDM origin for the GC excess. Indeed, this strategy has been successfully employed within the context of other, similar DDM scenarios~\cite{MeVDDM}. Such an analysis would of course require improved data on the gamma-ray spectrum at energies $E_\gamma \lesssim m_{\pi^0}/2$. However, several proposals for instruments which would provide significant improvements in energy resolution within that energy range have been advanced~\cite{Boggs:2006mh,ASTROGAM}. Thus, we see that ``contamination'' from the showering/hadron-decay contribution to the differential photon flux that arises in this DDM scenario when $c_g \sim \mathcal{O}(1)$ may actually be an asset rather than a hurdle in the effort to distinguish this scenario from other models for the origin of this excess. One final comment is in order. In particular, we stress that although the gamma-ray excess observed in the Fermi-LAT data is {\it consistent}\/ with an energy duality of the kind we have discussed in this paper, there are significant uncertainties in the spectral shape of the excess which, at present, preclude any more definitive statements along these lines. These include not only statistical uncertainties, but also systematic uncertainties in the astrophysical foregrounds/backgrounds in the vicinity of the GC and uncertainties resulting from the energy resolution of the the Fermi-LAT instrument. Moreover, the preferred value for the self-dual energy $E_\ast \sim \mathcal{O}(1\mathrm{~GeV})$ is very close to the lower limit of the energy range for which reliable data exists. As a result, current data does not yet permit us to distinguish between the annihilating/decaying dark-matter scenario we have described here and other possible explanations of the CG gamma-ray excess. However, there are new astronomical instruments, both planned and under consideration, which are far better equipped to investigate whether the gamma-ray spectrum from the GC indeed exhibits such an energy-duality. For example, GAMMA-400 is expected to have a better energy resolution than Fermi-LAT in the $E_\gamma \sim 1~\gev$ regime. A variety of instruments designed to study the gamma-ray spectrum in the $10~\mev\lesssim E_\gamma \lesssim 1~\gev$ regime, such as ASTROGAM~\cite{ASTROGAM}, have also recently been proposed, often with energy resolutions far superior to those of similar experiments past or present. High-statistics data from such experiments could potentially definitively rule out or else lend significant credence to our scenario. Indeed, this illustrates that even when an excess of photons observed at indirect-detection experiments has the form of a broad continuum bump, precision measurements of the spectral shape of this bump can prove crucial for our understanding of the underlying physics. \bigskip | 16 | 9 | 1609.09104 |
1609 | 1609.00025_arXiv.txt | We present Karl G. Janksy Very Large Array (VLA) 1.3 cm, 3.6 cm, and 6 cm continuum maps of compact radio sources in the Orion Nebular Cluster. We mosaicked 34 square arcminutes at 1.3 cm, 70 square arcminutes at 3.6 cm and 109 square arcminutes at 6 cm, containing 778 near-infrared detected YSOs and $~190$ {\it HST}-identified proplyds (with significant overlap between those characterizations). We detected radio emission from 175 compact radio sources in the ONC, including 26 sources that were detected for the first time at these wavelengths. For each detected source we fit a simple free-free and dust emission model to characterize the radio emission. We extrapolate the free-free emission spectrum model for each source to ALMA bands to illustrate how these measurements could be used to correctly measure protoplanetary disk dust masses from sub-millimeter flux measurements. Finally, we compare the fluxes measured in this survey with previously measured fluxes for our targets, as well as four separate epochs of 1.3 cm data, to search for and quantify variability of our sources. | 16 | 9 | 1609.00025 |
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1609 | 1609.02485_arXiv.txt | { Based on recent \emph{Herschel} results, the ortho-to-para ratio (OPR) of NH$_2$ has been measured towards the following high-mass star-forming regions: W31C (G10.6-0.4), W49N (G43.2-0.1), W51 (G49.5-0.4), and G34.3+0.1. The OPR at thermal equilibrium ranges from the statistical limit of three at high temperatures to infinity as the temperature tends toward zero, unlike the case of H$_{2}$. Depending on the position observed along the lines-of-sight, the OPR was found to lie either slightly below the high temperature limit of three (in the range 2.2-2.9) or above this limit ($\sim3.5$, $\gtrsim 4.2$, and $\gtrsim 5.0$). In low temperature interstellar gas, where the H$_{2}$ is para-enriched, our nearly pure gas-phase astrochemical models with nuclear-spin chemistry can account for anomalously low observed NH$_2$-OPR values. We have tentatively explained OPR values larger than three by assuming that spin thermalization of NH$_2$ can proceed at least partially by H-atom exchange collisions with atomic hydrogen, thus increasing the OPR with decreasing temperature. In this paper, we present quasi-classical trajectory calculations of the H-exchange reaction \mbox{NH$_2$ + H}, which show the reaction to proceed without a barrier, confirming that the H-exchange will be efficient in the temperature range of interest. With the inclusion of this process, our models suggest both that OPR values below three arise in regions with temperatures $\gtrsim20-25$~K, depending on time, and values above three but lower than the thermal limit arise at still lower temperatures.} | Hydrides play a crucial role in astrochemistry as initial building blocks of the chemistry in both diffuse and dense clouds. The new spectroscopic window opened by the {\it Herschel} Space Observatory in the submillimeter and in the far-infrared (FIR) has allowed the detection of some fundamental and excited rotational transitions of simple neutral or ionized hydrides in different types of sources and especially towards the cold interstellar medium, in either the envelopes of low- and high-mass star-forming regions or in more distant regions along the lines-of-sight to these objects. In addition to the detection of new hydrides such as ND \citep{bacmann2010}, \ce{HCl+} \citep{deLuca2012}, and \ce{ArH+} \citep{barlow2013}, doubly and triply hydrogenated hydrides in their ortho and para forms such as \ce{H2Cl+} \citep{lis2010}, \ce{H2O+} \citep{ossenkopf2010,schilke2013,gerin2013}, \ce{NH2} and \ce{NH3} \citep{hilyblant2010nh,persson2010,persson2012,persson2016}, and \ce{H2O} \citep{emprechtinger2013} have been detected. Some of the ortho-to-para ratios (OPRs) were found to be consistent with their thermal values and some such as the cases of water \citep{lis2013,flagey2013}, \ce{H3+} \citep{crabtree2011}, \ce{NH3} \citep[][]{persson2012} and \ce{NH2} \citep{persson2016} were not. Observing and being able to constrain OPRs in such environments can bring crucial information about the prevailing physical conditions, such as the temperature of the gas, and also, on the other hand, can yield new constraints on the interstellar chemistry occurring in these milieu. For instance, the \ce{H2} OPR has been suggested as a chemical clock in cold molecular gas \citep{flower1984,flower2006op,pagani2009,pagani2011,pagani2013,brunken2014}. Comprehensive analysis of how these OPRs arise involves a deep study of the interstellar chemistry of these simple polyatomic species, which is often poorly known, especially concerning the processes and rates governing {\it (i)} the formation of ortho and para species and {\it (ii)} the ortho-to-para conversion. As an example, interest in the interstellar chemistry of nitrogen-hydride species has arisen as a result of observations of the lowest rotational transitions of the nitrogen hydrides \ce{NH}, \ce{NH2} and \ce{NH3} at far-IR wavelengths towards cold interstellar sources \citep{bacmann2010,hilyblant2010nh,persson2010,persson2012}. During the last decades, different gas-phase and gas-grain models were developed in an attempt to reproduce the observational data of these hydrides but with at most moderate success \citep{meyer1991,millar1991,wagenblast1993,oneill2002,hilyblant2010nh}. Based on new theoretical and experimental data \citep[\eg][]{flower2006n,hugo2009,honvault2011,honvault2012,rist2013,daranlot2013,daranlot2012,jorfi2009NO}, \cite{legal2014a} revised the understanding of nitrogen chemistry by focussing on the study of the basic gas-phase processes for the specific physical conditions of cold molecular gas. They revisited the low temperature kinetics of the nitrogen-bearing species and provided a rigorous nuclear-spin chemistry of the N-hydrides \citep{rist2013}. The result was a nearly pure gas-phase astrochemical model, which does not consider any grain reactions except the formation of \ce{H2} and charge exchange reactions. The network includes nuclear-spin selection rules for the hydrogenated nitrogen molecules and also for the hydrogen chemistry, which plays a crucial role in the synthesis of nitrogen hydrides \citep{lebourlot1991,dislaire2012}. This model was able to reproduce the abundances and abundance ratios of \ce{NH}, \ce{NH2}, and \ce{NH3} observed towards the envelope of the protostar IRAS 16293-2422. The model was also the first to lead to an explanation of the puzzling measurements of the OPRs of ammonia in cold diffuse gas ($T \sim 30$~K), which were found to be \citep[$\approx 0.5 - 0.7 \pm 0.1$,][]{persson2012,faure2013}, below the thermal value of unity. More recently, \cite{persson2016} investigated the non-thermal observational values of the \ce{NH2} OPR measured towards the high-mass star-forming regions W31C (G10.6-0.4), W49N (G43.2-0.1), W51 (G49.5-0.4), and G34.3+0.1 including translucent clouds in front of these sources. These authors, focussing their study on the \ce{NH2} OPR, did not investigate in detail the column densities of ortho and para \ce{NH2}. The \ce{NH2}-OPR values observed by \cite{persson2016} towards the different sources are displayed in Figure~\ref{fig:NH2_OPR_ture_thermal_obs} with the lowest values of the observed temperature ranges. The uncertainties in the observed temperatures are very large; if the upper values were utilized in the figure, the OPR ratio corresponding to these values should never exceed the statistical value of three so that the OPR measurements with values greater than three would likely have much larger uncertainties than reported. Although the use of nuclear-spin selection rules in an improved model leads to the reproduction of most of the observed OPR values below three at reasonable temperatures, it was necessary to find a mechanism that can at least partially thermalize the OPR at particularly low temperatures where the thermal OPR exceeds three and goes to infinity as the temperature goes to 0 K, as shown in black in Figure~\ref{fig:NH2_OPR_ture_thermal_obs}. This pattern occurs in the opposite sense from the H$_{2}$ thermal OPR, which is depicted in Figure~\ref{fig:H2_OPR_ture_thermal}. The difference arises because of the additional asymmetry of the ground electronic state of NH$_{2}$. \begin{figure} \centering \resizebox{\hsize}{!}{\includegraphics{figure1}} \caption{NH$_2$ OPR computed as a function of temperature at thermal equilibrium in black along with the observed OPR values from \cite{persson2016} in blue, cyan and pink at the lowest values of the observed temperature ranges \citep{persson2016}. For the sake of clarity the uncertainties of the observed temperatures are omitted. The NH$_2$-OPR statistical value of 3 is represented by the dashed red line.} \label{fig:NH2_OPR_ture_thermal_obs} \end{figure} \begin{figure} \centering \resizebox{\hsize}{!}{\includegraphics{figure2}} \caption{\ce{H2} OPR computed as a function of temperature at thermal equilibrium.} \label{fig:H2_OPR_ture_thermal} \end{figure} This need led \cite{persson2016} to consider the poorly studied NH$_{2}$ - H atom-exchange reaction as a mechanism to interconvert NH$_{2}$ between its ortho and para forms, previously omitted in models. With the assumption that this exchange occurs rapidly in both directions, \cite{persson2016} were able to explain the large values of the OPR observed in selected cold sources. But the H-exchange reaction between H and NH$_2$ had not been studied in detail. Weak experimental evidence from the saturated three-body reaction to produce ammonia indicates at most a small barrier \citep{pagsberg_pulse_1979}, while theoretical calculations indicate a more substantial barrier and a conical intersection \citep{mccarthy_dissociation_1987,zhu2012}. In this paper, we present a new theoretical calculation developed to determine if this reaction is a plausible efficient para-to-ortho conversion pathway for NH$_{2}$. We also determine how our calculated OPR is affected by recent updates concerning gas-phase reactions between NH$_{2}$ and the abundant oxygen and nitrogen atoms. In Section~\ref{sec:OPR}, we discuss the OPR of \ce{NH2} at low temperatures and how different non-thermal OPR values can be obtained, while emphasizing the role of the \ce{H + NH2} H-exchange reactions. Section 3 begins with a presentation of the theoretical work developed in order to study the \ce{H + NH2} H-exchange reaction proposed in \cite{persson2016}. This section continues with our astrochemical results derived from six new nearly pure gas-phase models, which differ from each other concerning: {\bf{\it (i)}} the inclusion of new destruction reactions for NH$_{2}$, {\bf{\it (ii)}} the initial form of hydrogen, {\bf {\it (iii)}} the cosmic-ray ionization rate, {\bf {\it (iv)}} the gas density, and {\bf {\it (v)}} the sulfur elemental abundance. In Section 4, we discuss the results, chiefly our new findings and how they fit the observational data, including the temperature range over which they can reproduce the measured OPR of NH$_{2}$. Finally, Section 5 contains a summary of our calculations and results. | \label{section:disc} The temperature ranges for which each model reproduces the NH$_{2}$-OPR values within their uncertainty ranges are tabulated in Table~\ref{tab:opr_results}. The table is constructed in the following format: the first column on the left lists the three types of sources observed and studied: molecular envelopes, dense and cold cores, and translucent gas. In each category, we tabulate the observed NH$_{2}$ OPR and range of temperatures associated with the sources in the regions W31C, W49N, W51, and G34. Unless the OPR values have very large uncertainties, the observed temperature ranges likely pertain to a diversity of regions not included in the OPR observations, which have a tight constraint on the temperature range, corresponding to the lowest values of the observed ranges. We then list the models that can reproduce the observed OPR values over some temperature range, which may or may not overlap with the observed range of temperatures. Several timescales are listed for each model. For example, consider the molecular envelope in W49N, which has an OPR of~3.5. This value can be matched by Models 1, 4, and 5, but none of these models can also match the high upper limit of the observed temperature ranges, which is as high as 120~K. On the other hand, consider the cold core in W51, where the OPR is measured to be 3.4. This can be matched by Models 1, 4, and 5. The observed temperature range for W51 is 10-30~K. All three models present smaller temperature ranges within this rather large observational range. Finally, we consider the translucent gas in W49N. Here the OPR value of $\ge$ 5.0 is matched by Models 1', 6, and 7. The observed temperature of $\le$ 15 K is matched well by the three models. Note that Model 1' is similar to Model 1 except that the density, ionization rate and sulfur abundance are those used for the translucent case, as described in Table~\ref{tab:models}. For those models with the thermalization of the OPR for NH$_{2}$ activated, we must also investigate how this activation changes the degree of agreement with OPR values for sources at higher temperatures, where these values are lower than three. As described in Table~\ref{tab:opr_results}, Models 2, 4, and 5 for dense sources, and Models 1', 6 and 7 for translucent sources generally fit the data, despite the fact that the range of the observed temperature uncertainties is large. Since the models presented in this study always give smaller temperature ranges than the observed ones, and towards the lower range of temperature, we are tempted to claim that the OPR measurements do not pertain to those portions of the observed clouds with temperatures high enough for the OPR to be only at the statistical limit of three or below, which is roughly 40 K. \begin{figure} \centering \resizebox{\hsize}{!}{ \includegraphics{figure10}} \caption{ NH$_2$ OPR computed with Model 2 for different cosmic-ray ionization rates. The \ce{NH2}-OPR values are plotted at two different times: $1\times10^5$~yr and at steady state.} \label{fig:NH2-ROP_ture_zeta_impact} \end{figure} \begin{figure} \centering \resizebox{\hsize}{!}{ \includegraphics{figure11}} \caption{Evolution of the abundances of H, N, O , \ce{p-NH2} (NH2 in cyan) and \ce{o-NH2} (NHH in yellow) computed at 20~K with Model 2 as functions of time and for different cosmic-ray ionization rates. The details of Model 2 can be found in Table~\ref{tab:models}.} \label{fig:H_N_O_NH2_time_zeta_impact} \end{figure} \begin{figure} \centering \resizebox{\hsize}{!}{ \includegraphics{figure12}} \caption{Calculated OPR of NH$_2$ computed as a function of temperature for a density of $n_\mathrm{H} = 1\times10^3$~cm$^{-3}$ for translucent sources. Shown are OPR values for thermal equilibrium and Models 6-7 at assorted times. The hatched boxes represent the OPR measurements from \cite{persson2016} within their formal errors, for the translucent gas towards W31C found at three different gas velocities. The dotted horizontal lines with arrows mark the lower limits in the translucent gas towards W31C in blue and W49N in pink. The temperature range is limited to the 5~K - 35~K even though the observed ranges can be as large as 100~K. See text.} \label{fig:NH2_ture_translucent_case} \end{figure} \begin{landscape} \begin{table}[\!htb] \centering \caption{Observed and calculated OPR values of NH$_{2}$ for assorted models with associated temperature ranges.} \begin{tabular} {|c|c|c|cc|cccccccccc} \hline\hline Source & OPR\tablefoottext{a} & $T_\mathrm{K}$ & \multicolumn{2}{c|}{$T_\mathrm{range}$ for model 1\tablefoottext{b}} &\multicolumn{2}{c|}{$T_\mathrm{range}$ for model 2}&\multicolumn{4}{c|}{$T_\mathrm{range}$ for model 4} &\multicolumn{2}{c|}{$T_\mathrm{range}$ for model 5}\\ &&&$5\times10^5$~yr &$t\gtrsim10^6$~yr&$(1-5)\times10^5$~yr&\multicolumn{1}{c|}{s.s.} &$1\times10^5$~yr &$2\times10^5$~yr&$5\times10^5$~yr&\multicolumn{1}{c|}{s.s.}&$1\times10^5$~yr&\multicolumn{1}{c|}{s.s.} \\ & & (K)& (K) & (K) & (K) & \multicolumn{1}{c|}{(K)}& (K) & (K) & (K) & \multicolumn{1}{c|}{(K)}& (K) & \multicolumn{1}{c|}{(K)} \\ \noalign{\smallskip} \hline \noalign{\smallskip} Molecular envelopes \\ \hline W31C & 2.5($\pm0.1$) &$30-50$\tablefoottext{c}&$28-35$&$32-35$&$\gtrsim$18&\multicolumn{1}{c|}{$17-23$}&--&--&--&\multicolumn{1}{c|}{$17-23$}&$\gtrsim22$&\multicolumn{1}{c|}{$25-29$}\\ W49N & 3.5($\pm0.1$)&$\sim$15-120\tablefoottext{d}&$5-12$&$24-25$&--&\multicolumn{1}{c|}{--}&$24-25$ & $15-17$&--&\multicolumn{1}{c|}{--}&--&\multicolumn{1}{c|}{$13-15$}\\ W51 & 2.7($\pm0.1$)&20-50\tablefoottext{e}&$23-28$&$29-32$&$\lesssim$18&\multicolumn{1}{c|}{$13-17$}&--&--&$\gtrsim$24&\multicolumn{1}{c|}{$13-17$}&$14-23$&\multicolumn{1}{c|}{$22-25$} \\ G34 & 2.3($\pm0.1$)&20-70\tablefoottext{e,f}&$\gtrsim$35&$\gtrsim$35 &$\gtrsim$35&\multicolumn{1}{c|}{$23-32$} &--&--&--&\multicolumn{1}{c|}{$23-32$} &$\gtrsim$35&\multicolumn{1}{c|}{$29-35$}\\ \noalign{\smallskip} \hline \noalign{\smallskip} Dense $\&$ cold core \\ \hline \noalign{\smallskip} W51 & 3.4($\pm0.1$) &10-30\tablefoottext{g}&$5-13$&$24-25$&--&\multicolumn{1}{c|}{--}&$24-26$ & $15-18$&--&\multicolumn{1}{c|}{--}&--&\multicolumn{1}{c|}{$14-16$}\\ \noalign{\smallskip} \hline \hline \noalign{\smallskip} Translucent gas & & & \multicolumn{2}{c|}{$T_\mathrm{range}$ for model 1'\tablefoottext{b}}&\multicolumn{4}{c|}{$T_\mathrm{range}$ for model 6}&\multicolumn{2}{c}{$T_\mathrm{range}$ for model 7} &\multicolumn{2}{|c|}{\multirow{7}{*}{\diagbox[width=35mm,height=27.5mm]{}{}}}\\ \hline \noalign{\smallskip} &&&$10^4$~yr &$\gtrsim 10^6$~yr &$10^4$~yr &$10^5$~yr &$2\tdix{5}$~yr &\multicolumn{1}{c|}{$\gtrsim 5\tdix{5}$~yr} &$10^4-5\tdix{5}$~yr&s.s. &\multicolumn{2}{|c|}{}\\ \noalign{\smallskip} \hline W31C & 2.2($\pm0.2$) & 30-100\tablefoottext{h} &$\gtrsim$27&$\gtrsim$34 &$\gtrsim$25&$\gtrsim$33 &$\gtrsim$34 &\multicolumn{1}{c|}{$\gtrsim$34} &$\gtrsim$35 &$\gtrsim$34 &\multicolumn{2}{|c|}{}\\ & 2.9($\pm0.2$) & 20-100\tablefoottext{h} &$5-16$& $25-29$ &$5-11$& $23-28$& $25-29$&\multicolumn{1}{c|}{$26-30$}& $27-31$ &$26-30$ &\multicolumn{2}{|c|}{}\\ & 2.6($\pm0.2$) & 25-75\tablefoottext{h} &$12-27$&$28-34$ &$5-25$& $27-32$& $28-33$& \multicolumn{1}{c|}{$29-34$}& $30-35$ &$29-34$&\multicolumn{2}{|c|}{}\\ W31C & $\gtrsim$4.2 & 30-85\tablefoottext{h} &$\lesssim17$&$\lesssim21$&--&$\lesssim17$&$\lesssim19$&\multicolumn{1}{c|}{$\lesssim20$}&$\lesssim21$&$\lesssim20$&\multicolumn{2}{|c|}{}\\ W49N & $\gtrsim$5.0 & $<$15\tablefoottext{h} &$\lesssim14$&$\lesssim19$&--&$\lesssim14$&$\lesssim16$&\multicolumn{1}{c|}{$\lesssim18$}&$\lesssim19$&$\lesssim18$&\multicolumn{2}{|c|}{}\\ \hline \end{tabular} \tablefoot{$T_\mathrm{K}$ and s.s. stand respectively for the observed temperatures and steady state. \tablefoottext{a} The tabulated errors are the formal errors from \cite{persson2016}. \tablefoottext{b} The models 1 and 1' are similar to the models b for dense and translucent cases, respectively, described in \cite{persson2016}. \tablefoottext{c}{\citet{fazio1978} and \citet{mueller2002}.} \tablefoottext{d}{\citet{vastel2001}.} \tablefoottext{e}{\citet{vandertak2013}.} \tablefoottext{f}{Derived from NH$_3$ rotational transitions \citep{hajigholi2016}.} \tablefoottext{g}{Derived from CN and NH$_3$ rotational transitions \citep{mookerjea2014}.} \tablefoottext{h}{The excitation temperature of the C\ion{I} 492~GHz line \citep{gerin2015}.} } \label{tab:opr_results} \end{table} \end{landscape} | 16 | 9 | 1609.02485 |
1609 | 1609.00163_arXiv.txt | The main advantage of Microwave Kinetic Inductance Detector arrays (MKID) is their multiplexing capability, which allows for building cameras with a large number of pixels and good sensitivity, particularly suitable to perform large blank galaxy surveys. However, to have as many pixels as possible it is necessary to arrange detectors close in readout frequency. Consequently KIDs overlap in frequency and are coupled to each other producing crosstalk. Because crosstalk can be only minimised by improving the array design, in this work we aim to correct for this effect a posteriori. We analysed a MKID array consisting of 880 KIDs with readout frequencies at 4-8 GHz. We measured the beam patterns for every detector in the array and described the response of each detector by using a two-dimensional Gaussian fit. Then, we identified detectors affected by crosstalk above -30 dB level from the maximum and removed the signal of the crosstalking detectors. Moreover, we modelled the crosstalk level for each KID as a function of the readout frequency separation starting from the assumption that the transmission of a KID is a Lorenztian function in power. We were able to describe the general crosstalk level of the array and the crosstalk of each KID within 5 dB, so enabling the design of future arrays with the crosstalk as a design criterion. In this work, we demonstrate that it is possible to process MKID images a posteriori to decrease the crosstalk effect, subtracting the response of each coupled KID from the original map. | \label{sec:intro} % In astronomy, several blank imaging surveys at different wavelengths have been carried out to study the formation and evolution of galaxies at different cosmic epochs \cite{Weiss2009,Geach2013}. A multi-wavelength approach is essential to have a more complete view of galaxy properties and, in particular, sub-millimetre observations are necessary to explore the dust component of galaxies. Microwave kinetic inductance detectors (MKID\cite{Day2003,Baselmans2012,Mazin2009}) are the ideal technology to built fast and large cameras, such as A-MKID\footnote{http://www3.mpifr-bonn.mpg.de/div/submmtech/bolometer/A-MKID/a-mkidmain.html} or NIKA\cite{Monfardini2010,Monfardini2011}, to carry out deep and large blank galaxy surveys in the sub-millimetre regime. The main advantage of this technology is the possibility to read out all detectors simultaneously throughout a single readout line. This is possible because each detector is tuned to a specific resonance frequency and they are read out by sending wave tones though the readout line. \par We used an array of 880 twin-slot antenna coupled hybrid MKIDs made for development and test in view of the SPACEKIDS\footnote{http://www.spacekids.eu} project. This technology has been already applied to similar array showing good efficiency and sensitivity\cite{Janssen2013}. % KIDs are tuned to absorb 350 GHz and have resonance frequencies between 4 GHz and 8 GHz with a design separation in readout frequency of 2.64-5.28 MHz and designed quality factors (Q-factors) around 40000. Detectors are organised in the array such that the nearest spatial neighbours are always separated by at least one other KID in readout frequency domain \cite{Yates2014}, in order to minimise the number of crosstalking KIDs. The only crosstalk in this array is therefore from the overlapping resonant dips of the KIDs themselves. This is both due by design, to maximise number of KIDs per readout line, and due to scatter in the KID placement due to lithographical and film thickness variations. \par The aim of this paper is to correct for the crosstalk a posteriori, both by describing the point spread function (PSF), as well as by deriving a theoretical model to predict the crosstalk as a function of the separation in readout frequencies of the KIDs from the resonance frequency and the quality factor of each KID. \par | \label{sec:end} In this work we characterised the level of crosstalk of an MKID array in order to correct images for crosstalk a posteriori. We measured beam maps for all KIDs in the array and we described the PSF as a two-dimensional Gaussian, for each KID with response above -30dB from the maximum. We subtracted the best-fit PSF of all coupled KID from the original beam map in order to remove the crosstalk present. Following this procedure, it is possible to correct astronomical images for crosstalk a posteriori. Analysing the level of crosstalk present in the array, we derived that about 72$\%$ of KIDs in the array crosstalk above -30 dB level while $\sim$48$\%$ crosstalk above -20 dB level. In the full array, 48$\%$ of all KIDs are coupled to another detector, 16$\%$ are coupled to other two detectors and 7$\%$ are coupled to other three KIDs. \par We estimated the resonance frequency and the quality factor of each KID by measuring the frequency sweep and describing the power of the transmission as a Lorenztian function. By using these parameters we derived a model by assuming that all KIDs have the same dip depth of the transmission and Q-factor. This model describes the expected general level of crosstalk as a function of the readout frequency separation of the detectors. This shows, both experimentally and using a simple model, that the rule of thumb is that KID-KID separation higher than 10 KID bandwidth corresponds to $\sim$-25 dB crosstalk. This can be taken as design criterion, as required, and as a way to estimate crosstalk for future arrays. | 16 | 9 | 1609.00163 |
1609 | 1609.00449_arXiv.txt | { Ring-like structures in the interstellar medium (ISM) are commonly associated with high-mass stars. Kinematic studies of large structures in giant molecular clouds (GMCs) toward these ring-like structures may help us to understand how massive stars form. }{ The origin and properties of the \object{ring-like structure G345.45+1.50} is investigated through observations of the $^{13}$CO(3-2) line. The aim of the observations is to determine the kinematics in the region and to compare physical characteristics estimated from gas emission with those previously determined using dust continuum emission. This area in the sky is well suited for studies like this because the ring is located 1$\circdot$5 above the Galactic plane at 1.8\,kpc from the Sun, thus molecular structures are rarely superposed on our line of sight, which minimizes confusion effects that might hinder identifying of individual molecular condensations. }{ The $^{13}$CO(3-2) line was mapped toward the whole ring using the Atacama Pathfinder Experiment (APEX) telescope. The observations cover 17$'$$\times$20$'$ in the sky with a spatial resolution of 0.2\,pc and an rms of $\sim$1\,K at a spectral resolution of 0.1\,km\,s$^{-1}$. }{ The ring is found to be expanding with a velocity of 1.0\,km\,s$^{-1}$, containing a total mass of 6.9$\times$10$^3$\,M$_\odot$, which agrees well with that determined using 1.2\,mm dust continuum emission. An expansion timescale of $\sim$3$\times$10$^6$\,yr and a total energy of $\sim$7$\times$10$^{46}$\,erg are estimated. The origin of the ring might have been a supernova explosion, since a 35.5\,cm source, J165920-400424, is located at the center of the ring without an infrared counterpart. The ring is fragmented, and 104 clumps were identified with diameters of between 0.3 and 1.6\,pc, masses of between 2.3 and 7.5$\times$10$^2$\,M$_\odot$, and densities of between $\sim$10$^{2}$ and $\sim$10$^{4}$\,cm$^{-3}$. At least 18\% of the clumps are forming stars, as is shown in infrared images. Assuming that the clumps can be modeled as Bonnor-Ebert spheres, 13 clumps are collapsing, and the rest of them are in hydrostatic equilibrium with an external pressure with a median value of 4$\times$10$^4$\,K\,cm$^{-3}$. In the region, the molecular outflow IRAS\,16562-3959 is identified, with a velocity range of 38.4~km~s$^{-1}$, total mass of 13~M$_\odot$, and kinematic energy of 7$\times$10$^{45}$~erg. Finally, five filamentary structures were found at the edge of the ring with an average size of 3\,pc, a width of 0.6\,pc, a mass of 2$\times$10$^2$~M$_\odot$, and a column density of 6$\times$10$^{21}$\,cm$^{-2}$. }{} | A wealth of observations have shown that ring- and shell-like structures in the interstellar medium (ISM) are ubiquitous \citep[e.g.][]{churchwell2006,churchwell2007,churchwell2008,heiles1979,martin-pintado1999,oka1998,schuller2009,wong2008,beaumont2010,sidorin2014,anderson2015}. In the Galaxy, these structures range up to 1.2\,kpc in radius, 2x10$^7$ M$_\odot$ in mass, and 10$^{53}$\,erg in kinetic energy. Their formation has been associated with the energy released by massive stars through stellar winds, supernova explosions, and ionizing radiation, or by collisions of high-velocity HI clouds with the Galactic disk \citep{tenorio-tagle1988}. Rings with a size smaller than 100\,pc have been related with OB associations and stellar clusters in their interiors \citep{tenorio-tagle1988}, and hot cores and molecular clumps at their edges \citep[e.g.][]{martin-pintado1999,wong2008}. In this study, the entire ring-like structure G345.45+1.50 was observed in the $^{13}$CO(3-2) line to study its properties including the kinematics. The ring is located $\sim$1$\circdot$5 above the Galactic plane, so there is little contamination with foreground and background molecular structures, which are mainly concentrated toward the Galactic plane, thus minimizing confusion effects in identifying individual condensations. Properties of the ring and of the identified clumps are estimated from $^{13}$CO(3-2) line observations, and compared with those found using the 1.2\,mm continuum emission. The molecular structure is also compared with infrared observations, from MSX\footnote{Downloadable from http://irsa.ipac.caltech.edu} and Spitzer$^1$. This ring is located at a distance of 1.8 kpc from the Sun, inside the \object{GMC G345.5+1.0}. This GMC is located approximately between 344\dotcirc5 and 346\dotcirc5 in Galactic longitude, between 0\dotcirc2 and 2\dotcirc0 in Galactic latitude, and between $-$33 and $-$2 km\,s$^{-1}$ in LSR velocity \citep{bronfman1989}. The ring-like structure contains a total mass of $\sim$4.0$\times$10$^3$\,M$_\odot$ estimated through 1.2\,mm continuum emission \citep{lopez2011}. The 1.2\,mm emission is fragmented at a spatial resolution of $\sim$0.2\,pc and composed of $\sim$54 clumps, which have an average mass of 75\,M$_\odot$. Nineteen clumps are associated with infrared sources identified in the Midcourse Space Experiment \citep[MSX;][]{price2001} and Spitzer observations \citep{benjamin2003}, including the source \object{IRAS\,16562-3959}, a massive star-forming region (MSFR) with a luminosity of $\sim$5.3$\times$10$^4$\,L$_\odot$ and associated with the most massive and dense clump in the ring. Within the ring lie also 35 cold clumps, that is, dense condensations that are not associated with an infrared counterpart. They might be regarded as candidates in which massive star formation has not yet started. \begin{table} \caption{List of observations with the ASTE telescope of the $^{12}$CO(3-2) line. Columns 1 and 2 show equatorial coordinates, Col. 3 the antenna temperature intensity peaks, and Col. 4 the estimated kinetic temperatures.} \begin{center} \label{tableASTE} \begin{tabular}{cccc} \hline\hline\\ R.A. & Dec.& $^1$$T^*_\mathrm{A}$ & $^1$$T_\mathrm{K}$ \\ \multicolumn{2}{c}{J2000} & K & K \\ \hline\hline\\ 16:59:08&-40:13:57&34&57\\ 16:59:41&-40:03:37&32&53\\ 16:59:44&-40:10:17&23&40\\ 16:59:46&-40:11:37&25&43\\ 16:59:20&-40:11:17&20&36\\ 16:59:13&-40:11:17&23&41\\ 16:59:08&-40:05:37&19&35\\ 16:59:32&-40:10:17&15&29\\ 16:59:35&-40:07:37&26&45\\ 16:59:04&-40:11:17&20&36\\ 16:59:23&-40:13:17&16&30\\ 16:59:35&-40:12:57&21&38\\ 16:59:53&-40:08:37&9.7&21\\ 16:59:20&-40:09:37&14&28\\ 16:59:11&-40:02:37&14&27\\ 16:59:28&-39:57:57&13&25\\ 16:59:23&-39:58:37&14&27\\ 16:59:51&-40:00:37&18&34\\ 16:59:44&-39:59:17&12&25\\ 16:59:32&-39:58:17&13&26\\ 16:59:16&-40:00:57&15&29\\ 16:59:18&-39:58:57&13&26\\ 16:59:30&-39:56:37&9.6&21\\ 16:59:06&-39:55:57&11&23\\ 16:58:50&-40:08:16&13&26\\ 16:59:56&-39:58:37&11&23\\ 16:59:49&-39:56:17&7.2&17\\ 16:58:52&-40:09:16&15&29\\ 16:59:58&-40:05:57&7.7&18\\ \hline \end{tabular}\\ $^1$ The error for $T_\mathrm{A}^*$ and $T_\mathrm{K}$ values is $\sim$0.2\,K. \end{center} \end{table} | \label{sectionConclusions} The ring G345.45+1.50 was observed in the $^{13}$CO(3-2) line. We list our conclusions below. \begin{itemize} \item The ring contains a total mass of 6.9$\times$10$^3$\,M$_\odot$, which agrees well with results from previous observations in the 1.2\,mm continuum emission. \item The ratio of the column density estimated from the 1.2\,mm continuum to that estimated from the $^{13}$CO(3-2) line varies between 0.1 and 10, with an average value of $\sim$1.0. The column density estimated from the 1.2\,mm continuum is higher than that estimated from the $^{13}$CO(3-2) line toward sites with high density. It is possible that the 1.2\,mm continuum emission is more optically thin than the $^{13}$CO(3-2) line toward dense regions. \item The ring is expanding with a velocity of 1.0\,km\,s$^{-1}$ and has an expansion timescale of 3$\times$10$^6$\,yr and a total energy of 7$\times$10$^{46}$\,erg. This expansion might have been produced by a supernova explosion. This hypothesis is supported by the presence of a 35.6\,cm source, J165920-400424, in the spatial center of the ring. This source does not have an infrared counterpart and might be a pulsar, that remained from the gravitational collapse of a massive star. This needs to be tested in future observations. \item From the $^{13}$CO(3-2) line, the ring is composed of 104 clumps with diameters of between 0.3 and 1.6\,pc, masses of between 2.3 and 7.5$\times$10$^2$\,M$_\odot$, and densities of between 10$^{2}$ and 10$^{4}$\,cm$^{-3}$. About 6\% of them show clear signs of star formation, which allowed us to estimate a typical lifetime for these clumps of $\lesssim$40\,Myr. \item Assuming that clumps can be modeled as Bonnor-Ebert spheres, $\sim$13\% of the clumps are gravitationally unstable, while most of the clumps require an external pressure with a median value of 4$\times$10$^4$\,K\,cm$^{-3}$ to be in hydrostatic equilibrium. \item In the region, the outflow associated with the MSFR IRAS\,16562-3959 was identified, with a velocity range of 38.4~km~s$^{-1}$, a mass of 13~M$_\odot$, and a kinetic energy of 7$\times$10$^{45}$~erg. \item At the edge of the ring, five filamentary structures are found with lengths of between 2.2 and 5.1\,pc and with masses of between 10$^2$ and 2.5$\times$10$^2$~M$_\odot$. \end{itemize} | 16 | 9 | 1609.00449 |
1609 | 1609.02350_arXiv.txt | In a Galactic core-collapse supernova (SN), axionlike particles (ALPs) could be emitted via the Primakoff process and eventually convert into \Grays in the magnetic field of the Milky Way. From a data-driven sensitivity estimate, we find that, for a SN exploding in our Galaxy, the \Fermi Large Area Telescope (LAT) would be able to explore the photon-ALP coupling down to $g_{a\gamma} \simeq 2 \times 10^{-13}$~GeV$^{-1}$ for an ALP mass $m_a \lesssim 10^{-9}$~eV. These values are out of reach of next generation laboratory experiments. In this event, the \FermiLAT would probe large regions of the ALP parameter space invoked to explain the anomalous transparency of the Universe to \Grays, stellar cooling anomalies, and cold dark matter. If no $\gamma$-ray emission were to be detected, \fermiLAT observations would improve current bounds derived from SN\,1987A by more than one order of magnitude. | 16 | 9 | 1609.02350 |
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1609 | 1609.07743_arXiv.txt | During its two years mission around comet 67P/Churyumov-Gerasimenko, ESA's Rosetta spacecraft had the unique opportunity to follow closely a comet in the most active part of its orbit. Many studies have presented the typical features associated to the activity of the nucleus, such as localized dust and gas jets. Here we report on series of more energetic transient events observed during the three months surrounding the comet's perihelion passage in August 2015. We detected and characterized 34 outbursts with the Rosetta cameras, one every 2.4 nucleus rotation. We identified 3 main dust plume morphologies associated to these events: a narrow jet, a broad fan, and more complex plumes featuring both previous types together. These plumes are comparable in scale and temporal variation to what has been observed on other comets. We present a map of the outbursts source locations, and discuss the associated topography. We find that the spatial distribution sources on the nucleus correlates well with morphological region boundaries, especially in areas marked by steep scarps or cliffs. Outbursts occur either in the early morning or shortly after the local noon, indicating two potential processes: Morning outbursts may be triggered by thermal stresses linked to the rapid change of temperature; afternoon events are most likely related to the diurnal or seasonal heat wave reaching volatiles buried under the first surface layer. In addition, we propose that some events can be the result of a completely different mechanism, in which most of the dust is released upon the collapse of a cliff. | The OSIRIS cameras on board ESA's Rosetta spacecraft have monitored the activity of comet 67P-Churyumov-Gerasimenko (67P) across varying heliocentric distances (4 AU to 1.24 AU) and different seasons on the nucleus (sub solar latitude between +45 and -55 degrees). Previous publications focused particularly on coma features usually referred to as jets: collimated streams of dust and gas arising from the nucleus. The foot prints of these features on 67P, their migration with the seasons and heliocentric distance, their relation to topography, their photometry, and possible formation mechanisms are described in details in \cite{lara2015, lin2015, lin2016} and \cite{vincent2016}. One of the striking discoveries of Rosetta has been the clockwork repeatability of jets from one rotation to the next. Jets are very dynamic by nature, depending on the complex hydrodynamics of the gas and dust streams interacting with the local topography, and controlled by local thermal conditions. They grow and fade with the solar illumination as the nucleus rotates, but the same exact features can be observed from one rotation to the next. Figure \ref{fig:repeated_jets} shows an example of this phenomenon. This, of course, put constraints on the thermophysics and volatile content of active areas, which need to ensure the sustainability and repeatability of the jets we observed. \begin{figure} \centering \includegraphics[width=0.49\hsize]{figures/NAC_20150809T120949.png} \includegraphics[width=0.49\hsize]{figures/NAC_20150810T002300.png} \caption{Example of two images acquired on 2015-08-09T12.09.49 (left) and 2015-08-10T00.23.00 (right), almost one rotation apart (Rotation period - images separation = 5min33s). Both images contrast is stretched to the same level (5\% of the same maximum brightness value). Field of view 1x1 degree, distance = 305 km, resolution = 5.7 m/px.} \label{fig:repeated_jets} \end{figure} Long-lasting repetitive features are however not the only manifestation of activity on comet 67P. In this paper, we report on another types of events, much more transient in nature, which were observed most frequently around the summer months of 67P's southern hemisphere, i.e. from July to September 2015, when the comet reached its perihelion (13 August 2015, 1.24 AU). These events are characterized by the sudden and short release of a dust, sometimes collimated but not necessarily. While the typical jets are relatively faint (about 10\% of the nucleus surface brightness), the plumes ejected by these outbursts are usually as bright as the nucleus, and they can be detected in our images without enhancing the contrast. Contrary to the jets that last for several hours, most transient events are observed only once, indicating a lifetime shorter than the cadence of our images (between 5 and 30 min, depending on the observing sequence). One sequence showing a transient event is presented in Fig. \ref{fig:perihel_out}. We report here our detection of these transient events, during a 3-months period surrounding the perihelion passage. In the following text transient events will alternatively be referred to as \textit{outbursts} to indicate their sudden and bright behavior, bearing in mind that they are many orders of magnitude fainter than typical cometary outbursts detected routinely by ground based observers for other comets. \begin{figure} \centering \includegraphics[width=0.32\hsize]{figures/NAC_20150812T170504.png} \includegraphics[width=0.32\hsize]{figures/NAC_20150812T173504.png} \includegraphics[width=0.32\hsize]{figures/NAC_20150812T180504.png} \caption{A transient event detected on the day of perihelion (12 August 2015). Images are separated by only 1/2 h, contrast not enhanced. Observations before and after the event show only faint activity from the outbursting area, while the image at 17:35 reveals a plume of material as bright as the nucleus, expanding at least 10km from the source. Field of view 1x1 degree, distance = 332 km, resolution = 6.1 m/px. Outburst \#14 in Table \ref{tab:outbursts}.} \label{fig:perihel_out} \end{figure} | We have presented a series of 34 transient release of gas and dust by the nucleus of comet 67P over the three months surrounding its perihelion passage mid-August 2015. We found that outbursts occur about every 2.4 nucleus rotation and last at most a few minutes. They are comparable in scale and temporal variation to what has been observed on other comets. The dust plumes released by these events can be classified into 3 main morphologies: narrow jets, broad plumes, or a combination of both. We produced a map of the source locations of these events, and discussed the associated local topography. We find that the spatial distribution of outbursts locations on the nucleus correlates well with morphological region boundaries, especially areas marked by steep scarps or cliffs. Outbursts occur either in the early morning or shortly after the local noon, possibly indicating two types of processes: Morning outbursts may be triggered by thermal stresses linked to the rapid change of temperature, while afternoon events are most likely related to the diurnal or seasonal heat wave reaching volatiles buried deeper in the nucleus than those responsible for the more typical jets. In addition, we propose that some events can be the result of a completely different mechanism, in which most of the dust is released upon the collapse of a cliff. This idea will be tested with a more detailed morphological study using forthcoming high resolution images, joint analysis with other instruments on board Rosetta, and numerical modeling. | 16 | 9 | 1609.07743 |
1609 | 1609.03356_arXiv.txt | The main spectral characteristic of Balmer-dominated shocks (BDSs) is a two-component \Ha\ line. Components are produced in radiative decays of excited atoms \citep{ckr80}, whereas cold hydrogen atoms (the ones overrun by the shock) produce a narrow component ($\sim$10\,\kms) and broad-neutrals generate a broad component ($\sim$1000\,\kms). Broad-neutrals are formed in a charge exchange (CE) process between a cold atom and a hot proton downstream of the shock. BDSs are an important diagnostic tool for shock parameters \citep{heng10}: narrow (broad) component width indicates the pre (post)-shock temperature; the shock velocity can be estimated from the broad line width and in combination with proper motions of optical filaments provides distance estimates; the electron-to-proton temperature ratio can be constrained from the components' widths and their intensity ratios. The \Ha-line profile is also influenced by possible shock precursors (emissions from the shock interacting with the pre-shock medium), for example cosmic-rays (CRs) and broad-neutral precursors. Heating in the CR precursor results in a narrow \Ha\ component broadened beyond the normal 10--20\,\kms\ gas dispersion \citep{mor13}. Furthermore, CRs transfer also momentum to the pre-shock neutrals introducing a Doppler shift between the pre- and post-shock gas \citep{lee07}. CE in the broad-neutral precursor introduces an additional component to the \Ha\ profile -- intermediate component -- with the width of $\sim$150\kms\ \citep{mor12}. The \Ha-line that is broader than the intrinsic 20\,\kms\ was previously measured in low-spatial resolution data of Tycho \citep{ghava00,lee07}, and interpreted as a strong indicator of CR production. The same studies reported also on the detection of the intermediate component. However, both analyses focused on the \Ha-bright, but very complex 'knot\,g', where multiple or distorted shock fronts can contribute to the measured broadening of the narrow and intermediate \Ha\ line. Just as spatial resolution is crucial to eliminate or reduce the artificial broadening effect of differential projection, extended spatial coverage of the filament can help to ascertain spatially varying shock (and ambient ISM) conditions. For the first time, we provide both of these in our forthcoming analysis. We expand on previous studies also by providing full posterior distributions instead of relying merely on best-fit parameters (line fluxes, centroids and widths), and provide a quantitative assessment of the significance of the intermediate line as well as possible multiple narrow lines (projected shock fronts). | We presented the narrow \Ha\ spectroscopic observations of Tycho's NE Balmer filaments. This study provides spectroscopic data that is for the first time spatially resolved (spectro-imagery), with large coverage that comprises and resolves the entire NE filament. Our analysis includes Bayesian model comparison that enables a quantitative, probabilistic and well-defined model comparison. We find that the broadening of the NL beyond 20\,\kms\ that was noted in previous studies was not an artifact of the spatial integration, and that it extends across the whole filament, not only the previously covered 'knot\,g'. Likewise, we confirm the suspected presence of an IL, and show it to be widespread. The first result points toward the evidence of heating in the CR precursor, while the second result reveals the presence of the broad-neutral precursor. | 16 | 9 | 1609.03356 |
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1609 | 1609.06745_arXiv.txt | The initial magnetic field of previous magnetorotational instability (MRI) simulations has always included a significant system-scale component, even if stochastic. However, it is of conceptual and practical interest to assess whether the MRI can grow when the initial field is turbulent. The ubiquitous presence of turbulent or random flows in astrophysical plasmas generically leads to a small-scale dynamo (SSD), which would provide initial seed turbulent velocity and magnetic fields in the plasma that becomes an accretion disc. Can the MRI grow from these more realistic initial conditions? To address this we supply a standard shearing box with isotropically forced SSD generated magnetic and velocity fields as initial conditions, and remove the forcing. We find that if the initially supplied fields are too weak or too incoherent, they decay from the initial turbulent cascade faster than they can grow via the MRI. When the initially supplied fields are sufficient to allow MRI growth and sustenance, the saturated stresses, large-scale fields, and power spectra match those of the standard zero net flux MRI simulation with an initial large scale vertical field. | Magnetic fields are ubiquitous in turbulent astrophysical plasmas. While the mechanism of origin of large-scale ordered magnetic fields in these systems is a subtle business, more generic and less controversial is the amplification of total magnetic energy by the fluctuation or small-scale dynamo (SSD). Here, turbulence in a conducting plasma, with even a modest magnetic Reynolds number ($\Rm$), leads to exponential growth of the field on the shortest eddy turn over time scales, which is usually much smaller than the age of the astrophysical system. The SSD is likely to be important for the early generation of magnetic fields in stars and galaxies/inter-stellar medium (ISM) \citep{Suretal10,gentetal13,BS13}. Such SSD generated fields would then be present in the plasma from stars or the ISM that source accretion disks. In previous studies of the magneto-rotational instability (MRI) in shearing box models of accretion disks, the initial condition is typically an ordered \textit{non-stochastic} field with net flux or zero-net flux. Following linear stability analysis, this kind of initial condition is the most natural to compare with minimalist analytic theory, but in reality one expects a more random field without necessarily much of a large-scale field. The question then arises as to whether and what kind of disorder and incoherence in the initial field can still lead to MRI and field growth. Moreover, in general the source of an initial magnetic field itself, which could trigger MRI in a disk is not known. Initial fields in disks may be supplied externally and this may also be generated by SSD. Here, we explore this latter possibility, and the condition under which large-scale fields could grow and be sustained through MRI-driven turbulence initiated by SSD fields. \begin{figure*} \includegraphics[width=0.495\textwidth, height=0.3\textheight]{brms-time-ssd-mri512.png} \includegraphics[width=0.495\textwidth, height=0.3\textheight]{stressLSDevol-ssdmri512.png} \caption{Evolution of $\Brms$ and $\Urms$ is shown in upper and lower left panels respectively for Run~SSD, Run~A0 and Run~STD. Evolution of energy in large-scale magnetic fields and sum of Maxwell and Reynold stresses is shown in upper and lower right panels respectively for Run~SSD, Run~A0 and Run~STD. The solid curves refer to $x$-$y$ averaged fields while the dashed curves refer to $y$-$z$ averaged fields. } \label{timeplot} \end{figure*} In the case of initial vertical fields of zero flux or net flux field, the MRI modes grow large-scale radial and azimuthal fields in the early linear growth phase \citep{BEB16, EB16} on orbital times scales eventually saturating nonlinearly to generate turbulence. It has been a long-standing topic of investigation to understand what determines the amplitude of the fields and stresses on MRI saturation, because this is thought to constrain the rate of angular momentum transport in accretion disks. The saturation amplitude of the stresses is not a constant across simulations but, for fixed resolution and box dimensions, depends on whether the initial field is of zero-net flux type or there is a uniform background field \citep{HGB95, fromang07, Guanetal2009, Shi2016}. Here we investigate whether random fields (without a significant large-scale component) can seed the MRI, and whether the MRI sustains. Previous work using initially random fields \citep{HGB96} had adopted a flat 1-D magnetic spectrum, and also with mostly large scale modes having $1 < k(L/2\pi) < 4$ and the field being $0$ outside that interval. Such an initial field has a significant large scale component, (with the box scale field comparable to smaller scale fields), unlike the initial conditions we adopt below. In particular, we show that using small-scale fields (which by definition implies the absence of significant power at large scales) from a small-scale dynamo (SSD) as an initial condition, subsequent shear and rotation act to trigger MRI and as a result, sustains the turbulent fields even after the SSD forcing is turned off. We also discuss that the MRI fails to sustain when Gaussian random noise is instead used as an initial field. Note that \citet{Riols13} generate an initial condition with random set of Fourier modes to study the transition to MRI turbulence, but with apparently large scale modes being energetically dominant. In our case, the Gaussian random initial fields are essentially noise-like with no large-scale component. In section 2, we compare the sustained MRI turbulence with SSD initial condition against the case with zero net flux initial condition, with respect to the standard MRI signatures. Section 3 discusses the role of coherence in initial random magnetic fields for MRI to grow. The nature of MRI starting with SSD generated fields is investigated by spectral analysis in Section 4. We conclude in Section 5. | We have shown via direct numerical simulations that the MRI can sustain MHD turbulence even when seeded with an initially random small-scale magnetic and velocity fields supplied by an SSD. In this case, the energy is strongly dominated by fields at small scales, implying that substantial power at low wave modes is not necessary for MRI growth There is however, still a minimum strength and coherence required for this growth that is determined by a comparison between the turbulent diffusion time and MRI growth time near the wave number of maximum MRI growth. When the latter time scale is shorter than the former, diffusion wins and the MRI does not grow. If the turbulent velocity forcing scale (a proxy for coherence) of initial fields is decreased or the field magnitude is decreased (by supplying an early unsaturated SSD state as an initial condition) then the fastest growing linear MRI mode moves to smaller scales where it has a harder time competing with diffusion. In short, if the initial coherence scales or strengths are too small, the MRI is quenched. Generally, when the initial conditions are supplied by a saturated SSD, we find the conditions to be favorable for growth but when the initial field has a Gaussian random noise field the MRI fails to grow. The SSD is a natural initial condition as it is likely to be common in sufficiently conducting plasmas as in stars or the galactic ISM that feed accretion flows. After the MRI takes over from the SSD supplied initial conditions, the saturated state of the magnetic and velocity fields in our simulations is essentially indistinguishable from that of the more commonly studied case of initial vertical fields of zero net flux. In particular the saturated amplitudes of the total magnetic and velocity fields (magnetic field being dominant), the accretion stresses, ratio of Maxwell to Reynolds stress, and the magnetic and kinetic power spectra are all very similar for the two aforementioned initial conditions. Finally, we hypothesize that if only the SSD produced tangled magnetic field were supplied as an initial condition with the velocity removed, the turbulent decay rate would be lowered because the pre-removal $U_{rms} \ge V_A$. In the absence of an initial velocity, the minimum decay time is an Alfv\'en crossing time, which is itself the MRI growth time at $k_{max}$. Thus the aforementioned minimum field strength and coherence requirements are, if anything, less stringent when the velocity is absent. Further exploration of this hypothesis, and the understanding of the large-scale dynamo seen in these simulations would be of interest to explore in future work. | 16 | 9 | 1609.06745 |
1609 | 1609.08791_arXiv.txt | We observed RZ LMi, which is renowned for the extremely ($\sim$19~d) short supercycle and is a member of a small, unusual class of cataclysmic variables called ER UMa-type dwarf novae, in 2013 and 2016. In 2016, the supercycles of this object substantially lengthened in comparison to the previous measurements to 35, 32, 60~d for three consecutive superoutbursts. We consider that the object virtually experienced a transition to the novalike state (permanent superhumper). This observed behavior extremely well reproduced the prediction of the thermal-tidal instability model. We detected a precursor in the 2016 superoutburst and detected growing (stage A) superhumps with a mean period of 0.0602(1)~d in 2016 and in 2013. Combined with the period of superhumps immediately after the superoutburst, the mass ratio is not as small as in WZ Sge-type dwarf novae, having orbital periods similar to RZ LMi. By using least absolute shrinkage and selection operator (Lasso) two-dimensional power spectra, we detected possible negative superhumps with a period of 0.05710(1)~d. We estimated the orbital period of 0.05792~d, which suggests a mass ratio of 0.105(5). This relatively large mass ratio is even above ordinary SU UMa-type dwarf novae, and it is also possible that the exceptionally high mass-transfer rate in RZ LMi may be a result of a stripped core evolved secondary which are evolving toward an AM CVn-type object. | SU UMa-type dwarf novae (DNe) are a class of cataclysmic variables (CVs) which are close binary systems transferring matter from a red dwarf secondary to a white dwarf, forming an accretion disk. In SU UMa-type dwarf novae, two types of outbursts are seen: normal outbursts and superoutbursts. Superoutbursts are defined by the presence of (positive) superhumps, which are humps with a period a few percent longer than the orbital period [for general information of CVs, DNe, SU UMa-type dwarf novae and superhumps, see e.g. \citet{war95book}]. RZ LMi is one of the most enigmatic SU UMa-type dwarf novae. This object was originally discovered as an ultraviolet-excess variable star (\cite{lip81FBSCV}; later given a designation of FBS 0948$+$344: \cite{abr93FBS}). \citet{lip81FBSCV} reported that the object had a spectral energy distribution of spectral classes O--B at maximum. When the object was fading or rising, strong emission lines appeared. \citet{lip81FBSCV} studied 16 objective prism plates and seven direct imaging plates and recorded strong variation with a range of 14--17 mag. The object was recorded at 17 mag on the Palomar Obervatory Sky Survey (POSS). Based on the colors and rapid variation within a few days, \citet{lip81FBSCV} suggested a dwarf nova-type classification. \citet{gre82PGsurveyCV} also selected it as an ultraviolet-excess object (PG 0948$+$344) and was spectroscopically confirmed as a CV. \citet{gre82PGsurveyCV} gave a variability range of $B$=14.4--16.8 without details (the minimum likely referred to the magnitude in POSS). Despite the finding by \citet{lip81FBSCV}, \citet{NameList67} gave a designation of RZ LMi as a novalike (NL) variable by referring to \citet{gre82PGsurveyCV}. \citet{kon84KUV2} also selected RZ LMi as an ultraviolet-excess object. The nature of the variability of this object remained unclear. Knowing the dwarf nova-type classification, one of the authors (TK) visually observed this object in 1987--1988 and found relatively stable $\sim$20~d cycle lengths but probably with additional short outbursts. This result was reported in the domestic variable star bulletin ``Henkousei'' (in Japanese). This report was probably the first to document the unusual cyclic variation of RZ LMi. Two teams photometrically observed this object and \citet{rob94rzlmi} reported an unusual long-term repetitive light curve with a stable period (it is now known as the supercycle) of 19.2~d. \citet{pik95rzlmi} reported that the behavior was different from a typical dwarf nova in that it showed short fading and brightening. \citet{pik95rzlmi} reported a possible period of 21.167~d or 23.313~d and classified it to be a VY Scl-type NL. A clue to understanding the unusual behavior of RZ LMi came from the discovery of the SU UMa-type dwarf nova ER UMa with an ultrashort supercycle (interval between superoutbursts) of 43~d in 1994 \citep{kat95eruma}. As orally presented in the conference held in Abano Terme, Italy, 20--24 June 1994, Robertson's team originally considered that the recurring NL-type bright state and the dwarf nova-type state in ER UMa and RZ LMi was the best example to show the consequences of recurring mass-transfer bursts from the secondary. In this conference after Robertson's presentation, however, Y. Osaki introduced Kato and Kunjaya's detection of superhumps in ER UMa and showed that the unusual behavior of ER UMa could be understood within the framework of the thermal-tidal instability (TTI) model \citep{osa89suuma} if one can allow an exceptionally high mass-transfer rate (cf. \cite{osa95eruma}). After this conference, Robertson's team published a paper following our interpretation (\cite{rob94rzlmi}; \cite{rob95eruma}).\footnote{ The dwarf nova-type nature was first clarified by \citet{iid94eruma}. \citet{mis95PGCV} also observed ER UMa in 1993--1994 and finally reached the same conclusion as \citet{kat95eruma}. } After the discovery of superhumps in ER UMa and the recognition of the similarity between ER UMa and RZ LMi, superhumps were naturally sought --- the competition was intense: both \citet{rob95eruma} and \citet{nog95rzlmi} observed the same superoutburst in 1995 and detected superhumps. These observations confirmed that RZ LMi does indeed belongs to the SU UMa-type dwarf novae. These objects, together with V1159 Ori, are usually called ER UMa-type stars (cf. \cite{kat99erumareview}). The mechanism of the ultrashort (19~d) supercycle and the unusually regular outburst pattern remained a mystery. \citet{rob94rzlmi} suspected a mechanism outside the disk in addition to the one in the disk. A later publication by \citet{ole08rzlmi} followed the former possibility and suggested a third body. In the standard TTI model, this short supercycle is difficult to reproduce, and \citet{osa95rzlmi} presented a working hypothesis that in RZ LMi the disk becomes thermally unstable during the superoutburst earlier than in other SU UMa-type dwarf novae. \citet{hel01eruma} suggested, following the interpretation in \citet{osa95rzlmi}, that in systems with very small mass ratios ($q$), the tidal torque is too small to maintain the superoutburst, and that there occurs a decoupling between the thermal and tidal instabilities. \citet{hel01eruma} suggested that one of the consequences of this decoupling can be found as the persistent superhumps after superoutbursts. According to the working hypotheses by \citet{osa95rzlmi} and \citet{hel01eruma}, it is strongly predicted that the $q$ is small in RZ LMi and that the disk radius after the superoutburst is larger than in other SU UMa-type dwarf novae. RZ LMi, however, has defied every attempt to determine the orbital period, since it mostly stays in the ``outburst'' state and it is difficult to make a radial-velocity study in short quiescence. The almost continuous presence of superhumps also made it difficult to detect potential orbital variations by photometric methods. Without the orbital period or a radial-velocity study, it remained impossible to observationally determine $q$ and the evolutionary status of RZ LMi remained unclear despite its unusual outburst properties. The situation dramatically changed after the detection of a possible change in the outburst pattern in the AAVSO observations (vsnet-alert 19524). We have conducted a world-wide campaign to observe RZ LMi during the 2016 season. The new development has also been helped by the classification of superhumps stage (A, B and C: \cite{Pdot}) and identification of stage A superhumps representing the growing phase of superhumps at the radius of the 3:1 resonance (\cite{osa13v344lyrv1504cyg}; \cite{kat13qfromstageA}). In this paper, key information and results are given in the main paper. The results not directly related to the conclusion of the paper, such as variation of superhump periods and variation of the profile of superhumps, are given in the Supplementary Information since they would provide useful information to expert readers. | \label{sec:results} \subsection{Outburst light curve and emergence of superhumps}\label{sec:outburst} Figure \ref{fig:rzlmihumpall2} illustrates the three consecutive superoutbursts in 2016 during which we obtained time-resolved photometric observations. The initial part of the light curve (upper panel) was the final fading part of the preceding superoutburst. There was only one normal outburst between the first and second superoutbursts. Although the initial part of the first superoutburst was not well sampled, we found the superoutburst lasted for 26~d. There were two normal outbursts between the second and the third superoutburst. The duration of the second superoutburst was 48~d, even longer than the first one. The intervals (supercycles) between the three superoutbursts were 32~d and 60~d. These values were 2--3 times longer than the historical supercycle (19~d) of this object (\cite{rob94rzlmi}; \cite{nog95rzlmi}). During the best observed second superoutburst, there was a shoulder (precursor outburst). Figure \ref{fig:rzlmihumpall2a} shows the initial part of this superoutburst. The precursor part is clearly seen between BJD 2457483 and 2457484. During this phase, superhumps rapidly grew to the maximum amplitude and they started to decay slowly (lower panell; see also figure \ref{fig:rzlmi2016rise} for an enlargement). This behavior is the same as in ordinary SU UMa-type dwarf novae [see e.g. figure 19 in \citet{Pdot4} for one of the best known SU UMa-type dwarf novae VW Hyi; see also figure 4 in \citet{osa13v1504cygKepler} for the corresponding part of the Kepler data of V1504 Cyg]. Similar precursors were also recorded in snapshot observations of the first superoutburst and less markedly (fewer observations) during the third superoutburst. These observations indicate that RZ LMi shows precursor--main superoutburst as in ordinary SU UMa-type dwarf novae at least in this season. \begin{figure*} \begin{center} \FigureFile(110mm,170mm){fig2.eps} \end{center} \caption{Light curve and $O-C$ diagram of superhumps in RZ LMi (2016). This figure is an enlargement of figure \ref{fig:rzlmihumpall2}. (Upper:) Light curve. The data were binned to 0.0023~d. The precursor part is clearly seen between BJD 2457483 and 2457484. During this phase, superhumps rapidly grew. (Middle:) $O-C$ diagram for the second superoutburst (filled circles). We used a period of 0.05955~d for calculating the $O-C$ residuals. (Lower:) Amplitudes of superhumps. The scale is linear and the pulsed flux is shown in a unit corresponding to 18 mag = 1. } \label{fig:rzlmihumpall2a} \end{figure*} \begin{figure*} \begin{center} \FigureFile(130mm,80mm){fig3.eps} \end{center} \caption{Growing superhumps and precursor at the start of the 2016 April superoutburst. The data were binned to 0.0023~d. Superhump grew between BJD 2457483 and 2457484. } \label{fig:rzlmi2016rise} \end{figure*} The variable supercycles now safely excludes the previously supposed stable clock (\cite{rob94rzlmi}; \cite{ole08rzlmi}) to produce regular superoutbursts. The present finding completely excludes the third body for producing such stable periodicity (cf. \cite{ole08rzlmi}). Furthermore, \citet{osa95eruma} already showed that the supercycle length once reaches the minimum as the mass-transfer rate ($\dot{M}$) increases, but that it lengthens again as $\dot{M}$ further increases. As $\dot{M}$ increases, the system eventually reaches the ``permanent outburst'' state (see figure 2 in \cite{osa95eruma}), bridging ER UMa-type objects and what is called permanent superhumpers (\cite{ski93bklyn}; \cite{pat99SH}). The current state of RZ LMi exactly reproduced this prediction, and it provides a strong support to explain the unusual short supercycles in ER UMa-type objects. \citet{osa95rzlmi} formulated the duration of a superoutburst ($t_{\rm supermax}$) as below: \begin{equation} t_{\rm supermax} \sim t_{\rm vis}[f_M/(1-1/e)][1-(\dot{M}/\dot{M}_{\rm crit})]^{-1/2}, \label{equ:tsuper} \end{equation} where $t_{\rm vis}$ is the viscous depletion timescale and $\dot{M}_{\rm crit}$ is the critical $\dot{M}$ to produce a hot, stable disk, respectively. The factor $f_M$ is the fraction of the disk mass accreted during a superoutburst. It is given by \begin{equation} f_M \simeq 1-(R_0/R_{\rm d,crit})^{3.0}, \label{equ:fracmass} \end{equation} where $R_0$ and $R_{\rm d,crit}$ represent the disk radius at the end of a superoutburst and at the start of a superoutburst (assuming that the disk critically reaches the radius of the 3:1 resonance at the start of a superoutburst), respectively. If we consider that $t_{\rm vis}$ and $R_0$ are the same between different superoutbursts of RZ LMi, we can estimate $\dot{M}$ during the current state. Using the parameters in \citet{osa95rzlmi}, $t_{\rm vis}$=11.2~d and an assumption of a large disk radius at the end of a superoutburst $R_0$=0.42$a$, where $a$ is the binary separation, the historical $t_{\rm supermax}$ of 6~d \citep{rob95eruma} is reproduced with $\dot{M}/\dot{M}_{\rm crit}$=0.5. The current $t_{\rm supermax}$ values of 26~d and 48~d require $\dot{M}/\dot{M}_{\rm crit}$=0.97 and 0.99, respectively. Although these estimates have uncertainties due to various assumptions, it is certain that the current state of RZ LMi is {\it critically} close to the stability border. If RZ LMi increases $\dot{M}$ further by 1\%, the object should become a permanent superhumper. As judged from these estimates, we have seen an almost complete transition from an ER UMa-type object to a permanent superhumper. There has been at least one case (BK Lyn) in which a permanent superhumper became an ER UMa-type object and then returned back (\cite{pat13bklyn}; \cite{Pdot4}) and the case of RZ LMi may not be special. \subsection{Growing (stage A) superhumps and post-superoutburst superhumps}\label{sec:stageA} As is best seen in the lower panel of figure \ref{fig:rzlmihumpall2a}, the amplitudes of superhumps rapidly grew in $\sim$20 cycles. During the initial $\sim$13 cycles, the $O-C$ values (middle panel) were negative and a characteristic kink around $E$=13 indicates that this object showed long-period (stage A) superhumps before entering the stable phase of stage B superhumps with a shorter period as in other SU UMa-type dwarf novae (see e-tables for the full list of times of superhump maxima). This identification appears particularly confident since it has been recently demonstrated that ER UMa, the prototype of ER UMa-type objects, showed the same pattern with a precursor outburst \citep{ohs14eruma}. The period of stage A superhumps from the times of maxima ($E \le$ 13) is 0.0602(4)~d. By using the data between BJD 2457483.01 and 2457484.40, we have obtained a period of 0.0601(1)~d with the phase dispersion minimization (PDM; \cite{PDM}) method (figure \ref{fig:rzlmi2016shapdm}). The errors are 1$\sigma$ estimated by the methods of \citet{fer89error} and \citet{Pdot2}. We consider that the result by the PDM method is more reliable than from superhump maxima since it gives a smaller error and adopted this period. \begin{figure} \begin{center} \FigureFile(80mm,110mm){fig4.eps} \end{center} \caption{Stage A superhumps in RZ LMi (2016). The data between BJD 2457483.01 and 2457484.40 were used. (Upper): PDM analysis. We analyzed 100 samples which randomly contain 50\% of observations, and performed PDM analysis for these samples. The bootstrap result is shown as a form of 90\% confidence intervals in the resultant PDM $\theta$ statistics. (Lower): Phase-averaged profile.} \label{fig:rzlmi2016shapdm} \end{figure} Although the second superoutburst was generally well observed, the observations were unfortunately relatively sparse around stage A. The 2013 April superoutburst of this object was relatively well observed in the same phase, although the brightness peak was missed and instead of a distinct precursor, there was a stagnation in the rising phase, which was likely an embedded precursor (figures \ref{fig:rzlmi2013humpall}, \ref{fig:rzlmi2013rise}; see e-table for the full times of superhump maxima). During the observation in 2013, a continuous light curve of growing superhumps (0$\le E \le$4) was well recorded. Using the data between BJD 2456381.41 and 2456381.88, we obtained a period of 0.0602(3)~d with the PDM method (figure \ref{fig:rzlmi2013shapdm}). Since the value is consistent with the 2016 one, we adopted an averaged value of 0.0602(1)~d as the period of stage A superhumps. \begin{figure} \begin{center} \FigureFile(80mm,110mm){fig5.eps} \end{center} \caption{$O-C$ diagram of superhumps in RZ LMi in 2013 April. (Upper:) Light curve. The data were binned to 0.0023~d. (Middle:) $O-C$ diagram (filled circles). We used a period of 0.05945~d for calculating the $O-C$ residuals. (Lower:) Amplitudes of superhumps. The scale is linear and the pulsed flux is shown in a unit corresponding to 18 mag = 1. } \label{fig:rzlmi2013humpall} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,70mm){fig6.eps} \end{center} \caption{Growing superhumps and stagnation (embedded precursor) at the start of the 2013 April superoutburst. The data were binned to 0.0023~d. Superhump grew between BJD 2456381.41 and 2456381.88. } \label{fig:rzlmi2013rise} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,110mm){fig7.eps} \end{center} \caption{Stage A superhumps in RZ LMi (2013 April). The data between BJD 2456381.41 and 2456381.88 were used. (Upper): PDM analysis. (Lower): Phase-averaged profile.} \label{fig:rzlmi2013shapdm} \end{figure} The dynamical precession rate, $\omega_{\rm dyn}$ in the disk can be expressed by (see, \cite{hir90SHexcess}): \begin{equation} \label{equ:precession} \omega_{\rm dyn}/\omega_{\rm orb} = Q(q) R(r), \end{equation} where $\omega_{\rm orb}$ and $r$ are the angular orbital frequency and the dimensionless radius measured in units of the binary separation $a$. The dependence on $q$ and $r$ are \begin{equation} \label{equ:qpart} Q(q) = \frac{1}{2}\frac{q}{\sqrt{1+q}}, \end{equation} and\footnote{ There was a typographical error in the second line of equation (1) in \citet{kat13qfromstageA}. The correct formula is $\frac{q}{\sqrt{1+q}} \bigl[\frac{1}{4}\sqrt{r} b_{3/2}^{(1)}\bigr]$. The results (including tables) in \citet{kat13qfromstageA} used the correct formula and the conclusion is unchanged. We used the correct formula in equation (\ref{equ:rpart}) in this paper. The same correction of the equation should be applied to \citet{kat13j1222}, \citet{nak13j2112j2037} and \citet{kat15wzsge}. } \begin{equation} \label{equ:rpart} R(r) = \frac{1}{2}\sqrt{r} b_{3/2}^{(1)}(r), \end{equation} where $\frac{1}{2}b_{s/2}^{(j)}$ is the Laplace coefficient \begin{equation} \label{equ:laplace} \frac{1}{2}b_{s/2}^{(j)}(r)=\frac{1}{2\pi}\int_0^{2\pi}\frac{\cos(j\phi)d\phi} {(1+r^2-2r\cos\phi)^{s/2}}. \end{equation} This $\omega_{\rm dyn}/\omega_{\rm orb}$ is equal to the fractional superhump excess in frequency: $\epsilon^* \equiv 1-P_{\rm orb}/P_{\rm SH}$, where $P_{\rm orb}$ and $P_{\rm SH}$ are the orbital period and superhump period, respectively. If $P_{\rm orb}$ is known, we can directly determine $q$ from the observed $\epsilon^*$ of stage A superhumps under the assumption that the period of stage A superhumps reflects the purely dynamical precession rate at the radius of the 3:1 resonance \citep{kat13qfromstageA}. Since the orbital period of RZ LMi is not known, we cannot directly apply the method in \citet{kat13qfromstageA} to dynamically determine $q$. We can instead use the period of post-superoutburst superhumps to constrain $q$ and the disk radius as introduced in \citet{kat13j1222}: \begin{equation} \label{equ:epsstagea} \epsilon^*({\rm stage A}) = Q(q) R(r_{\rm 3:1}) \end{equation} and \begin{equation} \label{equ:epspost} \epsilon^*({\rm post}) = Q(q) R(r_{\rm post}), \end{equation} where $r_{\rm 3:1}$ is the radius of the 3:1 resonance \begin{equation} \label{equ:radius31} r_{3:1}=3^{(-2/3)}(1+q)^{-1/3}, \end{equation} $\epsilon^*({\rm post})$ and $r_{\rm post}$ are the fractional superhump excess and disk radius immediately after the outburst, respectively. By solving equations (\ref{equ:epsstagea}) and (\ref{equ:epspost}) simultaneously, we can obtain the relation between $r_{\rm post}$ and $q$. If we have knowledge about $r_{\rm post}$, as determined in other systems in \citet{kat13qfromstageA}, we have a more stringent constraint. On 2013 March 20, the object was observed in quiescence just following a superoutburst starting on March 6. The object displayed post-superoutburst superhumps and there was a continuous run covering six cycles. A PDM analysis of this continuous run yielded a period of 0.0594(2)~d (figure \ref{fig:rzlmi2013shpostpdm}). Since this variation was also present after one more normal outburst, and since the decaying amplitudes of this variation exclude the possibility of the orbital variation, we identified this period as that of post-superoutburst superhumps. We combined the quiescent data on March 20 and 24 and obtained a period of 0.05969(2)~d (figure \ref{fig:rzlmi2013shpostpdmall}, assuming that the superhump phase and period did not change during a normal outburst). The shorter one-day alias ($\sim$0.0584~d) is too close to the supposed orbital period (discussed later) and this alias is unlikely. Using the periods of stage A superhumps and post-superoutburst superhumps, the relation between $r_{\rm post}$ and $q$ is shown in figure \ref{fig:qrpost}. The measurements of $r_{\rm post}$ in SU UMa-type dwarf novae using the same method are within the range of 0.30 and 0.38 \citep{kat13qfromstageA}. The smaller values represent the values for WZ Sge-type dwarf novae with multiple rebrightenings (after such rebrightenings), and it is extremely unlikely the case for RZ LMi. If $r_{\rm post}$ is around 0.38, $q$ is estimated to be 0.06(1). It would be noteworthy that \citet{osa95rzlmi} assumed $r_{\rm post}$=0.42 for this particular object. If it is indeed the case, $q$ needs to be as large as 0.10(2). This consideration suggests that the expectations by \citet{osa95rzlmi} and \citet{hel01eruma} that RZ LMi has a very small $q$ and the large disk radius immediately following a superoutburst are not true at the same time: either $q$ is higher or the disk radius is smaller. This is the important conclusion from the present observation. Since stage A superhumps were not very ideally observed in the present study, further observations are needed to refine the result. \begin{figure} \begin{center} \FigureFile(80mm,110mm){fig8.eps} \end{center} \caption{Post-superoutburst superhumps in RZ LMi (2013 March 20). The data between BJD 2456371.5 and 2456371.9 were used. (Upper): PDM analysis. (Lower): Phase-averaged profile.} \label{fig:rzlmi2013shpostpdm} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,110mm){fig9.eps} \end{center} \caption{Post-superoutburst superhumps in RZ LMi (2013 March 20 and 24). (Upper): PDM analysis. (Lower): Phase-averaged profile.} \label{fig:rzlmi2013shpostpdmall} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,70mm){fig10.eps} \end{center} \caption{Relation between $q$ and $r_{\rm post}$ using the periods of stage A superhumps and post-superoutburst superhumps. Dashed curves represent the range of 1$\sigma$ errors.} \label{fig:qrpost} \end{figure} \subsection{Possible negative superhumps and orbital period}\label{sec:negSH} We used least absolute shrinkage and selection operator (Lasso) method (\cite{lasso}; \cite{kat12perlasso}), which has been proven to yield very sharp signals. We used the two-dimensional Lasso power spectra as introduced in the analysis of the Kepler data such as in \citet{kat13j1924}; \citet{osa13v344lyrv1504cyg}. These two-dimensional Lasso power spectra have been proven to be very effective in detecting signals in non-uniformly sampled ground-based data (see e.g. \cite{ohs14eruma}) since Lasso-type period analysis is less affected by aliasing than traditional Fourier-type power spectra. This characteristic has enabled detection of negative superhumps in ground-based observations (e.g. \cite{ohs14eruma}; \cite{Pdot6}). The result for RZ LMi in 2016 is shown in figure \ref{fig:rzlmi2016lasso}. The strong persistent signal around 16.8 cycle/day (c/d) is superhumps. A weaker signal around 17.50--17.55 c/d between BJD 2457510 and 2457530 is possible negative superhumps. Since the second superoutburst lasted very long, it may have been possible that negative superhumps were excited during this long-lasting standstill-like phase, which gives the condition almost the same as in permanent superhumpers (see subsection \ref{sec:outburst}). A PDM analysis of the data between BJD 2457510 and 2457530 yielded a period of 0.05710(1)~d. The superhump period in this interval was 0.059555(4)~d. It has been widely accepted that absolute fractional superhumps excesses ($\epsilon_+$ for positive superhumps, $\epsilon_-$ for negative superhumps, where $\epsilon \equiv P_{\rm SH}/P_{\rm orb}-1$) are tightly correlated when both signals are simultaneously seen in novalike CVs (e.g. \cite{pat97v603aql}; \cite{mon09diskprecession}). The empirical relation is $\epsilon_+ \simeq 2 |\epsilon_-|$. If it is also the case for RZ LMi, $P_{\rm orb}$ is expected to be around 0.05792~d. This period is labeled as Porb in figure \ref{fig:rzlmi2016lasso}. If this is indeed the orbital period, the consequence is important (cf. subsection \ref{sec:stageA}). The measured period of stage A superhumps [0.0602(1)~d] gives $\epsilon^*$=0.038(2), which is equivalent to $q$=0.105(5) (see table 1 in \cite{kat13qfromstageA}). This value is higher than the typical ones for WZ Sge-type dwarf novae, which have similar $P_{\rm orb}$ as RZ LMi (\cite{kat13qfromstageA}; \cite{kat15wzsge}; \cite{Pdot8}). This $q$ measurement invalidates the suggestion by \citet{osa95rzlmi} and \citet{hel01eruma} that the unusual properties of RZ LMi is a result of the very low $q$. The $q$, however, is consistent with the relation derived from the period of stage A and post-superoutburst superhumps (subsection \ref{sec:stageA}, figure \ref{fig:qrpost}) assuming the large disk radius at the end of a superoutburst. If this interpretation is correct, the disk radius at the end of a superoutburst is large as required by \citet{osa95rzlmi}, but the origin of the large disk (i.e. superoutbursts terminate earlier than in ordinary SU UMa-type dwarf novae) cannot be attributed to an exceptionally small $q$. In recent years, some SU UMa-type dwarf novae show early termination of superoutbursts (such as V1006 Cyg, \cite{kat16v1006cyg}). In \citet{kat16v1006cyg}, the termination may be associated with the supposed appearance of stage C superhumps (late-stage superhumps with a shorter constant period). The origin of stage C superhumps is still poorly known, the reason of premature termination of superoutbursts which may be related to the evolution of stage C superhumps still needed to be clarified. There is supporting evidence for a large $q$: the duration of stage A superhumps, which is considered to reflect the growth time of the 3:1 resonance, is very short (less than 1~d) in RZ LMi. From the theoretical standpoint, this growth time is expected to be proportional to $1/q^2$ \citep{lub91SHa}, and this relation has been confirmed in WZ Sge-type dwarf novae \citep{kat15wzsge}. As judged from the rapid growth of superhumps in RZ LMi, it appears extremely unlikely that RZ LMi has $q$ as small as in WZ Sge-type dwarf novae. \begin{figure} \begin{center} \FigureFile(80mm,95mm){fig11.eps} \end{center} \caption{Two-dimensional Lasso period analysis of RZ LMi (2016). (Upper:) Light curve. The data were binned to 0.02~d. (Lower:) Lasso period analysis. The strong persistent signal (disturbed during the starting phase of the superoutburst due to the non-sinusoidal profile and rapidly varying periods) around 16.8 c/d is superhumps. A weaker signal around 17.50--17.55 c/d between BJD 2457510 and 2457530 is possible negative superhumps. $\log \lambda=-6.4$ was used. Porb represents the orbital period we suggest. The width of the sliding window and the time step used are 18~d and 0.8~d, respectively. } \label{fig:rzlmi2016lasso} \end{figure} \subsection{Evolutionary status}\label{sec:evol} If the $q$ derived from the possible negative superhumps and suggested by the rapid growth of superhumps (subsection \ref{sec:negSH}) is correct, RZ LMi cannot be an object close to the period minimum or a period bouncer --- the idea conjectured by \citet{ole08rzlmi} that some ER UMa-type are period bouncers. The $q$ we suggested is similar to or even larger than those of ordinary SU UMa-type dwarf novae having similar $P_{\rm orb}$ (figure \ref{fig:qall5add}). We know at least another object GALEX J194419.33$+$491257.0 in the Kepler field which has a very short $P_{\rm orb}$ of 0.0528164(4)~d and very frequent outbursts (normal outbursts with intervals of 4--10~d) \citep{kat14j1944}. This object has an unusually high $q$=0.141(2) measured using stage A superhumps. These properties are somewhat similar to those of RZ LMi. \citet{kat14j1944} suggested the possibility that GALEX J194419.33$+$491257.0 may be a CV with a stripped core evolved secondary which are evolving toward AM CVn-type CVs. Such a condition might be a fascinating possibility to explain why RZ LMi has an exceptionally high $\dot{M}$ for its $P_{\rm orb}$. \begin{figure} \begin{center} \FigureFile(80mm,70mm){fig12.eps} \end{center} \caption{Location of RZ LMi on the diagram of mass ratio versus orbital period. The dashed and solid curves represent the standard and optimal evolutionary tracks in \citet{kni11CVdonor}, respectively. The filled circles, filled squares, filled stars, filled diamonds represent $q$ values from a combination of the estimates from stage A superhumps published in four preceding sources (\cite{kat13qfromstageA}; \cite{nak13j2112j2037}; \cite{Pdot5}; \cite{Pdot6}; \cite{Pdot7}; \citet{Pdot8}), known $q$ values from quiescent eclipses or radial-velocity study (see \cite{kat13qfromstageA} for the data source). In addition to the references listed in \citet{Pdot8}, we supplied the data for SDSS J143317.78$+$101123.3 in \citet{her16j1433}.} \label{fig:qall5add} \end{figure} \subsection{Secular variation of supercycle}\label{sec:secular} As discussed in subsection \ref{sec:outburst}, it is likely that RZ LMi changed $\dot{M}$ by a factor of $\sim$2 in the last two decades. In recent years, a transition from the novalike (permanent superhumper) state to the ER UMa-type state was discovered in BK Lyn (\cite{pat13bklyn}; \cite{Pdot4}). \citet{pat13bklyn} proposed, based on the potential identification with an ancient classical nova in 101, that ER UMa stars are transitional objects during the cooling phase of post-eruption classical novae [the idea was not new and it was already proposed in \citet{kat95eruma} and \citet{osa95eruma}]. Following this interpretation, \citet{otu13suumacycle} studied ER UMa-type objects and found a secular increase of the supercycle in most of typical ER UMa-type objects. RZ LMi was included, and \citet{otu13suumacycle} gave $\dot{P}$ of the supercycle of (5.0$\pm$1.9)$\times 10^{-4}$ in 18 yr. We should note that \citet{zem13eruma} also studied variation of supercycles in ER UMa and reported that supercycles vary in a range of 43.6--59.2~d with shorter time-scales of 300--1900~d. A secular increase of the supercycle was also statistically meaningful \citep{zem13eruma}. The conclusion of \citet{otu13suumacycle}, however, was only based on the variation of supercycles, and \citet{otu13suumacycle} disregarded the possibility that the supercycle can also increase if $\dot{M}$ increases toward $\dot{M}_{\rm crit}$ since there is a minimum in the supercycle as $\dot{M}$ increases (cf. the right branch of figure 1 in \cite{osa95eruma}; see also subsection \ref{sec:outburst}). Since it is apparently the case for RZ LMi, we studied secular variation of supercycles in RZ LMi. We summarized the result in table \ref{tab:rzlmisupercycle} [the data in \citet{Pdot4} also used AAVSO observations]. When raw data are not available, we measured the fraction of superoutburst and duty cycle by eye from the figures in the papers. The data for 1986--1989 (Shugarov) were too sparse to determine the supercycle and only the duty cycle was estimated. The data for 1987--1988 (Kato) were visual observations. The supercycle was determined from seven well-defined bright outbursts. The duty cycle was probably underestimated due to the insufficient detection limit of visual observations. For AAVSO observations, the fraction of superoutburst could be determined only to 0.05 since most of the data were randomly sampled snapshot observations. It has became apparent that the supercycle was not stable as had been supposed in \citet{rob95eruma} and \citet{ole08rzlmi}. The supercycle at the time of \citet{rob95eruma} probably reached the historical minimum. It was likely the supercycle once lengthened as long as to $\sim$24~d before 2013, but it again returned to 19--20~d. This behavior probably gave the impression that the supercycle is globally constant if seen in long time scales, such as several years. The situation changed in 2016 when the supercycle strongly increased. This increase was associated with the increase of the fraction of superoutburst and the duty cycle, indicating that $\dot{M}$ increased despite the lengthening of the supercycle (on the contrary to the conclusion by \cite{otu13suumacycle}). Although the situation before 2016 was less clear, the changing supercycle suggests fluctuating $\dot{M}$ within time scales of a year or two. The outburst pattern was generally regular in the second best observed season in 2013 (figure \ref{fig:rzlmi2013all}). The duration of the superoutburst was appreciably longer between BJD 2456380 and 2456400 (17~d, in contrast to 12--13~d in other superoutbursts in 2013; a mean value of supercycle in 2013 is given in table \ref{tab:rzlmisupercycle}). It was likely that $\dot{M}$ temporarily increased also in 2013. In the 2013--2014 season, the durations of superoutbursts again decreased to shorter than 12~d (figure \ref{fig:rzlmi2014all}) while supercycle lengths remained short ($\sim$21~d). \begin{figure} \begin{center} \FigureFile(80mm,70mm){fig13.eps} \end{center} \caption{Overall light curve of RZ LMi in 2013. The data were binned to 0.03~d. } \label{fig:rzlmi2013all} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,100mm){fig14.eps} \end{center} \caption{Overall light curve of RZ LMi in 2013--2014. The data were binned to 0.03~d. } \label{fig:rzlmi2014all} \end{figure} \begin{table*} \caption{Supercycles of RZ LMi}\label{tab:rzlmisupercycle} \begin{center} \begin{tabular}{ccccccc} \hline Year & JD range\commenta & supercycle (d) & fraction of superoutburst & duty cycle\commentb & nights & source \\ \hline 1984 & 45699--45851 & 19.6(1) & 0.4 & 0.53 & 30 & this work (Shugarov) \\ 1984--1985 & 46019--46230 & 19.5--21.5 & 0.35: & 0.39 & 31 & this work (Shugarov) \\ 1986--1989 & 46468--47682 & -- & -- & 0.6 & 7 & this work (Shugarov) \\ 1987--1988 & 47151--47318 & 22.5(9) & -- & 0.39 & 52 & this work (Kato) \\ 1992--1995 & 48896--49723 & 18.87 & 0.4 & 0.5 & -- & \citet{rob95eruma} \\ 2004--2005 & 53027--53519 & 19.07(4) & 0.4 & 0.4 & -- & \citet{ole08rzlmi} \\ 2005--2006 & 53644--53884 & 19.7(1) & 0.45 & 0.58 & 89 & AAVSO \\ 2006--2007 & 54041--54254 & 20.75(4) & 0.40 & 0.55 & 113 & AAVSO \\ 2007--2008 & 54374--54626 & 21.3(1) & 0.45 & 0.55 & 100 & AAVSO \\ 2008--2009 & 54749--54988 & 21.7(1) & 0.40 & 0.31 & 96 & AAVSO \\ 2009--2010 & 55146--55334 & 24.3(1) & 0.35 & 0.52 & 58 & AAVSO \\ 2012 & 55985--56070 & 21.61(2) & 0.55 & 0.46 & 76 & \citet{Pdot4} \\ 2012--2013 & 56273--56473 & 22.83(1) & 0.50 & 0.62 & 149 & this work \\ 2013--2014 & 56563--56819 & 20.8(1) & 0.40 & 0.48 & 151 & this work (Neustroev), AAVSO \\ 2014--2015 & 56930--57178 & 19.9(1) & 0.45 & 0.55 & 96 & AAVSO \\ 2016 & 57416--57451 & 35 & 0.60 & 0.50 & 28 & this work \\ 2016 & 57451--57483 & 32 & 0.81 & 0.74 & 31 & this work \\ 2016 & 57483--57543 & 60 & 0.80 & 0.83 & 54 & this work \\ \hline \multicolumn{5}{l}{\commenta JD$-$2400000.} \\ \multicolumn{5}{l}{\commentb Object brighter than 15.0 mag.} \\ \end{tabular} \end{center} \end{table*} As far as RZ LMi is concerned, $\dot{M}$ is not secularly {\rm decreasing} as in the scenario by \citet{pat13bklyn}. Among other ER UMa-type objects, BK Lyn returned back to its original novalike state in late 2013 and the ER UMa-type state was only transiently present \citep{Pdot4}. As discussed in \citet{Pdot4}, the hypothetical cooling sequence from novalike objects -- ER UMa-type dwarf novae -- ordinary SU UMa-type dwarf novae after nova eruptions as proposed by \citet{pat13bklyn} is not very consistent with observational statistics, since such a cooling sequence would predict a large number of intermediate (having much more slowly cooling white dwarfs) objects between ER UMa-type dwarf novae and ordinary SU UMa-type dwarf novae, while observations could not confirm a large number of such objects. As seen in ER UMa, BK Lyn and RZ LMi, the $\dot{M}$ variations look more irregular with time-scales of several years and it looks like that the majority of variations in outburst activities in these systems does not reflect the secular $\dot{M}$ variation as proposed by \citet{pat13bklyn}. As this paper has shown, the high activity of RZ LMi may be a result of a rare evolutionary condition with a relatively massive secondary. In such a case, we may not require a nova eruption to produce an object like RZ LMi. We observed RZ LMi, which is renowned for the extremely ($\sim$19~d) short supercycle and is a member of a small, unusual class of cataclysmic variable called ER UMa-type dwarf novae, in 2013 and 2016. In 2016, the supercycles of this object substantially lengthened in comparison to the previous measurements to 35, 32, 60~d for three consecutive superoutbursts. The durations of superoutbursts also lengthened and they composed 60--81\% of the supercycle. Such long durations of superoutbursts have never been observed in this object, and we consider that the object virtually experienced a transition to the novalike state (permanent superhumper). This observed behavior extremely well reproduced the prediction of the thermal-tidal instability model by \citet{osa95eruma}. We estimated that in 2016, the mass-transfer rate of RZ LMi reached 97--99\% of the thermal stability limit, and that it was about two times larger than the value in the past. We detected a precursor in the 2016 superoutburst and detected growing (stage A) superhumps in 2016 and in 2013. We estimated their period to be 0.0602(1)~d. This makes the second case in ER UMa-type dwarf novae in which growing superhumps were confidently recorded during the sequence of the precursor and the main superoutburst. We also detected post-superoutburst superhumps immediately after a superoutburst in 2013 with a period of 0.05969(2)~d. Since both stage A superhumps and post-superoutburst superhumps can be considered to reflect the dynamical precession rates, we can derive the relation between the mass ratio and the radius of the post-superoutburst. The result suggests that the mass ratio is not as small as in WZ Sge-type dwarf novae, having orbital periods similar to RZ LMi. By using least absolute shrinkage and selection operator (Lasso) two-dimensional power spectra, we detected possible negative superhumps with a period of 0.05710(1)~d. By combination with the period of ordinary superhumps, we estimated the orbital period of 0.05792~d. The period of stage A superhumps, combined with this orbital period, suggests a mass ratio of 0.105(5). This relatively large mass ratio is consistent with the rapid growth of superhumps. This mass ratio is even above ordinary SU UMa-type dwarf novae, and it is also possible that the exceptionally high mass-transfer rate in RZ LMi may be a result of a stripped core evolved secondary in a system evolving toward an AM CVn-type object. An analysis of historical records of supercycles in this system suggests that the variation of the outburst activity is sporadic with time-scales of years and it does not seem to reflect the secular variation caused by evolution. | 16 | 9 | 1609.08791 |
1609 | 1609.08875_arXiv.txt | {We use radio-continuum all-sky surveys at 1420 and 408 MHz with the aim to investigate properties of the Galactic radio source Lupus Loop. The survey data at 1435 MHz, with the linear polarization of the southern sky, is also used. We calculate properties of this supernova remnant: the brightness temperature, surface brightness and radio spectral index. For determining borders and calculation of its properties, we use the method we have developed. The non-thermal nature of its radiation is confirmed. The distribution of spectral index over its area is also given. A significant correlation between the radio spectral index distribution and the corresponding polarized intensity distribution inside the loop borders is found, indicating that the polarization maps could provide us information about the distribution of interstellar medium, and thus could represent one additional way to search for new Galactic loops.} \resumen{} \addkeyword{ISM: supernova remnants} \addkeyword{radiation mechanisms: non-thermal} \addkeyword{radio continuum: ISM} \begin{document} | Radio surveys of the region in vicinity of the supernova of 1006 A.D. revealed a plateau or spur running out from the galactic plane near $l$ = 330$^\circ$ which contains two shell-like objects: Lupus Loop and SN1006 \citep{miln71}. The general diffuse appearance of the Lupus Loop suggested that it is the remnant of a very old supernova (which is consistent with its low surface brightness) and these two remnants are not associated in any way \citep{miln74}. It was indicated earlier \citep{spoe73} that expanding sphere of the supernova remnants (SNRs) leads to compression of the interstellar magnetic field, which results in an observable radio source, and that the spatial orientation of the loops contains information on the direction of the magnetic field of the undisturbed medium outside the shell. \citet{spoe73} treated the Lupus Loop as an object similar to the main Galactic loops, and found that it indicates magnetic field direction parallel to field found from Loop I. Radio continuum observations of this source are given in \citet{miln71,miln74}, radio line observations in \citet{colo82}, $X$-ray observations and its spectrum can be found in \citet{toor80,leah91,ozak94,kapl06} and references therein, while far $UV$ observations are presented in \citet{shin06}. The analysis on the filamentary structure observed in polarization by WMAP (Wilkinson Microwave Anisotropy Probe) satellite is given in \citet{vida15,vida16}. It is described there that most of the polarized emission (at high latitudes) comes from individual filamentary features, and some of these structures are the well-known continuum radio loops. Using WMAP data at 23, 33 and 41 GHz, they studied the diffuse polarized emission over the entire sky, and they obtained the (average) polarization spectral indices which are consistent with synchrotron radiation. A catalogue of Galactic SNRs, with some statistics of their parameters, is presented in \citet{gree14a}, along with more detailed web-based version \citep{gree14b}. The current version of the catalogue contains 294 SNRs, and is based on research in the published literature up to the end of 2013. Lupus Loop, with catalogue name G330.0+15.0, is listed there, and some of its parameters are given. This low surface brightness loop has been observed in radio and $X$-ray wavelength range \citep{gree14b}. Our aim is to study the properties of this remnant, and to calculate radio spectral index using the method we have previously developed. Our method of calculation is explained in detail in \citet{bork12b} and references therein: we investigated Galactic Loops I-VI in papers \citet{bork06a,bork06b,bork07}, \newline \citet{bork08a,bork10,uros11}, and we investigated smaller remnants in \citet{bork08b,bork09a,bork09b,bork10,bork11}, \newline \citet{bork12a}. This method is applicable to extragalactic radio sources as well \citep{bork12c}. In this paper we also want to investigate the nature of its radiation and to study how spectral index varies across the face of the remnant. Besides, we want to analyze and study the polarization of this SNR and to investigate its connection with the spectral index. \begin{figure} \centering \includegraphics[width=0.45\textwidth]{fig1a.eps} \hspace{0.5cm} \includegraphics[width=0.45\textwidth]{fig1b.eps} \caption{Lupus Loop area with brightness temperature contour at 1420 MHz \textbf{(left)} and 408 MHz \textbf{(right)}. The contour is representing $T_{min}$ as given in the Table 1. Below, the temperature scales are given (in mK).} \label{fig01} \end{figure} \begin{figure} \centering \includegraphics[width=0.95\textwidth]{fig2.eps} \caption{The 3D map of the Lupus Loop and its surrounding at 1420 MHz. The brightness temperature is given in mK.} \label{fig02} \end{figure} \begin{figure} \centering \includegraphics[width=0.48\textwidth]{fig3a.eps} \hspace{0.2cm} \vspace{0.5cm} \includegraphics[width=0.48\textwidth]{fig3b.eps} \includegraphics[width=0.48\textwidth]{fig3c.eps} \hspace{0.2cm} \includegraphics[width=0.48\textwidth]{fig3d.eps} \caption{\textbf{Left:} 1420 MHz temperature profiles for area containing Lupus Loop at galactic longitude $b$ = 15$^\circ$ (top) and 17$^\circ$ (bottom). \textbf{Right:} 408 MHz temperature profiles for $b$ = 15$^\circ$ (top) and 17$^\circ$ (bottom).} \label{fig03} \end{figure} | As we showed earlier (\citet{bork12b} and references therein), the method for defining a loop border and for determining the values of brightness temperature and surface brightness, which we developed for main Galactic Loops I-VI, could be applicable to all SNRs. Here we use this method in order to: \begin{itemize} \item determine brightness temperature borders of the Lupus Loop at 1420 and 408 MHz, \item calculate the mean radio spectral index between the specified frequencies, as well as the distribution of indices across the face of this remnant, \item study the correlation between the radio spectral index distribution and the corresponding polarized intensity distribution within the given borders. \end{itemize} In the frequency range under consideration synchrotron radiation dominates the spectrum. We used the radio spectral index to study the radiation mechanism of this radio source. The value obtained for the spectral index (which is $>$ 0.1) confirmed non-thermal emission of radiation for this source. The main disagreement in the measured values can probably be caused by differences in the chosen area for Lupus Loop border. These new observations yielded value of $\alpha$ greater than \citet{miln74}. Besides the nature of the radiation, we also showed how spectral index varies across the face of the remnant. Taking into account that SNRs radiate non-thermal (synchrotron) radiation which is mainly caused by the magnetic field, which on the other hand is also responsible for the polarization of radiation, we supposed that there exists the connection between the polarization and radio spectral index $\alpha$, which we then showed. We can conclude that spectral index significantly varies across the Lupus Loop. Over the time, the synchrotron spectral index becomes steeper (gets greater value, i.e. loops steepen as they age). The boundary of the Lupus Loop is not well defined and the ISM is rather inhomogeneous. That is why there are significant variations in spectral indices $\alpha$ over the loop area. \textbf{Acknowledgments.} This research is part of the project 176003 ''Gravitation and the large scale structure of the Universe'' supported by the Ministry of Education, Science and Technological Development of the Republic of Serbia. | 16 | 9 | 1609.08875 |
1609 | 1609.04576_arXiv.txt | We describe general features that might be observed in the line spectra of relic cosmological particles should quantum nonequilibrium be preserved in their statistics. According to our arguments, these features would represent a significant departure from those of a conventional origin. Among other features, we find a possible spectral broadening (for incident photons) that is proportional to the energy resolution of the recording telescope (and so could be orders of magnitude larger than any intrinsic broadening). Notably, for a range of possible initial conditions we find the possibility of spectral line `narrowing' whereby a telescope could observe a spectral line which is narrower than it should conventionally be able to resolve. We briefly discuss implications for the indirect search for dark matter. | \label{sec_1} In the de Broglie-Bohm pilot-wave formulation of quantum theory \cite{deB28,BV09,B52a,B52b,Holl93}, the Born probability rule has been shown to arise spontaneously through quantum `relaxation' \cite{AV91a,AV92,AV01,VW05,SC12,EC06,TRV12,ACV14}--a dynamical process which is broadly similar to thermal relaxation in classical physics. The resulting relaxed or `equilibrium' state obeys the Born rule so that, following relaxation, the theory becomes experimentally indistinguishable from conventional quantum theory. If one is to regard this relaxation--which occurs without the need for additional postulates--as the ultimate cause of conventional quantum probabilities, the question arises as to what preceded the relaxation. As such, pilot-wave theory allows for arbitrary ensemble probabilities \cite{AV91a,AV91b,AV92,AV96,AV01,AV02,AV07,AV08,AV09,AV10,AVPwtMw,PV06} and consequently may be regarded as a more general theory of which standard quantum theory is a special equilibrium case. A straightforward corollary to such a viewpoint is that `quantum nonequilibrium'--defined generically as nonconformance to the Born rule--may have existed in the early universe, prior to relaxation \cite{AV91a,AV91b,AV92,AV96}. If so, then such primordial quantum nonequilibrium may have left traces that are still observable today, for example in the cosmic microwave background \cite{AV07,AV10,CV13,CV15}. Potentially it is also possible that quantum nonequilibrium may have survived in the statistical properties of some species of relic cosmological particles \cite{AV01,AV07,AV08,AV10,UV15}. In a previous article \cite{UV15} possible avenues through which nonequilibrium in particles could persist to this day were given and it was argued that detection of such relic nonequilibrium is in principle possible. This opens up the prospect that quantum nonequilibrium may play a role in contemporary experimentation--for example in the indirect search for dark matter. As yet however, it is unclear how this role would play out. The purpose of this article is to present a field theoretical account of the behaviour of quantum nonequilibrium under measurement that, whilst still far from an accurate description of the true workings of contemporary experiments (telescopes), takes at least a small step in this direction. Specifically, and for the reasons described below, we will take the example of telescopes tasked with the indirect detection of dark matter through spectral measurement of astrophysical photons, of which perhaps the best known is the Fermi-LAT \cite{LAT09}. We will draw comparisons to such telescopes throughout and section \ref{implications} will be devoted to a discussion of the possible implications for the indirect search for dark matter. To structure our discussion, and to provide explicit calculations demonstrating our arguments, we present a model that we hope captures some of the general features that we might expect to observe should nonequilibrium indeed persist. On the one hand, the model--which is particularly simple and parameter free--performs the same ultimate function and shares some of the key characteristics of contemporary experiments that are potentially the most likely to observe relic nonequilibrium. On the other, we emphasise that the model is not a realistic representation of such experiments, and that the comparisons we will make are intended merely to provide context. Rather, our purpose is only to begin the discussion of the qualitative phenomena may ultimately be observed if relic nonequilibrium does indeed exist in the statistics of some particle species. The phenomena that we will discuss are something of a departure from those of a classical origin. For instance, if the model is taken at face value, any broadening of a spectral line will take place on a lengthscale corresponding to the energy resolution of the telescope used. In addition, lines may acquire double or triple bumps, or as we shall discuss, more exotic profiles. There also exists, for a variety of nonequilibrium distributions, the possibility of spectral line `narrowing' in so much that a spectral line profile observed by a telescope would appear narrower than the energy resolution of the telescope could conventionally allow. It has been argued \cite{AV01,AV07,AV08,AV10,UV15} that quantum nonequilibrium could in principle have survived for some species of relic cosmological particles. In speculating on such a possibility, two questions immediately arise. Firstly--what requisite properties must a particle species possess in order to have been created in a state of nonequilibrium, and in order to retain at least a portion of it to the present day? An answer to this question, and an identification of a viable candidate in current particle physics models would presumably allow one to assess the likelihood that relic nonequilibrium could be discovered. It has been argued \cite{UV15,AV10}, for instance, that the decay of a nonequilibrium inflaton field may be a plausible mechanism for the production of such relics. A pilot-wave treatment of the inflaton field on the Bunch-Davies vacuum has been shown to exhibit trajectories that do not allow relaxation at all \cite{AV07,AV10}. Rather, any initial nonequilibrium in the inflaton field is simply scaled in proportion to physical wavelengths. In this respect, nonequilibrium from a pre-inflationary phase (if such a era existed) may have been conserved until (p)reheating, in which most of the matter in the universe is currently understood to have been created. Additionally, during inflation, field modes that are initially below the Planck scale are stretched to the scale at which current particle theories become meaningful. One could therefore conjecture that these so called trans-Planckian modes may have transferred exotic gravitational effects to conventional particle physics lengthscales \cite{AV10,UV15}. Of course, any discussion along these lines is highly speculative. An estimate of the likelihood of nonequilibrium surviving in particles to this day would depend delicately upon, not only the details of an assumed primordial cosmology, but also the particular particle physics model employed\footnote{Ref.\ \cite{UV15} discusses an illustrative scenario for the gravitino ($\tilde{G}$), which arises in supergravity theories and is a proposed dark matter candidate.}. As these are two of the most uncertain areas of contemporary physics, in this article we will address a different question. For the purposes of this article, it will suffice to say that within whatever cosmological and particle physics theories ultimately stand the test of time, there may exist a reasonable window of opportunity for quantum nonequilibrium to survive to this day for some species of relic particle. We refer the reader to \cite{UV15} for further discussion on this point. The question we shall address is as follows. If quantum nonequilibrium did indeed exist in the statistics of a relic particle, then how might such nonequilibrium relics manifest themselves in present-day experiments? A basic requirement to avoid relaxation is that the particle must be only very minimally interacting. The particle would therefore almost certainly come under the heading dark matter, whether as the whole of the observed matter deficit or as part of a larger dark sector. Additionally, whilst it may be possible to synthesise a candidate species in a particle accelerator, the parent particles used would already have relaxed and so the products would necessarily display equilibrium statistics. We must therefore concern ourselves with astrophysical sources, which might conceivably contain particles that have not yet undergone complete relaxation \cite{UV15}. In our discussion it will be useful to take, for the purposes of comparison and illustration, gamma-ray space telescopes concerned with the indirect detection of astrophysical candidates for dark matter generally referred to as weakly interacting massive particles (WIMPS). For example EGRET \cite{EGRET93}, Fermi-LAT \cite{LAT09}, DAMPE \cite{DAMPE}, GAMMA-400 \cite{GAMMA400}. It is hoped that such experiments may observe photons from the annihilation or decay of WIMPs. If these WIMPs were themselves in a state of quantum nonequilibrium, then as argued in ref.\ \cite{UV15} we might reasonably expect some of this nonequilibrium to remain in the statistics of the photons produced. Were such nonequilibrium photons to enter a telescope, then we might expect to see alterations to the spectrum observed, although the characteristics of these spectral alterations--the subject of this article--are not yet known. By definition, dark matter does not interact directly with the electromagnetic field, and so the production of the photons that could be observed by these telescopes would happen at loop level, or through intermediary particle production. Despite the suppression that typically results from loop level interactions, it has long been argued \cite{BS88,Rudaz89,B97,B04,Bertone12} that the detection of photon emission from annihilating or decaying dark matter would provide excellent evidence for a dark matter candidate. Primarily this is because, in cold dark matter models, annihilation would produce two back-to-back photons of energy $E_\gamma=m_{\text{WIMP}}$ and with only minimal intrinsic broadening (see for instance ref.\ \cite{Berg12}). Hence, observation of a spectral line would yield both the spatial location and the mass of the annihilating WIMP particles. Of course, the telescopes are not perfectly precise in their measurements. A single reading of a photon of energy $E_\gamma$ would satisfy what is termed the energy dispersion probability density function $D(E|E_\gamma)$, with a spread characterised by the energy dispersion $\Delta E/E_\gamma$\footnote{See for instance section 7 and figure 67 in \cite{LAT12}.}. Many individual readings on an ensemble of photons with an actual spectrum $\rho_\text{act}(E_\gamma)$ would produce an observed spectrum \begin{align} \rho_\text{obs}(E)=\int D(E|E_\gamma)\rho_\text{act}(E_\gamma)dE_\gamma,\label{eq1} \end{align} that is convolved by the energy dispersion function $D(E|E_\gamma)$. In the context of spectral lines, it is pertinent to consider the relative width of the $D(E|E_\gamma)$ and $\rho_\text{act}(E_\gamma)$ distributions. We may regard equation \eqref{eq1} in two separate regimes. Firstly, in the case of a higher resolution telescope where the width of $D(E|E_\gamma)$ is much smaller than any intrinsic spread in an actual line spectrum $\rho_\text{act}=\rho_\text{line}$, then we may approximate $D(E|E_\gamma)\rightarrow\delta(E-E_\gamma)$, and the observed spectrum would approximate the actual spectrum $\rho_\text{obs}(E)\approx\rho_\text{line}(E)$. In other words, a telescope of sufficient resolution may resolve the profile of the spectral line. Secondly, in the case of a lower resolution telescope, where the width of $ D(E|E_\gamma)$ is much larger than the intrinsic spread in the actual line spectrum, we may approximate $\rho_\text{act}=\rho_\text{line}(E_\gamma)\rightarrow\delta(E_\text{line}-E_\gamma)$ and the observed spectrum would instead approximate the energy dispersion function $\rho_\text{obs}(E)\approx D(E|E_\text{line})$. Hence, a telescope with an energy resolution that is inadequate to resolve the profile of an actual physical line, will instead observe a line whose profile is a function of the interaction between the telescope and the incident photon. In cold dark matter models of WIMPs, conventional line broadening occurs primarily due to the Doppler effect and is expected to produce an annihilation line with an intrinsic spread of 0.1\% of $E_\text{line}$ \cite{Berg12}. By comparison, the EGRET instrument aboard the Compton Gamma Ray Observatory that collected data from 1991-2000 achieved an energy dispersion of $\sim$20\% \cite{EGRET93}. The Large Area Telescope aboard the Fermi Gamma-ray Space telescope currently achieves around $\sim$10\% \cite{LAT12}. The DAMPE telescope, which was launched in December 2015, achieves an energy dispersion of $\sim$1.5\% \cite{DAMPE}. The GAMMA-400 and HERD telescopes are proposed to be launched in the early 2020s and reach an energy resolution of $\sim$1\% \cite{GAMMA400,HERD}. All of these telescopes have energy dispersions that are appreciably larger than the expected $0.1\%$ width of a WIMP annihilation line, and so could not be expected to resolve this conventional broadening. Instead, if a WIMP annihilation line were discovered, the observed line profile would closely approximate $D(E|E_\text{line})$--a property of the telescope itself. In contrast to the conventional broadening, and as we shall discuss, quantum nonequilibrium may be thought more properly to affect the interaction between the telescope and the photon, rather than the actual energy of the individual photons. The effect of quantum nonequilibrium is to alter the energy dispersion function $D(E|E_\gamma)$, rather than the actual spectrum $\rho_\text{act}(E_\gamma)$. This is important as, naively, one might expect a higher resolution telescope (with smaller $\Delta E/E_\gamma$) to be preferential for detection of quantum nonequilibrium, but it appears that this may not be the case. Our analysis indicates that in the presence of quantum nonequilibrium, whilst higher resolution telescopes will remain more favourable for the discovery of a sharp spectral line, nonequilibrium signatures may be more apparent in spectra observed by lower resolution telescopes. Nonequilibrium will be most evident when the width of $D(E|E_\gamma)$, the energy dispersion $\Delta E/E_\gamma$, is larger than any intrinsic energy spread in $\rho_\text{line}(E_\gamma)$--which we shall take as a working definition of lower resolution. Many of the current generation of telescopes are certainly within this regime. As such, if one were to accept these arguments and those we shall develop through the model below, then for many of the current generation of telescopes nonequilibrium line effects could in principle dominate conventional line broadening. Our paper is organised as follows. In section \ref{QM_model} we present an idealised and parameter-free field-theoretical model of a spectral measurement of the electromagnetic field. This will be sufficiently simple as to permit an explicit solution to the (functional) Schr\"{o}dinger equation. In section \ref{noneq_spectral_lines} we present the pilot-wave description of the model, we discuss how nonequilibrium may affect spectral lines, and we provide some explicit calculations. In section \ref{implications} we comment on the limitations of the model and discuss possible implications for the indirect search for dark matter. | 16 | 9 | 1609.04576 |
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1609 | 1609.08044_arXiv.txt | { Gaia is a space mission currently measuring the five astrometric parameters as well as spectrophotometry of at least 1 billion stars to $G = 20.7$~mag with unprecedented precision. The sixth parameter in phase space --- radial velocity --- is also measured thanks to medium-resolution spectroscopy being obtained for the 150 million brightest stars. During the commissioning phase, two fields, one around each ecliptic pole, have been repeatedly observed to assess and to improve the overall satellite performances as well as the associated reduction and analysis software. A ground-based photometric and spectroscopic survey was therefore initiated in 2007, and is still running in order to gather as much information as possible about the stars in these fields. This work is of particular interest to the validation of the Radial Velocity Spectrometer (RVS) outputs. } { The paper presents the radial velocity measurements performed for the Southern targets in the 12 - 17 $R$ magnitude range on high- to mid-resolution spectra obtained with the GIRAFFE and UVES spectrographs. } { Comparison of the South Ecliptic Pole (SEP) GIRAFFE data to spectroscopic templates observed with the HERMES (Mercator in La Palma, Spain) spectrograph allowed a first coarse characterisation of the 747 SEP targets. Radial velocities were then obtained by comparing the results of three different methods. } {In this paper we present an initial overview of the targets to be found in the 1-square-degree SEP region that was observed repeatedly by Gaia from its commissioning on. In our representative sample, we identified 1 galaxy, 6 LMC S-stars, 9 candidate chromospherically active stars, and confirmed the status of 18 LMC Carbon stars. A careful study of the 3471 epoch radial velocity measurements, led us to identify 145 RV constant stars with radial velocities varying by less than 1~\kps. 78 stars show significant RV scatter, while nine stars show a composite spectrum. As expected the distribution of the RVs exhibits 2 main peaks corresponding to Galactic and LMC stars. By combining \feh\, and \logg\, estimates, and RV determinations we identified 203 members of the LMC, while 51 more stars are candidate members. } { This is the first systematic spectroscopic characterisation of faint stars located in the SEP field. During the coming years, we plan to continue our survey and gather additional high- and mid-resolution data to better constrain our knowledge on key reference targets for Gaia. } | 16 | 9 | 1609.08044 |
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1609 | 1609.03817_arXiv.txt | Ellerman Bombs are signatures of magnetic reconnection, which is an important physical process in the solar atmosphere. How and where they occur is a subject of debate. In this paper we analyse \sunrise{}/IMaX data together with 3D MHD simulations that aim to reproduce the exact scenario proposed for the formation of these features. Although the observed event seems to be more dynamic and violent than the simulated one, simulations clearly confirm the basic scenario for the production of EBs. The simulations also reveal the full complexity of the underlying process. The simulated observations show that the Fe~I~525.02 nm line gives no information on the height where reconnection takes place. It can only give clues about the heating in the aftermath of the reconnection. The information on the magnetic field vector and velocity at this spatial resolution is, however, extremely valuable because it shows what numerical models miss and how they can be improved. | One of the important physical processes that has a major effect on the energy budget in the solar atmosphere is magnetic reconnection. It is a mechanism behind a myriad of dynamic atmospheric phenomena, starting from small scale flux cancellation in the solar photosphere to the largest solar disruptions - flares. Recently, a number of studies concentrated on explaining one of these phenomena called Ellerman Bombs \citep[EBs;][]{Ellerman1917} that pose a considerable challenge to models because of their specific characteristics. They are defined as transient brightenings of the extended wings of the H$\alpha$ line, but leave signatures also in Ca~II~H and Ca~II~IR~854~nm \citep{Gregal:2013,Reza2015} and sometimes even in observables which should sample orders of magnitude higher temperatures than the former \citep{Vissers2015,Tian2016}. Observations show that EBs are almost exclusively formed in young emerging active regions, usually between spots/pores where the emergence of serpentine field lines takes place. Series of EBs usually appear aligned along the orientation of the active region in so-called Bald Patches (BPs) - dips in magnetic field lines \citep{Pariat:etal:2004,Pariat:2006}. This scenario is also supported by MHD simulations \citep{Isobe2007,Archontis2009}. Because EBs assume the shape of a flame that seems to be rooted in the intergranular lanes \citep{Matsumoto2008,Watanabe2011}, it is suggested that their formation begins very low, near the surface. However, one- and two-dimensional modelling strongly suggests that only a temperature increase starting at heights of a few hundred km above the solar surface can produce the observed H$\alpha$ line profile without continuum brightening \citep{Kitai1983,Fang:2006,Bello:etal:2013,Berlicki:2014,Fang:2006,Hong:2014}. \cite{Nelson2013} analysed observations together with similar events in 3D MHD simulations and inferred that EBs occur co-spatially with regions of strong opposite-polarity magnetic field at locations where the Fe~I~630.25~nm line core intensity increases. According to them this gives enough evidence that EBs are in fact signatures of photospheric magnetic reconnection. Later it was established that most of the features observed by \cite{Nelson2013} were actually 'pseudo-EBs', just strong-field magnetic concentrations \citep{Rutten2013}. \cite{Reid2016} imposed a more discriminating threshold for detecting EBs and corrected the previously made mistake. Their detailed analysis of a large number of these features goes beyond just looking at the Fe~I~630.25~nm line's core intensity. Using inversions of the spectropolarimetric data they retrieved temperature enhancements of 200~K at the EBs' footpoints. This value is below the lowest number required by H$\alpha$ line modelling. \cite{Reid2016} attributed this to the low formation height of Fe~I~630.25~nm line. In this paper we analyse the photospheric response to an EB observed by the Imaging Magnetograph eXperiment \citep[IMaX,][]{MartinezPillet2011} on board of \sunrise{}, a balloon-borne solar observatory \citep{Solanki2010,Barthol2011,Berkefeld2011,Gandorfer2011}. Similar to \cite{Nelson2013} we compare the observations with MHD simulations, but unlike them we carry out an appropriate numerical experiment. While they used a run where cancellation of weak network was simulated, we reproduce the exact scenario proposed for the formation of EBs, i.e. the emergence of serpentine magnetic flux. Furthermore, we take into account the spectral and spatial resolution of our instrument, as well as its polarimetric sensitivity and apply the same inversion strategy as in the case of the observations. We pinpoint possible errors that this kind of analysis could produce and finally show which characteristics of EBs numerical simulations fail to reproduce and why. | During the evolution of NOAA AR 11768, an event at the confluence of two flux emergences produced a significant brightness increase in the SDO/AIA 1700 channel. Given that AIA 1700 brightness was above the $5\sigma$ threshold set by \cite{Gregal:2013}, we claim that this event can be classified as an Ellerman bomb. This claim is further reinforced by the underlying field configuration recovered by \cite{Centeno:etal:this}. They demonstrate that the event was produced during subsequent appearance of two emerging flux regions where the field is aligned in a way that points to a serpentine-like field topology. Observations strongly suggest that this is an environment where EBs inevitably appear \citep{Georgoulis:2002,Pariat:etal:2004,Pariat:2006}. In this work we analyse in detail photospheric signatures of the event recorded by \sunrise/IMaX. We compare the results with a numerical experiment that reproduces the proposed flux emergence scenario. The simulated event results in a temperature increase which is in agreement with previous EB models \citep{Fang:2006,Bello:etal:2013,Berlicki:2014,Fang:2006,Hong:2014} and produces the typically observed morphology of EBs in the wings of H$\alpha$. Simulations show the expected field topology at the onset of the EB-like feature. The reconnection seems to start in the photosphere, some 200~km above the surface. As new material and fresh magnetic flux emerges, the current sheet between the two opposite polarities of the two emerging features extends further up and the location where post-reconnection U-loops begin is shifted to above 1~Mm. During the whole evolution, no actual information on the location where the reconnection takes place is contained in the Fe~I~525.02~nm line. As the energy is deposited in the upper photosphere, the temperature rapidly changes first there and then it gets modified also further down by the reconnection aftermath, e.g. downward propagating reconnection jets as they collide with the local plasma. Because of this, the formation height of Fe~I~525.02~nm line is continually shifted downwards till the moment when the line samples only a very shallow layer near the solar surface. This is in agreement with observations that suggest that these events produce sufficient energy to ionize the neutral metals \citep{Rutten2015} and explains the very low temperature increase found by \cite{Reid2016}. Simulations suggest that not all locations where temperature increases are necessarily related to the event itself. Some of those appear at the footpoints of rapidly emerging loops as they expand in the horizontal direction upon reaching the surface. While material is drained, the footpoints are squeezed in between already developed magnetic features and fast-moving emerging material. As a result fast downflows are generated and footpoints quickly evacuated. These locations might be similar to the high-temperature points presented by \cite{Tortosa2009} and might have similar characteristics from a slanted viewing angles as the flows studied by \cite{Bellot2009} and \cite{Vitas2011}. They also resemble in some ways locations of convective collapse. Although they show qualitative similarities with the observations, the simulations fail to produce equally high temperatures. Comparison of the field strength and especially the velocities points to the fact that the choice of the initial magnetic field setup is not ideal. To produce an event that is as dynamic and violent as the observation suggest, more magnetic buoyancy is needed that would launch magnetic field into solar atmosphere more efficiently. Also, persistent upflows that last at least three times longer seem to be essential. For this, instead of an embedded thin flux sheet, one needs to insert into simulation domain more flux or a flux sheet over extended period of time. | 16 | 9 | 1609.03817 |
1609 | 1609.08933_arXiv.txt | The results of three-dimensional numerical simulations of the gas dynamics of the atmosphere of a ``hot Jupiter'' exoplanet during the passage of a coronal mass ejection (CME) from the central star are presented. These computations assumed the parameters for the stellar wind and the CME to be typical of the solar values. The characteristic variations of the flow pattern are considered for quasi-closed and closed (but appreciably distorted by the gravitational influence of the star) gaseous envelopes of the exoplanet. It is shown that a typical CME is sufficient to tear off the outer part of an asymmetric envelope that is located beyond the Roche lobe and carry it away from the exoplanet. This leads to a substantial increase in the mass-loss rate from the exoplanet envelope during the passage of CMEs. The mass-loss rate grows by about a factor of 11 for a closed envelope, and by about a factor of 14 for a quasi-closed envelope. Possible evolutionary consequences of the loss of part of the atmosphere during the passage of CMEs are discussed. | “Hot Jupiter” exoplanets have masses comparable to the mass of Jupiter and orbital semi-major axes not exceeding 0.1\,AU. The proximity of these objects to their central stars leads to a supersonic regime for the stellar wind flowing around the planet, and therefore to the formation of an outgoing shock followed by a contact discontinuity -- a boundary separating the matter of the wind from the gas of the exoplanetary atmosphere~\citep{Baranov-1977}. According to~\citep{Bisikalo-2013b, Bisikalo-2015}, the gaseous envelopes around hot Jupiters can be divided into three types. First and foremost, we must distinguish closed envelopes, where the head-on collision point (HOCP) corresponding to the minimum distance from the contact discontinuity to the planet is located inside the Roche lobe. Depending on the degree of filling of the Roche lobe, such envelopes may deviate from a spherical shape, but they are undoubtedly characterized by an absence of a significant outflow of matter, $\dot{M} < 10^{9}$\,\gs~\citep{Cherenkov-2014}. If the HOCP, and consequently part of the atmosphere, is located outside the Roche lobe, a significant outflow from the vicinity of the inner Lagrange point \Lp1 is present, and the envelope becomes appreciably asymmetrical. Some parts of the envelope may have restricted dimensions, since the propagation of the flow arising when the Roche lobe overflows may be stopped by the dynamical pressure of the stellar wind. In this case, a quasi-closed, stationary envelope with a complex shape forms in the system~\citep{Bisikalo-2013a}, with corresponding mass-loss rates of $\dot{M} < (3-5) \times 10^{9}$\,\gs~\citep{Cherenkov-2014, Bisikalo-2013c}. If the wind is not able to stop the flow from \Lp1, an open, aspherical envelope forms in the system. Closed and quasi-closed envelopes are of the most interest for our study, since they are stationary formations with relative low mass-loss rates. Open envelopes are characterized by high mass-loss rates~\citep{Cherenkov-2014} and short lifetimes, and can probably exist only during specific, short-lived stages in the evolution of the atmospheres of hot Jupiters. Previous results on gas-dynamical simulations of the envelopes of hot Jupiters have assumed that the stellar wind does not vary with time. At the same time, it is known from observations that even a relatively quiescent star such as the Sun has a wind in the form of a quasi-neutral plasma that can vary by factors of tens (see, e.g.,~\citep{Zastenker-1999, Liu-2014}). Beginning in the 1990s, it was known that the main source of the perturbed solar wind is giant ejections of matter from the solar corona, so-called coronal mass ejections (CMEs). According to~\citep{Zastenker-1999, Johnstone-2015, Howard-1985}, CMEs are characterized by a mean mass of plasma ejected into the interplanetary medium $\sim 10^{13}\,\mathrm{kg}$, a mean ejection energy $\sim 10^{31}\,\mathrm{erg}$, and ejection velocities that vary from $100$ to $3000\,\mathrm{km\,s^{-1}}$; i.e., the motion is obviously supersonic, so that the propagation of the ejection is accompanied by the formation of a shock. Moreover, a CME often differs from the ordinary solar wind by an up to 10-15\% increase in the abundance of helium ions. Note that, even for the Sun, the frequency of CMEs is very high, and varies from 0.5 to 2.5 per day, depending on the time separation from the maximum of the solar activity cycle~\citep{Zastenker-1999}. This means that an exopanet located even around a quiescent star will fairly often (no less than twice per month, as is the case for the Earth) be subject to the action of CMEs, which distort the flow patterns in its gaseous envelope. The main aim of our present study is to investigate the action of CMEs having fairly typical parameters on the structure of flows in the gaseous envelopes of hot Jupiters. The discovery in~\citep{Bisikalo-2013a, Bisikalo-2013b} of the possibility of forming extended gaseous envelopes around hot Jupiters not only fundamentally changes approach to interpreting observational data, but also substantially influences consideration of the physical processes occurring in these envelopes. In particular, it is very obvious that, if the size of the envelope exceeds the size of the Roche lobe, matter that is formally part of the planetary atmosphere will be only weakly gravitationally bound to it. In this case, even a small external action (such as a CME) is sufficient to tear off the outer part of the envelope from the planet, thereby appreciably increasing the mass loss by the exoplanet. Another important consequence of an extended envelope is the possibility of neglecting the influence of the planet's magnetic field on the dynamics of the outer layers of the envelope, and using gas-dynamical equations to describe the flows in the system. Indeed, even with the maximum estimates of the possible magnetic fields on hot Jupiters (up to a tenth of the magnetic moment of Jupiter~\citep{Kislyakova-2014}), the radius of the magnetosphere lies within the Roche lobe, so that the influence of the planet’s magnetic field on the dynamics of outer parts of the envelope is negligible. We used a modification of the computational code of~\citep{Bisikalo-2013b, Cherenkov-2014}, based on the solution of a system of gas-dynamical equations, for our new computations. We specified the characteristics of the plasma incident on the exoplanet based on the parameters of solar CMEs measured at the Earth’s orbit~\citep{Farrell-2012}. Further, we assumed that the relative variations of the parameters of the CME plasma at the orbit of the hot Jupiter match with the same as at the Earth, and occur at the same moments in time. The absolute values of the parameters were specified assuming that the parameters of the unperturbed wind correspond to the solar values at the orbit of the exoplanet. Computations were carried out for both closed and quasi-closed exoplanet atmospheres. Special attention was given to the variations in the mass-loss rate of the planetary atmosphere under the action of the CME. The paper is organized as follows. Section~\ref{sec:problem} presents our formulation of the problem. Section~\ref{sec:results} presents the results of our gas-dynamical simulations, and Section~\ref{sec:conclusion} summarizes the main results of our study. | \label{sec:conclusion} \begin{wrapfigure}{Lht!}{0.5\textwidth} \centering \epsfig{trim={2cm 2cm 0 5cm},clip,width=7cm,file=collage_7500_3_2.eps} \caption{Same as Fig.~7 for the fourth phase of the CME.} \end{wrapfigure} We have considered the gas dynamics of the interaction of a coronal mass ejection from a star with the atmosphere of a hot Jupiter exoplanet. These computations assumed that the parameters of the stellar wind and the CME corresponded to their solar values. Computations were carried out for both closed (but appreciably distorted by the gravitational influence of the star) and quasi-closed gaseous envelope of an exoplanet. With the adopted parameters for the atmosphere, an appreciable fraction of the envelope is located beyond the Roche lobe in both models. During the interaction with the CME, this part of the envelope, and even some of the envelope that is located within the Roche lobe, becomes gravitationally unbound from the planet, and is ejected from the system. The results of our three-dimensional numerical simulations have established that the total mass lost during the passage of the CME in the case of a closed envelope is $\Delta M \simeq 5 \times 10^{15}$\,\g. The mass lost in the case of a quasi-closed envelope is $\Delta M \simeq 1.0 \times 10^{16}$\,\g. These masses exceed the masses lost in the stationary solutions for the closed and quasi-closed envelopes over the same time intervals by factors of 10.8 and 14.1, respectively. Let us consider the possible evolutionary consequences of tearing off part of the envelope of a hot Jupiter during the passage of a CME. Let us suppose that the star displays solar-type activity; i.e., the rate at which CMEs distort the atmosphere of the planet is $\sim\!2$ per month. We will also suppose that the main parameters of the CMEs (the characteristic variations in the density, velocity, temperature, and the durations of the various phases) are equal to those for solar CMEs. In this case, according to our computations, the total mass loss over a year will be increased by factors of 2.7 and 3.5 for the closed and quasi-closed envelopes, compared to their equlibrium values (obtained for a stationary wind). Thus, even for weakly active stars, taking into account the influence of CMEs leads to approximately a threefold decrease in the characteristic lifetime of the atmosphere of a hot Jupiter. Note that the probability of a hot Jupiter having an extended quasi-closed envelope is approximately one-third~\citep{Bisikalo-2015}, so that our estimates should appreciably influence the overall evolution of the atmospheres of such exoplanets. Equally important, the probability and intensity of CMEs increase for young stars (compared to the Sun), further decreasing the lifetimes of extended envelopes around hot Jupiters. Summarizing, the role of CMEs in the evolution of the gaseous envelopes of hot Jupiters is very important, and must be included when considering the evolution of the atmospheres of such exoplanets. | 16 | 9 | 1609.08933 |
1609 | 1609.04716_arXiv.txt | \noindent We revisit the nonthermal gravitino production at the (p)reheating stage after inflation. Particular attention is paid to large field inflation models with a $\mathbb{Z}_2$ symmetry, for which the previous perturbative analysis is inapplicable; and inflation models with a stabilizer superfield, which have not been studied non-perturbatively. It is found that in single-superfield inflation models (without the stabilizer field), nonthermal production of the transverse gravitino can be cosmologically problematic while the abundance of the longitudinal gravitino is small enough. In multi-superfield inflation models (with the stabilizer field), production of the transverse and longitudinal gravitinos is significantly suppressed, and they are cosmologically harmless. We also clarify the relation between the background field method used in the preheating context and the standard perturbative decay method to estimate the gravitino abundance. | \subsection{Introduction} Supersymmetric (SUSY) models are well-motivated as a physics beyond the standard model, since it provides a successful gauge coupling unification, dark matter candidate, a great reduction of the hierarchy problem etc. In supergravity, however, there is a cosmological problem associated with the gravitino, the superpartner of the graviton, called the gravitino problem~\cite{Pagels:1981ke, Weinberg:1982zq, Khlopov:1984pf,Ellis:1984eq}. If the gravitino is not the lightest SUSY particle (LSP), it can decay into lighter SUSY particles. The lifetime of the gravitino is given by~\cite{Moroi:1995fs} \begin{align} \tau_{3/2} =\left(\frac{3}{8\pi} \frac{(m_{3/2}^0)^3}{\Mpl^2}\right)^{-1} \simeq 3\times 10^{-2}\,{\rm sec}\left( \frac{100\,{\rm TeV}}{m_{3/2}^0} \right)^3, \end{align} where $\Mpl$ is the reduced Planck scale and $m_{3/2}^0$ denotes the present gravitino mass.\footnote{ Throughout this paper, we distinguish the ``present gravitino mass'' $m_{3/2}^0$ and ``gravitino mass'' $m_{3/2}$, since the notion of gravitino and its mass is time-dependent in a cosmological evolution. The former corresponds to the gravitino mass in the present universe and it is just a constant while the latter is time-dependent. } Here we have assumed that the gravitino decays into only gaugino plus gauge boson pairs. If other decay modes are open, the lifetime becomes slightly shorter. Therefore, if the gravitino is much lighter than $100\,$TeV, it decays after the beginning of big-bang nucleosynthesis (BBN) and may affect light element abundances~\cite{Moroi:1995fs,Jedamzik:2004er, Kawasaki:2004yh, Kawasaki:2004qu, Jedamzik:2006xz, Kawasaki:2008qe}. If it is heavier, the decay itself does not affect BBN but LSPs produced by the gravitino decay can be overabundant compared with the observed dark matter abundance. If the gravitino is LSP and $R$-parity is conserved, the gravitino itself contributes to the dark matter abundance. In any case, there is a strict upper bound on the gravitino abundance. There are several processes that produce gravitinos in the early universe. One of the unavoidable production mechanisms is thermal production: in the high-temperature universe, scatterings of high-energy particles produce gravitinos~\cite{Bolz:2000fu,Pradler:2006qh,Rychkov:2007uq}.\footnote{ Ref.~\cite{Ellis:2015jpg} discussed the gravitino production by the scatterings of energetic inflaton decay products during the process of thermalization~\cite{Harigaya:2013vwa} and found that it is subdominant compared with the standard thermal production. } The abundance of thermally produced gravitinos is proportional to the reheating temperature after inflation $T_{\rm R}$, and hence we obtain an upper bound on $T_{\rm R}$ to avoid the gravitino problem. Gravitinos can also be produced nonthermally. Nonthermal gravitino production by the direct decay of inflaton was extensively studied in a series of works~\cite{Endo:2006zj,Nakamura:2006uc,Kawasaki:2006gs, Asaka:2006bv,Dine:2006ii,Endo:2006tf,Kawasaki:2006hm,Endo:2007ih,Endo:2007sz}. It was found that the inflaton generally decays into the gravitino pair with the partial decay rate~\cite{Kawasaki:2006gs, Asaka:2006bv, Endo:2006tf,Kawasaki:2006hm} \begin{align} \Gamma(\phi\to \psi\psi) \simeq \frac{1}{64\pi}\left( \frac{\left<\phi\right>}{\Mpl} \right)^2 \frac{m_\phi^3}{\Mpl^2}, \label{inf_dec} \end{align} where $m_\phi$ is the inflaton mass and $\left<\phi\right>$ is the vacuum expectation value (VEV) of the inflaton. It gives a stringent constraint on inflation models, although there are some loopholes~\cite{Endo:2007cu,Nakayama:2012hy}. This expression for the decay rate is valid for small-field inflation models such as new inflation or hybrid inflation.\footnote{ Chaotic inflation without a $\mathbb{Z}_2$ symmetry also leads to a similar expression. } On the other hand, large field inflation models~\cite{Linde:1983gd} attract lots of attentions in view of recent developments on successful inflation model building in the framework of supergravity~\cite{Kawasaki:2000yn,Kallosh:2010ug,Kallosh:2010xz,Ferrara:2010in,Nakayama:2010kt,Nakayama:2013txa,Kallosh:2013hoa,Kallosh:2013yoa,Galante:2014ifa}. It is interesting because it can be tested with on-going/future $B$-mode polarization experiments. In large field inflation models with a $\mathbb{Z}_2$ symmetry in which the inflaton field oscillates around the origin $\phi=0$ after inflation, we cannot use the expression (\ref{inf_dec}) as a gravitino production rate. This is simply because the calculations in Refs.~\cite{Endo:2006tf,Kawasaki:2006hm,Endo:2007ih,Endo:2007sz} assumed the perturbative decay of inflaton around its VEV. In inflation models with the $\mathbb{Z}_2$ symmetry, however, there is no such decay process due to the $\mathbb{Z}_2$ symmetry.\footnote{Another assumption there was that the SUSY is dominantly broken by the Polonyi field at the end of reheating so that the definition of ``gravitino'' at that epoch is the same as the present-day gravitino. This assumption is valid as long as we are interested in the gravitino production at the end of reheating. In large field inflation models, however, the gravitino production is often dominated at the epoch just after inflation (preheating) and hence this assumption is not justified, as we will see later.} This does not mean that the inflaton cannot decay into gravitinos as well as other light particles. The inflaton coherent oscillation affects the masses or kinetic terms of coupled particles. The coupled particles, including gravitinos, ``feel'' the rapid inflaton oscillation and it affects the evolution of their wave functions. It is known that this leads to particle production, often in the context of preheating~\cite{Dolgov:1989us,Traschen:1990sw,Shtanov:1994ce,Kofman:1994rk,Kofman:1997yn}. Therefore, even if the inflaton has the $\mathbb{Z}_2$ symmetry, its coherent oscillation necessarily transfers its energy to the coupled particles. The question we would like to address is: what amount of gravitinos is produced during the preheating? Production of gravitinos during the preheating was first discussed in Refs.~\cite{Kallosh:1999jj,Giudice:1999am} in a single-superfield case, in which only one inflaton chiral superfield was introduced. There it was found that in the preheating stage, longitudinal gravitinos are efficiently produced. Later it was recognized that the theory of gravitino preheating is much more involved due to the subtlety of the notion of ``gravitino''~\cite{Kallosh:2000ve,Nilles:2001ry,Nilles:2001fg}. The gravitino becomes massive by ``absorbing'' the goldstino, but the definition of goldstino is time-dependent in a cosmological background. In the early universe, the inflaton oscillation energy dominantly breaks SUSY and hence the goldstino is almost the inflatino, the fermionic superpartner of the inflaton. At late time, the Polonyi field\footnote{ In this paper we call the present-day SUSY breaking field as Polonyi field. } dominantly breaks SUSY and its fermionic component, Polonyino, becomes the goldstino. Thus the composition of goldstino changes with time. Refs.~\cite{Nilles:2001ry,Nilles:2001fg} noticed that it is essential to include (at least) two chiral superfields, inflaton and Polonyi, to correctly deal with this problem and concluded that what the preheating efficiently produces eventually becomes the inflatino, which is less harmful than the gravitino. Still, however, a quantitative/comprehensive analysis of the nonthermal gravitino production rate in such a case is missing. Although Refs.~\cite{Nilles:2001ry,Nilles:2001fg} revealed that the gravitino production is suppressed than previously thought, it is highly non-trivial how we can extrapolate their numerical results into more realistic setups and parameters both in the inflaton and SUSY breaking sector. Thus we would like to provide general analytic formulae for the nonthermal gravitino abundance that are applicable to any realistic models. \subsection{Brief summary} In this paper, we revisit the theory of nonthermal gravitino production in a comprehensive and unified manner. Our purposes and results are summarized below. \begin{itemize} \item We derive nonthermal gravitino production rates and their resulting abundances quantitatively with useful formulae. We find that in single-superfield inflation models, the production of transverse gravitino is significant and cosmologically problematic, while the production of longitudinal gravitino is less important. This aspect of the nonthermal gravitino production has been overlooked in previous literatures except for a few~\cite{Maroto:1999ch}. \item Recent realistic large field inflation models introduce an additional chiral superfield, called a stabilizer~\cite{Kawasaki:2000yn,Kallosh:2010ug,Kallosh:2010xz,Ferrara:2010in,Nakayama:2010kt,Nakayama:2013txa,Kallosh:2013hoa,Kallosh:2013yoa,Galante:2014ifa}. We find that in models with the stabilizer field, the production of transverse gravitino is significantly suppressed and it is cosmologically harmless. For the longitudinal component, the production rate is similar to the single-superfield case. \item We show the equivalence between the background field method developed in Refs.~\cite{Kallosh:2000ve,Nilles:2001ry,Nilles:2001fg} and the perturbative decay method developed in Refs.~\cite{Endo:2006tf,Kawasaki:2006hm,Endo:2007ih,Endo:2007sz} for evaluating the gravitino abundance in some sense. The former can deal with a broad class of models including $\mathbb{Z}_2$-symmetric large field models, while in models without the $\mathbb{Z}_2$ symmetry, it gives the same result as the perturbative decay method. \end{itemize} This paper is organized as follows. In Sec.~\ref{sec:grav}, we review the structure of the gravitino Lagrangian to set the stage of discussion in the subsequent Sections. We formulate a general setup to discuss multi-superfield case. Gravitino production in the single-superfield inflation models and in the multi-superfield inflation models are studied in Sec.~\ref{sec:single} and Sec.~\ref{sec:multi}, respectively. The analyses include both the transverse and longitudinal components of gravitino. Our conclusion is in Sec.~\ref{sec:conclusion}, and the gravitino abundance is summarized in Fig.~\ref{fig:Y}. App.~\ref{sec:notation} summarizes our notations and conventions. The background field method to evaluate the fermion production rate, often used in the preheating context, is reviewed in App.~\ref{sec:fermion}. We also briefly cover the gravitino production in small-field inflation models in App.~\ref{sec:small-field} in order to show the equivalence between our method and the perturbative decay method. In App.~\ref{sec:dyn}, we review the multi-field scalar dynamics to discuss the induced oscillation of the Polonyi field in the main text. Calculations of mass eigenvalues are given in App.~\ref{sec:mass}, which are used in Secs.~\ref{sec:long_single} and \ref{sec:long_multi}. | \label{sec:conclusion} We have studied the nonthermal gravitino production during (p)reheating paying particular attention to the case of $\mathbb{Z}_2$ symmetric large field inflation models. The result crucially depends on inflation models. In single-superfield inflation without a stabilizer field, production of the transverse gravitino is efficient and it can cause cosmological problems depending on the power law index of the inflaton potential. The longitudinal gravitino production is safely neglected. On the other hand, in multi-superfield inflation models with a stabilizer field, the transverse gravitino production is significantly suppressed and nonthermal gravitino production plays no important role in cosmology. Fig.~\ref{fig:Y} shows the gravitino abundance as a function of the reheating temperature $T_{\rm R}$ for the single-superfield inflation models (left) and multi-superfield inflation models with a stabilizer field (right). The solid line shows a contribution from thermal production~\cite{Bolz:2000fu,Pradler:2006qh,Rychkov:2007uq}, while dashed (dotted) lines show nonthermally produced ones for the inflaton potential $V\propto \phi^p$ with $p=2$ ($p=4$). Here we have taken $H_{\rm inf}=10^{13}\,$GeV (left) and $m_{3/2}^0=10^6$\,GeV (right). If the inflaton potential changes from $\phi^4$ to $\phi^2$ at some point, the prediction lies between these two lines. It is clearly seen that inflation models with a stabilizer predicts negligibly small nonthermal gravitino abundance. Therefore inflation models with a stabilizer field is motivated not only from the viewpoint of model building, but also from the requirement to avoid the nonthermal gravitino overproduction. Note that in this plot we have ignored some other nonthermal gravitino production processes such as those from Polonyi/inflatino decay since they are rather model-dependent~\cite{Nakayama:2012hy,Nilles:2001my}. However, inclusion of them does not much affect this conclusion. Some comments are in order. In the most part we assumed the (nearly) minimal K\"ahler potential for the inflaton superfield for simplicity. The production rate is significantly modified if there is a $\mathbb{Z}_2$-symmetric K\"ahler potential of the form \begin{align} K \sim \frac{1}{\Mpl^2}|\phi|^2 zz +{\rm h.c.}~~~{\rm or}~~~ \frac{1}{\Mpl^2}X^\dagger\phi zz + {\rm h.c.}, \end{align} for the single-superfield inflaton and multi-superfield inflaton case, respectively. Since these terms directly give the large oscillating Polonyino mass like $m_{\tilde z} \sim m_\phi \phi^2/{\Mpl^2}$, the longitudinal gravitino production rate is expected to be significantly enhanced to the same level as the transverse production rate, if the coefficients of these terms are $\mathcal O(1)$. Moreover, some inflation models, especially those constructed from the Jordan frame action, have a nonminimal K\"ahler potential of the inflaton sector itself which can potentially induce violent phenomena~\cite{Ema:2016dny}. Also we assumed that the inflaton is a gauge singlet: for a gauge non-singlet inflaton, the structure becomes more complicated. We will come back to these issues in future. \begin{figure}[t] \begin{center} \includegraphics[scale=1.3]{Ygrav.pdf} \caption{\small Gravitino abundance $n_{3/2}/s$ as a function of reheating temperature $T_{\rm R}$ for the single-superfield inflation model without a stabilizer field (\textbf{left}) and multi-superfield inflation model with a stabilizer field (\textbf{right}). The solid line shows a contribution from thermal production, while dotted (dashed) lines show nonthermally produced ones for the inflaton potential $V\propto \phi^p$ with $p=2$ ($p=4$). }\label{fig:Y} \end{center} \end{figure} | 16 | 9 | 1609.04716 |
1609 | 1609.04520_arXiv.txt | By monitoring a large number of stars in the Local Group galaxies, we can detect nanolensing events by sub-lunar mass compact objects (SULCOs) such as primordial black holes (PBHs) and rogue (free-floating) dwarf planets in the Milky Way halo. In contarst to microlensing by stellar-mass objects, the finite-source size effect becomes important and the lensing time duration becomes shorter ($\sim 10^{1-4}\,\textrm{s}$). Using stars with $V<26$ in M33 as sources, for one-night observation, we would be able to detect $10^{3-4}$ nanolensing events caused by SULCOs in the Milky Way halo with a mass of $10^{-9}\ms$ to $10^{-7}\ms$ for sources with S/N$>5$ if SULCOs constitute all the dark matter components. Moreover, we expect $10^{1-2}$ events in which bright blue stars with S/N$>100$ are weakly amplified due to lensing by SULCOs with a mass range of $10^{-11}\ms$ to $10^{-9}\ms$. Thus the method would open a new window on SULCOs in the Milky Way halo that would otherwise not be observable. | Currently, there has not been stringent observational constraint on the abundance of SUb-Lunar mass Compact Objects (SULCOs) with a mass of $10^{-13}\ms \le M \le10^{-9}\ms$ as the dark matter candidates\citep{carr16}. SULCOs can be either small planets, satellites, or primordial black holes (PBHs)\citep{carr74, carr75, hawking82, yokoyama97, inoue03}. Microlensing tests such as the MACHO and EROS collaborations ruled out the possibility that the compact objects with a mass of $10^{-7}M_{\odot}\le M\le 10^{-3}M_{\odot} $ constitute the Milky Way halo\citep{paczynski86, alcock97, alcock00, tisserand07}. On the other hand, femtolensing test of Gamma-ray bursts (GRBs) ruled out the mass range $10^{-16}M_{\odot}\le M\le 10^{-13}M_{\odot}$ assuming that the GRBs are at cosmological distance so that the angular size of the GRB source is sufficiently small~\citep{marani99, barnacka12}. The SULCOs also induce picolensing of GRBs but the observational limit is very weak~\citep{marani99}. The dynamical constraint on the amount of SULCOs in the dark halo is also less stringent\citep{carr99,carr16}. Recent lensing analyses based on the time variability of stars in the Milky Way disk using the data from the Kepler satellite, have yielded a constraint on the mass range $2 \times 10^{-9}\ms \le M \le 10^{-7}\ms$\citep{griest11,griest13,griest14}. However, the constraint is limited to local SULCOs at a distance $<4 \,\tr{kpc}$. Since the size of the Milky Way halo is much larger, it is important to constrain the abundance of SULCOs that reside at distance $>4\, \tr{kpc}$ as well. The neutron-star capture constraint for a mass range of $10^{-15}\ms \le M \le 10^{-9}\ms$ \citep{capera13} depends on the assumption that the PBHs reside in globular clusters, thus the limit is uncertain\citep{carr16}. SULCOs may be free-floating or rogue dwarf planets that have been ejected from developing or developed planetary systems\citep{ma16, smullen16}. Although the mass scale of the detected rogue planets\citep{penaramirez16}, typically a Jupiter mass scale, are much larger than the SULCO mass scale, a large number of rogue dwarf planets may reside in the Milky Way halo. For instance, the gravitational perturbation by passing stars may cause destruction of extrasolar planetary systems that may correspond to the Kuiper belt objects or the Oort cloud comets. In this paper, we propose a method using gravitational lensing to constrain the abundance of SULCOs in the Milky Way halo: By monitoring a large number of individual bright stars in the Local Group galaxies such as M33\citep{abe00}, we can obtain stringent constraint on the abundance of SULCOs as they induce amplification of the background source stars. The key factors are the source size and the scale of lensing time duration. In order to constrain compact objects with a sub-lunar mass via gravitational lensing, we need to have sources that are more distant than LMC or SMC for which the angular source size is smaller than the angular Einstein radius. Since the Einstein angular radius of SULCOs (typically $\lesssim 10^{-9}$ arcsec) is comparable to the radius of source stars, we need to consider the finite-source size effect \citep{witt95}. Even if the angular source size is larger than the Einstein angular radius of SULCOs, we can still detect the weak amplification of source stars if they are bright enough. Note that the finite-source size effect and the feasibility of observation with short time duration have not been explored in \citet{abe00}. Moreover, the scale of lensing time duration becomes shorter as the lens mass decreases. Thus it is important to assess whether currently available telescopes can probe SULCOs with a reasonable observation time. In what follows, for simplicity, we assume that the surface brightness of source stars is constant (limb darkening is not taken into account) and circular. | In this paper, we have discussed that nanolensing events by dark sub-lunar mass compact objects (SULCOs) such as PBHs and free-floating dwarf planets can be detected by monitoring $10^{5-7}$ stars in the Local Group galaxies such as M33. The typical lensing time scale is $\sim 100\,\tr{s}$. The exposure time for each snapshot must be $\sim 100\,\tr{s}$, which is the typical time scale of nanolensing variability for SULCOs with a mass of $\sim 10^{-9}\,\tr{M}$. Using a 8\,m class telescope, $10^{3-4}$ events per night would be detected if SULCOs constitute all the dark matter. Moreover, we expect $10^{1-2}$ events from weak amplification of very bright stars caused by SULCOs with a mass range of $10^{-11}\ms$ to $10^{-7}\ms$ though detection of such change may be a challenging task. Our method would provide a stringent constraint on the abundance of SULCOs at the distance $0.1-100\,\tr{kpc}$ from us. As a source galaxy, we have considered M33 as it has a relatively large number of blue main sequence stars and much diffuse spiral arms in comparison with M31. These features are important for observing weak amplification (which was not studied in \citet{abe00}) by SULCOs from time variability of very bright source stars. However, for the purpose of detecting nanolensing events, M31 would be much suitable as the number of available source stars is larger than M33 though the effect of blending due to neighbouring stars may be stronger as it is not face-on. For detecting SULCOs via weak amplification of source stars, other galaxies at a farther distance would be suitable if 30-m class telescope is available in the coming decade. We have taken into account the finite source-size effect, which is important for estimating the weak amplification caused by SULCOs. However, the effect of spatial variability in the source brightness such as limb darkening has not been taken into account. Such an effect becomes important when the impact parameter of the source is comparable to the radius of the star. More precise treatment is necessary for the cases in which the angular Einstein radius is approximately equal to the angular source size. A part of SULCOs may consist of free-floating dwarf planets. Detection of these objects in the intergalactic space is a challenging task. Like MACHOs, these small unbounded objects may constitute a large portion of baryonic masses in the disk or halo of our galaxy. Our method would open a new window on these small objects in the Milky Way halo that would otherwise not be observable. | 16 | 9 | 1609.04520 |
1609 | 1609.01419_arXiv.txt | {} {We investigate the nature of the compact, and possibly variable nuclear radio source in the centre of \wise, the proposed host galaxy of fast radio burst, \frb.} {We observed \wise\ at 5.0 GHz with the European VLBI Network four times between 2016 March 16 and June 2. At three epochs, we simultaneously observed the source with e-MERLIN at the same frequency.} {We detected a compact source in the EVN data in each epoch with a significance up to $\sim 8\sigma$. The four epochs yielded consistent results within their uncertainties, for both peak surface intensity and positions. The mean values for these quantities are $I_\mathrm{peak}=(115\pm9)$ \ujb\ and r.a.\ = 07$^{\rm h}$ 16$^{\rm m}$ 34.55496(7)$^{\rm s}$, dec.\ = $-19^\circ$ 00\arcmin\ 39.4754(8)\arcsec, respectively. The e-MERLIN data provided $\sim 3-5\sigma$ detections, at a position consistent with those of the EVN data. The presence of emission on angular scales intermediate between the EVN and e-MERLIN is consistent with being null. The brightness temperature of the EVN core is $T_{\rm b} \gtrsim 10^{8.5} {\rm K}$, close to the value required by Akiyama \& Johnson (2016) to explain the radio properties of \wise\ in terms of interstellar induced short-term variability.} {Our observations provide direct, independent evidence of the existence of a nuclear compact source in \wise, a physical scenario with no evident connection with \frb. However, the EVN data do not show indication of the variability observed with the VLA.} \clearpage | Fast radio bursts (FRBs) are transient episodes characterised by short (sub-ms) duration and large dispersion measure (DM). After the initial discovery by \citet{Lorimer2007}, several new such events have been discovered \citep[e.g.,][]{Thornton2013,Champion2016}, triggering debate about their nature. It is possible that they are due to young, highly magnetized neutron stars, as suggested for the repeating FRB\,121102 \citep{Spitler2016}, or to cataclysmic events. Both Galactic and extragalactic origins have been proposed. An extragalactic origin is preferred based on the large DM, $\simgt (0.5-1) \times 10^{3}$ cm$^{-3}$ pc typically found; however, only a precise localisation and a measurement of the redshift could be conclusive. For this reason, the reported localisation of \frb\ to the elliptical galaxy, WISE J071634.59$-$190039.2 (hereafter, \wise) by \citet[][hereafter, \citetalias{Keane2016}]{Keane2016} has attracted interest in the community. The precise redshift determination for this FRB has immediate implications for the system's energetics, thus possible FRB progenitors \citep{Liu2016,Zhang2016}, as well as applications to probe fundamental physics \citep{Bonetti2016,Tingay2016}. \begin{table*} \centering \caption{Log of observations, image parameters and model fit results. \label{t.log}} \begin{tabular}{lcccccccccc} \hline \hline \multicolumn{2}{c}{Epoch} & \multicolumn{4}{c}{EVN data} & \multicolumn{4}{c}{e-MERLIN data} & \\ Date in & & HPBW & $I_{\rm peak}$ & $I_{\rm noise} $& $S_{\rm 5.0, JMFIT}$ & HPBW & $I_{\rm peak}$ & $I_{\rm noise} $& $S_{\rm 5.0, JMFIT}$ & $\Delta S_{\rm 5.0}$\\ 2016 & MJD & (mas $\times$ mas, $^\circ$) & \multicolumn{2}{c}{(\ujb)} & (\uj) & (mas $\times$ mas, $^\circ$) & \multicolumn{2}{c}{(\ujb)} & (\uj)\\ \multicolumn{1}{c}{(1)} & \multicolumn{1}{c}{(2)} & \multicolumn{1}{c}{(3)} & \multicolumn{1}{c}{(4)} & \multicolumn{1}{c}{(5)} & \multicolumn{1}{c}{(6)} & \multicolumn{1}{c}{(7)} & \multicolumn{1}{c}{(8)} & (9) & (10) & (11) \\ \hline March 16 & 57463.8 & $10.1 \times 6.2, 3.9$ & 123 & 18 & $125\pm22$ & \dots & \dots & \dots & \dots \\ May 10 & 57518.6 & $9.7 \times 6.1, 8.7$ & 113 & 14 & $137\pm20$ & $261\times25, 12$ & 169 & 55 & $176\pm58$ & $40\pm60$ \\ May 31 & 57539.6 & $10.9 \times 6.1, -7.5$ & 107 & 16 & $117\pm20$ & $231\times27, 11$ & 145 & 48 & $158\pm51$ & $40\pm55$ \\ June 2 & 57541.6 & $9.3 \times 5.3, 1.3$ & 133 & 20 & $125\pm32$ & $212 \times 28, 10$ & 246 & 52 & $264\pm59$ & $140\pm70$ \\ \hline \end{tabular} \tablefoot{Cols.~(1, 2): observation date; Cols.~(3--6): EVN half-peak beam width (HPBW), naturally weighted image peak brightness and $1\sigma$ noise level, and results of a 2-d Gaussian fit to the image brightness distribution; Cols.~(7--10): same as Cols.~(3--6), for e-MERLIN data; Col.~(11): flux density difference between e-MERLIN and EVN data.} \end{table*} The proposed identification of \frb\ with \wise\ was based on the prompt detection (beginning 2 hrs after the FRB discovery) with the Australia Telescope Compact Array (ATCA) of a fading radio source within one beam of the 21-cm Parkes multi-beam receiver. Optical photometric and spectroscopic follow-up observations with the Subaru telescope identified an elliptical galaxy at $z=0.492\pm0.008$ consistent with the radio source within the $\sim1\arcsec$ positional uncertainty. In more detail, the radio transient emission was observed in only the first two epochs of ATCA follow-up separated by 6-days (flux densities $\sim 0.2$ mJy at 5.5 GHz) with three subsequent detections of essentially steady emission ($\sim 0.1$ mJy at 5.5 GHz) attributed to the emission from the host galaxy. This association of the prompt ms-duration emission from the FRB with the variable ATCA source has been questioned. \citet[][hereafter \citetalias{Williams2016a}]{Williams2016a} argued instead that \wise\ is consistent with being a random active galactic nucleus (AGN) found within the Parkes beam based on the known rate of variable (rather than transient) radio sources, that the steady radio emission component implies a large luminosity more typical of an AGN, and that the radio light curve is inconsistent with the evolution of a standard afterglow. Further criticism of the proposed association came from the results of a Karl G.\ Jansky Very Large Array (VLA) observing campaign almost a year after the FRB: in 10 total observations at 5.5 and 7.5 GHz spanning 35 days, \citet{Williams2016a,Williams2016b} found variable radio emission at an enhanced level with respect to the previously observed steady $\sim 0.1$ mJy source. \citet[][hereafter \citetalias{Vedantham2016}]{Vedantham2016} also reported a single epoch multi-frequency VLA observation over the 1-18 GHz frequency range that showed a flat spectrum radio source consistent with an AGN. Finally, numerical simulations by \citet{Akiyama2016} indicate that the reported light curve is consistent with scintillating radio emission from an AGN core with $T_\mathrm{b} \simgt 10^9 \, \mathrm{K}$. A final confirmation of the AGN scenario, plus a relevant contribution from refractive interstellar scintillation, can be obtained from high angular resolution Very Long Baseline Interferometry (VLBI) observations. In this Letter, we thus report on the results of European VLBI Network (EVN) e-MERLIN observations of \wise. In the following, we describe the observations in Sect.\ 2, present the results in Sect.~3, and discuss them in Sect.~4. | The presence of an AGN within \wise\ was already implicit in the results presented by \citetalias{Keane2016}. They reported an upper limit to the H$\upalpha$ luminosity associated with a star-formation rate of $\le 0.2 M_\odot \, {\rm yr}^{-1}$. Based on \citet{Condon1992}, this value corresponds to a radio luminosity $L_{\rm 5.0\, GHz} \le 10^{21}$\,W\,Hz$^{-1}$, about two orders of magnitude lower than that observed at the ATCA quiescence level. Our observations now provide a firm proof of the presence of a compact radio source at the centre of \wise, with a bolometric radio luminosity of $\nu L_\nu = 5.6 \times 10^{39}$ erg\,s$^{-1}$ and not variable, within the uncertainties. At first sight, this result supports the association proposed by \citetalias{Keane2016} between \frb\ and the subsequent episode of variable radio emission. However, the VLA data \citepalias[][angular resolution of $\sim 8\arcsec \times 3\arcsec$ at 5.0 GHz]{Williams2016a} indicate that variability is still present in the source more than one year after \frb: in Fig.~\ref{f.lc}, we show the light curve at 5.0 and 5.5 GHz over the time range 2016 February 27 to June 2, obtained with the VLA, EVN, VLBA, and e-MERLIN. The VLA measurements are generally higher and more variable than the EVN ones: the mean flux density and the variability index are $\langle S_{\rm VLA}\rangle = 195$ \uj, $V_{\rm VLA} = 0.49$ for the VLA, and $\langle S_{\rm VLBI}\rangle = 127$ \uj, $V_{\rm VLBI} = 0.08$ for the VLBI data (i.e.\ both EVN and VLBA). We exclude that this discrepancy is due to the presence of a secondary variable component in addition to the core imaged with EVN and e-MERLIN. Causality forbids variability on $\sim$ day time scales from diffuse emission resolved out by the EVN and e-MERLIN baselines (scale of $> 0.2\arcsec \sim 1.3$ kpc), which would also be inconsistent with the star formation rates determined by \citetalias{Keane2016}. The EVN data themselves do not show evidence for any secondary compact component either, in particular around MJD 57463, when the EVN and VLA data are nearly simultaneous, and differ by 150 \ujb; no non-nuclear sources are known to reach such a large luminosity. This requires us to explore the time, rather than the spatial, domain. It is possible, although unlikely, that the discrepancy is a chance coincidence: a K-S test on the distribution of the VLA and VLBI flux densities provides a probability that the two are drawn from the same distribution as low as 0.011. There is one significant factor to take into account: due to the different sensitivity of the two instruments, VLA data are obtained on much shorter time scales (typically, 30 minutes) than the EVN's (many hours). We can thus hypothesise that the parsec-scale source varies on short ($<$ hr) time scales, so that the VLA-based light curve resolves the variations, while they are averaged out by the longer EVN observations. The above scenario does also present some challenges. Intrinsic sub-hour time-scale variability from AGNs requires extremely large brightness temperature, exceeding the inverse Compton catastrophe limit. On the other hand, \wise\ is located at low Galactic latitude ($b=-3^\circ.2$), indicating that radio waves are subject to significant refractive scintillation in the ionized interstellar medium of the Milky Way. \citet{Akiyama2016} have argued that few-day time-scale variability of \wise\ could indeed be extrinsic, if the source has a $T_{\rm b} \gtrsim 10^9 {\rm K}$, which is consistent with our result. Scintillation has so far been studied mostly in blazars and little is known about the variability properties of weak sources; however, very rapid variations in \wise\ would at least be in agreement with the trend of increased variability found for lower flux density sources \citep{Lovell2008}. | 16 | 9 | 1609.01419 |
1609 | 1609.08012_arXiv.txt | {We contribute another anisotropy study to this field of research using Type Ia supernovae (SNe Ia). In this work, we utilise the power spectrum calculation method and apply it to both the current SNe Ia data and simulation. Using the Union2.1 data set at all redshifts, we compare the spectrum of the residuals of the observed distance moduli to that expected from an isotropic universe affected by the Union2.1 observational uncertainties at low multipoles. Through this comparison we find a dipolar anisotropy with tension of less that $2\sigma$ towards $l =171^\circ \pm 21^\circ$ and $b=-26^\circ \pm 28^\circ$ which is mainly induced by anisotropic spatial distribution of the SNe with $z > 0.2$ rather than being a cosmic effect. Furthermore, we find a tension of $\sim 4\sigma$ at $\ell=4$ between the two spectra. Our simulations are constructed with the characteristics of the upcoming surveys like the Large Synoptic Survey Telescope (LSST), which shall bring us the largest SNe Ia collection to date. We make predictions for the amplitude of a possible dipolar anisotropy that would be detectable by future SNe Ia surveys.} \begin{document} | \label{sec:intro} The observational investigation of the validity of the cosmological principle (CP) and its domain of validity is one of the main questions of modern cosmology. The isotropy of the cosmic microwave background radiation (CMB) suggests that the Universe is isotropic on very large scales (of order 100$h^{-1}$Mpc). Despite this fact, we observe cosmic structures such as voids and super-clusters in the nearby Universe. One then wonders where the transition between these two states occurs. While we have data on very large scales and also in our neighbourhood, we still require further data on the intermediate scales. This is where Type Ia supernovae (SNe Ia), which are the subject of our study in this work, could come useful. On the other hand, beside the important question of checking the validity of the CP, there are some known anomalies which are related to the possible anisotropy of the Universe. Accordingly, the search for a possible preferred axis, and hence anisotropy, in the cosmos has more critical motivations than to just look for a transition scale of the CP. These are as follows: \begin{itemize} \item \textit{Large scale velocity flows}: The scale of large scale (of order $100^{-1}$Mpc or larger) bulk flows is observed to be greater than what is expected in the standard $\Lambda$CDM cosmology \cite{Kash,Watkins,Lavaux}. \item \textit{The alignment of the CMB power spectrum low multipoles}: The directions of the normals to planes of the octopole and quadrupole moments and the dipole moments in the observed CMB map seem to point to a unified direction \cite{Eriksen,Schwa}. \item \textit{The CMB power asymmetry}: There is an indication of a power asymmetry in WMAP \cite{Eriksen:2007pc} and Planck data\cite{Ade:2013nlj}. \item \textit{Large scale alignment of quasar optical polarisation data}: It turns out that the quasar polarisation vectors point towards a common direction in the sky \cite{Hutse1,Hutse2}. \end{itemize} \vspace{2mm} If there is indeed anisotropy discovered in the Universe which is trustable and not due to systematic effects, various proposed physical effects could be responsible for such a signal. First and foremost the founding assumption upon which the standard cosmological paradigm is constructed would not hold any longer. Another possibility would be a dark energy with an anisotropic equation of state \cite{Koivis}. And also early universe models can introduce anisotropies \cite{Erickcek,Abolhasani}. In using SNe Ia as our anisotropy probe, there could occur various events which could disguise as anisotropy signals. Such these effects are the following: The intrinsic scatter of the SNe Ia (due to them not being perfect standard candles), scatter due to the location of the SNe Ia within the host galaxy and the type of the galaxy, extinction due to dust in the host galaxy, intergalactic medium and our own galaxy and finally, gravitational lensing along the line of sight to the SNe Ia, which could alter the light coming from the source \cite{Aman,Baghram}. But since most of the effects mentioned happen on the galactic scale they will be averaged out in a statistically large enough sample \cite{Kolatt}. Various methods have been employed over the years for investigating possible anisotropies in the Universe at varying scales. These are either model independent searches or ones assuming certain anisotropic models \cite{Alnes,Jain,Cai}, for a specific example we can mention dark energy models \cite{Koivisto,Blomqvist1,Blomqvist2,Camp}. With regards to the model independent approaches, many works employ the hemisphere comparison method, like \cite{Schwarz,Anton,Cai2,Tsagas,Kalus,Javan}. Other works such as that of Colin et al.'s in \cite {Colin} employ the statistical tool of `residual' they developed and analyse the data tomographically. Also the dipole anisotropy is studied \cite{Cooke,Lin} as well. All of these works find low significance anisotropies, which are mostly attributed either to systematic effects or low numbers of SNe Ia. While obviously these two sources of uncertainty are to be rectified in different ways, it is noteworthy to point out that as the number of data points (SNe Ia) increases in future data sets, the sources of systematics will be identified more clearly and therefore we will know how to try and improve upon our methodology and/or technologies we utilise at the moment. In this work we use the method of angular power spectrum calculation and apply it to the deviation of SNe Ia distance moduli from those in the standard isotropic model in order to probe anisotropy in the dipole and higher multipole moments. We also demonstrate that this method is potentially a strong probe to investigate anisotropy with future larger data sets. We should also mention that there exist studies that are pursued in the same direction \cite{Kolatt,Zhao,Carvalho}. One such important study is the work of Bengaly et al. \cite{Bengaly}. These authors consider only the low-redshift regime in order for their analysis to be model independent. Making use of the two most recent SNe Ia data sets available namely the Union2.1 and the JLA data set \cite{Betoule}, Bengaly et al. conclude that they cannot discard the possibility of the existence of a genuine anisotropy in the recent Universe. We point out here that in the present work, in contrast to recent similar studies, we probe the angular power spectrum of SNe Ia at higher redshifts, mainly with the motivation of bulk flow measurement of possibly cosmological origin. This would allow us to probe any possible discrepancy with the standard cosmological model prediction as discussed in the recent work by Colin et al. \cite{Colin:2017juj}. We should therefore observe some degree of dipolar anisotropy in the SNe Ia data, which is also found in many of the previous works albeit at low significance \cite{Colin,Schwarz,Anton}. Also all of the power spectrum investigations cited above utilise the SNe Ia data sets only, whereas we use the data as well as simulations in order to try and make predictions regarding what future data could tell us in this regard. We aim to determine a threshold value for the amplitude of a possible dipole anisotropy seeable by future supernova surveys such as the LSST \cite{LSST}. The paper is organized as follows. In section 2 we describe the methods we employed in our investigations and analyses. In section 3 we introduce the observational data we made use of in our work and investigate any detectable anisotropies out of the data. In section 4 we talk about our simulations for future surveys including dipolar anisotropy. Lastly, in section 5 we shall conclude. | In this paper we have studied the method of power spectrum calculation in order for identifying possible anisotropies in the Universe using the data of Type Ia supernovae (SNe Ia). For this, we considered using the Union2.1 SNe Ia data set (including all redshifts) and also simulating the future more abundant SNe Ia data of e.g. the LSST. To analyse the Union2.1 data set, we focused on the spectral values at low multipoles since the data were scarce. Furthermore, as discussed in section 3.1, we adopted a smoothing method which would assist us with our power spectrum calculation for the discontinuous field of our residuals. We then compared our data power spectrum with the equivalent for an isotropic universe. For this we created 1000 sky realisations for the same spatial distribution of the Union2.1 data and computed the mean spectrum of the distance modulus residuals corresponding to an isotropic universe affected by the observational uncertainties of the Union2.1 data. Through this comparison, we found $\sim 4\sigma$ tension at $\ell=4$ between the two spectra. An indication of a dipolar anisotropy exists with a significance of smaller than $2\sigma$ towards $l =171^\circ \pm 21^\circ$ and $b=-26^\circ \pm 28^\circ$. We did not detect any significant evidence for dipolar anisotropy at redshifts smaller than 0.2. We found that the observed dipole could hence be an artifact of the anisotropic spatial distribution of the high-redshift SNe Ia rather than being the result of cosmological effects. We then addressed the anisotropy search for future surveys with the aim of investigating how the increase in the number of observed SNe Ia in the future can affect the detectable threshold of a possible dipolar anisotropy in SNe Ia luminosity. To do so, we considered from 10,000 to 50,000 SNe Ia with a step size of 5000. For each set of SNe Ia, we made 700 realisations of the isotropic universe and computed the corresponding mean angular power spectrum. Then we created anisotropic simulations by adding a dipole term to the distance modulus relation (as shown in eq. \eqref{anismu}) and calculated the resulting anisotropic power spectra corresponding to each SNe Ia set. As for the uncertainties, we used the photometric uncertainties as predicted for the LSST survey in addition to adopting a conservative systematic error of 0.15 magnitudes. We calculated threshold values with 5$\sigma$ significance for possible dipole anisotropies that could be seen using the power spectrum method for increasing numbers of SNe Ia in the future data sets. As would be expected, such a dipole amplitude decreased in amplitude as the number of observed SNe Ia increased. As discussed at length in section 4, our simulations were simplistic versions of a proper one which would take into account every factor that could possibly affect the final results. As such we did not mask the Galactic plane and therefore, we can say that the dipole amplitudes we have calculated in this investigation would correspond to every direction other than the those towards the Galactic plane. Moreover, we compared the isotropic power spectra affected by the Union2.1 data observational uncertainties to the one with the predicted uncertainties of future surveys (like the LSST) with the same number and coordinates of SNe Ia as in the Union2.1 data set. Through this comparison, we showed that at low multipoles the constraining of a dipolar anisotropy can be improved up to about an order of magnitude for future surveys depending on the systematics. In this work we considered only the lowest possible moment of anisotropy. But of course as the SNe Ia data sets increase in size and accuracy of observations, possible anisotropies in higher multipole moments i.e. at smaller angular scales could also be resolved in principle. | 16 | 9 | 1609.08012 |
1609 | 1609.07465_arXiv.txt | \noindent Recently, a fully covariant version of the theory of $F(T)$ torsion gravity has been introduced \cite{ref:covariant}. In covariant $F(T)$ gravity the Schwarzschild solution is not a vacuum solution for $F(T)\neq T$ and therefore determining the spherically symmetric vacuum is an important open problem. Within the covariant framework we perturbatively solve the spherically symmetric vacuum gravitational equations around the Schwarzschild solution for the scenario with $F(T)=T + (\alpha/2)\, T^{2}$, representing the dominant terms in theories governed by Lagrangians analytic in the torsion scalar. From this we compute the perihelion shift correction to solar system planetary orbits as well as perturbative gravitational effects near neutron stars. This allows us to set an upper bound on the magnitude of the coupling constant, $\alpha$, which governs deviations from General Relativity. We find the bound on this nonlinear torsion coupling constant by specifically considering the uncertainty in the perihelion shift of Mercury. We also analyze a bound from a similar comparison with the periastron orbit of the binary pulsar \neut as an independent check for consistency. Setting bounds on the dominant nonlinear coupling is important in determining if other effects in the solar system or greater universe could be attributable to nonlinear torsion. | } The extended teleparallel gravity is a theory of gravity which is based purely on torsion, as opposed to curvature as in general relativity. In this theory the fundamental field is not the metric but instead is the tetrad. The teleparallel equivalent of general relativity (TEGR) is the one where the Lagrangian density is linear in the torsion scalar whereas the extended teleparallel gravity's Lagrangian density also possesses nonlinear terms. The action for the extended theory can be written as\footnote{Latin indices are orthonormal Lorentz indices whereas Greek indices are spacetime indices.} \begin{equation} S=\int \left\{\frac{F(T)}{2\kappa}+\mathcal{L}_{\mbox{\tiny{m}}}\right\}\,\mbox{det}[h^{a}_{\;\mu}]\,d^{4}x\,, \label{eq:gravaction} \end{equation} where $\mathcal{L}_{\mbox{\tiny{m}}}$ represents the matter Lagrangian density, $\kappa=8\pi$, $h^{a}_{\;\mu}$ is the tetrad, and $F(T)$ is a function of the torsion scalar, $T$. When $F(T)$ is linear in $T$, and the matter coupling is minimal \cite{ref:DandS}, \cite{ref:DandS2}, the theory is equivalent to general relativity. However, when $F(T)$ possesses terms non-linear in $T$ the theory can be radically different from the corresponding $F(R)$ theory. One interesting property of $F(T)$ theory not shared by $F(R)$ theory is that the resulting field equations remain second-order regardless of the form of $F(T)$. There has been much interesting work done in $F(T)$ theory, mainly in the realm of cosmology (see \cite{ref:cosmo1} - \cite{ref:cosmoend} and references therein). To a lesser extent, black holes have also been studied in $F(T)$ theory \cite{ref:bh1} - \cite{ref:bhend}. One issue that has been problematic with $F(T)$ gravity is that, in the usual formulation, the theory is not locally Lorentz invariant when $F(T)\neq T$. Therefore, inequivalent equations of motion are obtained by Lorentz transforming the tetrad field. Because of this, researchers were forced into searching for what are known as ``good'' (vs ``bad'') tetrads \cite{ref:goodbadtets}. It was not always clear what constituted a good tetrad. In general, one could appeal to the following criteria \cite{ref:goodbadtets},\cite{ref:bh2}: \begin{itemize}\label{page:criteria} \item The tetrad should not place restrictions on the form of $F(T)$. That is, the tetrad needs to retain acceptable equations of motion regardless of the function $F(T)$ and not just work well for certain functions. In the past the following conditions were proposed for acceptable equations of motion: \item The tetrad must produce equations of motion which are compatible with a symmetric stress-energy tensor $\mathcal{T}_{\mu\nu}=\mathcal{T}_{\nu\mu}$. \item The resulting equations of motion should not produce peculiar physics. For example, in spherical symmetry there should be no energy transport in the angular directions. In a static scenario, there should be no energy flux from one location to the other, etc. \end{itemize} It was realized that the reason for the lack of local Lorentz invariance in the non covariant theory was related to the fact that the theory did not contain simply the gravitational degrees of freedom, but also ones which depend on the particular Lorentz frame one is in \cite{ref:earlycov}, \cite{ref:earlycov2}, \cite{ref:davood}. This comes in via the inertial part of the spin connection\footnote{We slightly abuse nomenclature in that ``spin connection'' always refers to $\omega$ and we refer to the transport connection of the theory explicitly as the Weitzenb\"{o}ck connection regardless of index character.} \begin{equation} \omega^{a}_{\;\;b \sigma}=\Lambda^{a}_{\;\;c}\partial_{\sigma}\Lambda_{b}^{\;\;\;c}, \label{eq:spincon} \end{equation} where $\Lambda^{\cdot}_{\;\cdot}$ are the usual Lorentz transformation matrix components. Note that the purely inertial spin connection above depends only on which Lorentz frame and it can be shown to yield no curvature. It can therefore contribute to curvatureless torsion-only theory such as $F(T)$ gravity, and its ambiguity (as there are infinitely many valid Lorentz frames) is a source of introducing inertial effects in non covariant $F(T)$ theory. Therefore, in the non-covariant theory one must choose a tetrad where a {purely inertial} spin connection (the specific spin connection we are interested in will be discussed below) vanishes, for the tetrad to be a truly good tetrad. In this particular (non covariant) case the connection is then the curvatureless Weitzenb\"{o}ck connection, $\Gamma^{\alpha}_{\;\;\beta\gamma}$, but without the spin connection, and the torsion tensor, $T^{\alpha}_{\;\;\beta\gamma}$, is defined by its commutator \begin{equation} T^{\alpha}_{\;\;\beta\gamma}=\Gamma^{\alpha}_{\;\;\gamma\beta}-\Gamma^{\alpha}_{\;\;\beta \gamma}= h_{a}^{\;\;\alpha}\left(\partial_{\beta}h^{a}_{\;\gamma}-\partial_{\gamma}h^{a}_{\;\beta}\right)\,. \label{eq:nospintorsion} \end{equation} In this non-covariant version of the theory the torsion scalar is defined by \begin{equation} T := \frac14 T_{\alpha\beta\gamma} T^{\alpha\beta\gamma} + \frac12 T_{\alpha\beta\gamma} T^{\gamma\beta\alpha} - T_{\alpha\beta}^{\;\;\;\alpha}T^{\gamma\beta}_{\;\;\;\;\gamma}\,, \label{eq:nospinT} \end{equation} and the action (\ref{eq:gravaction}) is varied with respect to the tetrad degrees of freedom to yield the extended teleparallel gravitational equations of motion \begin{equation} h^{-1} h^{a}_{\;\rho} \partial_{\mu} \Big( h \frac{F(T)}{dT} S_{a}^{\;\nu\mu} \Big) - \frac{dF(T)}{dT} T_{\alpha\beta\rho} S^{\alpha\beta\nu} + \frac{1}{2} F(T) \delta_{\rho}^{\;\nu} = 8\pi \mathcal{T}_{\rho}^{\;\nu}\,. \label{eq:nonceoms} \end{equation} Here $\mathcal{T}_{\rho}^{\;\nu}$ is the usual stress-energy tensor and \begin{equation} S_{\alpha\beta\gamma} := K_{\beta\gamma\alpha} + g_{\alpha\beta} \, T_{\sigma\gamma}^{\;\;\;\sigma} - g_{\alpha\gamma} \, T_{\sigma\beta}^{\;\;\;\sigma} \label{eq:stensor} \end{equation} with \begin{equation} K_{\alpha\beta\gamma} := \frac12 \left( T_{\alpha\gamma\beta} + T_{\beta\alpha\gamma} + T_{\gamma\alpha\beta} \right)\,. \label{eq:ktensor} \end{equation} The tensors $S_{\alpha\beta\gamma}$ and $K_{\alpha\beta\gamma}$ are known as the modified torsion and contorsion tensors respectively. As long as the tetrad yields zero for the inertial spin connection (to be discussed), the tetrad is a ``good'' one and the above theory is robust. However, sometimes the criteria listed earlier for a good tetrad are not sufficient to yield a tetrad which makes this spin connection vanish. In such a case one is inadvertently introducing non-gravitational effects into the equations of motion. Historically this issue of choosing an acceptable tetrad has been a rather difficult one, as isolating the correct spin connection and formulating the correct, fully covariant, equations of motion is rather difficult. Therefore the non-covariant theory is generally used and the price to pay is that there is a chance one may accidentally introduce inertial effects, even if the tetrad chosen satisfies the criteria mentioned earlier. Recently there has been a breakthrough with the situation of covariance in $F(T)$ gravity. Kr\v{s}\v{s}\'{a}k and Saridakis have shown how to identify the correct inertial spin connection to use and how to incorporate it consistently in the equations of motion \cite{ref:covariant} to yield a fully (spacetime and orthonormal-Lorentz) covariant theory. Briefly, in \cite{ref:covariant} the curvatureless spin connection (\ref{eq:spincon}) to be chosen is constructed as follows: \begin{itemize} \item Choose any metric compatible tetrad, $h^{a}_{\;\mu}$. \item In this tetrad ``turn off'' gravity by taking $G\rightarrow 0$. This resulting tetrad is called $\hr$. \item Calculate the torsion tensor with $\hr$, including the inertial spin connection, and set it to zero. That is, form the following equation: \begin{equation} T^{a}_{\;\;\mu\nu}(\hr,\omega^{a}_{\;\;b\sigma})=\partial_{\mu}h^{a}_{\mbox{\tiny{(r)}}\nu} - \partial_{\nu}h^{a}_{\mbox{\tiny{(r)}}\mu} + \omega^{a}_{\;\;b\mu}h^{b}_{\mbox{\tiny{(r)}}\nu} - \omega^{a}_{\;\;b\nu}h^{b}_{\mbox{\tiny{(r)}}\mu} = 0 \label{eq:torinertial} \end{equation} and solve for the components of $\omega^{a}_{\;\;b\sigma}$. \item The above $\omega^{a}_{\;\;b\sigma}$ is the one to use in the covariant formulation of $F(T)$ gravity in the equations that follow. \end{itemize} With the appropriate spin connection chosen as above, the full torsion tensor, which is the commutator of the Weitzenb\"{o}ck connection now with non-zero spin connection, is calculated via \begin{equation} T^{\alpha}_{\;\;\beta\gamma}=h_{a}^{\;\;\alpha}\left(\partial_{\beta}h^{a}_{\;\gamma}-\partial_{\gamma}h^{a}_{\;\beta}\right) + h_{a}^{\;\;\alpha} \omega^{a}_{\;\;b\beta}h^{b}_{\;\;\gamma} - h_{a}^{\;\;\alpha} \omega^{a}_{\;\;b\gamma}h^{b}_{\;\;\beta} \label{eq:torsionproper} \end{equation} and the torsion scalar, the modified torsion tensor, and the contorsion tensor are all calculated as in (\ref{eq:nospinT}), (\ref{eq:stensor}) and (\ref{eq:ktensor}) using (\ref{eq:torsionproper}). The resulting equation of motion are \cite{ref:covariant} \begin{equation} \Scale[0.95]{h^{-1} h^{a}_{\;\rho} \partial_{\mu} \Big( h \frac{F(T)}{dT} S_{a}^{\;\nu\mu} \Big) - \frac{dF(T)}{dT} T_{\alpha\beta\rho} S^{\alpha\beta\nu} + \frac{1}{2} F(T) \delta_{\rho}^{\;\nu} + \frac{dF(T)}{dT} S_{a}^{\;\;\alpha\nu} h^{b}_{\;\;\rho}\omega^{a}_{\;\;b\alpha} = 8\pi \mathcal{T}_{\rho}^{\;\nu}}\,. \label{eq:eoms} \end{equation} (Some factors differ from \cite{ref:covariant} due to an overall multiplicative factor difference in our expression for $T$. Equations (\ref{eq:eoms}) were derived from scratch using our conventions to ensure that they are compatible with the quantities as defined here.) In section \ref{sec:vac} we perturbatively solve the vacuum version of the above equations of motion and obtain the correction to the Schwarzschild vacuum in $T+(\alpha/2)\;T^{2}$ gravity. We perturb about the Schwarzschild solution since it is known the even in the regime where gravity is not very weak, the Schwarzschild solution provides an excellent approximation to the spherically symmetric field. (In fact, perturbations about Schwarzschild are commonly used in general relativity studies to take into account deviations from spherical symmetry due to higher polar moments.) This allows us to not just limit our study to solar system observations but also in the arena where deviations from Minkowski spacetime may not be small but deviations from the Schwarzschild vacuum is still small, such as in the vicinity of neutron stars. In the non-covariant theory there are tetrads in the literature which yield the Schwarzschild or related Kottler vacuum as a solution to $F(T)$ gravity even for the case $F(T)\neq T$ \cite{ref:schw1}, \cite{ref:schw2}. Interesting non-Schwarzschild vacua have also been obtained in \cite{ref:schw3}. This is because, before the covariant theory, any tetrad compatible with criteria such as those listed on page \pageref{page:criteria} could be considered acceptable as it was difficult to discern what was a truly gravitational tetrad and which incorporated Lorentz frame effects. The tetrads leading to the Schwarzschild vacuum turn out to not be compatible with vanishing inertial spin connection derived via the solution to (\ref{eq:torinertial}), as required for the non-covariant theory to coincide with the covariant theory. Given the difficulty in finding the appropriate tetrad in the non-covariant theory, it is understandable that those tetrads were considered. It can be shown that in the covariant theory, the Schwarzschild solution is not a vacuum solution for $F(T)\neq T$ by plugging in the Schwarzschild solution into (\ref{eq:eoms}) (with any metric compatible tetrad now since all yield the same equations of motion in the covariant theory). In section \ref{sec:orbits} we study some properties of test-particle orbits in the vacuum and derive the perihelion shift of orbits. A comparison of this perihelion shift with the observed value and its uncertainty for Mercury and in the vicinity of neutron stars allows us to set an upper limit on the value of the magnitude of the non-linear torsion coupling, $\alpha$. The bounds on the coupling set here differ from those in previous studies (\cite{ref:solar1}-\cite{ref:solarend}) due to the fact that the current study deals with the Lorentz covariant version of the theory, which differs from the standard formulations of $F(T)$ theory when the spin connection is non-zero. The rotated tetrads utilized in \cite{ref:solar3},\cite{ref:solarend} meet the set of criteria set out on page \pageref{page:criteria} and therefore produce reasonable physics. However, those rotated tetrads do not yield zero components for the inertial spin connection, and hence the theories there differ from the Lorentz covariant theory studied here due to the fact that the torsion tensors differ. This results in slightly different torsion scalars and modified torsion tensors (\ref{eq:stensor}). One can see how the resulting equations of motion are affected by noting how these quantities enter the equations of motion, and so if these quantities differ, different equations of motion can ensue. We hasten to add that this is by no means an implication that previous works are incorrect, it is simply that they are working in a different theory than the Lorentz covariant one considered here. Interestingly, in the literature, rotated spherically symmetric tetrads which do yield zero spin connection components can be found in \cite{ref:covariant}, \cite{ref:tamgood}, \cite{ref:tampres}, and it can be checked that such rotated tetrads do produce equations of motion equivalent to the covariant version of $F(T)$ gravity. Solar system tests turn out to provide a higher bound to the value of the nonlinear coupling, due to the extremely weak nature of torsion in this regime. Given the precision of solar system measurements it is worthwhile noting this bound. In an arena where torsion is stronger, but still very close to general relativity so that the Schwarzschild solution is expected to dominate, such as in the vicinity of neutron stars, a similar bound on the nonlinear coupling can be achieved for consistency, which we also study in section \ref{sec:orbits}. We choose \neut as the system to study since there is fairly accurate data on its orbital properties. It also has a mass ratio slightly more favorable to the approximations made here than many other such binaries for which good data is available. Finally, we make some concluding remarks in section \ref{sec:conc}. | }\label{sec:conc} We have perturbatively solved for the spherically symmetric vacuum in the recently developed Lorentz covariant $F(T)$ gravity theory. The spherically symmetric vacuum is arguably one of the most interesting arenas of physical study in gravitational physics and therefore knowing its structure, even if at the perturbative level, is of relevance. Assuming that for small values of the torsion the first nonlinear term in the gravitational Lagrangian is of order $T^{2}$, we set a bound on the magnitude of the nonlinear coupling constant. The bound was set from considering the observational uncertainty in the perihelion shift of Mercury, as well as another check by considering the periastron precession of the binary system \neut. Both these calculations are consistent with an upper bound on the coupling, $\alpha$, of approximately $10^{20}$km$^{2}$. Any phenomena which one wishes to attribute to nonlinear torsion should therefore not require an $\alpha$ of greater value than this order, unless one is dealing in the strong torsion regime where $\mathcal{O}(T^{3})$ effects are expected to play an important role in the expansion of the action. It should be noted that this ``weak field'' regime actually encompasses a rather wide range of gravitational phenomena. The value of the torsion scalar $T$, even in the vicinity of \neut, is approximately of order $|T|\approx 10^{-24}$km$^{-2}$. Therefore it would require rather extreme gravitational conditions before one would be in what could be called the ``strong field'' regime. | 16 | 9 | 1609.07465 |
1609 | 1609.00209.txt | %We investigate the properties and evolution of star particles in numerical simulations of two isolated spiral galaxies and two spiral galaxies from cosmological simulations. Unlike previous numerical work, where typically each star particle represents one `cluster', for the isolated galaxies we are able to model features we term `clusters' with groups of particles. We compute the spatial distribution of stars with different ages, and cluster mass distributions, comparing our findings with observations including the recent LEGUS survey. We find that spiral structure tends to be present in older (100's Myrs) stars and clusters in the simulations compared to the observations. This likely reflects differences in the numbers of stars or clusters, the strength of spiral arms, and whether the clusters are allowed to evolve. Where we model clusters with multiple particles, we are able to study their evolution. The evolution of simulated clusters tends to follow that of their natal gas clouds. Massive, dense, long-lived clouds host massive clusters, whilst short-lived clouds host smaller clusters which readily disperse. We note that embedded clusters may be less inclined to disperse in simulations in a galactic environment with continuous accretion of gas onto the clouds than isolated clouds, and correspondingly, massive young clusters which are no longer associated with their natal clouds are difficult to form in the simulations. On average though, clusters appear to disperse fairly quickly, in basic agreement with observational measures of the number of clusters of different ages. However we caution that both the relaxation times and initial densities of clusters in the simulations are underestimated compared to real clusters, whilst both the ages and age spreads of clusters will be subject to how stellar feedback behaves in the simulations. We investigate the properties and evolution of star particles in two simulations of isolated spiral galaxies, and two galaxies from cosmological simulations. Unlike previous numerical work, where typically each star particle represents one `cluster', for the isolated galaxies we are able to model features we term `clusters' with groups of particles. We compute the spatial distribution of stars with different ages, and cluster mass distributions, comparing our findings with observations including the recent LEGUS survey. We find that spiral structure tends to be present in older (100s Myrs) stars and clusters in the simulations compared to the observations. This likely reflects differences in the numbers of stars or clusters, the strength of spiral arms, and whether the clusters are allowed to evolve. Where we model clusters with multiple particles, we are able to study their evolution. The evolution of simulated clusters tends to follow that of their natal gas clouds. Massive, dense, long-lived clouds host massive clusters, whilst short-lived clouds host smaller clusters which readily disperse. Most clusters appear to disperse fairly quickly, in basic agreement with observational findings. We note that embedded clusters may be less inclined to disperse in simulations in a galactic environment with continuous accretion of gas onto the clouds than isolated clouds and correspondingly, massive young clusters which are no longer associated with gas tend not to occur in the simulations. Caveats of our models include that the cluster densities are lower than realistic clusters, and the simplistic implementation of stellar feedback. | The evolution and properties of stellar clusters is a huge area of interest in astronomy. To date most work in this area has been observational, looking at the age and mass functions, types and distributions of star clusters in nearby galaxies (see e.g. recent reviews by \citealt{Adamo2015} and \citealt{Longmore2014}). There has been relatively little input in this area from galactic numerical simulations. Here we use galaxy simulations to study the properties of clusters (which are represented by one or many star particles depending on resolution) and compare with observations. Where clusters are represented by many star particles, we also study their evolution including in relation to their natal molecular clouds. Clusters evolve both in the context of their formation within a molecular cloud (i.e. at the embedded stage), and in terms of their dissociation from the gas of their natal cloud. Observational measures of cluster behaviour typically focus on the age and mass distributions of clusters. Cluster age distributions can be used to predict cluster evolution, since obviously a sharply declining distribution will suggest that clusters disperse much quicker than those with distributions with a flat profile. Recent observations show a constant power law profile followed by a clear downturn around 100 \citep{Silva-Villa2014} or 200 Myr \citep{Baumgardt2013} in the cluster age distributions in nearby galaxies. This appears to indicate that a significant population of clusters disperse on timescales of around 100 Myr. There is some discrepancy between results, with some groups finding steeper mass distributions \citep{Chandar2006,Chandar2014} compared to others \citep{Gieles2007a,Silva-Villa2014} even for the same galaxy. There does however appear to be a genuine environmental dependence, some galaxies (typically more quiescent galaxies such as the SMC) having flatter distributions, with less indication of cluster disruption, than others \citep{Adamo2015}. There are a number of possible mechanisms whereby clusters may disperse, or become unbound over time (see e.g. \citealt{Adamo2015}) and only a brief outline is included here. The first mechanism is so called `infant mortality', whereby the loss of gas (e.g. by stellar feedback processes) from a stellar cluster still associated with a molecular cloud changes the gravitational potential experienced by the cluster \citep{Lada1984}. Other processes include tidal effects within the cloud \citep{Elmegreen2010,Kruijssen2011, Kruijssen2012}. Other mechanisms are relevant instead to the evolution of the cluster independent of, or after the gas of the natal molecular cloud has disappeared. Again tidal effects in the galaxy may cause a cluster to become unbound, whilst the interactions of clusters with other molecular clouds may also play a role (e.g. \citealt{Spitzer1958}). Finally, N-body effects, and longer term stellar evolution may also have an effect on the cluster evolution. Another question regarding cluster evolution is the spatial evolution of the stars. Young stars are preferentially formed in the spiral arms of galaxies, but as clusters evolve they may tend to be less associated with a spiral pattern. Again, it is of interest over what timescale this occurs; the distribution of stars of different ages may vary according to the nature of the spiral arms in galaxies, whether they are density waves, short-lived transient arms, or arms induced by a bar or tidal perturbation (\citealt{Dobbs2009,Chandar2011}, Chandar al. 2016 in prep.). The distribution of stars is also often described as hierarchical (e.g. \citealt{Larson1995, Elmegreen1997, Bonnell2003, Elmegreen2014, Goul2015}), at least for young stars in clusters, often thought to be a consequence of the turbulent nature of the ISM. \citet{Grasha2015} recently find that in NGC 628 this hierarchical nature tends to disappear for clusters older than 40 Myr. Other studies also indicate that hierarchical structure reduces over timescales of $\sim 70 $ Myr \citep{Gieles2008,Goul2015}, though longer timescales are found for the LMC \citep{Bastian2009}. It is not currently computationally feasible to fully model cluster evolution in a galaxy. Even the evolution of isolated clusters is difficult to follow, with fully resolved studies of cluster formation and evolution limited to only small clusters \citep{Moeckel2010,Moeckel2012}, and these still require input regarding how much and how quickly gas disperses when a cluster forms. Although not able to follow individual stars, \citet{Fujii2016}, model more massive clusters, again combining hydrodynamic and N-body physics, and assuming the removal of gas after stars form. They find that the cluster evolution is largely dependent on the initial conditions; lower mass clouds tend to form open clusters, higher mass clouds form dense massive clusters, and lower density high mass clouds form `leaky clusters' \citep{Pfalzner2009}, clumpy clusters which more gradually evolve into less dense clusters as gas disperses. On galaxy scales, one approach is to model a cluster using an N-body code, subject to a galactic potential \citep{Baumgardt2003,Hurley2008,Renaud2013,Renaud2015,Rossi2015}, which allows the study of galactic tides, galaxy collisions and other such large scale processes on the cluster. \citet{Kruijssen2011} perform an N-body$+$SPH model of galaxies, allowing for the consistent formation of clusters in high gas density regions, and including a sub-grid model for stellar cluster evolution including mass loss from stellar evolution, two-body relaxation, and tidal shocks. The resolution of these models however is such that a single particle represents a full cluster, and the ISM and GMCs are not well resolved. In addition to simulations on galaxy scales or smaller, nowadays cosmological simulations have sufficient resolution in order for a single particle to represent a stellar cluster and at least give an indication of the spatial distribution of star clusters, although they do not follow cluster evolution. Previous numerical work has also considered the evolution of GMCs, which may well be linked with the evolution of star clusters. \citet{Dobbs2013} showed that the evolution of GMCs is highly complex, with multiple cloud mergers (see also \citealt{Dobbs2015}) leading to the formation of GMCs, and likewise the dispersal of GMCs into multiple clumps. \citet{Dobbs2013} find typical lifetimes of 10 Myr (as measured by the time over which GMCs retain at least half their mass), with longer lifetimes for the most massive clusters. GMCs are dispersed by feedback and shear. \citet{Hopkins2012} find lifetimes typically of a few 10s of Myrs. Observational estimates find lifetimes of 20-30 Myr \citep{Kawamura2009,Meidt2015}, although \citet{Whitmore2014} estimate that gas is expelled form clusters within $\sim$ 10 Myr. \citet{Dobbs2014} also investigate stellar age ranges in GMCs. In particular, they note that long-lived inter-arm clouds typically contain a larger age spread compared to GMCs in the spiral arms, which tend to predominantly contain young stars. In this paper, we primarily focus on isolated galaxy simulations, which have sufficient resolution to follow `clusters' as groups containing multiple stars. We note throughout that although we term these features `clusters' they may range from massive dense clusters to unbound associations. We also utilise cosmological simulations to compare the spatial distribution of clusters in quite different simulations, over longer time periods and with more realistically induced spiral arms. We consider a number of properties of the clusters, including spatial distribution, mass distributions and rough estimates of the age distributions. We also follow the more detailed evolution of the clusters in the isolated galaxy simulations. We do not include stellar evolution (other than stellar feedback) of the clusters, or mass loss other than the result of star particles dispersing, and likewise cannot model two-body relaxation. Instead the evolution of clusters in the isolated galaxy simulations is just driven by the immediate gas structure dynamics, gas and stellar gravity and, in relation to these, stellar feedback from the star particles. | We have examined stellar clusters in isolated galaxy and cosmological galaxy simulations. We first considered the spatial distributions of star particles and clusters in these simulations and compared them with observational surveys. Most notably, the simulated clusters (both in the case of the isolated galaxies and cosmological galaxies) display clearer spiral structure in older clusters compared to the observations. We identified a couple of possible reasons for this difference. Firstly there are few older clusters present in the observations compared to the simulations. Related to this is the fact that in the observations, clusters disperse, whereas in the cosmological simulation there is no cluster evolution, so the number of clusters does not decrease with age. Likewise the same occurs for the star particles in the isolated galaxy simulations, although if grouping stars into clusters, there are smaller numbers of older clusters and the distribution of clusters better resembles the observations. A second effect may be the global galactic structure. The simulations show particularly clear spiral structure, and likewise the observations of NGC 1566 with the strongest spiral arms show the clearest spiral structure in the older stars. We then considered the mass distributions of clusters for the isolated galaxy simulations, grouping stars together into clusters using a friends of friends algorithm. With different levels of feedback, the distributions show the same slope, but are shifted up and down. This is similar to the behaviour seen for clouds \citep{Dobbs2011new}. We also saw simply from the spatial distribution of star particles that the lower feedback model contained more massive clusters, since the feedback is less able to disperse the clouds and prevent continuing star formation and further cluster growth. The cluster distributions are slightly steeper than that for clouds, and a little shallower compared to observations. We then studied the evolution of clusters in the isolated galaxy simulations, again grouping star particles into clusters using a friends of friends algorithm. Some clusters are shown to resemble what we typically think of as clusters, i.e. groups of stars that form mostly together (although we also see indications that clusters may merge during cloud-cloud collisions) and those that are simply associations of stars that happen to be spatially coincident. We perform the analysis with different restrictions on the age spreads to try to deselect the latter examples. We compared age distributions from the clusters in the simulations with results both from observations and the theoretical models of \citet{Elmegreen2010}. We find a gradient of slope $\zeta \sim -1$ to $-0.5$, indicative of moderate to rapid dispersal, dependent on the strictness of the cluster definition. Values at the less steep end of this range are in good agreement with observations, though correspond to our least strict definition of clusters, suggesting that probably our clusters disperse faster than observed clusters. The evolution of the star clusters in the simulations appears to largely follow the evolution of the GMCs themselves -- most dispersing relatively quickly, some surviving longer. Thus the timescales for cluster dispersal (10s of Myrs) we find are not dissimilar to cloud lifetimes. They are longer than the lifetimes typically found in \citet{Dobbs2013} but there we used a fairly strict definition of lifetime that at least half the gas contained by the cloud must be the same over its lifetime. However our resolution is such that we may well be underestimating the time clusters survive after the gas disperses. This may be reflected in the timescales we find compared to observed clusters. Achieving such resolution to study clusters fully (and in particular achieve high cluster densities) is not yet viable in galaxy simulations. A second point is that massive clusters with small age spreads tend not to occur in the simulations, as massive clouds tend to build up over longer timescales (i.e. 10s of Myrs). In the simulations, continuous gas accretion onto the clouds hinders gas dispersal and cluster dissociation compared to simulations or analysis of isolated clouds. Again though, resolution may be an issue, particularly in regards to how feedback is modelled and how effective stellar feedback is in dispersing the gas. In our discussion section, we have mentioned a number of further caveats and uncertainties in our interpretation of the simulations. In particular, we have only used a very simple approach to identifying clusters, and due to our resolution they tend to be less concentrated or dense than real clusters. Again obtaining the densities of real clusters would require higher resolution that is higher than typically feasible in these type of simulations at present. We also make no attempt to follow processes such as mass loss of individual stars, or two-body effects of individual stars in the simulations which in reality may also effect cluster evolution. | 16 | 9 | 1609.00209 |
1609 | 1609.05460_arXiv.txt | {Heavy elements, even though its smaller constituent, are crucial to understand Jupiter formation history. Interior models are used to determine the amount of heavy elements in Jupiter interior, nevertheless this range is still subject to degeneracies due to uncertainties in the equations of state.} {Prior to Juno mission data arrival, we present Jupiter optimized calculations exploring the effect of different model parameters in the determination of Jupiter's core and heavy element's mass. We perform comparisons between equations of state published recently.} {The interior model of Jupiter is calculated from the equations of hydrostatic equilibrium, mass and energy conservation, and energy transport. The mass of the core and heavy elements is adjusted to match Jupiter's observational constrains radius and gravitational moments.} {We show that the determination of Jupiter interior structure is tied to the estimation of its gravitational moments and the accuracy of equations of state of hydrogen, helium and heavy elements. The location of the region where Helium rain occurs as well as its timescale are important to determine the distribution of heavy elements and helium in the interior of Jupiter. We show that differences find when modeling Jupiter's interior with recent EOS are more likely due to differences in the internal energy and entropy calculation. The consequent changes in the thermal profile lead to different estimations of the mass of the core and heavy elements, explaining differences in recently published Jupiter interior models.} {Our results help clarify differences find in Jupiter interior models and will help the interpretation of upcoming Juno data.} | S_{H}=\frac{1}{X}\Big(S_{MH13}-Y~S_{SCvH,He}\Big) \end{equation} \begin{figure*}[ht] \begin{center} \subfigure{\label{1}\includegraphics[angle=0,width=.45\textwidth]{UvsRho-H-paper.pdf}}\subfigure{\label{2}\includegraphics[angle=0,width=.45\textwidth]{UvsRho-He-paper.pdf}} \end{center} \caption{Specific internal energy as a function of density at different temperatures for hydrogen (left panel) and helium (right panel), using two different equations of state. SCvH is shown in green solid lines and the values shown in blue dotted lines correspond to the u in REOS.3 plus $\Delta u$ ($\Delta u_H=1590.12135$ for hydrogen and $\Delta u_{He}=1843.06795$ for helium) or REOS3b.} \label{internal-energy} \end{figure*} \begin{figure*}[ht] \begin{center} \subfigure{\includegraphics[angle=0,width=.45\textwidth]{sRho-H-paper.pdf}}\subfigure{\includegraphics[angle=0,width=.45\textwidth]{sRho-He-paper-2.pdf}} \end{center} \caption{Specific entropy vs. density at different temperatures for hydrogen (left panel) and helium (right panel). For hydrogen we show a comparison between the entropy calculated with REOS3b (blue), the one published in SCvH (green) and MH13+SCvH values (red). Since MH13+SCvH is a pure hydrogen table, the right panel shows a comparison between REOS3b and SCvH only. } \label{entropy} \end{figure*} with $\rho_H$ and S$_H$ the density and entropy of the pure hydrogen equation of state we extracted from MH13 table, $\rho_{SCvH,He}$ and S$_{SCvH,He}$ the density and entropy in the SCvH helium table, X$_{MH13}$, $\rho_{MH13}$ and S$_{MH13}$ the hydrogen mass fraction, density and entropy in MH13, respectively. Eq. \eqref{entropy-mixture} neglects the entropy of mixing. Detailed calculations using the SCvH EOS with and without this entropy of mixing show that this is a much smaller effect than the uncertainties on the EOSs themselves discussed here. We call this new hydrogen table MH13+SCvH (shown in appendix \ref{appendix}). \subsubsection{Entropy calculation for hydrogen and helium using REOS.3}\label{REOS3-changes} REOS.3 is a density-temperature equation of state with pressure and specific internal energy that covers a large range in pressure and temperature (figure \ref{Fig:H-phase}, for hydrogen). To allow comparisons between the tables and avoid errors in the entropy calculation, we changed the zero point of the specific internal energy in the REOS.3 tables to make them coincide in the ideal gas regime with the SCvH EOS (N. Nettelmann and A. Becker private communication). Since the difference between the specific internal energy of REOS.3 and SCvH equations of state at T=60 K and $\rho=10^{-3}$ g/cm$^3$ is $\Delta u_H=1590.12135$ for hydrogen and $\Delta u_{He}=1843.06795$ for helium, we added these values to all the specific internal energies in the REOS.3 H and He tables, respectively. Figure \ref{internal-energy} shows a comparison between the internal energies of SCvH and REOS.3 + $\Delta_u$. The entropy is a necessary parameter in internal structure calculations. The two layers considered in the model follow an adiabat, therefore the ratio between the derivatives of the entropy with respect to pressure and temperature gives us the temperature gradient in the planet's interior. We calculate the specific entropy, s, for each point of the REOS.3 table through thermodynamic relations between the published u, P,T and $\rho$ \citep{ne12}. From the definition of the Helmholtz free energy: \begin{equation}\label{F} F=U-TS \end{equation} it follows, \begin{equation}\label{s} s(T,V) = \frac{u(T,V)}{T}-\frac{1}{M}\Big(\frac{F(T,V)}{T}-\frac{F(T_0,V_0)}{T_0}\Big)+s_0 \end{equation} Since, \begin{equation}\ \frac{1}{M}\Big(\frac{F(T,V)}{T}-\frac{F(T_0,V_0)}{T_0}\Big)=\frac{1}{M}\int^{T,V}_{T_0,V_0}d\Big(\frac{F(T',V')}{T'}\Big) \end{equation} and \begin{equation}\label{FT} d\Big(\frac{F(T',V')}{T'}\Big)=\frac{dF}{T}-\frac{F}{T^2}dT \end{equation} from Eq. \eqref{F} it follows, \begin{equation} \frac{dF}{T}=\frac{d(U-TS)}{T}=\frac{dU}{T}-dS-\frac{S}{T}dT \end{equation} and \begin{equation} \frac{F}{T^2}dT=\frac{(U-TS)}{T^2}dT \end{equation} then Eq. \eqref{FT} can be written as: \begin{equation} d\Big(\frac{F(T',V')}{T'}\Big)=\frac{dU}{T}-dS-\frac{U}{T^2}dT \end{equation} using that \begin{equation} \frac{dU}{T}=-\frac{P}{T}dV + dS \end{equation} then \begin{equation} d\Big(\frac{F(T',V')}{T'}\Big)=-\frac{P}{T}dV -\frac{U}{T^2}dT \end{equation} Now, going to $\rho$ and T plane \begin{equation} \frac{1}{M}d\Big(\frac{F(T',\rho')}{T'}\Big)=\frac{P}{T}\frac{1}{\rho^2}d\rho -\frac{u}{T^2}dT \end{equation} Finally, \begin{equation} \frac{1}{M}\int^{T,\rho}_{T_0,\rho_0}d\Big(\frac{F(T',\rho')}{T'}\Big)=\int^{\rho}_{\rho_0}\frac{P(T_0,\rho')}{T_0}\frac{1}{\rho'^2}d\rho'-\int^{T}_{T_0}\frac{u(T',\rho)}{T'^2}dT' \end{equation} and going back to Eq. \eqref{s}: \begin{equation}\label{eq-entropy} s(T,\rho) = \frac{u(T,V)}{T} - \Big[\int^{\rho}_{\rho_0}\frac{P(T_0,\rho')}{T_0}\frac{1}{\rho'^2}d\rho'-\int^{T}_{T_0}\frac{u(T',\rho)}{T'^2}dT'\Big] + s_0 \end{equation} The specific entropy at each point is calculated from Eq. \eqref{eq-entropy}, using the trapezoid rule for the numerical integration and cubic splines interpolation to add temperature and density points to improve the numerical calculation. Figure \ref{entropy} shows a comparison of the entropy calculated at different temperatures with other equations of state. These new equations of state with entropy and internal energies that coincide with SCvH at T=60 K and $\rho=10^{-3}$ g/cm$^3$ are called REOS3b (see appendix \ref{appendix}). \begin{figure*}[ht] \begin{center} \includegraphics[angle=0,width=.8\textwidth]{Hugoniot.pdf} \end{center} \vspace*{-45mm} \caption{Principal Hugoniot of hydrogen (left panel) and helium (right panel). The curves were calculated for an initial state of $\rho_0= 0.0855$ g/cm$^3$ and $T_0=20$ K for hydrogen and $\rho_0= 0.123$ g/cm$^3$ and $T_0=4$ K for helium. Experimental results are shown with different point styles for comparison. We included recent estimations by \citet{br15} who presented corrections of previously published data on He \citep{eg08,ce10}, H$_2$ and D$_2$ \citep{lo12} based on a better understanding of shocked compressed SiO$_2$.} \label{figure-hugoniot} \end{figure*} \subsection{Comparison with experiments} The original equations of state MH13 and REOS.3 experienced some changes such as the creation of a pure hydrogen table and the extension of such table for a large pressure and temperature range (MH13+SCvH, section \ref{MH13-changes}), the change of the u$_0$ and entropy calculation (REOS3b section \ref{REOS3-changes}), and interpolation to add more points and make a pressure-temperature table (MH13+SCvH and REOS3b). In order to test our final tables, we make comparisons with high pressure experiments. A lot of attention has been devoted to experiments designed to understand the properties of hydrogen (or deuterium) and helium at high densities \citep{ne83,ne84,ho95,co98,be02,bo03,gr04,kn04,eg08,hi09,ce10,lo12}. In these experiments a gas at rest with an initial thermodynamic state ($u_0,\rho_0,P_0$) is exposed to an abrupt change in pressure, temperature and density. Applying the laws of conservation of mass, momentum and energy at both sides of this shock wave, we derive a relation between the state of the gas before and after the shock, called the Rankine-Hugoniot equation: \begin{equation}\label{Hugoniot} H(\rho,P)=u-u_0+\frac{1}{2}\bigg(P+P_0\bigg)\bigg(\frac{1}{\rho}-\frac{1}{\rho_0}\bigg) \end{equation} where $\rho,P,u$ are the density, pressure and internal energy of the final shocked gas. Equation \ref{Hugoniot} defines all states on the (u,$\rho$,P) surface that can be reached from the initial condition by a single shock. \subsubsection{Hugoniot-curve calculation from P, T $\rho$ and s} The Hugoniot curve, H($\rho$,P), is defined by: \begin{equation}\label{H} H(\rho,P)=0 \end{equation} Since our EOS tables give us P, T, $\rho$ and s we want to write Eq. \eqref{H} as a function of these variables. If we differentiate Eq. \eqref{H} we obtain: \begin{equation}\label{eq-dH} dH=du+\frac{1}{2}\bigg[\bigg(\frac{1}{\rho}-\frac{1}{\rho_0}\bigg)dP-\bigg(\frac{1}{\rho^2}(P+P_0)d\rho\bigg)\bigg]=0 \end{equation} Now we know that: \begin{equation}\label{eq-du} du=-PdV+Tds \end{equation} where $V=\frac{1}{\rho}$ and therefore, \begin{equation}\label{eq-dv} dV=-\frac{d\rho}{\rho^2} \end{equation} Using Eq. \eqref{eq-du} and \eqref{eq-dv} in \eqref{eq-dH}: \begin{equation}\label{eq-Hs} dH=\frac{1}{2}\bigg(\frac{1}{\rho}-\frac{1}{\rho_0}\bigg)dP+\frac{1}{2}\frac{(P-P_0)}{\rho^2}d\rho+Tds=0 \end{equation} to integrate in the P,T plane, we use: \begin{equation} d\rho(P,T)=\frac{\partial\rho(P,T)}{\partial P} dP+\frac{\partial\rho(P,T)}{\partial T} dT \end{equation} \begin{equation} ds(P,T)=\frac{\partial s(P,T)}{\partial P} dP+\frac{\partial s(P,T)}{\partial T} dT \end{equation} Equation \eqref{eq-Hs} is written as: \begin{equation}\label{eq-finaldH} \begin{split} dH(P,T)=\frac{1}{2}\bigg(\frac{1}{\rho(P,T)}-\frac{1}{\rho_0}\bigg)dP+\frac{1}{2}\frac{(P-P_0)}{\rho(P,T)^2}\frac{\partial\rho(P,T)}{\partial P} dP+\\ \frac{1}{2}\frac{(P-P_0)}{\rho(P,T)^2}\frac{\partial\rho(P,T)}{\partial T} dT+T\frac{\partial s(P,T)}{\partial P} dP+T\frac{\partial s(P,T)}{\partial T} dT \end{split} \end{equation} Integrating Eq. \eqref{eq-finaldH} between an initial point and the final state, we get the Hugoniot curve as a function of the variables present in our EOS tables: \begin{displaymath} H(P,T)-H_0=\frac{1}{2}\int_{P(H_0)}^{P}\bigg(\frac{1}{\rho(P,T(H_0))}-\frac{1}{\rho_0}\bigg)dP+\\ \end{displaymath} \begin{displaymath} \frac{1}{2}\int_{P(H_0)}^{P}\frac{(P-P_0)}{\rho(P,T(H_0))^2}\frac{\partial\rho(P,T(H_0))}{\partial P} dP+ \end{displaymath} \begin{displaymath} \int_{P(H_0)}^{P}T(H_0)\frac{\partial s(P,T(H_0))}{\partial P} dP+ \end{displaymath} \begin{equation}\label{eq-HugoniotCurve} \frac{1}{2}\int_{T(H_0)}^{T}\frac{(P-P_0)}{\rho(P,T)^2}\frac{\partial\rho(P,T)}{\partial T} dT+ \int_{T(H_0)}^{T}T\frac{\partial s(P,T)}{\partial T} dT \end{equation} To find the zeros in Eq. \eqref{eq-HugoniotCurve} we calculate $H(P,T)$ at each P and T in the EOS table and when it changes sign we do a cubic spline interpolation in P and T to find the exact values of P,T, $\rho(P,T)$ and s(P,T) that will give us $H(P,T)=0$. Figure \ref{figure-hugoniot} shows Hugoniot curves for hydrogen and helium obtained when using different equations of state and compared with experimental data. \subsection{Heavy elements} \label{section:heavies} Hydrogen and helium are the most relevant species, but an accurate description of Jupiter's interior needs a definition of the heavy elements equation of state. In our model heavy elements are water and rocks, and we use three different equations of state to test their sensitivity. Following \citet{sg04} we use for rocks the equation of state for a mixture of silicates called "dry sand" in SESAME \citep{lj92}. For water we use the SESAME EOS \citep{lj92}, and a more recent equation of state calculated in \citet{va13}, which combines an equation of state for water at high temperatures (T>1000K) \citep{fr09} with results taken from NIST database \citep{sw89}. | Jupiter reservoir of heavy elements is key to understand the origin of our Solar system. Nevertheless, the distribution and amount of heavy elements in its interior is difficult to constrain and degeneracies arise depending on assumed observational constrains and model parameters in interior structure calculations. We present Jupiter optimized models, where the mass of the core and the mass of heavy elements are adjusted to reproduce Jupiter's radius, $J_2$ and $J_4$. We show how our solutions change drastically with the EOS for hydrogen and helium and also explore the sensitivity to heavy elements equations of state, separation between metallic and molecular envelope and distribution of heavy elements in Jupiter's interior. We adopt two different models for Jupiter, both scenarios consider helium phase separation and correspondingly different helium abundance in the outer and deeper layer. The difference is in the heavy elements distribution: one scenario has an homogeneous distribution of heavy elements and its mass mixing ratio is adjusted according to the observables. In the second scenario, Jupiter has different compositions of heavy elements in the two layers and the difference in the abundance in the outer and deeper envelope ($\Delta Z$) is adjusted to find solutions that best reproduce Jupiter observational data. Allowing a change in heavy elements between the two layers adds a degree of freedom to the problem, which grants more solutions in the M$_Z$-M$_{core}$ space. The pressure at which the separation between the two envelope layers occurs affects the solutions. This separation occurs between 0.8 and 4 Mbar, according to \citet{mo13} helium rain studies. We find that M$_Z$ decreases and M$_{core}$ increases when P$_{sep}$ moves from high to low pressures. Based on the works by \citet{SCvH95,MH13,REOS3}, we explored hydrogen and helium equations of state and show that significant differences remain in these EOSs, although they match experimental data obtained by compression experiments along a Hugoniot. Some of the differences come from internal energy and entropy calculations. We show how small changes in the internal energy lead to differences in the entropy calculated which in turn affect the thermal profile and the estimation of the mass of the core and heavy elements. This explains differences seen in recently published interior models of the planet. Jupiter internal structure has a much large temperature when using REOS3b than with SCvH. For densities $\rho>0.22246$ g/cm$^3$, MH13+SCvH leads to much lower temperatures than the other two EOS. This differences in the thermal structure lead to differences in the derived M$_{core}$ and M$_Z$. MH13+SCvH allows larger M$_{core}$ and smaller M$_Z$ while REOS3b has larger M$_{core}$ but similar M$_Z$ than results find with SCvH. In our baseline simulations, MH13+SCvH leads to M$_{core}$ between 11 and 17 M$_{Earth}$, in agreement with results by \citet{MH13} and the preferred model of \citet{hm16}. REOS3b leads to M$_{core}$ between 7 and 16 M$_{Earth}$, larger than estimations by \citet{ne12} and \citet{REOS3}. While their preferred model has P$_{sep}\ge$ 4 Mbar, our models put the separation between Z$_{atm}$ and Z$_{deep}$ in the same place as the helium phase transition, between 0.8 and 4 Mbar \citep{mo13} and the baseline simulations have P$_{sep}$=2 Mbar. When comparing the results at P$_{sep}$= 4 Mbar we find a lower limit for the mass of the core of 4 M$_{Earth}$, consistent with the small core hypothesis showed by \citet{ne12} and \citet{REOS3} for the same case. Other small differences are due to different model parameters such us the temperature at the 1 bar limit, equation of state used for solids and differences in entropy calculation. The equation of state for the heavy elements is also relevant. We study three different equations of state for rocks and water. Dry sand SESAME \citep{lj92} allows smaller M$_{core}$, while M$_Z$ increase when using $H_2O$ SESAME \citep{lj92} when compared with solutions obtained with hot water NIST EOS \citep{va13}. Our results help in the interpretation of Jupiter observational data. Its gravitational moments changed from the first pre-Juno data \citep{cs85} to the constrains we have today (Jacobson, 2013). They also change according to the dynamics and rotation of Jupiter adopted in the model. Given the relatively large scatter in the gravitational moments of Jupiter inferred between 1985 and today, in our baseline simulations we chose to use conservative $2\sigma$ error bars based on the published value of \citet{cs85} which encompass all of these values. We also show how different Js lead to different estimations of the core and heavy elements masses having a difference of up to 4M$_{Earth}$ in M$_{core}$ and $\sim$6M$_{Earth}$ in M$_Z$ for REOS3b and SCvH. Our preferred results have larger $J_6$ than the ones currently published. Juno mission will provide more accurate data, improving our knowledge of Jupiter internal structure. | 16 | 9 | 1609.05460 |
1609 | 1609.08632_arXiv.txt | {Future large scale structure surveys will provide increasingly tight constraints on our cosmological model. These surveys will report results on the distance scale and growth rate of perturbations through measurements of Baryon Acoustic Oscillations and Redshift-Space Distortions. It is interesting to ask: what further analyses should become routine, so as to test as-yet-unknown models of cosmic acceleration? Models which aim to explain the accelerated expansion rate of the Universe by modifications to General Relativity often invoke screening mechanisms which can imprint a non-standard density dependence on their predictions. This suggests density-dependent clustering as a `generic' constraint. This paper argues that a density-marked correlation function provides a density-dependent statistic which is easy to compute and report and requires minimal additional infrastructure beyond what is routinely available to such survey analyses. We give one realization of this idea and study it using low order perturbation theory. We encourage groups developing modified gravity theories to see whether such statistics provide discriminatory power for their models.} \arxivnumber{1609.08632} \begin{document} | The observation that the expansion rate of the Universe is accelerating is one of the most puzzling aspects of our current cosmological model. Two classes of explanation have been investigated, one based on a modification of the contents of the Universe \cite{Wei13,PDG14} and one based on a modification of gravity (see e.g.~Refs.~\cite{JaiKho10,Cli12,Joy15,Hut15} for recent reviews). At present there are no theoretically consistent, observationally allowed models which provide cosmic acceleration through modifications to gravity. Therefore, observational constraints on modifications to general relativity (GR) often focus on `generic' features that some as-yet-to-be-determined models might be expected to have. Large-scale structure surveys typically provide constraints on the distance scale and rate-of-growth of fluctuations through studies of baryon acoustic oscillations (BAO) and redshift-space distortions (RSD) \cite{Wei13,PDG14}. In combination with gravitational lensing surveys (of galaxies or the cosmic microwave background) a number of tests of GR on linear scales can be constructed \cite{Wei13,PDG14}. It is reasonable to assume that such analyses will be an integral part of analyses of future surveys as well, improving some tests of dark energy and modified gravity models. A question then arises what other analyses should `routinely' be performed on future surveys, so that observational constraints are available for theorists and phenomenologists seeking to constrain next-generation models? Ideally these analyses should be simple to perform and report, while at the same time providing information beyond the standard analyses currently published. In the absence of a specific theoretical framework this is a difficult question to answer. The tightest constraints will come from a model-by-model analysis, but more generic constraints can also be useful when investigating wide classes of models. One frequently encountered phenomenon in modified gravity models is a screening mechanism that forces model predictions to approach those of GR in regions of high density or strong gravitational potential. Conversely, signatures of modified gravity will show up in regions where gravity is weak. Screening is a property of many modified gravity models that offers potentially distinctive observational signatures. Here we advocate the use of the density-marked correlation function \cite{WhiPad09} as an easy-to-compute statistic which may test future modified gravity models. Computation of the marked correlation function requires minimal modification to existing analysis frameworks, and requires no further infrastructure beyond that which is routinely available for studying BAO and RSD. By weighting the pairs of galaxies by a `mark' which depends on a local density estimate (e.g.~increasing the weight of low density regions) it provides density-dependent information from the survey which may be of use in constraining future theories. This paper introduces the inverse-density-marked correlation function and provides an exploration of its properties using low order Lagrangian perturbation theory. The latter is primarily for convenience -- such statistics can also be calculated and explored using N-body simulations or mock catalogs which will give access to smaller scales where the signal may be larger. We shall return to this point in the conclusions. The outline of this paper is as follows. In section \ref{sec:background} we introduce the marked correlation function and quickly review Lagrangian perturbation theory and how to compute the marked correlation function within this framework. We present some results to build intuition on the marked correlation function in section \ref{sec:results}. Finally we conclude in section \ref{sec:conclusions}. | \label{sec:conclusions} The growth of large-scale structure as observed in modern cosmological surveys offers one means of testing general relativity on the largest scales. Constraints on the distance scale, growth rate and deflection of light (from BAO, RSD and lens modeling respectively) have become standard and such analyses are likely to be performed on all future surveys. In the absence of a single, compelling model of modified gravity it is difficult to know how to augment these `standard' analyses so as to best constrain modifications. Many models which modify gravity invoke a screening mechanism that forces the predictions to become equal to those of GR in regions of high density or strong gravitational potential. This suggests that generic tests for the density dependence of the growth of structure could be added to our list of `standard' analyses and might provide useful in constraining models of modified gravity in the future. In this paper we have pointed out that an inverse-density-marked correlation function, $\mathcal{M}(r)$, is very easy to compute and report, requiring little computational overhead, code development or additional survey products. To illustrate the form that $\mathcal{M}(r)$ takes in the standard theory, we present the lowest order Lagrangian perturbation theory calculation. In keeping with our perturbative approach we have emphasized large scales, such as can be easily probed with upcoming surveys aimed at BAO and RSD. Of course it may be that modifications to gravity are more easily seen on smaller scales with density fields estimated from denser samples of galaxies (or other objects). While this likely invalidates the perturbative calculation presented earlier, it is relatively straightforward to generate predictions on such scales from simulations given a suitably refined model for bias (e.g.~populating halos in N-body simulations with galaxies using an HOD model \cite{CooShe02} or a semi-analytic model \cite{Bau06}). Indeed, even if such statistics prove unconstraining for the modified gravity models of the future, the density information they encode can be very valuable for validating and refining the bias model used to model BAO and RSD statistics, e.g.~breaking degeneracies \cite{WhiPad09} or testing for assembly bias. Should the population of galaxies in a survey depend on properties of a dark matter halo beyond its mass (e.g.~formation time or concentration) which correlates with large-scale environment we expect to see a departure from the simplest theoretical models which neglect such effects. The nature of this departure can be understood from the theory above and parameterized in a very flexible way. Since the GR predictions are so well known, as long as modifications to gravity are not fully degenerate with these effects, we should still be able to disentangle them. There are several obvious lines of development. First, this statistic could be computed on existing modified gravity simulations to get a sense for the size of the effect as a function of scale and the ideal redshift and number density of tracer for this test. This would help to calibrate expectations, but cannot be taken as definitive since existing models are either ruled out or do not explain acceleration with modified gravity and we do not know how the predictions would differ in a model which could explain our Universe. Second, a study should be undertaken of the best density estimate and to what extent noise in this estimate adversely affects the results. Third, if further investigation warrants, this model can be extended to higher order in perturbation theory or to include a dependence on the derivatives of the density or the dimensionality of the structure. Within the Lagrangian framework it is relatively straightforward to include marks which can be expressed in terms of initial density and its derivatives. Finally, it is worth investigating whether marked statistics for auto- and cross-correlations of imaging and spectroscopic surveys could yield other, valuable constraints on modifications to GR. It is straightforward \cite{Mat08b,CLPT,WanReiWhi14} to modify the formulae in this paper to account for cross-correlations of biased and marked tracers, thus opening the possibility to predict a range of other statistics. | 16 | 9 | 1609.08632 |
1609 | 1609.01713_arXiv.txt | Here we utilise recent measures of galaxy total dynamical mass and X-ray gas luminosities (L$_{X,Gas}$) for a sample of 29 massive early-type galaxies from the SLUGGS survey to probe L$_{X,Gas}$--mass scaling relations. In particular, we investigate scalings with stellar mass, dynamical mass within 5 effective radii (R$_e$) and total virial mass. We also compare these relations with predictions from $\Lambda$CDM simulations. We find a strong linear relationship between L$_{X,Gas}$ and galaxy dynamical mass within 5 R$_e$, which is consistent with the recent cosmological simulations of Choi et al. that incorporate mechanical heating from AGN. We conclude that the gas surrounding massive early-type galaxies was shock heated as it fell into collapsing dark matter halos so that L$_{X,Gas}$ is primarily driven by the depth of a galaxy's potential well. Heating by an AGN plays an important secondary role in determining L$_{X,Gas}$. | Massive early-type galaxies (ETGs) reveal halos of hot gas. This X-ray emitting diffuse gas varies in luminosity by a factor of $\sim$100-1000 for a given optical luminosity, or equivalently, stellar mass (O'Sullivan, Forbes \& Ponman 2001; Fabbiano 2006; Boroson, Kim \& Fabbiano 2011). It has been suggested that stellar mass loss dominates the source of this gas (Sun et al. 2007; Sarzi et al. 2013). However, in a large X-ray study of ETGs, Goulding et al. (2016) concluded that the data are not consistent with a simple stellar mass loss picture. An additional source of gas is indicated by the cold dark matter (CDM) paradigm for galaxy formation. Here infalling pristine gas in massive halos is shock-heated to X-ray emitting temperatures as these halos collapse. The hot gas slowly cools while emitting X-rays, with the X-ray luminosity directly related to the depth of the potential well (White \& Frenk 1991). In the GIMIC simulations of Crain et al. (2010), the density of the hot gas is less concentrated than that of the canonical dark matter density profile. This leads to much lower X-ray luminosities than predicted by White \& Frenk (1991), which is then closer to the X-ray luminosity as observed from late-type galaxy halos. However, for early-type galaxies, AGN heating of hot gas is also potentially important. Recently, Choi et al. (2015) included AGN in cosmological models of massive ETGs. They found better agreement with the X-ray luminosities of ETGs than when using models without AGN feedback, although some discrepancies remained. These cosmological simulations all suggest that the key parameter determining the luminosity of the diffuse X-ray gas (L$_{X,Gas}$) is the galaxy mass (which determines the depth of the potential well) with AGN perhaps playing a secondary role in ETGs. Combining new {\it Chandra} data with results from the literature, Kim \& Fabbiano (2013; KF13) have shown that the dynamical (i.e. baryonic plus dark matter) mass may indeed be the key parameter in determining L$_{X,Gas}$ surrounding an ETG (see also Mathews et al. 2006). In particular, they showed that measurements of L$_{X,Gas}$ correlated strongly with the dynamical mass within 5 R$_e$ (where R$_e$ is the effective radius of the galaxy optical light) for a sample of 14 early-type galaxies. The total mass measurements within 5 R$_e$ came from the work of Deason et al. (2012) who used a combination of planetary nebulae (PNe) and globular cluster (GC) kinematics from the literature, and the assumption of power-law mass and tracer density profiles. Here, we use a homogeneous analysis of dynamical masses from GC kinematics within 5 R$_e$ for 29 ETGs (Alabi et al. 2017). This analysis is based on data from the SLUGGS survey (Brodie et al. 2014) and supplemented by data from the literature. X-ray gas luminosities mostly come from the compilation of Kim \& Fabbiano (2015; hereafter KF15) and are extracted from radii which vary from 30 to 240 arcsec ($\sim$ 2-5 R$_e$.) We investigate the trends of the X-ray gas luminosity for our sample with stellar mass, dynamical mass within 5 R$_e$ and with extrapolated virial mass. We compare these new observations with simulated galaxy halos from Crain et al. (2010) and Choi et al. (2015), and discuss whether the hot gas halos around ETGs are consistent with an origin, and heating mechanisms, as described by the $\Lambda$CDM paradigm. | Using new dynamical masses we find a strong linear relationship between the diffuse gas X-ray luminosity for a sample of 29 nearby massive early-type galaxies and their total dynamical mass within 5 R$_e$. This result supports the earlier analysis by KF13 based on 14 galaxies. We also show that the cosmological simulations of Choi et al. (2015), which incorporate AGN mechanical heating, now agree with observations of the X-ray luminosity and mass within 5 R$_e$. This good agreement with model galaxies in a cosmological framework supports the idea that the diffuse X-ray luminosity is primarily driven by the depth of a galaxy's potential well, with a significant contribution from AGN mechanical heating. Supernova heating may contribute in lower mass galaxies. | 16 | 9 | 1609.01713 |
1609 | 1609.03520_arXiv.txt | Knowledge of the stellar content near the Sun is important for a broad range of topics ranging from the search for planets to the study of Milky Way structure. The most powerful method for identifying potentially nearby stars is proper motion (PM) surveys. All old optical surveys avoid, or are at least substantially incomplete, near the Galactic plane. The depth and breadth of the ``Vista Variables in V\'ia L\'actea'' (VVV) near-IR survey significantly improves this situation. Taking advantage of the VVV survey database, we have measured PMs in the densest regions of the MW bulge and southern plane in order to complete the census of nearby objects. We have developed a custom PM pipeline based on VVV catalogues from the Cambridge Astronomy Survey Unit (CASU), by comparing the first epoch of $JHK_{\rm S}$ with the multi-epoch $K_{\rm S}$-bands acquired later. Taking advantage of the large time baseline between the 2MASS and the VVV observations, we also obtained 2MASS-VVV PMs. We present a near-IR proper motion catalogue for the whole area of the VVV survey, which includes 3003 moving stellar sources. All of these have been visually inspected and are real PM objects. Our catalogue is in very good agreement with the proper motion data supplied in IR catalogues outside the densest zone of the MW. The majority of the PM objects in our catalogue are nearby M-dwarfs, as expected. This new database allow us to identify 57 common proper motion binary candidates, among which are two new systems within 30~pc of the Sun. | A complete census of stars within the solar neighbourhood out to a specified distance will inform us about the stellar mass function, star formation, and the kinematics of the Galaxy and of young, nearby clusters and moving groups. The main difficulty in constructing a volume-limited sample is identification of nearby, low-mass objects because of their low luminosity. Also, accurate distance measurements for these stars are not easy to obtain. The most powerful method for identifying potential nearby stars comes from PM surveys. PM surveys continually improve, as longer time baselines increase the accuracy of the measurement. A uniform census of nearby stars allows characterisation of the relative occurrence rates of different types of stars, and allows relationships between intrinsic properties of those stars, including absolute magnitude and colour, to be examined. The first attempts at large surveys for high proper-motion (HPM) stars began in the early 20th century with works by \citet{1915PASP...27..240V}, \citet{1919VeHei...7..195W} and \citet{1939AJ.....48..163R}. Later, additional surveys were completed: e.g., \citet{1971lpms.book.....G, 1978LowOB...8...89G}. The first all-sky search exploiting these initial photographic surveys for nearby stars was by Luyten, who published two PM catalogues: the Luyten Half-Second catalogue \citep[LHS:][]{1979lccs.book.....L}, and the New Luyten Two-Tenths catalogue \citep[NLTT:][]{1979nlcs.book.....L, 1979NLTT..C01....0L, 1980nltt.bookQ....L, 1980nltt.bookR....L}. After these early works many papers concerning proper motion studies were published. \cite{2005AJ....129.1483L} compiled a list of 61977 stars in the northern hemisphere with $\mu > 0.15$ arcsec yr$^{-1}$, identifying over 90\% of those stars down to a limiting magnitude of V$\approx$19.0, excluding the Galactic plane, and southern hemisphere (\citet{2008AJ....135.2177L}, $0.45 \le \mu \le 2.0\,\arcsec{\rm yr^{-1}}$). \citet{2005AJ....130.1247L} reported the discovery of 182 southern stars ($\delta < -30^\circ$) with proper motion $0.45 < \mu < 2.0\,\arcsec{\rm yr^{-1}}$. \cite{2005AJ....130.1658S} reported the discovery of 152 new high proper motion systems ($\mu\ge 0.4\,\arcsec {\rm yr^{-1}}$) in the southern sky ($\delta=-47^\circ$ to $00^\circ$) brighter than UKST plate $\rm R_{59F}=16.5$ via their SuperCOSMOS-RECONS (SCR) search. \cite{2008AJ....135.2177L}, completing their SUPERBLINK proper motion survey in the southern celestial hemisphere, found 170 additional new stars with proper motions $0.45 < \mu < 2.0\,\arcsec{\rm yr^{-1}}$. This final part of their search covers 11,600 deg$^{-2}$ in the declination range $-30^\circ < \delta < 0^\circ$ and in low Galactic latitude areas south of $\delta$ = -30$^\circ$ which had not been covered in earlier data releases. Most of the new discoveries were found in densely populated fields along the Milky Way, toward the Galactic bulge/center. Their total list of high proper motion stars recovered by SUPERBLINK in the southern sky contains 2228 stars with proper motions $0.45 < \mu < 2.0\,\arcsec{\rm yr^{-1}}$. \cite{2011AJ....142...10B}, as a continuation of the SCR search in the southern sky, presented 2817 new southern proper motion systems with $0.18 < \mu < 0.40\,\arcsec{\rm yr^{-1}}$ and declinations between $-47^\circ$ and $00^\circ$. Subsequently, \cite{2011AJ....142...92B} published 1584 new southern proper motion systems with $\mu > 0.18\,\arcsec{\rm yr^{-1}}$ and $16.5 > R_{59F} \geq 18.0$. This search complemented the six previous SCR searches of the southern sky for stars within the PM motion range, but shallower than $R_{59F} = 16.5$. \cite{2011AJ....142..138L} published an all-sky catalogue of M-dwarfs with apparent infrared (IR) magnitude $J<10$. The 8889 stars were selected from the on-going SUPERBLINK survey of stars with $\mu>40$~mas yr$^{-1}$, supplemented on the bright end with the Tycho-2 catalogue. Completeness tests suggest that this catalogue represents $\sim$75\% of the estimated $\sim$11900 M-dwarfs with $J<10$ expected for the entire sky. This catalogue is, however, significantly more complete for the northern sky ($\approx$90\%) than it is for the south ($\approx$60\%). \citet{2011AJ....141...97W} presented a spectroscopic catalogue of 70841 visually inspected M-dwarfs from the seventh release of the Sloan Digital Sky Survey. Recently, \citet{2013AN....334..176L}, again using the SUPERBLINK PM survey, reported a catalogue of $\sim$200,000 M-dwarfs in the northern sky. They presented a new census of $\sim$100,000 M-dwarfs located within 100 pc. The new census is 95\% complete to 50 pc, and $>$75\% complete to 100~pc. It was followed by the spectroscopic catalogue of the brightest ($J$$<$9) M dwarf candidates in the northern sky \citep{2013AJ....145..102L}. \citet{2013MNRAS.435.2161F} used the Position and Proper Motion Extended-L (PPMXL) catalogue \citep{2010AJ....139.2440R} and applied optical and near-IR colour cuts together with a reduced proper motion (RPM) cut to find bright M-dwarfs for future exoplanet transit studies. Recently, \citet{2014MNRAS.437.3603S,2014MNRAS.443.2327S} presented two new infrared PM catalogues based on the UKIDSS Large Area Survey and Galactic Plane Survey. \citet{2014ApJ...781....4L} used multi-epoch astrometry from the Wide-field Infrared Survey Explorer (WISE) to identify 762 high proper motion objects, 761 of which were detected also by the Two Micron All Sky Survey (2MASS). \citet{2014ApJ...783..122K}, using the AllWISE processing pipeline, have measured motions for all objects detected on WISE images taken between 2010 January and 2011 February. They found 22445 objects that have significant AllWISE motions, of which 3525 have motions that can be independently confirmed from earlier 2MASS images, yet lack any published motions in SIMBAD. M-dwarfs are the most abundant inhabitants of our Galaxy and also are probably the most common sites of planet formation \citep{2006ApJ...640L..63L}. They account for over 70\% of stellar systems in the solar neighbourhood \citep{1997AJ....114..388H}. In addition, the single star fraction -- a crucial statistic for giant planet formation \citep[e.g.][]{2012ApJ...745...19K} -- decreases from $\sim$60-70\% for M-dwarfs \citep{1992ApJ...396..178F, 2010A&A...520A..54B} to $\sim$54\% for solar-type stars \citep{1991A&A...248..485D, 2010ApJS..190....1R} to near 0\% for the most massive stars \citep{1999NewA....4..531P}, further separating M-dwarfs from AFGK stars as the most numerous potential planet hosts of all the stellar classes \citep{2006ApJ...640L..63L}. About 25\% of all Doppler-confirmed planets with $M \sin i < 30$~$M_\oplus$ are orbiting M-dwarfs. Large exoplanet surveys have now started to monitor sizable numbers of M-dwarfs, such as the M2K program which is targeting some 1600 M-dwarfs for radial velocity (RV) monitoring \citep{2010PASP..122..156A} and the MEarth project \citep{2014arXiv1409.0891I} which is designed to detect exoplanet transits in nearby late-type M-dwarfs. The principal methods of exoplanet detection, RV and transits, are both more sensitive to planets around stars of lower and substellar mass. The brightest M-dwarfs are the ideal (highest priority) targets for high precision RV searches, the latest M-dwarfs are the most suitable for transit surveys, and the youngest M-dwarfs are preferable targets for Adaptive Optics (AO) imaging. The above-mentioned factors make searches for new, nearby, and young low-mass stars and substellar objects highly valuable. All old optical surveys avoid or are at least substantially incomplete near to the Galactic plane. In the Subsection\,\ref{sec:catalog} we present a detailed comparison with previous PM studies covering the VVV area. There are two modern surveys that make an exception: The INT Photometric H$_\alpha$ Survey of the Northern Galactic Plane (IPHAS) \citep{2005MNRAS.362..753D} and The VST Photometric H$_\alpha$ Survey of the Southern Galactic Plane and Bulge (VPHAS+) \citep{2014MNRAS.440.2036D}. Nevertheless, this region is still referred to as the ``zone of avoidance'' as it contains the highest stellar densities down to faint limiting magnitudes in addition to regions with dark molecular clouds, nebulosity, and current star formation, which produces substantial confusion. Nevertheless, the Galactic plane and the Galactic bulge offer considerable latent potential for new discoveries of nearby low-mass stars and ultra-cool dwarfs (UCDs) from deeper searches. Such discoveries may contain young, unusual, and nearby/bright examples of these objects, and will also complement those made at higher Galactic latitudes \citep{2012MNRAS.427.3280F}. Good recent examples of nearby ($<$10 pc) interesting discoveries at low Galactic latitude are those of UGPS~J0722-05 \citep[T9:][]{2010MNRAS.408L..56L}, DENIS~J081730.0-615520 \citep[T6.5:][]{2010ApJ...718L..38A}, and the amazing discoveries of the nearest brown dwarfs WISE~J104915.57-531906.1AB \citep{2013ApJ...767L...1L} and WISE~J085510.83-071442.5 \citep{2014ApJ...786L..18L} at $\sim$2~pc from the Sun. These discoveries near to the Galactic plane highlight the important serendipitous nature in which new IR surveys like VVV \citep{2010NewA...15..433M} can improve on the incomplete Solar neighbourhood census of low-mass stellar and sub-stellar systems. A positive aspect of the high stellar densities encountered in the Galactic plane and bulge is the plethora of bright reference stars for good AO tip-tilt low-order correction. This will facilitate high-Strehl imaging measurements to identify very low-mass brown dwarf/planetary-mass, companions for studying multiplicity, and also measuring dynamical masses. Moreover, high-Strehl AO studies of newly identified binary moving group members could also provide good age and composition constrains, as well as dynamical masses, enabling direct feed-back to evolutionary models. The crowded fields are also perfect for time series and astrometric studies because they provide many suitable reference stars. The VVV near-IR ($ZY JHK_{\rm S}$) survey \citep{2010NewA...15..433M, 2012A&A...544A.147S} covers 562 deg$^2$ of the Galactic bulge and the Southern Galactic disk and provides accurate photometry, and multi-epoch $K_{\rm S}$-band imaging, enabling us to discover a meaningful sample of new nearby cool and ultra-cool dwarfs (UCDs: spectral types $>$M6) with a higher completeness than has previously been achieved in the low Southern Galactic latitudes (e.g., from 2MASS and Deep Near Infrared Survey of the Southern Sky (DENIS)), from a PM search. The first study of the PM objects using VVV data was made by \citet{2013AA...560A..21I}. The common proper motion method was used and seven new co-moving companions around known HPM stars were discovered. \citet{2013A&A...557L...8B} reported the discovery of the first VVV brown dwarf -- VVV\,BD001 -- a new member of the 20 pc sample with well defined proper motion, distance, and luminosity. For our initial search we limited our sample selection to the brightest ($K_{\rm S}<13.5$ magnitude) objects only. A more complete and deeper catalogue will be published later \citep[Smith et al., in prep]{smith2015a}. Here we present the first VVV HPM catalogue limited to $K_{\rm S}\le13.5$ using the VVV databases. | This work has identified 3003 PM stars (mostly late K- and M-dwarfs) with magnitudes of $K_{\rm S}$<13.5 and PM$<$30~mas\,yr$^{-1}$ from the VVV catalogues. The completeness in its brighter part ($K_{\rm S}$<9) is comparable with the completeness of the similar catalogues outside the ``zone of avoidance'' near to the MW bulge and disk. We found 57 wide CPMBs all dK+dM or dM+dM binaries. We started an intense spectral follow-up of the most interesting candidates \citep{gromadzki2015}. Low-resolution spectra will confirm the spectral types of the objects and higher-resolution spectra will provide constraints on the RV. Such observations would allow prioritisation of bright younger M-dwarfs for light-curve follow-up and transit searches and also for AO imaging for searching for nearby companions. The incompleteness of our sample becomes significant at PM$\leq$100 mas\,yr$^{-1}$. One hundred thirty-four of the catalogue stars have PM$>$300 mas\,yr$^{-1}$ and 42 of them are newly found HPM objects. 382 catalogue stars have PM$>$200 mas\,yr$^{-1}$ and 179 of them are new HPM objects, 1576 stars showing PM>100 mas\,yr$^{-1}$ -- 1247 are new ones. The star with the highest proper motion in the VVV area (except $\alpha$~Centauri) is HD~156384C with PM=1175 mas\,yr$^{-1}$ and the new found star with the highest PM motion is VVV~J180414.62-312937.18 with PM=810 mas\,yr$^{-1}$. Here we limited our search only for the brightest objects. The three thousand HPM stars found are only the tip of the iceberg. Most tiles in the VVV database have limiting magnitudes $K_{s,{\rm lim}}\sim17$--18 \citep[see the VVV DR1 paper][]{2012A&A...544A.147S}, and we expect that the final catalogue of VVV HPM stars will contain $>$10$^5$ objects. | 16 | 9 | 1609.03520 |
1609 | 1609.02185.txt | A current issue in the study of planetary nebulae with close binary central stars is the extent to which the binaries affect the shaping of the nebulae. Recent studies have begun to show a high coincidence rate between nebulae with large-scale axial or point symmetries and close binary stars. In addition, combined binary-star and spatio-kinematic modeling of the nebulae have demonstrated that all of the systems studied to date appear to have their central binary axis aligned with the primary axis of the nebula. Here we add two more systems to the list, the central stars and nebulae of NGC~6337 and Sp~1. We show both systems to be low inclination, with their binary axis nearly aligned with our line-of-sight. Their inclinations match published values for the inclinations of their surrounding nebulae. Including these two systems with the existing sample statistically demonstrates a direct link between the central binary and the nebular morphology. In addition to the systems' inclinations we give ranges for other orbital parameters from binary modeling, including updated orbital periods for the binary central stars of NGC~6337 and Sp~1. | The shaping of planetary nebulae (PNe) has been a matter of interest for some time with the problem being approached from a number of different but complimentary methods \citep{kwi14}. One approach has been to identify binary central stars (CSs) where the companions are close enough to have interacted in the past and determine whether the interaction might have produced the observed morphologies. Since most known binary CSs are post common envelope (CE) binaries \citep[for a recent review of the CE binary interaction see ][]{iva13}, it is the CE interaction that has been primarily under scrutiny. Searches for close binary central stars of planetary nebulae (CSPNe) have been successful in discovering these systems, showing that approximately 10--20\% of all CSPNe appear to have a binary companion with an orbital period of less than a few days \citep{bon00,mis09a,dem09}. Many of those systems were discovered through photometric variability and while most are likely to be real binaries, additional confirmation is necessary for some of them \citep[e.g., Kn~61;][]{dem15}. Along with studies confirming the binarity of several of these systems \citep[e.g.][]{shi08,hil15a,hil16}, discoveries of additional close binary CSPNe are helping us to better understand the nature of these systems. In addition, studies of the CS can be linked to kinematic studies of the nebulae to determine whether a causal link exists between the interaction and the nebular morphology and kinematics. Surprisingly, while there are several suggestions that the CE interaction is the cause for the shape of post-CE PNe \citep{mor81,bon90,zij07,dem09,mis09b} there has never been a quantification of the link. The reason has been a lack of data. Binary modeling of known and newly discovered binary CSPNe along with spatio- kinematic modeling of the PNe have shown that of the systems studied, all seem to show an alignment between the central binary axis and the primary geometrical axis of the PN. Including the two systems we provide values for in this paper, there are now eight known PNe with binary CSs for which both inclinations are known. Using these data we are here, for the first time, demonstrating a correlation between post-CE CSs and their PN shapes, which we argue below implies causation. | We have shown conclusively that the CSs of both NGC~6337 and Sp~1 are short period irradiated binary systems with cool companions in low inclination orbits. The companions in both systems appear to be larger than a main sequence counterpart of the same mass, as is typical in these systems. The companions are also typically hotter than expected for a main sequence counterpart. However, we find that in the case of Sp~1 the companion may be roughly the same temperature, or cooler. It is unclear why this is the case here. The secondary {\it may} be slightly evolved, though there is no evidence to support this possibility. The companion in NGC~6337 does show a typically higher temperature. The modeled inclinations of both binary systems also align with the inclinations of the surrounding nebulae. And we show that these two systems, now along with six other systems, make a sample of eight PNe for which the nebular inclination and binary inclination of a close binary CS are known. In all eight cases the two inclinations agree with one another, within the uncertainties. We demonstrate that the likelihood that all eight of these systems are aligned merely by chance is vanishingly small. All other known parameters considered, the conclusion of a causal link between binarity and the axial symmetry of the PN is now on solid statistical grounds. Post-CE PNe have already been tentatively associated with bipolar morphologies \citep{zij07,dem09,mis09b} although the link is not clear cut because some of the PNe do not show the distinctive bi-lobal structure \citep[e.g., M2-29;][]{haj08}. From CE simulations \citep[e.g.,][]{san98,pas12} it is clear that the ejection of the CE happens preferentially on the equatorial plane. This equatorial ejection is a torus with a very large scale height. The scale height is likely a function of the companion mass because it depends on the amount of angular momentum injected which is larger for a more massive companion. A more massive companion also induces a stronger tide on the giant, which in turn results in more distant companions being captured into a CE interaction by the AGB star. These more massive and more distant companions carry more angular momentum into the envelope at the time of in-spiral. Magnetic fields in CE interactions likely play an important dynamical role in the ejection \citep[e.g.,][]{reg95}, and can be investigated by jets observed in PNe. Some post-CE PNe have jets. Jet masses and kinematics have been measured \citep[e.g.,][]{jon14} and exploited to determine the likely magnetic field strength and geometry at the time of CE \citep{toc14}. They also allow us to determine the elusive timescales of common envelope events, such as the length of the pre-in-spiral phase or the formation of post-CE accretion disks after the in-spiral by a new phase of Roche lobe overflow or fall-back of envelope gas \citep{kur16}. Once the in-spiral takes place, it is over very quickly (of the order of a dynamical timescale of the giant which is between a month and a year for the giants of interest here). On the assumption that the envelope is actually ejected during the dynamical in-fall \citep[something that is currently debated;][]{nan15,iac16} we can assume that the orbital separation of the post-in-spiral giant is approximately the same as we see today. This means that the giant star is quickly changed to a star with a radius smaller than today's Roche lobe radius. Since the luminosity is fully dependent on the core mass, it does not change. As a result the temperature of the CS must incur a relatively similar increase. The least massive companions would, at least in principle sink deeper into the envelope of the giant to eject it. This may generate a correlation between orbital separations and companion mass, although one may have to account for more massive primaries likely needing more orbital shrinkage to eject the massive envelope. In conclusion, a number of complex correlations are expected between stellar, binary and PN parameters in these objects, something that we will be able to test thanks to continuous characterization of these objects increasing the number of accurate parameters available. Soon PNe will become the best testing ground of the common envelope interaction. But for now we have demonstrated observationally a statistically significant connection between close binary CSs and their surrounding PN. | 16 | 9 | 1609.02185 |
1609 | 1609.03466_arXiv.txt | {% Treating dark matter at large scales as an effectively viscous fluid provides an improved framework for the calculation of the density and velocity power spectra compared to the standard assumption of an ideal pressureless fluid. We discuss how this framework can be made concrete through an appropriate coarse-graining procedure. We also review results that demonstrate that it improves the convergence of cosmological perturbation theory. } | The calculation of the power spectrum of cosmological fluctuations for stochastic initial conditions is an important and challenging task. It provides the basis for analyzing data on the large-scale structure of the universe and can lead to constraints on the parameters of the cosmological model. The spectrum can be computed through $N$-body simulations. However, these are CPU-intensive, while they do not provide an intuitive understanding of the underlying physical processes. The alternative option is an analytical treatment, which amounts to solving the collisionless Vlasov-Poisson equation perturbatively in the density contrast for appropriate classes of initial conditions \cite{bernardeau0}. The use of perturbative techniques in this context is complicated by the growth of non-linearities at short length scales. Within the strongly non-linear regime, velocity dispersion and virialization becomes important. The inability of analytical methods to describe the short scales reliably suggests to reformulate cosmological perturbation theory based on an effective description, applicable only above some length scale. The description should contain effective parameters which absorb the effect of short-scale perturbations that are `integrated out'~\cite{baumann,effective,Pietroni:2011iz}. In the conventional description of dark matter as a pressureless ideal fluid, only the lowest moments of the phase-space distribution are taken into account. Enlarging this framework by including higher moments results in viscous transport coefficients that parametrize a non-ideal stress tensor for the dark matter. In the following we present heuristic arguments that support the suggestion that an efficient fluid description of dark matter must allow for such non-ideal terms. The role of the effective viscosity and sound velocity is to account for the interaction of large-scale fluctuations (with wavenumbers $k < k_m$) with the short-scale ones ($k > k_m$) that are not followed explicitly in a fluid description limited to $k<k_m$. In our analysis the UV contributions are incorporated in the large-scale theory at the level of a one-loop approximation. Going beyond one-loop can be achieved through renormalization-group techniques \cite{Max1,floerchinger}. The approximation that we review here provides an intuitive and rather simple framework, while it still resolves the main deficiency of standard perturbation theory (SPT), namely the strong dependence on the short-scale dynamics that are out of the reach of analytical techniques. We review the calculation of the density power spectrum presented in ref. \cite{blas} in the context of the coarse-grained theory. We also present results for the velocity and cross spectra within the same framework. | We conclude that the description of dark matter as a viscous fluid discussed here yields a robust framework for predicting power spectra for $k\lesssim 0.2\, h/$Mpc at $z=0$ without the need to adjust any free parameters. The framework can be constructed and extended through a formal analysis based on the renormalization group. We refer the reader to ref. \cite{floerchinger} for details. \begin{figure*} \centering \includegraphics[width=6cm]{A-P22OverP11Viscz0}\qquad \includegraphics[width=6cm]{A-P22OverP11SPTz0} \caption{Power spectrum of the velocity-divergence, normalized to the density power spectrum, for the viscous theory (left) and standard perturbation theory (right). The various lines show the dependence on $k_m$ and on $\Lambda$, respectively, as in fig.\,\ref{fig-2}. The red lines are $N$-body results ((J12) \cite{Jennings:2012ej}, (HAA14) \cite{Hahn:2014lca}), and the shaded region is the quoted uncertainty of the velocity power spectra extracted from $N$-body data. } \label{fig-3} \end{figure*} \begin{figure*} \centering \includegraphics[width=6cm]{A-P12OverP11Viscz0}\qquad \includegraphics[width=6cm]{A-P12OverP11SPTz0} \caption{Cross density-velocity power spectrum, corresponding to the same approximations as in fig.\,\ref{fig-3}.} \label{fig-4} \end{figure*} | 16 | 9 | 1609.03466 |
1609 | 1609.04841_arXiv.txt | Precise measurements of eclipsing binary parameters and statistical studies of young clusters have suggested that some magnetically active low-mass dwarfs possess radii inflated by $\sim$5--15\% relative to theoretical expectations. If true, this effect should be pronounced in young open clusters, due to the rapid rotation and strong magnetic activity of their most extreme members. We explore this possibility by determining empirical radii for 83 members of the nearby Pleiades open cluster, using spectral energy distribution fitting to establish \fbol\ with a typical accuracy of $\approx$3\% together with color and spectro-photometric indices to determine \teff. We find several Pleiades members with radii inflated above radius-\teff\ models from state-of-the-art calculations, and apparent dispersions in radii for the K-dwarfs of the cluster. Moreover, we demonstrate that this putative radius inflation correlates strongly with rotation rate, consistent with inflation of young stars by magnetic activity and/or starspots. We argue that this signal is not a consequence of starspot-induced color anomalies, binarity, or depth effects in the cluster, employing {\it Gaia} DR1 distances as a check. Finally, we consider the lithium abundances of these stars, demonstrating a triple correlation between rotation rate, radius inflation, and enhanced lithium abundance. Our result---already significant to $\sim$99.99\% confidence---provides strong support for a magnetic origin of the inflated radii and lithium dispersion observed in young, low-mass stars. | \label{sec:intro} Precise knowledge of the fundamental parameters of stars is of considerable importance to understanding both their exact nature, their exo-planets, and the timescales of diverse astrophysical phenomenon such as star formation and circum-stellar disk evaporation. Standard stellar theory has been extremely successful at predicting stellar properties throughout the wide and varied life cycles of stars, but their fidelity with regard to the radii of low mass ($M < 1$\msun) stars has been called into question by a consistent trend emerging from direct measurements: young, active stars appear to have radii that are larger than standard predictions by $\sim 5-15$\%. This phenomenon has been claimed in eclipsing binaries \citep[e.g][]{popper97,torres02,lopez-morales05}, statistical studies of open clusters \citep[e.g.][]{jackson16}, on both sides of the fully-convective boundary of $0.35$~\msun\ \citep[e.g.][]{clausen09,Stassun:2012}, and on both the pre-main sequence and main sequence \citep[e.g.][]{torres10,feiden12,torres14,Stassun:2014}. Additionally, corresponding anomalies in the \teffs\ of the afflicted stars have been noted in several instances \citep[e.g.][]{stassun06,dupuy16}. While the underlying cause remains controversial, it has been shown that in some cases the degree of radius inflation appears correlated with the strength of surface magnetic activity \citep[e.g.][]{torres06,lopez-morales07,morales08,Stassun:2012}. Newer generations of theoretical models have begun to consider such effects, and ongoing research is studying their impact on the fundamental parameters, abundances, and evolutionary timescales of young stars \citep[e.g.][]{mullan01,chabrier07,macdonald10,feiden13,feiden14,jackson14a,jackson14b,Somers:2014,Somers:2015b,Somers:2015a}. Results have been particularly promising for young stars: the well-known lithium-rotation correlation in pre- and zero-age main sequence clusters \citep[e.g.][]{soderblom93,messina2016,bouvier2016} is a direct prediction of an activity-radius connection on the pre-main sequence \citep[e.g.][]{ventura98,Somers:2015b,Somers:2015a,jeffries2016}. If an activity-radius connection truly exists, young ($\lesssim 200$~Myr) main sequence open clusters present a valuable test bed to uncover the nature of the correlation. Young clusters generally contain members with extraordinarily rapid rotation, as the magnetic stellar winds which efficiently break stars on the main sequence have not yet significantly depleted the initial stellar angular momentum \citep[e.g.][]{pinsonneault89,gallet2015,somers16b}. Furthermore, young clusters host large dispersions in rotation rate at fixed mass due to the diversity of rotative initial conditions in star-forming regions, and consequently show a range of activity levels and starspot coverage from moderate to extreme \citep[e.g.][]{soderblom93,odell95,gallet2015,fang16}. Such clusters have been the target of several recent statistical studies examining the fundamental parameters of stars \citep{littlefair11,cottaar14,jackson14a,jackson14b,jackson16,lanzafame2016}. In several cases, these studies have concluded both that the average stellar radius is larger compared to standard predictions, and that dispersions likely exist in radius at fixed \teff. In particular, \citet{jackson14a} recently examined the well-studied open cluster, the Pleiades, for signs of radius inflation. By measuring the projected radii of a large sample of its members, and statistically analyzing their results, they found that the average Pleiad radius is $\sim 10$\% larger than theoretical predictions below $1$~\msun. The Pleiades is an advantageous laboratory for such experiments, given its proximity \citep[$d \sim 136$~pc,][]{melis14}, near-solar composition of [Fe/H]~$\sim 0.03$ \citep{soderblom09}, and young age of 125~Myr \citep{stauffer1998}. Furthermore, its members more massive than $\sim 0.6$~\msun\ have reached the main sequence, meaning that the harrowing uncertainties dogging pre-main sequence models can be avoided in comparisons with theory \citep[e.g.][]{Stassun:2014}. In this paper, we apply a distinct method to this same open cluster to search for corroborating signs of radius inflation and dispersion. Our method involves fitting the broadband spectral energy distributions (SEDs) of individual members to establish their bolometric fluxes (\fbol), and combining this result with the known distance of the cluster and an estimate of the \teff, to calculate the stellar radius. This approach is attractive because, in principle, individual stars can be tested for radius inflation, and correlations with non-standard physical effects like activity and rotation can be explored. We devote considerable discussion to the accuracy of our \teffs, as the active and spotted nature of young stars complicates simple extrapolation from photometry. In the end, we find a clear connection between rotation rate and apparent radius inflation above a putative \teff-radius relation. The paper is organized as follows. In $\S$\ref{sec:methods}, we describe our sample selection, our methods for deriving \fbol\ and \teff\ for our stars, and our procedure for deriving the corresponding radii. In $\S$\ref{sec:results}, we present these radii and compare them with predictions from theoretical models, looking particular at the influence of rotation on the agreement. In $\S$\ref{sec:discussion}, discuss the possible contaminating influence of starspots, binaries, and extinction, and describe the relationship between inflated radii and lithium abundance in this cluster. Finally, we summarize our conclusions in $\S$\ref{sec:summary}. | \label{sec:summary} Previous studies have reported anomalously large radii among low-mass stars in the Pleiades and other young clusters, using ensemble averages of projected rotation velocity measurements. Here, we have measured the radii of several Pleiades cluster members via the Stefan-Boltzmann law, combining (i) \teffs\ determined through color and spectro-photometric techniques, (ii) bolometric fluxes determined by summing the observed spectral energy distributions, and (iii) the known cluster distance. Our sample specifically includes stars with previously determined rotation periods and lithium abundances. We compare our radius measurements to literature isochrones, calibrated on older, inactive stars, and find that in many cases the Pleiades radii can be larger by 10--20\% compared to expectations. We further show that this over-radius correlates with rapid rotation at greater than 99.99\% confidence, strongly suggesting a magnetic origin. We discuss whether this radius-rotation correlation could be a spurious artifact brought on by poorly calculated radii due to the distorted SEDs of rapidly rotating stars. A very large systematic offset in \teff, afflicting only stars which rotate with periods shorter than 1.5~days, would be required to reconcile the rapidly rotating Pleiads with the model expectations. However, our quantitative measures of SED distortion find that this is principally a function of \teff, not rotation, with the coolest stars tending to show modestly distorted SEDs that might be better fit by a two-temperature model. We conclude that the most likely explanation is magnetically-driven radius inflation amongst the most rapidly rotating Pleiads. | 16 | 9 | 1609.04841 |
1609 | 1609.01463_arXiv.txt | Electron acceleration in the solar corona is often associated with flares and the eruption of twisted magnetic structures known as flux ropes. However, the locations and mechanisms of such particle acceleration during the flare and eruption are still subject to much investigation. Observing the exact sites of particle acceleration can help confirm how the flare and eruption are initiated and how they evolve. Here we use the Atmospheric Imaging Assembly to analyse a flare and erupting flux rope on 2014-April-18, while observations from the Nan\c{c}ay Radio Astronomy Facility allows us to diagnose the sites of electron acceleration during the eruption. Our analysis shows evidence for a pre-formed flux rope which slowly rises and becomes destabilised at the time of a C-class flare, plasma jet and the escape of {\color{black}$\gtrsim$75\,keV} electrons from rope center into the corona. As the eruption proceeds, continued acceleration of {\color{black}electrons with energies of $\sim$5\,keV} occurs above the flux rope for a period over 5 minutes. % At flare peak, one site of electron acceleration is located close to the flare site while another is driven by the erupting flux rope into the corona at speeds of up to 400\,km\,s$^{-1}$. Energetic electrons then fill the erupting volume, eventually allowing the flux rope legs to be clearly {\color{black} imaged from radio sources at 150--445\,MHz}. Following the analysis of \citet{joshi2015}, we conclude that the sites of energetic electrons are consistent with flux rope eruption via a tether-cutting or flux cancellation scenario inside a magnetic fan-spine structure. In total, our radio observations allow us to better understand the evolution of a flux rope eruption and its associated electron acceleration sites, from eruption initiation to propagation into the corona. | Flares and coronal mass ejections (CMEs) are thought to result from magnetic energy release in the solar corona, often involving the destabilisation of a twisted magnetic structure known as a flux rope \citep{chen2011, webb2012}. This activity may be accompanied by the acceleration of energetic particles \citep{lin2000a, kahler2007, lin2011}. However, there is ongoing debate on exactly where, when and how the particle acceleration occurs during flaring and eruption. Observing the locations of {\color{black}energetic electrons} during an eruptive event may help confirm how the electrons are accelerated, how the eruption proceeds, and also help in identifying which models of solar eruptive activity are correct. The observation of the sites of electron acceleration during flaring or eruptive activity in the corona has traditionally been made using radio observations (see \citet{pick2008} for a review). Some of the longest known signatures of particle acceleration in the corona are type III radio bursts \citep{wild1959}, now believed to be from energetic electrons causing plasma emission as they {\color{black}propagate through} the corona \citep{paesold2001, yan2006, chen2013}. Type IIIs are amongst various types of radio bursts generated by energetic electrons that are accelerated during both flares and small scale eruptions in the form of plasma jets \citep{aurass1994, kundu1995, nitta2006, klassen2012, chen2013a}. Sites of energetic electrons (radio sources) are also known to be closely associated with larger scale eruptions such as plasmoids and sigmoids \citep{kundu2001, khan2002, marque2002}. Such eruptive activity often shows the {\color{black}sites of energetic electrons} to be located close to the underlying active region or moving with the {\color{black}erupting structure} itself \citep{pick2005, bain2014}, originally observed as flare continua and moving type IV bursts \citep{robinson1975, pick1986}. Electron acceleration sites may also be located on the boundaries of CMEs, located at its nose or flanks. \citep{zimovets2012, bain2012, carley2013, zucca2014, salas2016}. In the later stages of an eruptive event, {\color{black}the radio emission from energised electrons} can be observed to be within the erupting structure, allowing sources of plasma emission at the CME legs to be imaged \citep{maia1999, huang2011}. In very rare cases, the energised electrons fill the entire volume of the CME and interact with the magnetic field to produce gyrosynchrotron emission in the metric/decimetric domain, allowing observation of what is generally known as a `radio CME' \citep{bastian2001, maia2007, demoulin2012}. Theoretical models of solar eruptive activity often include a variety of sites of magnetic reconnection and shocks \citep{chen2011}, implying a variety of possible particle acceleration sites during eruption. The models may be unique in where they predict these sites to occur in relation to the erupting flux rope. For example, the magnetic breakout model specifically predicts a site of reconnection above a flux rope \citep{antio99, lynch2004}, while the tether-cutting or flux cancellation models may both build and release the flux rope via reconnection quite close to the rope center \citep{vanboo1989, moore2001}. During the final stages of propagation of the erupting structure, nearly all models predict the development of a current sheet below the main body of the eruption. This is generally known as the `standard' or CSHKP model, and predicts reconnection, shocks and particle acceleration in this current sheet \citep{carmichael1964, sturrock1966, hirayama1974, kopp1976}. Elsewhere during the eruption, models also predict particle acceleration from interchange reconnection or shocks driven at the outer boundaries of the erupting structure as it propagates into the corona \citep{kozarev2011, schmidt2012, masson2013}. While theoretical models predict a variety of possible electron accelerations sites during flux rope eruption, the radio observations provide a means for detecting theses sites. However, observing the sites of electron acceleration simultaneously with flux rope observations has proven difficult in the past. It is only recently that the Atmospheric Imaging Assembly \citep[AIA;][]{lemen2012} has made available high spatial and temporal resolution observations of flux rope signatures in the corona e.g., twisted sigmoids with temperatures of $\gtrsim$10\,MK \citep{zhang2012, cheng2015, joshi2015, song2015}. Hence, there is now the possibility to combine the high time and spatial resolution flux rope observations with radio imaging and dynamic spectra to explore where, when and how the electron acceleration occurs during the eruption of such a body and to compare this to what theoretical models predict. In this paper, we examine an eruptive event from 2014-April-18. This event has previously been studied by \citet{joshi2015} and \cite{cheng2015} using X-ray and EUV imaging and UV spectroscopy, respectively. They identify a flux rope and multiple hypothesized sites of reconnection (potentially associated with electron acceleration) during eruption. Here we attempt to identify these sites of electron acceleration using radio imaging from the Nan\c{c}ay Radioheliograph \citep[NRH;][]{kerdraon1997} combined with metric and decimetric radio spectrography. We show strong observational evidence of {\color{black} energized electrons produced from an eruptive mechanism which closely resembles a tether-cutting or flux cancelation model. This is followed by electron acceleration from reconnection above the erupting structure in the surrounding magnetic environment.} Sites of electron acceleration are then both associated with the flare and driven by the {\color{black}erupting structure} itself. We also show how electrons then fill the erupting volume, making visible at radio wavelengths the CME legs. In total, we reveal in unprecedented detail the sites and kinds of electron acceleration occurring throughout this event, from eruption initiation to propagation into the corona. \begin{figure}[!t] \begin{center} \includegraphics[scale=0.35, trim=2.5cm 0.8cm 2cm 1.0cm]{carley_figure1.pdf} \caption{ (a) GOES X-ray light curves showing the initial C-class flare at $\sim$12:35\,UT. This is followed by an M7.3 class flare peaking at $\sim$13:00\,UT. (b) RHESSI X-ray flux observations from 3-100\,keV. Data gaps due to RHESSI night and the South Atlantic Anomaly are indicated by blue and brown lines, respectively. (c) FERMI GBM light curves showing emission from 4.5-101.6\,keV. (d) Radio dynamic spectra from NDA and Orf\'{e}es covering 10--1000\,MHz. Based on the observations below, the dynamic spectrum is split into five periods, indicated by the vertical dashed lines.} \label{fig:xray_radio} \end{center} \end{figure} | In this study, {\color{black}we have presented analysis of radio sources which indicate the regions of accelerated electrons at each stage of a flare and flux rope eruption}, from its initiation to its propagation into the corona. While the flux rope was observed in AIA, the sites of electron acceleration were identified using multiple radio frequency images of the Nan\c{c}ay Radioheliograph. This combined with high time and frequency resolution dynamic spectral observations from Nan\c{c}ay's new Orf\'{e}es instrument allowed us to identify when and where the electron acceleration took place during the event. Our observations and analysis reveal the following properties of electron acceleration sites during flux rope eruption in this event: \begin{enumerate} \item At the time of flux rope eruption, tether-cutting or flux cancelation-type reconnection takes place at flux rope center resulting in the expulsion of a plasma jet and the acceleration of electron beams to {\color{black}$\gtrsim$75\,keV} which escape into the corona and produce a type IIId radio burst. The escape of these beams requires the null point of a fan-spine structure to be located close to the rope center, with electrons propagating along the spine. At the same time, the presence of a type IIIn burst indicates $\gtrsim$52\,keV electrons propagating to interplanetary distances. However, it is unclear as to the relationship between the type IIId and type IIIn. The two different radio bursts may be from two separate populations of electrons, so there may only be an indirect link between the population of electrons accelerated along with the jet, and those escaping into the heliosphere. \item As the flux rope erupts, reconnection takes place above the rope, resulting in repeated acceleration of {\color{black}electrons of energies of 5\,keV} for a period of up to 5 minutes. The site of reconnection is likely at a null point in a large fan-spine structure above the flux rope in the corona. The electrons accelerated during this time escape into the corona along the spine of this structure and produce type III radio bursts. This observation is in support of the hypothesised points of reconnection (implying electron acceleration) at the fan-spine null-point outlined in \citet{joshi2015}. \item During flare peak, the majority of the electron acceleration takes place close to the flaring active region. Simultaneously, reconnection driven above the rope as it erupts into the corona results in a site of further electron acceleration which propagates outwards at the same speed as the rope ($\sim$400\,km\,s$^{-1}$). Following this, electron acceleration continues in the active region and electrons are injected onto the loops of the rope. \item {\color{black}Electron acceleration} continues close to the flare site. During this time we find evidence of energetic electrons beginning to fill the erupting volume. The growth of this volume results in electrons being contained on the magnetic fields that make up the legs of the CME, allowing the legs to be clearly imaged at multiple frequencies high in the corona. In future, low frequency imaging spectroscopy, such as that now provided by instruments like the Low Frequency Array \citep[LOFAR;][]{vanHaarlem2013}, may reveal just how far CME legs may be imaged in the corona e.g., as low as 30\,MHz. \end{enumerate} | 16 | 9 | 1609.01463 |
1609 | 1609.02183_arXiv.txt | This work reports new experimental radiative lifetimes and calculated oscillator strengths for transitions of astrophysical interest in singly ionized cobalt. More precisely, nineteen radiative lifetimes in Co$^+$ have been measured with the time-resolved laser-induced fluorescence technique using one- and two-step excitations. Out of these, seven belonging to the high lying 3d$^7$($^4$F)4d configuration in the energy range 90697 -- 93738 cm$^{-1}$ are new, and the other twelve from the 3d$^7$($^4$F)4p configuration with energies between 45972 and 49328 cm$^{-1}$ are compared with previous measurements. In addition, a relativistic Hartree-Fock model including core-polarization effects has been employed to compute transition rates. Supported by the good agreement between theory and experiment for the lifetimes, new reliable transition probabilities and oscillator strengths have been deduced for 5080 \ion{Co}{ii} transitions in the spectral range 114 -- 8744 nm. | The final stage of exothermal elemental production in stars is the iron-group elements. The even-$Z$ atoms are produced by consecutive capture of helium nuclei, and named $\alpha$-elements. The production of the odd-$Z$ elements is not as well constrained, and does not follow the abundance trends of the $\alpha$-elements, indicating non-common production sites. As a result of this nucleosynthesis in the interior of stars, the even-$Z$ nuclei such as Ca, Ti, Cr, and Fe have a higher cosmic abundance compared to the odd-$Z$ nuclei located in between. However, the astrophysical interest for the odd-$Z$ iron-group elements has increased in recent years. In the present project, we target atomic data for \ion{Co}{ii} (Z=27). Cobalt is believed to be produced primarily in type II supernova, and also to a lesser extent in type Ia (Woosley \& Weaver 1995, Bravo \& Martinez-Pinedo 2012, Battistini \& Bensby 2015). Abundance determinations in stars serve as important tests of the stellar evolution and supernova explosion models (Pagel 2009). Furthermore, high-excitation spectral lines have additional diagnostic value, since they can be used to benchmark non local thermodynamical equilibrium (non-LTE) modelling of stellar atmospheres. Along with the development of 3D model atmospheres, a trustworthy non-LTE treatment is the current challenge for accurate stellar abundances. High-precision atomic data for selected lines is important for this development (Lind {\it et al} 2012). In the case of cobalt, a big effort was recently made by Lawler {\it et al} (2015) to provide improved oscillator strengths for about 900 lines belonging to the first spectrum (\ion{Co}{i}). These were deduced from emission branching fractions measured from hollow cathode lamp spectra recorded using a Fourier transform spectrometer and a high-resolution echelle spectrograph combined with radiative lifetimes determined by the time-resolved laser-induced fluorescence (TR-LIF) technique. Concerning the second spectrum (\ion{Co}{ii}), several papers have reported experimental determinations of radiative data. Lifetime measurements have been performed by Pinnington {\it et al} (1974) and S\o rensen (1979) using the beam-foil technique and by Salih {\it et al} (1985) and Mullman {\it et al} (1998) using TR-LIF. Experimental transition probabilities were reported for 41 spectral lines by Salih {\it et al} (1985), Crespo Lopez-Urrutia {\it et al} (1994) and Mullman {\it et al} (1998) by combining branching fraction measurements with available lifetimes. Theoretical data have been published by Raassen {\it et al} (1998) who computed oscillator strengths for a large number of \ion{Co}{ii} lines using the method of orthogonal operators. However, all these studies were limited to radiative decays from the odd-parity 3d$^7$4p configuration to lower states belonging to the 3d$^8$ and 3d$^7$4s even-parity configurations. More extensive calculations are reported by Kurucz (2011). The main goal of the present work is to extend the knowledge of radiative data to higher energy states and to provide a consistent set of transition rates for a large number of spectral lines in singly ionized cobalt. More precisely, new experimental lifetime measurements were performed by TR-LIF for 12 energy levels belonging to the 3d$^7$4p odd-parity configuration and 7 energy levels belonging to the 3d$^7$4d even-parity configuration using one- and two-step excitations, respectively. In addition, transition probabilities and oscillator strengths were computed for 5080 \ion{Co}{ii} lines in a wide spectral region, from ultraviolet to infrared, using a pseudo-relativistic Hartree-Fock model including core-polarization effects. | Transition probabilities and oscillator strengths have been obtained for 5080 spectral lines in \ion{Co}{ii} using the pseudo-relativistic Hartree-Fock method including the most important intravalence correlation and core-polarization effects. The accuracy of the new data has been assessed through detailed comparisons with previously published experimental and theoretical radiative rates together with new lifetime measurements performed in the present work using the laser-induced fluorescence technique with one- and two-step excitations. In view of the overall agreement obtained between all sets of results, it is expected that the new $gA$- and $gf$-values should be accurate to a few percent for the strongest transitions and to within 20--30\% for weaker lines. | 16 | 9 | 1609.02183 |
1609 | 1609.05595_arXiv.txt | In modern high precision optical instruments, such as in gravitational wave detectors or frequency references, thermally induced fluctuations in the reflective coatings can be a limiting noise source. This noise, known as coating thermal noise, can be reduced by choosing materials with low mechanical loss. Examination of new materials becomes a necessity in order to further minimize the coating thermal noise and thus improve sensitivity of next generation instruments. We present a novel approach to directly measure coating thermal noise using a high finesse folded cavity in which multiple Hermite-Gaussian modes co-resonate. This method is used to probe surface fluctuations on the order $10^{-17} \mrtHz$ in the frequency range \SI{30-400}{Hz}. We applied this technique to measure thermal noise and loss angle of the coating used in Advanced LIGO. | Dielectric coatings used in high precision optical instruments consist of alternating layers of materials with low and high index of refraction. Thermal noise in these coatings arises from mechanical dissipation in the coating materials described by the fluctuation dissipation theorem. This noise limits the sensitivity of the current gravitational wave detectors~\cite{0264-9381-24-2-008, PhysRevD.78.102003, PhysRevLett.116.061102}, of the best frequency references~\cite{Ludlow:07}, and of macroscopic quantum measurements~\cite{1367-2630-11-7-073032, Poot2012273}. Further improvement of these instruments calls for reduction of the coating thermal noise. The materials presently in use belong to the class of amorphous glassy oxides including SiO$_2$, Ta$_2$O$_5$, ZrO$_2$, Nb$_2$O$_5$, HfO$_2$ and Al$_2$O$_3$. The search for the new high reflectivity surfaces with low mechanical loss explores a wide range of possibilities: from new amorphous coatings produced with conventional ion beam sputtering techniques \cite{0264-9381-27-8-084030, PhysRevD.91.042002}, to crystal coatings \cite{nphoton.2013.174} and grating reflectors \cite{PhysRevD.88.042001}. The loss angle of new coating materials is most frequently obtained based on the measurement of the mechanical quality factor. The techniques presented in the literature include, among others, suspended disks~\cite{crooks_blades, harry_blades}, clamped cantilevers~\cite{pierro_cantilevers}, and the gentle nodal suspension~\cite{cesarini_nodal}. The level of coating thermal noise is then calculated from the measured parameters, such as mechanical loss angles, Poisson ratio, and Young's modulus. However, due to uncertainties in the multilayer parameters a robust experimental setup is necessary to directly measure coating thermal noise of a particular sample. Such a measurement is complicated by multiple noise sources such as table vibrations, laser frequency and amplitude noise, and various readout noises. In the past, direct measurements of the coating thermal noise have been accomplished using suspended free-space Fabry-Perot cavities~\cite{numata_ctn, black_ctn}. Seismic motion limits the sensitivity of these experiments below 100\,Hz. On the other hand, fixed-spacer cavities with optically contacting mirrors were recently developed to observe coating thermal noise below 100\,Hz~\cite{tara_ctn}. However, the readout of this experiment is located in transmission of the cavities. This sets an upper limit on the reflectivity of the measured sample. This paper describes a novel technique for the direct observation of the coating thermal noise which uses only one free-space Fabry-Perot cavity, and in which there is no upper limit on the sample reflectivity. Multiple transverse electromagnetic modes (TEM) co-resonate in the cavity: 00, 02 and 20. These modes have orthogonal spacial profiles, and probe different areas of the sample coating, while other displacement noises of the cavity are common to all resonating modes. Coating thermal noise is extracted from the frequency difference between the two higher order modes. In Sec.~\ref{ctn} we describe analytical calculations of the coating thermal noise for the fundamental and higher order modes in the linear and folded cavities. Sec. \ref{exp} describes our experimental setup. We have used it to measure the coating thermal noise of an Advanced LIGO~\cite{design_aligo} witness sample. In Sec. \ref{res} we discuss the sensitivity of our experiment, measured coating thermal noise of Advanced LIGO sample and the estimation of \ttan:\tant\ loss angle. | We presented a novel experiment for the broadband direct measurements of the coating thermal noise. The sensitivity of \SI{10^{-17}}{\mrtHz} has been achieved in the frequency band \SI{30 - 1000}{Hz}. This is made possible by our novel measurement technique, in which TEM00, 02 and 20 spatial modes all co-resonate in a folded cavity. As a first application of this technique, we measured the coating thermal noise from Advanced LIGO coating and estimated the loss angle of \ttan:\tant. Our results are broadly consistent with the previous estimations, but give a 20\% higher coating thermal noise compared to the published Advanced LIGO noise estimates~\cite{0264-9381-32-7-074001, den_nb}. With the ever increasing sensitivity of precision optical measurements, coating thermal noise has become a significant obstacle. In terms of the gravitational wave interferometers and some macroscopic quantum measurement experiments, it is essential to reduce this noise in order to reach and surpass the standard quantum limit. Our experiment design will allow for rapid testing new coatings, thereby helping to reduce the coating thermal noise in the future generation of gravitational wave detectors, frequency references and quantum measurements. | 16 | 9 | 1609.05595 |
1609 | 1609.00379_arXiv.txt | {We present a complete framework for numerical calculation of the power spectrum and bispectrum in canonical inflation with an arbitrary number of light or heavy fields. Our method includes all relevant effects at tree-level in the loop expansion, including (i) interference between growing and decaying modes near horizon exit; (ii) correlation and coupling between species near horizon exit and on superhorizon scales; (iii) contributions from mass terms; and (iv) all contributions from coupling to gravity. We track the evolution of each correlation function from the vacuum state through horizon exit and the superhorizon regime, with no need to match quantum and classical parts of the calculation; when integrated, our approach corresponds exactly with the tree-level Schwinger or `in--in' formulation of quantum field theory. In this paper we give the equations necessary to evolve all two- and three-point correlation functions together with suitable initial conditions. The final formalism is suitable to compute the amplitude, shape, and scale dependence of the bispectrum in models with $|\fNL|$ of order unity or less, which are a target for future galaxy surveys such as Euclid, DESI and LSST. As an illustration we apply our framework to a number of examples, obtaining quantitatively accurate predictions for their bispectra for the first time. Two accompanying reports describe publicly-available software packages that implement the method. \par\vspace{6mm}\mbox{}\hfill \includegraphics[scale=0.1]{Logos/LOGO-ERC} \quad \includegraphics[scale=0.5]{Logos/LOGO-STFC} \quad \includegraphics[scale=0.3]{Logos/RS}} | \label{sec:introduction} In the inflationary scenario, quantum-mechanical processes seeded an early distribution of gravitational potential wells. As matter sank into these wells it formed the largest structures we observe, with the result that any observable tracing this structure can be used to infer details of the seeding process. To do so we require measurements on cosmological scales, which continue to improve at a remarkable rate---and, soon, we can expect the cosmic microwave background temperature and polarization anisotropies to be joined as a precision probe by the galaxy density field, intrinsic alignments, weak lensing shear maps and perhaps others. The raw materials for these analyses are the correlation functions describing the primordial distribution of potential wells, and to calculate them we need the methods of quantum field theory in curved spacetime. We will discuss these methods in detail below, and explain why such calculations are challenging and why analytic methods are limited. But two of the reasons are easy to state and especially difficult to overcome: first, the algebraic complexity arising when even the simplest models are coupled to gravity; and second, the occurrence of large hierarchies that compensate for the smallness of any natural expansion parameters and render na\"{\i}ve perturbation theory useless. The most straightforward solution to both these issues is to switch to a numerical method. Exactly this approach has been adopted in other areas of physics---% including collider phenomenology, the paradigmatic example of extracting observational predictions from quantum field theory---% when the same difficulties are encountered. The tools available to assist cosmologists in these calculations are substantially less sophisticated than in collider phenomenology, where not only the numerical computations are implemented by software---% it is possible to combine powerful computer packages that partially \emph{automate} the calculation of LHC observables directly from a Lagrangian~\cite{Degrande:2011ua}. Examples include the Feynman diagram generators \href{http://theory.sinp.msu.ru/~semenov/lanhep.html}{\LanHEP}~\cite{Semenov:1998eb,Semenov:2008jy} or \href{https://feynrules.irmp.ucl.ac.be}{\FeynRules}~\cite{Christensen:2008py,Christensen:2009jx} combined with \href{http://theory.sinp.msu.ru/~pukhov/calchep.html}{{\CompHEP}/{\CalcHEP}}~\cite{Boos:2004kh,Pukhov:1999gg,Pukhov:2004ca}, \href{http://madgraph.hep.uiuc.edu}{\MadGraph}~\cite{Alwall:2014hca} or \href{http://www.feynarts.de/formcalc/}{\FormCalc}~\cite{Hahn:1998yk,Hahn:2006qw}. Numerical tools for inflationary calculations have been developed, but typically they work on a case-by-case basis where derivatives of the potential must be obtained by hand and supplied as subroutines. Examples include the Fortran codes \href{http://theory.physics.unige.ch/~ringeval/fieldinf.html}{\FieldInf} \cite{Ringeval:2005yn,Martin:2006rs,Ringeval:2007am}, \href{http://modecode.org}{\ModeCode} and \href{http://modecode.org}{\MultiModeCode} \cite{Mortonson:2010er,Easther:2011yq,Norena:2012rs,Price:2014xpa}, and the Python package \href{http://pyflation.ianhuston.net}{\PyFlation}~\cite{Huston:2009ac,Huston:2011vt,Huston:2011fr}, which are all solvers for the two-point function in models of varying generality. For the three-point function, the only public code of which we are aware is the Fortran90 solver \href{https://sites.google.com/site/codecosmo/bingo}{\BINGO}~\cite{Hazra:2012yn,Sreenath:2014nca} which is restricted to single-field canonical models. Providing derivatives of the potential by hand becomes burdensome when the potential is complex or there are many fields---unfortunately, precisely the cases where numerical methods have most value. The situation is worse where a nontrivial field-space metric or kinetic structure means that further derivatives are required, such as the field-space Riemann tensor. It would be preferable, as in collider phenomenology, to automate these calculations. But automated tools have advantages beyond mere convenience, providing a fair basis for comparison between models by dropping simplifying assumptions and enabling researchers whose primary interest may be model building (rather than the calculation of correlation functions as an end in themselves) to obtain observable predictions close to the state-of-the-art in technical sophistication. More ambitiously, once $n$-point functions are available there is no need to stop: we can extend the analysis to include automated calculation of late-universe observables such as the CMB angular correlation functions, the dark matter or galaxy correlation functions, estimates of scale-dependent bias, and so on. In Ref.~\cite{Dias:2015rca}, three of us presented a Mathematica \href{http://transportmethod.com}{`transport' code} that can automate (in the sense just described) calculation of the inflationary two-point function in a multiple-field model with nontrivial field-space metric, given only symbolic expressions for the metric and potential. Like those described above this code is a two-point function solver. But present-day datasets already have some sensitivity to inflationary \emph{three}-point functions, and this sensitivity is expected to improve as large-scale galaxy surveys such as Euclid, DESI and LSST become available. To enable model-building theorists to compare their scenarios with these datasets it would be very convenient to automate calculation of the three-point functions. In this paper we describe a numerical method for doing so---% currently, applied to an arbitrary multifield model with canonical kinetic terms---% and collect a number of results showcasing its utility. The method generalizes to nontrivial kinetic terms but the calculations necessary to implement it have not yet been performed. We intend to return to this in the future. For those wishing to replicate our analyses, or apply our methods to their own models, we have made our computer codes available. The transport code described in Ref.~\cite{Dias:2015rca} was implemented in Mathematica, which trades speed and flexibility for a certain kind of convenience inherited from access to Mathematica's symbolic engine and visualization capabilities. However, because numerical calculation of three-point functions is substantially slower than for two-point functions---and the complexity of implementation and maintenance substantially higher---% it is not clear that Mathematica will continue to provide a suitable platform. To accommodate this our bispectrum codes are mostly written in compiled languages and optionally can be parallelized. They have been tested against each other but do not share code, and have different specializations and use cases: \begin{itemize} \item {\PyTransport} is developed by a team at Queen Mary, University of London. It has a {\CC} core but is intended to be used through a Python interface, and uses the \href{http://www.sympy.org/en/index.html}{\SymPy} package to provide its symbolic algebra support. It has minimal prerequisites, supports rapid development, and its Python interface means that it is easily scriptable. It can be used with libraries such as \href{http://matplotlib.org}{\Matplotlib} or \href{http://code.enthought.com/projects/mayavi/}{\Mayavi} for visualization and can be parallelized using \href{http://pythonhosted.org/mpi4py/}{\MpiForPy} or similar packages. \item {\CppTransport} is developed by a team at the University of Sussex. It is a pure {\CC} platform using the \href{http://www.ginac.de}{\GiNaC} library (originally developed as part of the one-loop particle physics project \href{http://wwwthep.physik.uni-mainz.de/~xloops/}{\xloopsginac}~\cite{Bauer:2001ig}) for symbolic algebra. It uses \href{http://www.mpi-forum.org}{\MPI} multi-process communication to parallelize calculations and scales from laptop-class hardware up to many cores on a HPC cluster, without requiring a shared-memory architecture. Results are stored as \href{https://www.sqlite.org}{\SQLite} databases, enabling sophisticated postprocessing using SQL database queries. For simple analyses a suite of built-in visualization and analysis tools are provided. \end{itemize} Both codes automate the calculation of two- and three-point field-space correlation functions (and hence two- and three-point $\zeta$ correlation functions) directly from a potential. They implement the same computational scheme, but in slightly different ways. In this paper our intention is to describe this scheme, illustrate its utility, and explain its relationship to earlier work. We refer to the codes when giving examples, but they will be described more fully elsewhere~\cite{CppTransportUserGuide, PyTransportUserGuide}. They can be downloaded by following the links given in~\S\ref{sec:automated-codes}. \para{Synopsis} This paper is divided into three parts. First, in~\S\ref{sec:why-automated} we review the current state-of-the-art in computing inflationary three-point functions. This depends on accurate calculation of vacuum fluctuations over a time-dependent cosmological background for a variety of masses and parameter regimes. In some regimes it is possible to find successful analytic approximations, but in others the success of such approximations has been meagre. We argue that dramatic short-term improvements in our ability to make analytic estimates are unlikely, making a numerical method essential. Second, in~\S\S\ref{sec:numerical-computation}--\ref{sec:gauge-transform} we describe our numerical approach. In~\S\S\ref{sec:numerical-computation}--\ref{sec:transport-equations} we show that it can be regarded as a reformulation of the Schwinger or `in--in' method to compute expectation values---as distinct from transition amplitudes, which are the output of the Feynman calculus described in particle physics textbooks. As with all formulations of field theory, ours has advantages and disadvantages. A major advantage is that it is very well suited to numerical evaluation. For example, challenging steps such as Wick rotation (that are best handled analytically) appear only in the calculation of initial conditions but \emph{not} in the subsequent evolution equations. Suitable initial conditions need be calculated only once and are universal for all models, no matter what parameter regime they inhabit. We describe this calculation in~\S\ref{sec:initial-conditions}, and the transition to $\zeta$ correlation functions in~\S\ref{sec:gauge-transform}. The major disadvantage of our method is that it obscures the clear physical interpretation of Feynman diagrams as processes occurring in spacetime~\cite{Seery:2009hs,Arkani-Hamed:2015bza}. But since we intend to apply our method for numerical evaluation this is not so important. Third, in~\S\S\ref{sec:automated-codes}--\ref{sec:performance} we briefly describe our implementations {\PyTransport} and {\CppTransport} before using them to showcase the utility of the method. In~\S\ref{sec:examples} we verify numerically that our formalism successfully tracks the evolution of the two- and three-point functions from sub- to super-horizon scales. We demonstrate that it can be used to extract very delicate variation with shape and scale, and produce examples that exemplify some of the physical processes discussed in~\S\ref{sec:why-automated}. In each case, our analysis yields results that can be obtained only using numerical techniques. In~\S\ref{sec:performance} we discuss the numerical characteristics of our method and its performance as implemented in {\PyTransport} and {\CppTransport}. In particular, we explain how the integration time and convergence properties scale with key adjustable parameters. Finally, we conclude in~\S\ref{sec:conclusions}. \para{Notation} We work in units where $c = \hbar = 1$ and use the reduced Planck mass $\Mp^2 = (8\pi G)^{-1}$, where $G$ is Newton's gravitational constant. Our metric signature is $(-,+,+,+)$. We work in the Heisenberg picture except where otherwise stated. In order to write concise expressions we use a number of different summation conventions. For details, see the discussion below Eq.~\eqref{eq:zeta-twopf-factorization}; the discussion around Eq.~\eqref{eq:extended-summation-example}; and the discussion above Eqs.~\eqref{eq:action-S2}--\eqref{eq:action-S3}. | \label{sec:conclusions} The major result of this paper is a complete formalism for numerical calculation of the tree-level correlation functions produced during an epoch of early-universe inflation. This formalism was described in~\S\S\ref{sec:numerical-computation}--\ref{sec:gauge-transform}. It does not require the slow-roll approximation except to obtain the estimates of initial conditions given in~\S\ref{sec:initial-conditions}. As explained in~\S\ref{sec:tree-level}, the tree-level approximation means that our formalism should produce accurate estimates unless multiparticle production channels make a significant contribution to the curvature perturbation $\zeta$, for example by nontrivial scattering processes or decays. In certain scenarios, such as warm inflation or trapped inflation, the tree approximation may also fail to capture processes by which energy is drained from the zero-mode into finite-wavenumber excitations. These restrictions should be carefully considered before applying our tools---or any others based on a tree-level approximation---to study some particular inflationary model. We are supplying two concrete implementations of this formalism, described in~\S\ref{sec:automated-codes} and available for download under open source licences. These are not just bare implementations of the evolution equations; instead, they both support \emph{automated} analysis of models from a high-level Lagrangian description, and can be used to produce immediate high-resolution numerical results for models whose bispectra were previously intractable. In~\S\S\ref{sec:numerical-computation}--\ref{sec:gauge-transform} we focused on the correlation functions generated by a system of canonically-normalized scalar fields (and their contribution to the curvature perturbation $\zeta$), and the two bispectrum-level implementations described in~\S\ref{sec:automated-codes} are currently restricted to scenarios of this type. However, this restriction is not necessary as a matter of principle. Extensions to more complex models, such as those with a nontrivial kinetic sector $G_{\alpha\beta}(\phi) \partial_a \phi^\alpha \partial^a \phi^\beta$ or interactions of Galileon-type, are limited only by algebraic complexity. All that would be required are replacements for the tensors ${u^{\ext{a}}}_{\ext{b}}$ and ${u^{\ext{a}}}_{\ext{b}\ext{c}}$, and appropriate initial conditions that account for the new interactions and kinetic structure. An older Mathematica-based implementation capable of handling models with nontrivial field-space metric is already available, although its output is limited to the two-point function~\cite{Dias:2015rca}. All three implementations, together with links to further resources, can be found at the website \href{https://transportmethod.com}{transportmethod.com}. In~\S\ref{sec:examples} we exhibited results for a selection of concrete models exemplifying the wide range of mass spectra and coupling constants that can be accommodated. Collectively, these demonstrate that our formalism successfully tracks highly nuanced features of the bispectrum amplitude and shape. Just as important, because it includes all relevant effects (especially from gravitational-strength couplings), it can be used to predict the complete scale- and shape-dependent bispectrum generated in models where the reduced bispectrum $\fNL(k_1, k_2, k_3)$ is order unity. These models are a target for the next generation of galaxy surveys, including Euclid, DESI and LSST~\cite{Alvarez:2014vva}. To obtain robust predictions with these low amplitudes it is not sufficient to rely on approximations that discard physical effects occurring on subhorizon scales or around horizon exit, or that do not accurately account for the effect of hierarchies among the wavenumbers appearing in each correlation function. Our formalism, and especially the reusable implementations we provide, supply a means for these models to be accurately analysed for the first time. Meanwhile, to achieve a suitable level of preparation for a Euclid-, DESI- or LSST-like survey it will be insufficient merely to improve the accuracy of primordial calculations. Reliable forecasts for models where $|\fNL(k_1, k_2, k_3)| \lesssim 1$ must at least account for gravitational evolution after horizon exit and the characteristics of the survey. These details are now well-understood; what is required is an integrated toolchain that links them all together. At the simpler level required by CMB experiments, an analysis such as that given in Ref.~\cite{Byrnes:2015dub}---% which relied on numerical bispectra produced using the {\CppTransport} platform---% demonstrates how accurate, high-resolution calculations of primordial correlation functions can be integrated into a numerical toolchain producing accurate, reliable results customized to a specific experiment. (In Ref.~\cite{Byrnes:2015dub} the customization was for a Planck-like CMB survey, but the point we are making is much more general.) As we stockpile datasets of ever-increasing accuracy there is a corresponding burden on theorists to generate predictions with matching refinement. No one tool, or single approach, will be sufficient---but we hope that that the software tools we are making available constitute a step towards this goal. | 16 | 9 | 1609.00379 |
1609 | 1609.06308.txt | We use the combined data from the COS-GASS and COS-Halos surveys to characterize the Circum-Galactic Medium (CGM) surrounding typical low-redshift galaxies in the mass range $\rm~M_*\sim~10^{9.5-11.5}~M_{\odot} $, and over a range of impact parameters extending to just beyond the halo virial radius ($\rm~R_{vir}$). We find the radial scale length of the distributions of the equivalent widths of the \Lya and \ion{Si}{3} absorbers to be 0.9 and 0.4 $\rm~R_{vir}$, respectively. The radial distribution of equivalent widths is relatively uniform for the blue galaxies, but highly patchy (low covering fraction) for the red galaxies. We also find that the \Lya and \ion{Si}{3} equivalent widths show significant positive correlations with the specific star-formation rate (sSFR) of the galaxy. We find a surprising lack of correlations between the halo mass (virial velocity) and either the velocity dispersions or velocity offsets of the \Lya lines. The ratio of the velocity offset to the velocity dispersion for the \Lya absorbers has a mean value of $\sim$ 4, suggesting that a given the line-of-sight is intersecting a dynamically coherent structure in the CGM rather than a sea of orbiting clouds. The kinematic properties of the CGM are similar in the blue and red galaxies, although we find that a significantly larger fraction of the blue galaxies have large \Lya velocity offsets ($>$~200~\kms). We show that - if the CGM clouds represent future fuel for star-formation - our new results could imply a large drop in the specific star-formation rate across the galaxy mass-range we probe. | } Galaxy growth is fundamentally connected to the cycle of accretion and ejection of matter into and out of galaxies. In the simplest picture, galaxies acquire gas that reaches the central regions via the circum-galactic medium (CGM). There it condenses into neutral and then molecular gas, some of which is then converted into stars. Young stars in turn drive strong winds, outflows, and radiation that deposit mass, metals, energy, and momentum to the CGM, thus significantly influencing its properties \citep[see review by][and references therein]{somerville15,fielding16}. These linked processes are commonly termed the baryon cycle. The CGM then lies at the heart of this cycle, as it is the interface between the stellar body of the galaxy and the intergalactic medium. It is the primary spatial pathway for the baryon cycle into and out of galaxies \citep[][and references therein]{ford16, tumlinson13, borthakur15, nielsen15, shen14, mitra15, brook14}. The CGM is also a reservoir of low-density gas that may have as much mass as the stellar component of the galaxy \citep{werk13,werk14, tumlinson13, peeples14, richter16}. It extends out from the stellar disk out to the virial radius of the galaxy \citep{chen01a, stocke13, borthakur13}. However, due to its low surface-brightness, we have not yet been able to directly image this vast baryonic reservoir. On the other hand, absorption-line spectroscopy provides an avenue to probe the physical conditions in this low-density gaseous medium. Rest-frame ultra-violet (UV) spectroscopy enables us to use various absorption-line transitions, including hydrogen and metal-line species spanning a broad range of ionization states. Mapping the CGM with the help of a large sample of sightlines probing a range of impact parameters is crucial for understanding its properties and its variations as a function of radius. The radial dependence in the properties of the neutral hydrogen in the CGM has been known for decades, based on observations of the \Lya absorption-line \citep[][and references therein]{lanzetta95, chen98, tripp98, chen01b, bowen02, prochaska11, stocke13, tumlinson13, liang14, borthakur15}. However, only recently, with the installation of Cosmic Origins Spectrograph (COS) aboard the Hubble Space Telescope (HST), has it become feasible to undertake detailed probes of the CGM properties as a function of other global properties of the central galaxy. One of the consequences of the accretion of gas passing through the CGM is that this provides the raw material to sustain the growth of the galaxy via star formation \citep[e.g.][]{bouche13}. Not all galaxies produce stars at the same rate \citep{brinchmann04, salim07, noeske07, daddi07, rodighiero11, speagle14,snyder15}. In particular, galaxies show two distinct populations in terms of their star-formation rate (SFR). While most low mass galaxies form stars at significant rates, most high mass galaxies produce stars at negligible levels. This was termed as the galaxy color bimodality defined in terms of ``blue" (star-forming) galaxies and ``red" (quiescent) galaxies \citep[e.g.][]{kauffmann03, blanton03, baldry04, brinch04, tully82}. About a decade back, cosmological hydrodynamical simulations revealed two distinct ways that galaxies accrete gas into their dark matter halo as a function of halo mass. The predominant mode of gas accretion for low mass galaxies is believed to be the ``cold" mode \citep{keres05, keres09, dekel09}, where gas falls into galaxies as streams or lumps at temperatures much less than that of the virial temperature. For the higher mass halos, the accretion process is expected to be in the ``hot" mode \citep{white91, fukugita06}, in which the incoming gas shock heats to the virial temperature. This broadly can explain why high-mass galaxies have little to no cold gas reservoirs to fuel star-formation \citep[see work on condensation in hydrodynamical simulations by ][]{kaufmann06, kaufmann09, sommerlarsen06}. Galaxies also recycle gas from previous generations of star-formation that is stored in their CGM \citep{ford13,fraternali15}. However, the process of how gas gets into the disk from the CGM is fairly complex. Nonlinear perturbations in the filamentary flows may help the cool accreting gas condense and add cold gas to the disk \citep{keres_hernquist09, joung12}. These condensing clouds may contain as much as 25\%-75\% of the cold gas in the CGM \citep{fernandez12}. %FIGURE1 \begin{figure*} \includegraphics[trim = 20mm 105mm 20mm 0mm, clip,scale=0.535,angle=-0]{f1a.ps} \includegraphics[trim = 20mm 105mm 20mm 0mm, clip,scale=0.535,angle=-0]{f1b.ps} \caption{Distribution of galaxy properties for the COS-GASS and COS-Halos samples. The left panel shows the stellar mass distribution and the right panel shows the virial radius distribution. The $\rm R_{vir}$ for both the samples were estimated using the prescription described by \citet{kravtsov14, mandelbaum16,liang14} as described in section 2.1.} \label{fig-sample_m_Rvir} \end{figure*} %FIGURE2 \begin{figure*} \hspace{-0.6cm} \includegraphics[trim = 0mm 0mm 0mm 0mm, clip,scale=0.535,angle=-0]{f2a.ps} \includegraphics[trim = 0mm 0mm 0mm 0mm, clip,scale=0.535,angle=-0]{f2b.ps} \caption{Variation of \Lya and Si~III equivalent width with normalized impact parameter (i.e. $\rm \rho/R_{vir}$) for a combined COS-GASS and COS-Halos sample. The colors blue and cyan indicate``blue" galaxies and red and yellow denote ``red" galaxies. The black thick line denotes the fits to the data using the Buckley-James method. The calculations were performed using the survival analysis software ASURV that takes into account the censored data. The parameters describing the best-fit lines are printed at the bottom left corner. Since the fits presented here take into account the censored data, the parameters of the best-fit in the left panel are slightly different from those published by \citet{borthakur15}.} \label{fig-W_rho_Rvir} \end{figure*} In addition to accretion, star-formation driven feedback may change the nature and properties of the gas in the CGM \citep{kauffmann16, liang16, nelson15, nelson16, marasco15}. Massive young stars inject energy and/or momentum into outflows \citep{veilleux05, heckman11, borthakur14, heckman15, heckman16} that may travel into the CGM, enriching it with metals, shock-heating the cooler CGM clouds, and possibly even expelling/unbinding the CGM \citep{borthakur13}. Therefore, if feedback provided by massive stars plays a role in the observed bimodality, then we should see a change in the structure, ionization state, and/or kinematics of the CGM as a function of SFR. To that end, we have selected a subsample of galaxies from the GALEX Arecibo SDSS Survey \citep[GASS;][]{catinella10,catinella12,catinella13} that have background UV-bright quasi-stellar objects (QSOs) located within a projected distance of 250~kpc in the rest-frame of the galaxy. This yielded the COS-GASS sample \citep{borthakur15} whose members were observed with COS using the G130M grating. This provided a spectral R$ =$ 20,000$-$24,000 (FWHM $\sim$ 12 to 15~\kms). We have multi-band data for these galaxies from the parent GASS survey: 21~cm \HI spectroscopic data obtained with the Arecibo telescope, optical images and spectroscopy from the Sloan Digital Sky Survey (SDSS), UV imaging with the Galaxy Evolution Explorer (GALEX), molecular gas data from IRAM \citep[COLD GASS;][]{saintonge11}, and long-slit optical spectroscopy \citep{moran12} for a portion of the sample. Therefore, we have the stellar mass, SFR, gas-phase metallicity, stellar morphology, and atomic and molecular gas masses for all the 45 galaxies from the COS-GASS sample. Here we present our study utilizing the combined COS-GASS \citep{borthakur15} and COS-Halos \citep{tumlinson13} samples. Detailed descriptions of our sample, the COS observations and data reduction are presented in Section~2. The results are presented in Section~ 3 and their implications are discussed in Section~4. Finally, we summarize our findings in Section~5. The cosmological parameters used in this study are $H_0 =70~{\rm km~s}^{-1}~{\rm Mpc}^{-1}$ (in between the two recent measurements of $\rm 73.24 \pm 1.74 ~ km~s^{-1}~Mpc^{-1}$ \citep{riess16} and $\rm 67.6_{-0.6}^{+0.7}~{\rm km~s}^{-1}~{\rm Mpc}^{-1}$ \citep{grieb16}) , $\Omega_m = 0.3$, and $\Omega_{\Lambda} = 0.7$. We note that varying the Hubble constant value from $\rm 65~ to~ 75 ~ km~s^{-1}~Mpc^{-1}$ does not affect the conclusions in the paper. %FIGURE3 \begin{figure} \hspace{-0.2cm} \includegraphics[trim = 0mm 0mm 0mm 0mm, clip,scale=0.5, angle=-0]{f3.ps} \caption{ The cross-correlation function between the galaxy systemic velocities and velocity centroids of \Lya absorbers. The cross-correlation function was calculated using the same data analysis criterion as the observations/measurements. For example, the absorbers were randomly distributed within the allowed velocity range of $\pm$600~\kms. Caution must be applied when comparing these results to those from blind surveys or surveys with different intrinsic resolutions for the spectrograph. } \label{fig-cc_lya} \end{figure} | } We have presented the analysis of a comprehensive data set combining the COS-GASS and COS-Halos samples to probe the CGM of low-z galaxies spanning a stellar mass range of almost two orders-of-magnitude centered on the characteristic mass ($\sim 10^{10.5} \rm M_{\odot}$) at which the galaxy population transitions from mostly blue, star-forming galaxies to red, quiescent ones. These two surveys cover similar ranges in stellar masses and dark halo virial radii ($\rm R_{vir}$). In addition, the combined sample uniformly samples a large range of radial distances from 0.02 to 1.3~$\rm R_{vir}$. The COS-GASS survey primarily samples the outer CGM and COS-Halos survey primarily samples the inner CGM. We characterized the CGM properties, including its radial profile, its kinematics, and its correlation with the global properties of the galaxies. In particular we have divided the sample into blue galaxies (with sSFR $\rm >~10^{-11}~ M_{\odot}yr^{-1}$) and red galaxies (with lower sSFR). In this analysis, we discussed the \Lya $\lambda$1215$\rm \AA$ and \ion{Si}{3} $\lambda$1206$\rm \AA$ transitions tracing intermediate ionization gas. \ion{Si}{3} was chosen as it is the strongest metal transition detected in the combined data set. The typical detection limit for the COS-GASS sample is $\sim \rm 50~m\AA$ corresponding to 3$\sigma$ uncertainty in the data. In the combined sample, the detection rates of \Lya and \ion{Si}{3} were 91\% and $\sim$50\% respectively. Based on the analysis of the combined sample we conclude the following: \begin{itemize} \item[1.] The radial distribution of the equivalent width of \Lya as a function of normalized impact parameter ($\rm \rho/R_{vir}$) can be expressed as an exponential. The scale-lengths are similar for the red and blue galaxies (0.75 and 0.72 $\rm R_{vir}$ respectively). The radial distribution of equivalent width of \ion{Si}{3} can also be expressed as an exponential with a scale-length of 0.36 (0.33) $\rm R_{vir}$ for the red (blue) galaxies. The detection rate of \ion{Si}{3} drops to almost zero beyond about 0.8 $\rm R_{vir}$. \item[2.] The blue galaxies show a relatively uniform radial distribution of \Lya absorbers, implying an areal covering fraction of nearly 100\% in the CGM. In contrast, the \Lya absorbers have a much less uniform radial distribution in the CGM of the red galaxies, suggesting a patchy distribution with smaller areal covering fractions. These differences are reflected in the overall normalization of the radial distribution of equivalent widths, which is higher for the blue galaxies (by 0.45 dex). Similar results were found for \ion{Si}{3}, but are restricted to the region interior to 0.8 $\rm R_{vir}$ (where \ion{Si}{3} is detected). \item[3.] We found a significant positive correlation between the equivalent width of \Lya and the star-formation rate (at the 99.8\% confidence level). The correlation is even more significant for normalized quantities: the impact- parameter-corrected equivalent with of \Lya ([Log W - $\rm \overline{Log~W}]_{Ly\alpha}$) and the specific SFR (SFR/M$_{\star}$) were found to correlate at the 99.99\% confidence level. Similar results were found for \ion{Si}{3}. \item[4.] We found the velocity distribution of the centroids of the majority of the \Lya and \ion{Si}{3} to generally lie within $\sim$150~\kms of the systemic velocity of the galaxy. These velocities are smaller than the escape velocity, thus suggesting the gas seen in absorption is the gravitationally bound within the halo. The metal-line transitions are also found mostly within $\pm$ 100~\kms of \Lya absorbers, although not all strong ($\rm >0.3\AA$) \Lya absorbers showed associated \ion{Si}{3}. \item[5.] We find that the velocity offset between the \Lya centroid and the systemic velocity ($\Delta v$) is usually significantly larger than the line-of-sight velocity dispersion of the \Lya line ($\sigma_{los}$). The mean ratio $\Delta v/\sigma_{los} \sim4$. \item[6.] We find no dependence of the kinematic properties of the CGM ($\Delta v$ or $\sigma_{los}$) on the galaxy halo mass (virial velocity). This is surprising, as the sample spans ranges of about $10^{2}$ in halo mass and $\sim$5 in $v_{vir}$. \item[7.] We found that the kinematic properties of the CGM are generally similar between the blue and red galaxies. However, while the majority of both the blue and red galaxies have $\Delta v <$ 100 \kms, the distribution of $\Delta v$ for the blue galaxies shows a pronounced tail out to values as high as 500 \kms. \item[8.] We found a significant change in the CGM kinematics at about a radius of 1.0 (0.7) $\rm R_{vir}$ for the blue (red) galaxies. In the outer CGM $\Delta v$ for the \Lya absorbers is always less than 150 \kms, while the distributions of $\Delta v$ show tails out to values as high as 500 \kms in the inner CGM. In addition, $\sigma_{los}$ is higher on average in the outer CGM for both the blue and red galaxies. These two results lead to a corresponding decrease in $\Delta v/\sigma_{los}$ in the outer CGM. \end{itemize} The combined COS-GASS and COS-Halos sample has allowed us to conduct a comprehensive study of the connection of the properties of the CGM with those of the stellar body of the galaxy. We think that three of the specific results from above are particularly noteworthy. First, the differences in the radial distributions of the \Lya {\it vs.} the \ion{Si}{3} absorbers, suggest that the inner CGM is being (or has at some time been) affected by feedback associated with massive stars and supernovae. This feedback has chemically-enriched the CGM. It is interesting that the kinematics of the inner and outer CGM also show differences, although it is not clear that these are related to feedback or to projection effects. Secondly, the fact that the typical ratio of the velocity offset to the line-of-sight velocity dispersion for the \Lya absorption-lines is so large is an important clue as to the structure of the CGM. It implies that a line-of-sight through the CGM does not intersect a whole sea of many clouds orbiting in the halo, but is rather passing through a coherent structure (cloud, sheet, filament). Model-dependent estimates imply a path-length of order 1 to 10 kpc for these structures. Thirdly, the independence of the kinematic properties of the warm CGM on the halo mass is quite remarkable. This implies that, even though the observed absorption-line systems are mostly gravitationally bound to the halo, simple gravitational forces alone do not adequately explain the CGM dynamics. For the massive red galaxies, the ``sub-virial" velocities could be understood if the absorbers represent material cooling and condensing out of (or suffering drag as they move through) a hot volume-filling phase that is supported hydrostatically against gravity. A major motivation of this study was to understand how and why galaxies in this stellar mass regime exhibit the color bimodality stemming from a suppression/cessation of star-formation in some of them. Since the CGM is the interface through which galaxies could exchange gas and energy that is required to form stars (or is expelled as a result of star-formation), the CGM properties could hold clues as to how this process of gas delivery may be disrupted leaving some galaxies deprived of fuel to form stars. We have explored a simple scenario in which the star-formation rate in a galaxy is proportional to the total mass of CGM clouds divided by an inflow time. We then show that the empirical results on the independence of CGM kinematic properties on halo mass and the smaller covering factor in the CGM in the red galaxies would imply a drop in the specific star formation rate by about a factor of 30 over the stellar mass range from 10$^{10}$ to 10$^{11.5}$ M$_{\odot}$. In any event, we believe the data we have presented here provide a valuable observational resource for on-going and future numerical simulations that try to reproduce CGM properties such as \Lya and metal-line column density profiles, covering fraction, and dynamics, such as line-widths and velocity spreads \citep[for example][and references therein]{hummels12, stinson12, shen13, ford13, ford14, ford16, liang16, kauffmann16, fielding16}. The ultimate goal is understanding the role of the CGM in the evolution of galaxies. \vspace{.5cm} | 16 | 9 | 1609.06308 |
1609 | 1609.05901_arXiv.txt | We calculate the sensitivity to a circular polarization of an isotropic stochastic gravitational wave background (ISGWB) as a function of frequency for ground- and space-based interferometers and observations of the cosmic microwave background. The origin of a circularly polarized ISGWB may be due to exotic primordial physics (i.e., parity violation in the early universe) and may be strongly frequency dependent. We present calculations within a coherent framework which clarifies the basic requirements for sensitivity to circular polarization, in distinction from previous work which focused on each of these techniques separately. We find that the addition of an interferometer with the sensitivity of the Einstein Telescope in the southern hemisphere improves the sensitivity of the ground-based network to circular polarization by about a factor of two. The sensitivity curves presented in this paper make clear that the wide range in frequencies of current and planned observations ($10^{-18}\ {\rm Hz} \lesssim f \lesssim 100\ {\rm Hz}$) will be critical to determining the physics that underlies any positive detection of circular polarization in the ISGWB. We also identify a desert in circular polarization sensitivity for frequencies between $10^{-15}\ {\rm Hz} \lesssim f \lesssim 10^{-3}\ {\rm Hz}$, given the inability for pulsar timing arrays and indirect-detection methods to distinguish the gravitational wave polarization. | Starting with the first stargazers, our knowledge of the heavens has come in the form of electromagnetic waves. The intensity and polarization of these massless messengers have been shown to contain a wealth of information about the physics and astrophysics of distant objects and the conditions along the line of sight, extending to the earliest moments after the primordial universe became transparent. Now the same is coming true for gravitational waves \cite{Abbott:2016blz}. In this paper we consider the frequency-dependent sensitivity of the most common gravitational wave detection techniques to a net circular polarization of an isotropic stochastic gravitational wave background (ISGWB). Since gravitational waves have two polarizations, any stochastic gravitational wave background can be expanded in terms of the standard four Stokes parameters: $I$, $Q$, $U$, and $V$. However, given the spin-2 nature of gravitational waves, the $Q$ and $U$ linear polarizations are only non-zero for anisotropic backgrounds (with the first non-zero contribution at the quadrupole). On the other hand both $I$ and $V$ are scalar quantities, and as such, may be non-zero for isotropic backgrounds. Since most stochastic gravitational wave backgrounds are predicted to be nearly isotropic we only consider the sensitivity of the most common techniques to the intensity, $I$, and level of circular polarization, $V$. The detection of a non-zero circularly-polarized ISGWB would indicate new fundamental physics \cite{PhysRev.106.388}. Leading examples consist of inflationary models in which the inflaton couples to the parity-odd Chern-Simons scalar of a U(1) vector field, as in Refs.~\cite{Anber:2009ua,Anber:2012du}, or the inflaton couples similarly to a non-Abelian SU(2) gauge field, as in chromo-natural inflation \cite{Adshead:2012kp,Adshead:2013qp} or gauge-flation \cite{Maleknejad:2011jw,Namba:2013kia} and its variants \cite{Maleknejad:2012fw,Cai:2016ihp,Dimastrogiovanni:2016fuu,Adshead:2016}. Through different mechanisms, these scenarios all generate a primordial spectrum of gravitational waves with a scale-dependent chiral asymmetry, whereby the spectra for left- and right-circular polarizations differ. An inflaton that couples directly to the gravitational Chern-Simons scalar \cite{Alexander:2004us,Alexander:2009tp} and other quantum gravity schemes \cite{Takahashi:2009wc,Contaldi:2008yz} will also generate an asymmetry. In the case of chromo-natural and gauge-flation, the spectrum is chirally-symmetric below a cutoff scale typically close to the present-day horizon radius, and chirally-asymmetric with a blue tilt on smaller scales. This opens the possibility of detection across a wide range of frequencies by the cosmic microwave background (CMB), satellite, and ground-based detectors. Numerous gravitational wave observatories are on line or in planning stages. LIGO and VIRGO are already taking data; Pulsar Timing arrays (PTAs) may expect to see a signal in the near future; LIGO India and KAGRA are under development; and a global network of ground-based interferometers has been proposed, under the name Einstein Telescope. Technology is already developing for space-based observatories such as the Evolved Laser Interferometer Space Antenna (eLISA) \cite{eLISA} (recently re-named LISA) and the Big Bang Observer (BBO) \cite{Crowder:2005nr}. These independent but complementary observatories are sensitive to different frequencies. On the largest scales (i.e., frequencies of $f \simeq 10^{-18}\ {\rm Hz}$) we may detect gravitational waves through their effects on the CMB. Observations of both the intensity (i.e., temperature) and linear polarization of the CMB give information about the properties of a stochastic gravitational wave background on scales equal to the size of the observable universe. The autocorrelation of the temperature and B-mode polarization provides an estimate of the intensity of a possible ISGWB whereas the cross-correlation of the temperature and B-mode polarization as well as the cross-correlation of the E and B-mode polarization provide estimates of the level of net circular polarization. Pulsar timing arrays (PTAs) are most sensitive to gravitational waves at frequencies $f \simeq 10^{-9}\ {\rm Hz}$. However, because of the effective geometry of the detector -- the fact that we measure each pulse time of arrival at the Earth -- PTAs are not sensitive to the circular polarization of an ISGWB. At moderate frequencies, $f \simeq 1\ {\rm Hz}$, space-based laser interferometers can be made to be sensitive to the circular polarization of an ISGWB by correlating the signals recorded by two independent observatories lying in different planes. Finally at high frequencies ($f \simeq 100\ {\rm Hz}$) the correlation between signals of ground-based laser interferometers are already sensitive to the circular polarization of the ISGWB. Previous work has considered the sensitivity to the circular polarization of the ISGWB for gound-based and space-based gravitational wave observatories. Refs.~\cite{Seto:2006hf,Seto:2006dz,Seto:2007tn,Seto:2008sr,Seto:2008sr2} consider the sensitivity of ground-based and space-based interferometers and Ref.~\cite{Gluscevic:2010vv} considers the sensitivity of measurements of CMB polarization. Much of this is included in the exhaustive review in Ref.~\cite{Romano:2016dpx}. We extend this work in several ways. First, this paper presents the sensitivity to circular polarization in a consistent framework for each observatory. This allows us to gain a clearer intuition for what type of observatory will provide useful information on the circular polarization as well as provide formulae which can be used to calculate the sensitivity of future observatories. Second, while previous work calculated the sensitivity to a flat spectrum (i.e. $\Omega_{\rm gw} = {\rm constant}$), we present the full sensitivity curves, which determines the frequency range associated with each observatory and allows a comparison of the sensitivity to non-flat spectra. Finally, we consider the sensitivity of several observatories (such as the Einstein Telescope) which were not included in previous work. % This paper is organized as follows: In Sec.~II we discuss the basic physics of a gravitational wave detector, present a calculation of the optimal signal to noise, discuss the properties of the ISGWB, and present the method we use to calculate the sensitivity curves. In Sec.~III we calculate the sensitivity curves for space-based observatories. In Sec.~IV we calculate the sensitivity curves for a network of ground-based observatories. In Sec.~V we calculate the sensitivity curves for observations of the CMB. In Sec.~VI we present our conclusions. | \begin{figure}[h!] \begin{center} \resizebox{!}{7.5cm}{\includegraphics{All_plot.pdf}} \caption{A collection of all of the sensitivity curves calculated in this paper with solid blue giving the sensitivity to the intensity and dashed orange to the level of circular polarization. We have also included the sensitivity to the ISGWB from pulsar timing array from Ref.~\cite{Ellis:2012zv} and indirect limits coming from CMB measurements of the radiative energy density of the universe \cite{Smith:2006nka,Pagano:2015hma}.} \label{fig:ALL} \end{center} \end{figure} As shown in this paper, most of the common techniques used to detect the ISGWB will be sensitive to both the intensity and level of circular polarization. We have summarized the sensitivity curves calculated in this paper in Fig.~\ref{fig:ALL}: the solid blue curves show the sensitivity to the intensity and the dashed orange curves show the sensitivity to the level of circular polarization. Space-based detectors will be sensitive to both the intensity and circular polarization as long as they utilize more than three inertial masses. We have considered the case where these detectors operate as a constellation of two equilateral triangles. The two triangles must be separated by some distance, and there is a distance at which the overall sensitivity to both the intensity and circular polarization are equal, in agreement with Ref.~\cite{Seto:2006dz}. In addition to this we found that this optimal distance has a strong dependence on the specifications of the observatory -- for LISA we found that the optimal distance $D \simeq 7L$ whereas for BBO $D \simeq 2L$. Ground-based detectors are sensitive to both the intensity and circular polarization as long as we correlate the signal from at least three widely separated sites. This means that the current collection of ground-based detectors (LIGO Hanford and LIGO Livingston) are not capable of separating out these two signals. However, with VIRGO and KAGRA soon to turn on, the ground-based network will become sensitive to both signals. We find that this total network sensitivity is greatly enhanced if we include the Einstein Telescope. Since the intrinsic sensitivity of the Einstein Telescope to the intensity is significantly better than current gravitational wave observatories, it has a disproportionate effect on the overall sensitivity to the intensity. However, it also significantly improves the network's sensitivity to the level of circular polarization. We also found that if we were to locate the Einstein Telescope in the southern hemisphere the improvement in the total sensitivity to the level of circular polarization further improves by another factor of two. Observations of the temperature and polarization of the CMB are sensitive to both the intensity and circular polarization of the ISGWB. The correlation between the CMB temperature and the $E$ and $B$ mode polarization can isolate the effects of the ISGWB intensity from those of the circular polarization. In agreement with Ref.~\cite{Gluscevic:2010vv} we find that the Planck satellite is equally (in)sensitive to the intensity and circular polarization of the ISGWB, but that a future CMB satellite dedicated to measuring the CMB polarization -- CMBpol -- will improve the sensitivity by three orders of magnitude for the intensity of the ISGWB and two orders of magnitude for the circular polarization. As opposed to reporting the sensitivity as a single number, the calculation of sensitivity curves gives a quantitative accounting of the frequency coverage by these various observatories. Looking at the combination of all of the observatories considered in this paper in Fig.~\ref{fig:ALL}, it is interesting to note the absence of any detector operating at frequencies between $10^{-15}\ {\rm Hz} \lesssim f\lesssim 10^{-3}\ {\rm Hz}$ which will be sensitive to the level of circular polarization. This 18 orders of magnitude is a wide swath of frequency space inside of which we do not have any known technique to detect the circular polarization of the gravitational wave background. This sensitivity desert calls out for new and creative ideas on how to detect the level of circular polarization in an ISGWB. | 16 | 9 | 1609.05901 |
1609 | 1609.00009_arXiv.txt | We explore the rates of tidal disruption events (TDEs) of stars by supermassive black holes (SBHs) in galactic nuclei formed in mergers followed by a formation and coalescence of a binary SBH. Such systems initially have a deficit of stars on low-angular-momentum orbits caused by the slingshot process during the binary SBH stage, which tends to reduce the flux of stars into the SBH compared to the steady-state value. On the other hand, a newly formed galactic nucleus has a non-spherical shape which enhances the mixing of stars in angular momentum and thus the TDE rate. In galaxies with relatively low SBH masses ($\lesssim 10^7\,M_\odot$), relaxation times are short enough to wash out the anisotropy in initial conditions, and for more massive SBH the enhancement of flux due to non-sphericity turns out to be more important than its suppression due to initial anisotropy. Therefore, the present-day TDE rates generally exceed conventional steady-state estimates based on a spherical isotropic approximation. We thus conjecture that the lower observationally inferred TDE rates compared to theoretical predictions cannot be attributed to the depletion of low-angular-momentum stars by SBH binaries. | Tidal disruption events (TDEs) are luminous short-living flares that are believed to happen in galactic nuclei containing supermassive black holes (SBHs): a close encounter between a star and SBH leads to the disruption of the star if the gradient of gravitational force is large enough \citep{Rees1988}. These events leave imprints in the entire electromagnetic spectrum: optical \citep{vanVelzenFarrar2014}, UV \citep{Gezari2008} or X-ray \citep{Donley2002, KhabibullinSazonov2014} observations imply an average rate of $10^{-4}-10^{-5}$ events per year per galaxy \citep[e.g.,][]{Komossa2015}. The loss-cone theory describing the flux of stars into a SBH was first developed in 1970s in the context of globular clusters \citep[e.g.,][]{FrankRees1976, LightmanShapiro1977}, and later applied to galactic nuclei \citep{SyerUlmer1999,MagorrianTremaine1999,WangMerritt2004,MageshwaranMangalam2015,Kochanek2016,Aharon2016}. Even though the TDE rates inferred from observations are of the same order as the theoretical estimates, the latter are systematically higher, as stressed by \citet{StoneMetzger2016},\defcitealias{StoneMetzger2016}{SM2016} hereafter SM2016. However, uncertainties are high on both sides, and the rate of \textit{observable} events depends both on their \textit{intrinsic} rate and the details of emission mechanisms in the process of tidal disruption itself. In this paper we deal only with the first factor, namely the stellar-dynamical estimate of the flux of stars into the loss cone of a SBH. The simplest and most commonly used estimate is based on the steady-state solution of the Fokker--Planck equation describing the diffusion of stars in angular momentum for a spherically-symmetric stellar system, driven by two-body relaxation. There are many physical effects that may complicate the model: resonant relaxation \citep{RauchTremaine1996,HopmanAlexander2006,Merritt2015,BarOrAlexander2016}, massive perturbers \citep{Perets2007}, non-spherical geometry \citep{MagorrianTremaine1999,MerrittPoon2004,HolleySigurdsson2006,VasilievMerritt2013,Vasiliev2014}, anisotropy in the initial conditions \citep[hereafter Paper I]{MerrittWang2005,LezhninVasiliev2015}\defcitealias{LezhninVasiliev2015}{Paper I}. All but the last one generally tend to increase the TDE rate compared to the `reference' spherical steady-state value (SSS), and thus would only increase the tension between theory and observations reported in \citetalias{StoneMetzger2016}. The last factor, however, may act in the opposite sense, if the initial distribution of stars was tangentially-biased, having a deficit of stars with low angular momentum. Such a gap could have arised if the galaxy previously contained a binary SBH, which ejected stars from the galactic core through the slingshot process on the way to its merger \citep{MilosMerritt2001}. In \citetalias{LezhninVasiliev2015}, we explored the effect of gap in the angular momentum distribution on the present-day TDE rates, solving the time-dependent Fokker--Planck equation in spherical geometry. We found that the gap is refilled after $\sim 10^{-2}\,T_\mathrm{rel}$, where $T_\mathrm{rel}$ is the local relaxation time measured at the radius of influence $r_\mathrm{infl}$ (the latter defined as the radius enclosing the mass of stars equal to twice the SBH mass $M_\bullet$). We used \citet{Dehnen1993} models with various values of power-law index $\gamma$ of the density profile and scaled them to real galaxies using the $M_\bullet-\sigma$ relation \citep{FerrareseMerritt2000,Gebhardt2000}, obtaining a one-parameter family of models for each $\gamma$. Using this scaling, the gap-refill time is longer than the Hubble time for $M_\bullet \gtrsim 10^7\,M_\odot$, thus for these galactic nuclei we might expect a reduction of TDE rates compared to the SSS value. However, the same galaxy merger that creates a binary SBH necessary for the formation of the gap, also leads to a significantly non-spherical shape of the merger remnant. Thus it is important to consider both the shape and velocity anisotropy together in order to obtain a more reliable estimate of TDE rates. The Fokker--Planck method is unsuitable for this task, as the existing implementations are restricted to spherical or at most axisymmetric geometry under certain simplifying assumptions: either that the distibution function depends only on the two classical integrals of motion (the energy $E$ and the conserved component of angular momentum $J_z$), as in \citet{Goodman1983,FiestasSpurzem2010}, or additionally on a third integral that exists only within the sphere of influence, as in \citet{VasilievMerritt2013}. On the other hand, relaxation processes around a SBH can also be studied with the Monte Carlo method, which has been used in spherical geometry by \citet{DuncanShapiro1983,FreitagBenz2002}, but can be easily extended to non-spherical systems \citep{Vasiliev2014}. The paper is organized as follows. In Section~\ref{sec:theory} we review the basics of the loss-cone theory and the methods used to compute the TDE rates, and in Section~\ref{sec:ic} we describe the initial conditions for our models. In Section~\ref{sec:fp_sph} we apply the spherical Fokker--Planck formalism of \citet{LezhninVasiliev2015} to a sample of galaxies from \citet{Lauer2007}, estimating the TDE rates in the presence of the angular momentum gap for a diverse collection of density profiles. Section~\ref{sec:mc} presents the main suite of our Monte Carlo simulations: we demonstrate that they agree with the Fokker--Planck approach in the case of spherical galaxies, and then consider non-spherical models with and without a gap in the initial distribution. We find that the enhancement of TDE rate due to non-sphericity outweighs the suppression due to the gap, which itself is much less pronounced in non-spherical cases. Section~\ref{sec:results} summarizes our results. | \label{sec:results} We considered the rate of disruption of stars by massive black holes in galactic centers, with a particular focus on the influence of anisotropic initial conditions resulting from a previously existing binary SBH, and on the role of geometry of the galactic nucleus (spherical, axisymmetric or triaxial). In doing so, we generalized the results of \citet{Vasiliev2014}, where different geometries were considered under the assumption of isotropic initial conditions, and of \citet{LezhninVasiliev2015}, where the effect of a gap at low angular momentum in the initial distribution of stars was studied in spherical galactic models. Our main tool for this study is the Monte Carlo method for dynamical evolution of stellar systems with arbitrary geometry, introduced in \citet{Vasiliev2015}. The initial conditions for the Monte Carlo simulations are taken from models of galactic nuclei with binary SBHs, in which the slingshot process ejects stars with low angular momenta. In order to isolate the effect of angular-momentum gap from the role of geometry, we constructed isotropized models with the same density profile and shape, but having a uniform distribution of stars in angular momentum. The results can be summarized as follows. \begin{itemize}[leftmargin=8mm] \setlength\itemsep{0pt} \item For spherical systems, we compared the evolution of TDE rates computed by Fokker--Planck and Monte Carlo methods, and found a satisfactory agreement, confirming our earlier results obtained with the former approach. \item Using the Fokker--Planck approach, we computed the expected TDE rates for a sample of galaxies from \citet{Lauer2007} in the presence of a gap in the initial distribution in angular momentum, and compared them with the steady-state estimates from \citetalias{StoneMetzger2016}. We find that the suppression of TDE rates is important for galaxies with $M_\bullet \gtrsim 10^{7.5}\,M_\odot$, again confirming our earlier findings with a more diverse collection of galaxy models. \item The depletion of stars with low angular momentum due to slingshot ejection during the binary SBH evolution is less pronounced in non-spherical systems. \item The difference between models with isotropic and anisotropic initial conditions and with different geometry starts to play role for $M_\bullet \gtrsim 10^{7.5}\,M_\odot$. Moreover, for non-spherical models, the influence of a gap is less dramatic than for spherical models, and in fact the TDE rate stays well above the SSS estimate (\mbox{$\sim 2\times 10^{-5}\,M_\odot$/yr} under our scaling) for all of them, regardless of the SBH mass and the initial conditions. \item Thus we conclude that the effect of tangential anisotropy could be only a minor factor in resolving the discrepancy between theoretical predictions and observationally inferred TDE rates, raised by \citetalias{StoneMetzger2016}, since the majority of TDEs are expected to occur in galaxies with lower SBH masses than our threshold value. The explanation must therefore lie elsewhere, for instance, in the properties of optical emission \citep[e.g.,][]{MetzgerStone2016}. \end{itemize} A number of caveats should be mentioned. Our models were scaled to physical units using Equation~\ref{eq:rinfl}, i.e., all properties are fixed by the SBH mass, while for real galaxies this is certainly not the case. Moreover, in the range of $M_\bullet$ considered in this paper, SSS estimates of TDE rates for the sample of galaxies in \citetalias{StoneMetzger2016} are up to an order of magnitude higher than our values $\sim 2\times10^{-5}\,M_\odot$/yr, which indicates that our adopted scaling (taken from a series of previous papers) underestimates the density of real galaxies in the low-mass end. Thus the TDE rates shown in Figure~\ref{fig:summary} should be interpreted as lower boundaries on the range of values possible in real galaxies, and a typical galaxy would have a TDE rate a factor of few higher. As illustrated by \citet{StoneVelzen2016}, some E+A galaxies may have TDE rates still an order of magnitude higher than average, owing to their high central density. This also means that the effects of non-sphericity and angular-momentum gap would be noticeable starting from a somewhat higher $M_\bullet$ than in our simulations, since a denser galaxy has a shorter relaxation time and thus is less affected by these factors. Our initial conditions are taken from simulations of merging binary SBHs conducted in \citet{VasilievAM2015}. The main suite of simulations were set up not in a cosmological context, but rather embedded an equal-mass binary SBH into a steady-state galaxy with a prescribed shape. This is of course not very realistic, but allows a more precise control on the role of geometry in both the evolution of the binary SBH, and the subsequent loss-cone dynamics around a single SBH. The models considered here and in that paper are only moderately non-spherical, with minor-to-major axis ratio $\gtrsim 0.75$ and intermediate-to-major axis $\gtrsim 0.9$. We also considered a series of `merger' models using initial conditions from a simulation that followed a merger of two galaxies followed by formation and merger of SBHs. The shape of this model was close to but not precisely axisymmetric. Nevertheless, even this small deviation brings qualitative changes to the structure of orbits around the SBH, and the TDE rates for these models were very similar to the triaxial ones. Our models also have a very shallow density profile in the center, resulting from the destruction of the pre-existing cusp during the merger. Galaxies with a steeper density profile ($\gamma \gtrsim 1$) probably either have not experienced a major merger in their history, or have re-grown the cusp due to a short relaxation time or due to dissipative processes (inflow of fresh gas and star formation). In all these cases we don't expect the angular-momentum gap to exist. To summarize, simple estimates of TDE rates under the assumption of spherical geometry, steady state and isotropic initial conditions, commonly used in the literature, should be regarded as a lower bound for the TDE rates expected in more complex systems resulting from galaxy mergers, because the effect of angular-momentum gap is more than compensated by a non-spherical shape of the merger remnant. We remind that the software for computing the steady-state and time-dependent TDE rates in spherical geometry for an arbitrary density profile, using the Fokker--Planck approach, is available at \url{http://td.lpi.ru/~eugvas/losscone}. | 16 | 9 | 1609.00009 |
1609 | 1609.05638_arXiv.txt | {} {In 2015, we initiated a survey of Scorpius-Centaurus A-F stars that are predicted to host warm-inner and cold-outer belts of debris similar to the case of the system HR~8799. The survey aims to resolve the disks and detect planets responsible for the disk morphology. In this paper, we study the F-type star HIP~67497 and present a first-order modelisation of the disk in order to derive its main properties.} {We used the near-infrared integral field spectrograph (IFS) and dual-band imager IRDIS of VLT/SPHERE to obtain angular-differential imaging observations of the circumstellar environnement of HIP~67497. We removed the stellar halo with PCA and TLOCI algorithms. We modeled the disk emission with the \texttt{GRaTeR} code.} {We resolve a ring-like structure that extends up to $\sim$450 mas ($\sim$50 au) from the star in the IRDIS and IFS data. It is best reproduced by models of a non-eccentric ring with an inclination of $80\pm1^{\circ}$, a position angle of $-93\pm1^{\circ}$, and a semi-major axis of $59\pm3$ au. We also detect an additional, but fainter, arc-like structure with a larger extension (0.65 arcsec) South of the ring that we model as a second belt of debris at $\sim$130 au. We detect 10 candidate companions at separations $\geq$1". We estimate the mass of putative perturbers responsible for the disk morphology and compare it to our detection limits. Additional data are needed to find those perturbers, and to relate our images to large-scale structures seen with HST/STIS.} {} | The Scorpius-Centaurus OB association (Sco-Cen) is the nearest site of recent massive star formation \citep{1999AJ....117..354D}. The proximity (d=90-200 pc) and young age ($\sim$11-17 Myr) of Sco-Cen all contribute to make it an excellent niche for direct imaging (DI) search of planets and disks. The planet imager instruments GPI \citep{2008SPIE.7015E..18M} and SPHERE \citep{2008SPIE.7014E..18B} have initiated surveys of the circumstellar environment of a few Sco-Cen stars at unprecedented contrasts and angular resolution. They have already resolved several new debris disks around those stars \citep[]{2015ApJ...807L...7C, 2015ApJ...812L..33K, 2015ApJ...814...32K, 2016A&A...586L...8L, 2016arXiv160502771D} with wing-tilt and ringed morphologies indicative of the presence of planets. % We initiated a DI survey with SPHERE to image new giant planets and circumstellar disks around a sample of Sco-Cen F5-A0 stars with infrared excess. The excesses can be modeled by 2 black-body components, each corresponding to a belt of debris \citep[][hereafter C14]{2014ApJS..211...25C}. This architecture is reminiscent of the Solar System, and of benchmark systems previously identified by DI like HR~8799 \citep{2008Sci...322.1348M, 2010Natur.468.1080M}, or HD~95086 \citep{2013ApJ...779L..26R} which is also part of Sco-Cen.\\ In the course of the survey, we resolved a disk around the F0-type star HIP~67497 (HD~120326). This intermediate-mass star (M$=1.6 M_{\odot}$) is located at a distance of $107.4^{+10.4}_{- 8.7}$ pc \citep{2007A&A...474..653V} and belongs to the 16 Myr old \citep{2002AJ....124.1670M} Upper Centaurus Lupus sub-group \citep{1999AJ....117..354D}. C14 modeled the infrared excess of the star with two belts of debris: a cold ($127\pm5$ K) bright ($L_{IR}/L_{*}=1.1\times10^{-3}$) belt at 13.9 au and a second colder ($63\pm5$ K) dimmer ($L_{IR}/L_{*}=1.4\times 10^{-4}$) belt at 116.5 au. The same team produced an alternative model of the excess requiring only one belt located at 8.82$\pm$1 au \citep{2015ApJ...808..167J}. Previous observations with adaptive optics systems did not reveal any companion or structure close to HIP~67497 \citep{2013ApJ...773..170J}. Our resolved observations therefore offer the opportunity to better constrain the radial distribution of the dust around this star and to look for the companions responsible for that distribution. We present the observations and data in \S~\ref{section:data}, and the models in \S~\ref{section:diskprop}. We discuss the morphology of the disk and the existence of putative perturbers (planets) in \S~\ref{section:conclusion}. | \label{section:conclusion} The disk around HIP~67497 has an infrared luminosity (IL) in the same range as the disks already resolved around the Sco-Cen A-F type stars HD 111520, HIP 79977, HD 106906, and HD 115600 ($L_{IR}/L_{*}\sim10^{-3}-10^{-4}$). This could explain why we were able to resolve it. For HD~106906, and HD~115600, C14 finds outer belts at shorter separations than observed. In the case of HIP~67497, the second belt corresponds roughly to the location of the coldest belt found by C14, but the flux ratio between the two belt models found in \S~\ref{sec_outer_ring} is of the same order as the ratio of the IL of the belts found by C14 (7.85). HST/STIS images of HIP~67497 show extended emission at large scales \citep{2016IAUS..314..175P}, which we compare to the SPHERE images in Fig. \ref{fig:HST}. The inner structures ($\leq$3'') seen with HST have an orientation compatible with the belt and the arc of the SPHERE images. Another asymmetric feature is seen on the east side up to 8 arcsec, but can not be easily related to the rest. Different asymmetries at different scales have already been noticed for HD106906 \citep{2015ApJ...814...32K} or HD 32297 \citep{2014ApJ...780...25E, 2014AJ....148...59S} for instance. One candidate companions (\#1, see App. \ref{section:companions}) could lie within the structures revealed by STIS and the arc seen with SPHERE. The present STIS data are unfortunately affected by blind zones caused by the position of the coronagraphic bars of STIS (only two roll angles). Several options exist to explain the morphology of the disk of HIP~67497. The observed ring-like structures may be caused by dust-gas interactions \citep{2013Natur.499..184L}. Unfortunately, the gas content of the disk remains unknown. Alternately, planets with different individual eccentricities and semi-major axes \citep{2016ApJ...827..125L} may also provide an explanation for the double-ring structure. Our observations are sensitive to 1.5 to 15 M$_{Jup}$ in-between the ring and the arc when accounting for the disk inclination (Fig. \ref{fig:dl}). We explore in App. \ref{AppC:perturbers} the case of one, two, or three perturbers on circular orbit, or one and two planets on eccentric orbits using numerical simulations. The predicted masses can reach $\sim$21 M$_{Jup}$ for the case of a single planet. But the mass estimate of the perturber(s) is sensitive to the eccentricity of the orbit(s). New observations with STIS and SPHERE are required to reveal the full morphology of the disk, improve the detection performances, and to clarify the nature of the CCs. \begin{figure} \begin{center} \includegraphics[width=\columnwidth]{Figures/Final/Fig_HST.jpg} \caption{Sketch showing the HST/STIS and the SPHERE images of the debris disk around HIP~67497.} \label{fig:HST} \end{center} \end{figure} | 16 | 9 | 1609.05638 |
1609 | 1609.00537_arXiv.txt | The Lockman Hole is a well-studied extragalactic field with extensive multi-band ancillary data covering a wide range in frequency, essential for characterising the physical and evolutionary properties of the various source populations detected in deep radio fields (mainly star-forming galaxies and AGNs). In this paper we present new 150-MHz observations carried out with the LOw Frequency ARray (LOFAR), allowing us to explore a new spectral window for the faint radio source population. This 150-MHz image covers an area of 34.7 square degrees with a resolution of 18.6$\times$14.7\,arcsec and reaches an rms of 160\,$\mu$Jy\,beam$^{-1}$ at the centre of the field. As expected for a low-frequency selected sample, the vast majority of sources exhibit steep spectra, with a median spectral index of $\alpha_{150}^{1400}=-0.78\pm0.015$. The median spectral index becomes slightly flatter (increasing from $\alpha_{150}^{1400}=-0.84$ to $\alpha_{150}^{1400}=-0.75$) with decreasing flux density down to $S_{150} \sim$10\,mJy before flattening out and remaining constant below this flux level. For a bright subset of the 150-MHz selected sample we can trace the spectral properties down to lower frequencies using 60-MHz LOFAR observations, finding tentative evidence for sources to become flatter in spectrum between 60 and 150\,MHz. Using the deep, multi-frequency data available in the Lockman Hole, we identify a sample of 100 Ultra-steep spectrum (USS) sources and 13 peaked spectrum sources. We estimate that up to 21\,per\,cent of these could have $z>4$ and are candidate high-$z$ radio galaxies, but further follow-up observations are required to confirm the physical nature of these objects. | \label{intro} Although the majority of the earliest radio surveys were carried out at very low frequencies (i.e. the 3rd, 6th and 7th Cambridge Surveys; \citealt{3c,3cr,6c,Hales2007} and the Mills, Slee and Hill (MSH) survey; \citealt{msh}), in more recent years large area radio surveys have primarily been performed at frequencies around 1\,GHz such as the NRAO VLA Sky Survey (NVSS; \citealt{nvss}), the Sydney University Molonglo Sky Survey (SUMSS: \citealt{sumss} and the Faint Images of the Radio Sky at Twenty Centimeters (FIRST) survey; \citealt{first}). With the advent of radio interferometers using aperture arrays such as the Low-Frequency ARray (LOFAR), we now have the ability to re-visit the low-frequency radio sky and survey large areas down to much fainter flux density levels, and at higher resolution, than these earlier surveys. LOFAR is a low-frequency radio interferometer based primarily in the Netherlands with stations spread across Europe \citep{lofar}. It consists of two different types of antenna which operate in two frequency bands: the Low-Band-Antennas (LBA) are formed from dipole arrays which operate from 10 to 90\,MHz, and the High-Band Antennas (HBA) are tile aperture arrays which can observe in the frequency range 110--240\,MHz. The long baselines of LOFAR allow us to probe this frequency regime at much higher spatial resolution than previously done; up to 5\,arcsec resolution for the longest dutch baseline, and up to 0.5\,arcsec resolution for the European baselines. The combination of LOFAR's large field-of-view, long baselines and large fractional bandwidth make it an ideal instrument for carrying out large surveys. The majority of sources detected in these surveys to date have been radio-loud AGN, with star-forming galaxies only beginning to come into the sample at lower flux densities ($S<5-10$\,mJy). Obtaining large samples of these objects allows us to study the source population in a statistically significant manner and investigate how the properties of radio galaxies evolve over cosmic time. In addition, large surveys allow us to search for rare, unusual objects in a systematic way. However, in order to maximise the scientific value of these large surveys, complementary multi-wavelength data are essential to obtain a comprehensive view of the source populations. One such field with extensive multi-band coverage is the Lockman Hole field. This field was first identified by \citet{lockman} who noted the region had a very low column density of Galactic H{\sc i}. This smaller amount of foreground H{\sc i} makes it an ideal field for deep observations of extragalactic sources, particularly in the IR due to the low IR background \citep{Lonsdale2003}. Because of this, there is extensive multi-band ancillary data available, including deep optical/NIR data from ground based telescopes (e.g. \citealt{Fotopoulou2012}), midIR/FIR/sub-mm data from the {\it Spitzer} and {\it Herschel} satellites \citep{Mauduit2012a,Oliver2012} and deep X-ray observations from {\it XMM-Newton} and {\it Chandra} \citep{Brunner2007,Polletta2006}. In addition, the Lockman Hole field has an extensive amount of radio data covering a wide range in frequency. This includes the 15-GHz 10C survey \citep{10c, Whittam2013}, deep 1.4-GHz observations over 7 square degrees observed with the Westerbork Synthesis Radio Telescope (WSRT; \citealt{Guglielmino2012}, Prandoni et al., 2016a, in preparation), 610-MHz GMRT observations \citep{Garn2008} and 345-MHz WSRT observations (Prandoni et al., 2016b, in preparation). In this paper we present LOFAR observations of the Lockman Hole field, which extends this multi-frequency information down to 150\,MHz, allowing us to study the low-frequency spectral properties of the faint radio source population. For a brighter sub-sample we were also able to perform a preliminary analysis of the spectral properties down to 60\,MHz. Studying the spectral index properties of the radio source population can provide insight into a range of source properties. For example, the radio spectral index is often used to distinguish between source components in AGN (i.e. flat-spectrum cores vs. steep-spectrum lobes or ultra-steep spectrum relic emission). Spectral information can also be used to derive approximate ages of the radio source based on spectral ageing models (see e.g. \citealt{Harwood2013, Harwood2015} and references therein), providing insight into the average life-cycle of radio-loud AGN. Previous studies that have looked at the average spectral index properties of large samples have reported conflicting results as to whether the median spectral index changes as a function of flux density \citep{Prandoni2006, Ibar2010, Randall2012, Whittam2013}. These studies have generally been carried out at GHz frequencies, with studies of low-frequency selected sources typically showing evidence for the median spectral index to become flatter at fainter flux densities \citep{Ishwara-Chandra2010, Intema2011, Williams2013}. However, most of these latter studies have been biased against detecting steep-spectrum sources at the fainter end of the flux density distribution due to the flux limits imposed by the different surveys used. The wide frequency coverage available in the Lockman Hole, along with the large area surveyed, also allows us to search for sources with more atypical spectral properties such as those with ultra-steep or peaked spectra. In this paper we present 150-MHz LOFAR observations of the Lockman Hole field. In Section 2 we discuss the observational parameters, data reduction and source extraction of the 150-MHz LOFAR observations, followed by a brief overview of additional 60-MHz LOFAR observations that are used for the spectral analysis. In Section 3 we present an analysis of the source sizes and resolution bias which are used to derive the 150-MHz source counts and in Section 4 we investigate the spectral index properties of low-frequency selected radio sources. Section 5 presents a deeper look at sources that exhibit more unusual spectral properties (e.g. ultra steep spectrum or peaked spectrum sources), providing insight into how many of these sources we might expect to find in the completed LOFAR all-sky survey. We conclude in Section 6. Throughout this paper we use the convention $S_{\nu} \propto \nu^{\alpha}$. | We present 150-MHz LOFAR observations of the Lockman Hole field, allowing us to study the spectral index properties of low-frequency radio sources. The field was observed for 10~hrs using the HBA array covering the frequency range 110--172~MHz. The final 150~MHz image covers an area of 34.7 square degrees, has a resolution of 18.6$\times$14.7\,arcsec and reaches an rms of 160\,$\mu$Jy\,beam$^{-1}$ at the centre of the field. This results in a catalogue of 4882 sources within 3 degrees from the pointing centre. From these data we derive the 150-MHz source counts, which are in good agreement with other observations at 150\,MHz (e.g. \citealt{Williams2016}) as well as with extrapolations from observations of the Lockman Hole at other frequencies (Prandoni et al., 2016a,b, in preparation). By crossmatching this LOFAR 150-MHz catalogue with the NVSS, WENSS and VLSSr all-sky surveys, we form a sample with which to study the spectral index properties of low-frequency radio sources. Due to the different flux limits and resolutions of each of these surveys, we only include sources with deconvolved sizes less than 40\,arcsec and with $S_{150}>40$~mJy. This results in a sample of 385 LOFAR sources which we term the Lockman--wide sample. As expected for low-frequency selected sources, the sample is dominated by steep spectrum sources, with a median spectral index of $\alpha_{150}^{1400}=-0.82\pm0.02$ and an interquartile range of [$-0.94$, $-0.70$]. Although we reach sub-mJy flux densities in the LOFAR 150-MHz observations, the spectral analysis is severely limited by the much higher flux limits of other surveys that cover the entire LOFAR field of view (i.e. NVSS, WENSS, VLSSr). To investigate the spectral index properties of mJy sources selected at 150~MHz, we crossmatch the LOFAR catalogue with deeper 1.4~GHz observations from WSRT (Prandoni et al., 2016a, in preparation). This deep WSRT mosaic covers 7 square degrees with resolution 11$\times$9\,arcsec and reaches an rms of 11$\mu$Jy. This results in a sample of 1302 matched sources which we refer to as the Lockman--WSRT catalogue. Investigating the spectral index properties of this larger sample, again we find that the sample is dominated by steep-spectrum sources (82.1\,per\,cent), with 11.3\,per\,cent classified as flat-spectrum and 6.6\,per\,cent classified as USS. The median spectral index of this much larger sample is $\alpha_{150}^{1400}=-0.78\pm0.015$ with an interquartile range of [$-0.95$, $-0.65$]. The median spectral indices become slightly flatter with decreasing flux density, increasing from $\alpha_{150}^{1400}=-0.84$ for sources above 50\,mJy to $\alpha_{150}^{1400}=-0.76$ at 5--10\,mJy. At fainter flux densities the median spectral index stays relatively constant. Crossmatching this Lockman--WSRT catalogue with other deep, multi-frequency radio surveys available in the Lockman Hole we form the Lockman--deep sample with which to study the spectral properties of low-frequency radio sources above a flux density threshold of $S_{150}=9$\,mJy. This again shows that the majority of sources exhibit straight, power-law spectra over the full frequency range from 150\,MHz to 15\,GHz. Preliminary 60-MHz LOFAR observations of the Lockman Hole field shows tentative evidence that a large fraction of sources begin to flatten between 60\,MHz and 150\,MHz, but a more comprehensive analysis of the absolute flux calibration is needed to confirm this result. Using this wide-frequency coverage we also search for sources with more extreme spectral properties such as Ultra-Steep Spectrum (USS) sources and peaked spectrum sources. We identify a total of 100 USS sources and 13 candidate peaked spectrum sources in the Lockman Hole. Using the $K$-$z$ relation we estimate that approximately 42\,per\,cent of these could be at $z>3$ and up to 21\,per\,cent at $z>4$ indicating that these could be candidate high-$z$ galaxies. However, further follow-up is required to confirm the physical properties of these objects. The depth reached in these observations closely matches that expected for the `Tier-1' all-sky LOFAR survey at 150~MHz. As such, the results obtained here provide insight into what we can expect to detect in the completed all-sky survey. However, this study also highlights the need for complementary, sensitive datasets across a wide range in frequency to maximise the scientific return of low-frequency radio surveys. | 16 | 9 | 1609.00537 |
1609 | 1609.05254_arXiv.txt | We present the first results from an ongoing survey for multiplicity among the bright stars using the Navy Precision Optical Interferometer (NPOI). We first present a summary of NPOI observations of known multiple systems, including the first detection of the companion of $\beta$ Scuti with precise relative astrometry, to illustrate the instrument's detection sensitivity for binaries at magnitude differences $\Delta$$m$ $\lessapprox$ 3 over the range of angular separation 3 - 860 milliarcseconds (mas). A limiting $\Delta$$m_{700}$ $\sim$ 3.5 is likely for binaries where the component spectral types differ by less than two. Model fits to these data show good agreement with published orbits, and we additionally present a new orbit solution for one of these stars, $\sigma$ Her. We then discuss early results of the survey of bright stars at $\delta$ $\geq$ -20$\arcdeg$. This survey, which complements previous surveys of the bright stars by speckle interferometry, initially emphasizes bright stars of spectral types F0 through K2. We report observations of 41 stars of apparent visual magnitude $m_V$ $\leq$ 4.30, all having been observed on multiple nights. Analysis of these data produces fitted angular separations, position angles, and component magnitude differences for six previously known visual binaries. Three additional systems were examined as possible binaries, but no conclusive detection could be made. No evidence of close stellar companions within our detection limit of $\Delta$$m$ $\approx$ 3 was found for the remaining 32 stars observed; however, uniform-disk angular diameters are reported for 11 of the resolved stars in this last group. | \label{intro} Knowledge of the frequency of multiplicity among stars is fundamental to furthering our understanding of many areas of astrophysics and has direct impact on the design of future experiments to detect and image extrasolar planets. However, multiplicity surveys of even the brightest stars using modern techniques were surprisingly incomplete until the 1990s \citep{har00}. Early surveys of the bright stars by speckle interferometry \citep{mca87,mca93}, although themselves incomplete, revealed that the frequency of visual binaries in the range of angular separations accessible to that technique was several times that previously known, including substantial numbers of wider binaries missed by classical visual observers. These imaging surveys, combined with radial velocity observations and other techniques, contributed to bright star multiplicity catalogs. \citet{etv08} present a catalog of bright, multiple (two or more components) star systems and discuss the implications of multiplicity upon several topics in astrophysics. Their catalog is derived from a number of input catalogs and observational techniques and consists of more than 4,500 stars. A catalog consisting of bright systems with three or more stars, the $\mathit{Multiple\ Star\ Catalog}$ \citep[MSC,][]{tok97}, also draws its sources from different observational techniques. The MSC is available online \footnote{Catalog J/A+AS/124/75 in the VizieR catalog access tool \citep{och00}.} and was updated in 2010. As noted by \citet{rag10}, continued multiplicity survey efforts using speckle interferometry \citep[e.g.,][]{mas98,hor08} have now resulted in nearly complete coverage, at least for bright, solar-type stars, down to the diffraction limit of large telescopes ($\sim$ 30 mas). These surveys, along with systematic, higher-precision radial velocity observations have also largely closed the historical gap in orbital period coverage between these techniques \citep{rag12}. \citet{mca92} points out, however, that only long-baseline optical interferometry, using multiple-telescope arrays with baseline lengths of hundreds of meters, offers a single technique for multiplicity detection throughout the angular-separation/period range from classical visual doubles to interacting binaries \citep{hrs09,zav10} and contact binaries. Optical interferometry not only provides the data products of speckle interferometry with improved precision, but at the narrower angular separations not accessible to speckle where most spectroscopic binaries reside, offers sensitivity for binary detection in a range of orbital inclinations complementary to spectroscopy. The combination of visual orbits from interferometry with spectroscopic orbits for SB2 systems can yield stellar masses and orbital parallax and, if one or more of the components are resolved, stellar angular and potentially linear diameters \citep{hum94}. High precision mass determinations are potentially possible for SB1 systems as well should GAIA parallaxes become available for such bright stars. Improving our knowledge of stellar multiplicity for both physical and ${\it optical}$ systems of small angular separation also has important implications for precision navigation, where the presence of stellar companions and their relative motions can affect the determination of the ``center of light'' by relatively low-resolution star trackers \citep{har00}. We report here the first results of what is anticipated to be an ongoing survey of the brighter stars using the Navy Precision Optical Interferometer (NPOI). We first describe the capabilities of the NPOI for such a survey in \S~\ref{npoi}. Based on these capabilities (${\it c.}$ 2004), we next discuss the selection of the targets for the initial survey in \S~\ref{targsel}. The standard observing procedures and data reduction are described in \S~\ref{seqred}, including the selection and observation of calibration stars, and use of the resulting data to produce the accurately calibrated fringe visibility data for the program stars upon which all subsequent source modeling depends. The remainder of \S~\ref{obs}, followed by \S~\ref{modeling}, present the results of models in three areas; first, the systematic examination of the calibrated data for evidence of binary systems among the program stars (\S~\ref{bindet}) and subsequent detailed astrometric and photometric modeling of the detected systems (\S~\ref{binBSS},~\ref{mardet}); second, the observation and modeling of previously known binary systems (\S~\ref{obslog},~\ref{binknown}), plus detected binaries among the program stars (\S~\ref{bindet},~\ref{binBSS}), to determine the maximum detected magnitude difference ($\Delta$$m$) of binary star pairs in our survey as a function of their angular separation (\S~\ref{limbindet}); and, third, the subsidiary result of accurate angular diameters for the resolved, single stars among the program sample (\S~\ref{diamBSS}). Plans for future stages of this survey are also discussed (\S~\ref{resdisc}). | \label{conclus} We have presented the first results from an ongoing survey for multiplicity among the bright stars using the NPOI. The initial source sample emphasized bright stars of spectral types F0 through K2. We report observations of 41 stars ($m_V$ $\leq$ 4.30), all having been observed on multiple nights. Observations of known binary systems among the program star sample, combined with additional observations of other known binaries, including the first ``visual'' detection of the secondary star in the $\beta$ Sct system at precisely measured separations and position angles, demonstrate the NPOI's sensitivity for binary detection over a wide range of angular separations (3 - 860 mas) at component magnitude differences $\Delta$$m_{700}$ $\lessapprox$ 3, and to perhaps $\Delta$$m_{700}$ $\sim$ 3.5 for binaries where the component spectral types differ by less than two. Fitted angular separations, position angles, and component magnitude differences for six previously known binaries from the program sample are presented, along with angular diameters for 11 resolved stars from the sample that have no detected stellar companions. In the light of the significant improvements made to the limiting sensitivity of NPOI observations ($m_V$ $\approx$ 6.0) in recent years, plans are being drafted to extend the initial survey to a complete, volume-limited sample of stars of spectral types F0 through G5. | 16 | 9 | 1609.05254 |
1609 | 1609.03969_arXiv.txt | We investigate the fossil magnetic field in the accretion and protoplanetary discs using the Shakura and Sunyaev approach. The distinguishing feature of this study is the accurate solution of the ionization balance equations and the induction equation with Ohmic diffusion, magnetic ambipolar diffusion, buoyancy and the Hall effect. We consider the ionization by cosmic rays, X-rays and radionuclides, radiative recombinations, recombinations onto dust grains, and also thermal ionization. The buoyancy appears as the additional mechanism of magnetic flux escape in the steady-state solution of the induction equation. Calculations show that Ohmic diffusion and magnetic ambipolar diffusion constraint the generation of the magnetic field inside the `dead' zones. The magnetic field in these regions is quasi-vertical. The buoyancy constraints the toroidal magnetic field strength close to the disc inner edge. As a result, the toroidal and vertical magnetic fields become comparable. The Hall effect is important in the regions close to the borders of the `dead' zones because electrons are magnetized there. The magnetic field in these regions is quasi-radial. We calculate the magnetic field strength and geometry for the discs with accretion rates $(10^{-8}-10^{-6})\,\rm{M}_{\odot}\,\rm{yr}^{-1}$. The fossil magnetic field geometry does not change significantly during the disc evolution while the accretion rate decreases. We construct the synthetic maps of dust emission polarized due to the dust grain alignment by the magnetic field. In the polarization maps, the `dead' zones appear as the regions with the reduced values of polarization degree in comparison to those in the adjacent regions. | Contemporary star formation takes place in the magnetic rotating cores of molecular clouds which we call protostellar clouds. The collapse of protostellar clouds leads to the formation of the protostars with protostellar discs. The discs form under the action of the electromagnetic and centrifugal forces. The protostellar discs are geometrically thick self-gravitating discs. After the formation of a young star, the protostellar disc transforms into the accretion disc, where matter accretes onto the star with the accretion rate $\dot{M}=(10^{-6}-10^{-8})\,\rm{M}_{\odot}\,\rm{yr}^{-1}$ in the case of T~Tauri stars. The accretion discs of the T~Tauri stars are geometrically thin structures with masses $\sim(0.001-0.1)\,\rm{M}_{\odot}$ and sizes $(100-1000)$~au (see review of \citet{williams11}). During evolution, the accretion discs transform to the protoplanetary discs similar to the protosolar nebula. There are some observational data concerning the magnetic fields in the accretion discs of young stellar objects. The measurements of the remnant magnetization of the meteorites indicate that the magnetic field strength in the protosolar nebula was $\sim(0.1-1)$ Gs~\citep{levy78, fu14}. \citet{donati05} reported about the measurement of the magnetic field strength of about 1~kGs in the inner regions of the FU~Ori system using the Zeeman splitting technique. \citet{girart06} made the measurements of the polarized dust emission from the low-mass protostellar system NGC~1333~IRAS~4A and found that the geometry of the magnetic field is `hour-glass' in this system. \cite{Rao_etal2014} detected linearly polarized 878 $\mu\rm{m}$ dust emission in the circumstellar disc around the IRAS~16293-2422~B protostar from the Submillimetre Array observations. The measurements indicate that the geometry of the magnetic field is complex in this system. \citet{stephens14} performed resolved measurements of the polarized 1.25~mm continuum emission from the HL~Tau disc using the Combined Array for Millimetre-wave Astronomy (\textsc{CARMA}). They suggested that the geometry of the magnetic field in the disc is complex with both toroidal and poloidal components. \citet{scox15} performed observations of the dust polarization from the circumstellar disc around the Class 0 protostar L1527 using the \textsc{CARMA}. They concluded that the geometry of the magnetic field is toroidal. Interpretation of (sub)millimetre polarization observations of the discs around young stellar objects can be complicated. The polarization of the disc dust emission is usually interpreted as a result of the alignment of non-spherical dust grains with magnetic field. But as it was demonstrated by \citet{Yang_etal2016} the (sub)mm continumm polarization could also be a result of scattering and disc inclination. Direct measurements of the magnetic field strength based on the Zeeman effect are yet not possible and the polarization observations are not able to detect the large-scale magnetic field geometry in details. Therefore, theoretical investigations of the magnetic field in young stellar objects are needed to explain the observations and to predict what magnetic field will be detected in future high-resolution measurements. There are two main conceptions of the origin of the magnetic field in young stellar objects: the dynamo theory and the theory of the fossil magnetic field. The latter is based on the numerical simulations and observations of the star formation in molecular cloud cores (see recent review by \citet{fmft}). One-and-half-dimensional magneto-gas-dynamic simulations of \citet{dudorov87} have shown that the initial magnetic flux of protostellar clouds is conserved partially during collapse and accretion. These conclusions have been confirmed by a number of multi-dimensional simulations (see review by \citet{i12}). Therefore, from the point of view of the discs formation theory it is natural to assume that the magnetic field of stars and their accretion discs is the {\it fossil} one, i.e. it is the remnant of the protostellar cloud's magnetic field. The strength and geometry of the magnetic field in the accretion discs is influenced by accretion, differential rotation, magnetic diffusion and different instabilities. Ohmic diffusion (OD) and magnetic ambipolar diffusion (MAD) are the main processes that constraint the generation of the magnetic field during the star and disc formation. OD of the magnetic field is caused by currents dissipation in gas with finite conductivity (see \citet{parker_book}). MAD is the process of plasma drift through the neutral gas under the action of the electromagnetic force. \citet{mestel56} first pointed that MAD will allow a protostellar cloud to contract across the field lines if the cloud is dense enough and the frictional coupling between plasma and neutral gas is small. The magnetic diffusion efficiency depends on the ionization fraction. \citet{hayashi81} estimated the conductivity of the protosolar nebula with the magnetic field taking into account ionization by cosmic rays and decay of radioactive elements. Hayashi concluded that OD prevents generation of the strong magnetic field in the regions of the terrestrial planet formation. \citet{lubow94} and \citet{agapitou96} have shown that significant inward dragging of the large-scale magnetic field lines occurs if the magnetic diffusivity is much less than the turbulent viscosity in the disc. \citet{rr96} pointed out that toroidal magnetic field can be generated in a rotating disc when magnetic diffusion is weak. In these papers, OD efficiency was described with the help of non-dimensional magnetic Prandtl number determined as the relation of the magnetic diffusivity to the turbulent viscosity. \citet{gammie96} introduced layered accretion scenario for the T~Tauri accretion discs. He pointed out that the ionization fraction can be very low, $x<10^{-13}$, near the mid-plane of the disc due to the attenuation of cosmic rays. Magneto-rotational instability \citep[MRI, ][]{velikhov59, chandra60, bh91} is not developed and magneto-hydrodynamic (MHD) turbulence is weakened in such regions of low ionization fraction. He called these regions `dead' zones. The accretion takes place only in the surface layers of the discs in this model. \citet{sano00} calculated the ionization fraction of the minimum mass solar nebula \citep[MMSN, ][]{cameron73, mmsn}. The `dead' zone was determined as the region where wavelength of most unstable MRI mode exceeds the accretion disc scale height. This critical wavelength depends on the magnetic field strength and the ohmic diffusivity. \citet{sano00} investigated dependence of `dead' zone size on the plasma parameter $\beta=(100,\,1000)$ and on the dust grain size. They have shown that the `dead' zone shrinks as the dust grains radius increases and $\beta$ decreases. The `dead' zones are located at the distances from 1~au to (10-20)~au inside the disc. \citet{mohanty13} calculated the characteristics of the `dead' zones in the frame of the MMSN model. The `dead' zones were treated as the regions where MRI growth rate is less than the dissipation rate. The authors considered damping of MRI by both OD and MAD. They concluded that MAD determines the outer boundary of the `dead' zone. Using local shearing-box simulations of the MMSN discs, \citet{bai13} has shown that MRI is completely suppressed in the region from $0.5$~au to $(5-10)$~au near the mid-plane due to OD and in the surface layers due to MAD. Global MHD simulations of accretion discs dynamics were performed in the ideal MHD approach \citep{fromang06, flock11, suzuki14} and in the resistive limit \citep{dzyu10}. \citet{gressel15} carried out global simulations of the protoplanetary discs using the \textsc{nirvana} III code. The calculations were done for the part of the disc between $1$~au and $5$~au in two dimensions taking into account OD and MAD and ionization by FUV and X-rays only. Buoyancy is another mechanism that influences the magnetic field in the accretion discs. \citet{parker_book} has shown that the magnetized plasma in gravitationally stratified fluid is buoyantly unstable. The magnetic field splits into flux tubes rising from the fluid because of the buoyancy force. Various aspects of the magnetic buoyancy in the accretion discs have been investigated, such as the dynamics of slender magnetic flux tubes \citep{sakimoto89, schram93, ziegler01}, the generation of turbulence \citep{rozyczka96}, the formation of hot corona and bursts activity \citep{miller00, machida00}, magnetic dynamo \citep{tout92, johansen08}. \cite{zhilkin12} performed three-dimensional MHD simulations of flows in the contact binary systems and have shown that buoyancy can constraint the generation of the toroidal magnetic field.The influence of the buoyancy on the fossil magnetic field strength has not been investigated yet. Conductivity is anisotropic in the magnetized plasma \citep{alfven_book, cowling_book}. Electrons drift causes the Hall current directed perpendicular to the magnetic field lines. \citet{urpin91}, \citet{shalybkov97} have shown that the Hall current can generate the toroidal field from the poloidal one in the conducting medium. \citet{vainshtein00} also have found that the Hall currents can lead to the exchange of energy between different components of the magnetic field in the stratified plasma. \citet{wardle99a} calculated the conductivity tensor for the molecular gas taking into account the Hall effect. They have found that the Hall effect contribution to the conductivity is important for the densities between $10^7$ to $10^{11}\,\rm{cm}^{-3}$ in the case of Mathis-Rumpl-Nordsieck (MRN) grain-size distribution. Analytical and numerical investigations have shown that the Hall effect can increase or decrease MRI growth rate \citep{wardle99b, balbus01, sano02} and `dead' zone size \citep{wardle12} depending on the direction of the magnetic field vector ${\bf B}$ with respect to the rotation axis direction ${\bf \Omega}$. \citet{lesur14} and \citet{bai14} performed local shearing-box simulations of the protoplanetary discs dynamics taking into account OD, MAD, and the Hall effect. They analysed the role of these effects at fixed radial distances from a star. Simulations have shown the artificial generation of the strong azimuthal magnetic field in the midplane of the disc in the case $\left({\bf B}\cdot{\bf \Omega}\right)>0$. This effect is caused by the symmetry features of the shearing-box approximation and artificial limitation of the diffusivity. Thus, the role of OD and MAD has been mainly investigated in application to MRI and `dead' zones characteristics. \citet{guilet14} developed the approach proposed by \citet{lubow94}. To investigate the poloidal magnetic field dragging in the accretion discs, they used the vertically averaged values of the advection and diffusion rates which take into account the back-reaction of the mean magnetic field on the flow. The effective turbulent diffusivity was characterized by \citet{ss73} $\alpha$ parameter provided that the magnetic Prandtl number equals 1. This approach allowed to obtain more efficient dragging of the poloidal magnetic field than in previous works. Assuming balance between the advection and diffusion of the magnetic field, \citet{takeuchi14a} derived the steady-state radial distribution of the vertical magnetic field $B_z(r)\propto r^{-2}$ reflecting the magnetic flux conservation law. Using this profile, the authors obtained the upper limit on the magnetic field strength, $0.1$~Gs at $r=1$~au and $\sim1$~mGs at $10$~au. \citet{takeuchi14b} investigated the relaxation of the poloidal magnetic field to this steady state taking into account Ohmic diffusion only. Both \citet{guilet14} and \citet{takeuchi14a} did not consider the toroidal magnetic field in the disc. \citet{fmfadys} (DK14, hereafter) developed the MHD model of the stationary geometrically thin low-massive accretion discs with the fossil magnetic field. The model for the accretion disc contains \citet{ss73} equations, the induction equation with OD and MAD, the ionization balance equations taking into account thermal ionization, shock ionization by cosmic rays, X-rays and radioactive elements, radiative recombinations, recombinations on various dust grains. The dust particles are considered to be well-mixed with the gas. DK14 concluded that OD and MAD constraint the generation of the magnetic field inside the `dead' zones, and the magnetic field there is quasi-poloidal. The magnetic field can be quasi-azimuthal or quasi-radial in the outer regions depending on the ionization rates and dust characteristics. Many previous works a priori determined the magnetic field strength and/or magnetic diffusion efficiency in a variety of ways, in particular assuming constant plasma parameter $\beta$ and/or constant magnetic Prandtl number, etc. Following DK14, we calculate the magnetic field strength and geometry on the basis of the theory of the fossil magnetic field, i.e. using the initial conditions of the disc formation. Currently, 2D and 3D MHD simulations of the dynamics of protostars with accretion discs are difficult because of the limitations on the temporal and spatial resolution. Our approach is semi-analytical. We do not perform full 3D simulations, but calculate all three components of the magnetic field taking into account many physical effects, such as ionization, recombinations, magnetic diffusion, and thermal effects. In this paper, we modify our basic accretion disc model in order to take into account the magnetic buoyancy and the Hall effect. In section \ref{Sec:BaseModel}, we describe the model of the accretion disc. Details can be found in the paper DK14. The modification of the model is carried out in section \ref{Sec:Modif}. We derive the induction equation with OD, MAD, buoyancy and Hall effect in section \ref{Sec:IndEq}. Estimation of the buoyancy velocity is given in section \ref{Sec:Vb}. The results of calculation of the magnetic field are presented in section \ref{Sec:Results}. We investigate influence of the buoyancy and the Hall effect on the fossil magnetic field in sections \ref{Sec:ResBu} and \ref{Sec:ResHall}, respectively. In section \ref{Sec:Mdot}, we study the intensity and geometry of the fossil magnetic field with different mass accretion rates. To demonstrate how the disc magnetic field can affect the (sub)mm continuum polarization, we performed radiative transfer calculations. First, we calculated the dust temperature, $T_{\rm{d}}$, distribution in the disc (Section~\ref{sec:hyperion}). This distribution was then used to calculate the intensity and polarization of the (sub)mm continuum emission (Section~\ref{sec:lime}). We summarize our findings in section \ref{Sec:Summary}. | \label{Sec:Summary} We investigated the fossil magnetic field in the accretion and protoplanetary discs of young stellar objects. The accretion disc model described in DK14 includes Shakura and Sunyaev equations, the induction equation with Ohmic diffusion and magnetic ambipolar diffusion, the equations of ionization-recombination balance. We take into account ionization by cosmic rays, X-rays, radiative recombinations, recombinations onto dust grains, and thermal ionization. We investigated the influence of the magnetic field buoyancy and the Hall effect on the magnetic field in the disk, which were not taken into account in our previous study (DK14). The stationary solution of the induction equation has the form, in which the buoyancy represents the additional mechanism of the magnetic flux loss. Such a modification allows us to investigate the fossil magnetic field both in the regions of the effective generation, and in the `dead' zones. We also take into account non-linearity of magnetic ambipolar diffusion (MAD) in the modified model. Our model is the useful tool for the investigation of dynamics of accretion and protoplanetary discs. In contrast to the other investigations, where the magnetic field strength and/or geometry are a priori determined, in our model the strength and geometry of the magnetic field in the discs are calculated. The predictions of the model are in agreement with observational constraints, as discussed below. The radial and azimuthal magnetic field components are zero at $z=0$ due to the equatorial symmetry of the disc. In recent shearing-box MHD simulations, the generation of the strong azimuthal magnetic field in the midplane due to the Hall effect was found \citep{bai14, lesur14, bai16}. This result is caused by the peculiarities of the shearing-box approximation which does not allow to distinguish between the odd and even symmetries of the magnetic field. Moreover, the artificial limitation of the diffusivity in the region of the ionization fraction minimum was used in these calculations, which leads to the overestimation of the Hall effect efficiency. Our calculations show that the buoyancy constraints the toroidal magnetic field generation. The strength of $B_{\varphi}$ is comparable with the vertical magnetic field strength in the inner region of the disc, $r\lesssim1$~au at $z\sim 0.5$~H. This result confirms simple estimates of DK14. The fossil magnetic field strength is nearly equal to the stellar magnetic field at the inner edge of the disc. In the outer region, the non-linear magnetic ambipolar diffusion and buoyancy constraint the growth of the toroidal magnetic field, so that the toroidal magnetic field $B_{\varphi}$ remains comparable with $B_z$. The Hall effect leads to the transformation of the azimuthal magnetic field to the radial one, and vice versa. The radial magnetic field becomes comparable with the azimuthal and vertical ones due to the Hall effect in the regions where electrons are magnetized. This happens near the borders of the `dead' zone. Thus, the magnetic field gains the quasi-radial geometry, $B_r \sim B_z$, in these regions at $z\sim 0.5$~H. The quasi-radial magnetic field promotes the generation of centrifugal wind \citep{blandford82}. DK14 have shown that magnetic ambipolar diffusion may prevent the generation of the quasi-radial magnetic field. On the basis of our new results, we conclude that the Hall effect is an important factor determining the possibility of the centrifugal wind launching in the accretion discs of young stars. We calculated the geometry and strength of the fossil magnetic field in discs with different accretion rates. We refer to these cases as the evolutionary sequence from the protostellar to the protoplanetary discs. The protostellar discs are characterized by the higher accretion rate, $\dot{M}=10^{-6}\,\rm{M}_{\odot}\,\rm{yr}^{-1}$. The case with $\dot{M}=10^{-7}\,\rm{M}_{\odot}\,\rm{yr}^{-1}$ corresponds to the accretion disc of a young T~Tauri star. In the case of the lower accretion rate $\dot{M}=10^{-8}\,\rm{M}_{\odot}\,\rm{yr^{-1}}$, the disc is considered to be protoplanetary. Calculations show that the geometry of the fossil magnetic field does not depend significantly on the accretion rate. In all cases, the magnetic field remains quasi-azimuthal in the inner region, quasi-vertical inside the `dead' zone, quasi-radial or quasi-azimuthal in the outer regions of the disc. The `dead' zone achieves its maximum size in the protostellar disc. The extent of the `dead' zone decreases with the decrease of the accretion rate. The inner and outer boundaries of the `dead' zone move closer to the star. The strength of the fossil magnetic field goes down at any given distance during the evolution. For example, the vertical magnetic field strength falls down from $\sim 0.1$~Gs ($\dot{M}=10^{-6}\,\rm{M}_{\odot}\,\rm{yr}^{-1}$, protostellar disc) to $\sim 0.01$~Gs ($\dot{M}=10^{-8}\,\rm{M}_{\odot}\,\rm{yr^{-1}}$, protoplanetary disc) at $r=3$~au. These values coincide with the remnant magnetic field strength inferred from the meteorites magnetization measurements \citep{levy78, fu14}. We constructed the synthetic maps of the dust emission polarized due to the alignment of the dust grains with the magnetic field direction. The combination of the quasi-azimuthal and the quasi-radial magnetic field geometries in the inner disc appear as the spiral magnetic field structure in the face-on disc map. The synthetic polarization map for the face-on disc shows that it will be possible to spatially resolve the `dead' zones in nearby accretion discs in observations with the instruments like The Atacama Large Millimeter/Submillimeter Array (\textsc{ALMA}). The `dead' zones will appear as a hole-like central regions where the value of polarization degree is small compared to those in the adjacent parts of the disc. Thus, the observations of the polarized emission can be a useful tool for the investigation of the planet formation region properties in the protoplanetary discs. Our synthetic maps are more detailed than the current observational data. In recent observations, it has been found that the magnetic field may be toroidal \citep{scox15} or the combination of toroidal and poloidal \citep{stephens14} in the protostellar and protoplanetary discs. We predict that we will see the different types of the magnetic field geometries in the different parts of the discs. For example, the quasi-radial magnetic field still has not been observed in the outer regions of protoplanetary discs. We hope that future polarization measurements will confirm our predictions. We stress out that the radiative transfer calculations presented in this work are not aimed to reproduce the observations and to perform the detailed simulations of disc observations because of the simplicity of our disc model. The model does not account the dust growth, settling or radial drift that can affect the dust size and spatial distribution and, thus, the polarization of the disc emission \citep[see e.g.][]{Cho_etal2007}. The detailed observations of the continuum emission polarization may be useful to investigate the dust dynamics and protoplanet formation process. Moreover, we do not consider the cases when the disc is embedded in a dense envelope. In the present work we haven't considered UV radiation from the central star and the role of dust particles charge. The stellar UV ionizes only thin surface layers of the disc. Our conclusions about the magnetic field inside the disc will not change if UV radiation will be taken into account. Examinations of more complex ionization model including different dust grain charges will be addressed to the future papers. We neglected the effect of the charged dust grains on the anisotropy of the conductivity tensor. Charged dust grains with mass $m_{\rm{g}}$ and charge $e$ can produce the anisotropy of the conductivity only in case when their magnetization parameter \begin{equation} \beta_{\rm{g}} = \frac{eB}{m_{\rm{g}}c}\frac{1}{\nu_{\rm{gn}}}\label{Eq:beta_g1} \end{equation} is more than unity. In (\ref{Eq:beta_g1}), $\nu_{\rm{gn}}=\langle\sigma v\rangle_{\rm{gn}}n_{\rm{n}}$ is the frequency of collisions of the grains with the neutrals, $\langle\sigma v\rangle_{\rm{gn}}\simeq \pi a_{\rm{d}}\sqrt{\frac{8kT}{\pi m_{\rm{n}}}}$ is the corresponding coefficient of the momentum transfer, $m_{\rm{n}}$ is the mass of the neutral particle. Assuming that the density of the dust grains $\rho_{\rm{g}}=2\,\rm{g}\,\rm{cm}^{-3}$, we derive from (\ref{Eq:beta_g1}) with the typical parameters at $1$~au in our model (see Figs.~\ref{Fig:1}, \ref{Fig:2}, \ref{Fig:4}) \begin{equation} \beta_{\rm{g}} = 7\times 10^{-17}\left(\frac{B}{0.1\,\rm{Gs}}\right)\left(\frac{a_{\rm{d}}}{0.1\,\mu\rm{m}}\right)^{-5}\left(\frac{T}{800\,\rm{K}}\right)^{-\frac{1}{2}}\left(\frac{n_n}{10^{14}\,\rm{cm}^{-3}}\right)^{-1}.\label{Eq:beta_g2} \end{equation} Therefore, the dust particles with sizes $a_{\rm{d}}\geq 0.1\,\mu$m (considered in our calculations) are too large to be magnetized and contribute to the anisotropy of conductivity. Expression (\ref{Eq:beta_g2}) shows that the dust grains with $a_{\rm{d}}<6\times 10^{-3}\,\mu$m will be magnetized with the adopted values of density, temperature, and magnetic field strength. It should be noted that the radius of the dust grains at $T=800$~K is 15 times smaller than the initial radius $a_{\rm{d}}=0.1\mu$m at $T<150$~K due to the evaporation of ices in our calculations. Such dust grains are also non-magnetized according to (\ref{Eq:beta_g2}). The evaporation of the dust grains in the inner regions of the discs should be taken into account in the conductivity calculations. The evaporation of the dust grains causes the reduction of the recombination rate which leads to the growth of the ionization fraction. Turbulence in our model is described in terms of \citet{ss73} $\alpha$ approximation. In order to determine the turbulent diffusivity and its influence on the magnetic field correctly, it is required to study MHD turbulence properties by solving the Reynolds equations \citep[e.g., ][]{ruzmaikin88}. Turbulence leads to the generation of the small scale magnetic field out of the large-scale one. The question whether the turbulence can transport the large-scale magnetic field requires further analysis. This complex issue is beyond the scope of current work. Our further investigations will focus on the dynamical influence of the magnetic field on the structure of the accretion discs. | 16 | 9 | 1609.03969 |
1609 | 1609.03472_arXiv.txt | Recent observations from both space-based GeV observatories and ground-based TeV observatories reveal $\gamma$-ray emission from several middle aged supernova remnants (SNRs) which are interacting with molecular clouds (MCs), e.g. W44 \citep{Abdo10a,Giuliani11,Uchiyama12}, IC443 \citep{Albert07,Acciari09,Abdo10b}, W28 \citep{Aharonian08,Abdo10c,Hanabata14} and W51C \citep{Abdo09,Aleksic12}. The observed $\gamma$-ray emission from these objects shows a smooth transition from the GeV band to the TeV band (if detected) and they all peak in the GeV band. The characteristic $\pi^0$-decay signature identified in IC 443 and W44 \citep{Giuliani11,Ackermann13} provides possible direct evidence for cosmic ray (CR) proton acceleration in SNRs, making middle aged SNRs detected in $\gamma$-rays an important class of objects for understanding CR acceleration in SNRs. So far two scenarios have been proposed to explain the observed spatial correlation between the $\gamma$-ray emission region and the MC region with hadronic origin: one is the escaping scenario \citep{Gabici09,Fujita09,LC10,Ohira11}, which focuses on the CR particles that escaped from the remnant, and the other one is the direct interaction scenario \citep{bykov00,Uchiyama10,Inoue10,TC14,TC15,Cardillo16}, which studies the CR particles accelerated at the remnant. In the escaping scenario, CR particles having escaped from the remnant together with the pre-existing CR in the ambient medium illuminate the nearby MCs, producing the $\gamma$-ray emission. A collision between the SNR and MC is not necessarily required and the resulting $\gamma$-ray emission can be external to the remnant. A good example for this scenario might be W28. In the direct interaction scenario, a collision between the remnant and MC occurs, as indicated by both observation and theory. Shocked clouds are identified in middle aged SNRs and are spatially coincident with the $\gamma$-ray emission region \citep[e.g.,][]{Abdo10a,Abdo10b,uchiyama11,Nicholas12}. Besides, the MC interaction creates a region with high density which then becomes an ideal site for $\pi^0$-decay emission. Diffusive shock acceleration (DSA) is believed to be the particle acceleration mechanism in most astrophysical environments involving shock waves \citep[e.g.,][]{Bell78,B&E87}. The theory naturally produces a power law spectrum of energetic particles in a steady state which is close to the observed CR spectrum after taking into account propagation effects. DSA is considered to be the most plausible particle acceleration mechanism at the SNR shock front for both scenarios. Here, we limit our discussion to the direct interaction scenario. In this scenario, re-acceleration of pre-existing CRs is taken into account, while particle injection through the thermal pool is neglected in view of the slow radiative shock \citep{B&C82,Uchiyama10,TC14}. If this picture is correct, then it implies a transition of seed particles for CR acceleration in SNRs from injected seed particles to pre-existing CRs. When does this transition occur and how would it affect the accelerated CR spectrum in SNRs could be interesting problems for future study. The standard DSA theory produces too flat a steady state particle spectrum compared to that indicated by observations, so it has been suggested that the observed particle spectrum in the energy range of interest has not reached a steady state yet. An exponential cutoff has been explored in \cite{Uchiyama10} and \cite{TC14} as the high energy cutoff for MC interaction models. It is found that for remnants like IC 443, W28 and W51C with TeV detection, the exponential cutoff tends to underproduce the TeV emission. In \cite{TC15}, we developed a time dependent DSA solution in the test particle limit for the case of re-acceleration of pre-existing CRs and show that the time dependent solution in combination with the MC interaction model in \cite{TC14} is capable of explaining the observed $\gamma $-ray emission in remnants like W44 and IC 443 from GeV to TeV bands. In the following section, we first discuss the properties of the time dependent DSA solution and then provide a simple physical picture to understand the spectral shape of the time dependent DSA solution. | If the observed $\gamma$-ray emission from those middle aged SNRs interacting with MCs is the result of time dependent DSA, according to eq.\ \ref{shift} the extension of the plateau region in the $\gamma$-ray spectrum reflects the logarithmic ratio $\ln(p_f/p_i)$ between the initial and final momentum. Thus it can be used to estimate the spatially averaged diffusion coefficient around the SNRs. In SNRs, the spatially averaged diffusion coefficient in the upstream region $\kappa_1$ is much larger than that in the downstream region $\kappa_2$, as $\kappa_1$ gradually approaches the interstellar value when particles move far away from the shock front. Hence, according to eq.\ \ref{shift} we have \begin{equation} \kappa_1 \approx\frac{(U_1-U_2)U_1t}{3\ln (p_f/p_i)}\approx 10^{25}{\rm cm^2s^{-1}}\frac{r-1}{r\ln(p_f/p_i)}\left(\frac{U_1}{100\kms}\right)^2\left(\frac{t}{10^4~ \rm yrs}\right), \end{equation} where $r$ is the shock compression ratio and $r=4$ for strong shocks. In middle aged SNRs $\ln(p_f/p_i)$ is of order unity based on the observed $\gamma$-ray spectrum. Hence the spatially averaged diffusion coefficient in these middle aged SNRs at the energy range around the plateau region ($\sim \rm GeV$) is approximately $10^{25}\rm cm^2s^{-1}$. The number is close to the value derived in \cite{TC15} through detailed fitting and is found to be between the Bohm limit and the interstellar value. \begin{figure}[htb] \center \includegraphics[width=0.8\textwidth]{schematic.pdf} \caption{Schematic figure for the time dependent DSA solution.} \label{schematic} \end{figure} \small % | 16 | 9 | 1609.03472 |
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1609 | 1609.08626_arXiv.txt | We investigate the response of self-interacting dark matter (SIDM) halos to the growth of galaxy potentials using idealized simulations, each run in tandem with standard collisionless Cold Dark Matter (CDM). We find a greater diversity in the SIDM halo profiles compared to the CDM halo profiles. If the stellar gravitational potential strongly dominates in the central parts of a galaxy, then SIDM halos can be as dense as CDM halos on observable scales. For extreme cases with highly compact disks core collapse can occur, leading to SIDM halos that are denser and cuspier than their CDM counterparts. If the stellar potential is not dominant, then SIDM halos retain constant density cores with densities far below CDM predictions. When a disk potential is present, the inner SIDM halo becomes {\em more flattened} in the disk plane than the CDM halo These results are in excellent quantitative agreement with the predictions of Kaplinghat et al. (2014). We also simulated a galaxy cluster halo with a central stellar distribution similar to the brightest central galaxy of the cluster A2667. A SIDM halo simulated with cross section over mass $\sigma/m = 0.1\ \cmg$ provides a good match to the measured dark matter density profile of A2667, while an adiabatically-contracted CDM halo is denser and cuspier. The cored profile of the same halo simulated with $\sigma/m = 0.5\ \cmg$ is not dense enough to match A2667. Our findings are in agreement with previous results that $\sigma/m \gtrsim 0.1 \, \cmg$ is disfavored for dark matter collision velocities in excess of about 1500 km/s. More generally, the predictive cross-talk between baryonic potentials and SIDM density distributions offers new directions for constraining SIDM cross sections in massive galaxies where baryons are dynamically important. | \label{sec:intro} The dark matter (DM) paradigm has been tremendously successful in explaining the large-scale structure of our universe \citep[see, e.g. ][]{PlanckCos15, SDSS16}, though the precise nature of dark matter remains unknown. The simplest example of cold dark matter (CDM), consisting of a single, collisionless particle with negligible primordial thermal velocity dispersion, can match the large-scale data remarkably well. Alternatively, dark matter could be more complex, with nongravitational coupling to standard model particles \citep[e.g.][]{BoehmSchaeffer05,Escudero15} and/or new dark sector particles \citep[e.g.][]{Feng10,Khlopov10,Lesgourgues16}; many models of this kind produce observable signatures in astronomical data sets \citep{Mangano06,Feng09,Racine15}. In this paper we consider the possibility that dark matter has strong elastic self-scattering interactions and explore the implications of such interactions on the dark matter distributions within individual galaxies. We specifically focus on the back-reaction associated with galaxy formation. Collisional or Self-Interacting dark matter (SIDM) was first explored in the context of galaxy formation by \citet{Spergel00}, who argued that SIDM models with cross-section over mass $\sigmam \sim 1\ \cmg$ should lead to observable constant density cores in galaxies, in agreement with observations at that time. While early estimates suggested that SIDM models of this kind would significantly reduce substructure counts compared to CDM, more recent numerical investigations have shown that the substructure differences are minimal \citep{Vogelsberger12,Rocha13}. However, the original expectation that SIDM halos should have constant-density cores has been demonstrated robustly in cosmological simulations \citep[][]{Dave2001,Rocha13,Zavala13}. SIDM cores are generated by energy-exchange interactions, which heat the halo center until it becomes isothermal. The radial extent of this core is set by the requirement that a typical dark matter particle will experience at least one interaction per Hubble time \citep{Rocha13}. This implies that larger SIDM cross sections produce larger isothermal cores. If the cross section is large enough, the isothermal region can extend beyond the peak in the halo's velocity dispersion profile; in this case, energy-exchange interactions could extract heat from the core leading to core-collapse, which increases the central density \citep{Kochanek2000,Balberg02,Colin02,Koda11,Vogelsberger12}. However, this effect is muted in cosmological simulations and \citet{Elbert15} used dark matter (only) zoom cosmological simulations to show that core-collapse behavior sets in only for very large cross sections $\sigmam \gtrsim 10\ \cmg$. \begin{table*} \centering \begin{tabular}{lccccccccccc} Name & $\mathrm{\mvir}$ & $c_{\rm v}$ & $N_\mathrm{p}$ & $\mathrm{r_{cut}}$ & $\epsilon$ & Convergence Radius & Potential Shape & $\mathrm{M_{\rm gal}}$ & $\mathrm{\textit{a}}$ & $\mathrm{\textit{b}}$ & $h$ \\ & ($10^{12} \msun$) & & ($10^6$) & (kpc) & (kpc) & (kpc) & & ($10^{10} \msun$) & (kpc) & (kpc) & (kpc) \\ \hline \hline MW & $1.0$ & 13 & $3$ & 230 & $0.4$ & 0.83 & MN Disk & $5.0$ & $1.5, 3.0, 6.0$ & $0.3$ & -- \\ \hline LSB & $0.2$ & 11.8 & $10$ & 170 & $0.19$ & 0.30 & MN Disk & $0.06$ & $2.2$ & $0.2$ & -- \\ \hline Elliptical & $1.8$ & 9.7 & $10$ & 300 & $0.37$ & 1.0 & Hernquist Sphere & $6.2$ & $ - $ & -- & 3.0 \\ \hline Cluster & $10^3$ & 3 & $50$ & 500 & $2$ & 3.4 & Hernquist Sphere & $210$ & $ - $ & -- & 28.5 \\ \hline \end{tabular} \caption{Summary of simulated halos. The first five columns list identifying names and general simulation properties: halo mass, NFW concentration, particle number, exponential cutoff radius and force softening. We define $\mvir$ following \citet{Bryan98} with a virial overdensity of $\Delta_{\rm v} = 97$ with respect to the critical density. The sixth column lists the convergence radius for the SIDM runs, which we set to $0.6$ times the \citet{Power03} radius for CDM as found in Elbert et al. (2015). The last four columns summarize the properties of the galaxy potentials grown in each case, where $M_{\rm gal}$ is the final galaxy mass and the other shape parameters are defined in Equations \ref{eqn:MN} and \ref{eqn:H}. Note that there are three separate disks of varying scale length for the Milky Way runs. We refer to these in the text and figures as `Compact', `Fiducial', and `Extended.'} \label{sims.tab} \end{table*} The tendency for SIDM models with $\sigmam \lesssim 10\ \cmg$ to produce constant-density cores with lower overall density is of special interest for comparisons to dwarf and low surface brightness (LSB) galaxies. This is because many of these galaxies are observed to have cores on roughly the scales expected in SIDM \citep{Flores1994,Moore94,deBlok96,Salucci00,deBlok01,Swaters03,Gentile04,Simon05,Spekkens05,KuziodeNaray08,deBlok08,Donato09,Oh11,Adams14} as opposed to the cusps predicted in dissipationless CDM simulations \citep{Dubinski1991,NFW}. SIDM cores also may provide a natural explanation for the unexpectedly low densities of local dwarf galaxies \citep{Vogelsberger12,Vogelsberger14,Elbert15}, a problem known as ``Too Big to Fail" (TBTF) \citep{MBK11,MBK12,Ferrero2012,Klypin2014,Papastergis14,GarrisonKimmelTBTF,Tollerud14}. There are many in the galaxy formation community who believe these issues may be resolved by baryonic processes such as supernova feedback \citep{Navarro96,Read05,Pontzen12,Governato12,DiCintio14,Onorbe15,Maxwell15,Dutton16,Katz16,Read2016} though not all authors necessarily agree \citep{Penarrubia12,SGK13,Pace16}. Tidal effects have been shown to solve TBTF in satellite galaxies \citep[see e.g.][]{Read06,Zolotov2012,Brooks2014,DelPopolo2014,Arraki14}, but the evidence for TBTF in the local field \citep{Kirby14,GarrisonKimmelTBTF} necessitates another solution for these galaxies. This ongoing debate and the lack of DM detections in direct, indirect and collider searches motivates a thorough exploration of the SIDM hypothesis. The goal of this paper is to investigate the effects of galaxy formation on SIDM halos, specifically the contraction of these halos due to the gravitational potential of the galaxy. To this end we use a set of N-body simulations similar to those initially used to examine contraction in CDM halos. The work is organized as follows: in \S \ref{sec:constraints}, we briefly describe the properties required of a viable SIDM model and in \S\ref{sec:motivation}, we sketch the physics of contraction of SIDM halos and motivate our work in this paper. We describe our simulations and analysis in \S\ref{sec:simulations}. We present our results in \S\ref{sec:results}, discussing our Milky Way analogue halos in \S\ref{ssec:MW} and our elliptical and lsb simulations in \S\ref{ssec:other}, while in \S\ref{ssec:analytic} we directly compare our simulations to the analytic model presented in \citet{Kaplinghat15b}. \S\ref{sec:cluster} shows the results of our cluster simulations, and compares these with the observations of \citet{Newman13b}. We summarize our results and conclude in \S\ref{sec:conclusions}. | \label{sec:conclusions} In this work we have investigated the combined effects of baryonic gravitational potentials and dark matter self-interactions on dark matter halos using idealized simulations of dark matter halos with galactic potentials. By simulating halos of various sizes with many different potentials we have found: \begin{itemize} \item SIDM halo shapes are not inherently more resilient to effects from baryons than their CDM counterparts. For a Milky Way halo hosting a Milky Way analogue disk, the SIDM halo is more compact along the disk axis than its CDM equivalent in agreement with the prediction of \citet{Kaplinghat14b}. For an elliptical galaxy, whose stellar potential is markedly more spherical, we expect the SIDM halo to be correspondingly more spherical. \item Halos that host substantial baryonic populations display few differences in spherically-averaged density profiles between CDM and SIDM models on observable scales. Even extended baryon populations can contract halos with respect to SIDM-only simulations, though these systems retain potentially observable constant-density cores and are less dense than CDM. In extreme cases, we find that potentials from dense baryonic structure can cause SIDM halos to core-collapse and become denser than their CDM counterparts. \item Halos that host less massive or highly diffuse stellar and gas disks display substantial differences between CDM and SIDM models. Thus, the original motivation for explaining the low densities observed in galaxies dominated by dark matter is still intact. Among galaxies, these are likely the best systems to measure or constrain the SIDM cross section. \item The densities of our contracted halos are in good agreement with the analytical predictions in \citet{Kaplinghat14b, Kaplinghat15b}, with the exception of the core-collapsing Compact Disk because it no longer obeys the isothermal assumption of the model. In particular, we show that the spherically-averaged density profiles are well approximated by the simple model in \citet{Kaplinghat15b}, which has an isothermal core and an undisturbed CDM outer profile. \item We simulated a cluster halo for 3 Gyr after the brightest cluster galaxy was fully in place to test against the mass measurements for A2667 \cite{Newman13b}. Our simulated CDM halo was denser than the observed central profile for A2667. On the other hand, SIDM with $\sigma/m \simeq 0.5\ \cmg$ was too low in density compared to the measurements. The choice of $\sigma/m \simeq 0.1\ \cmg$ was in much better agreement with the measured normalization (and inner slope) of the A2667 density profile. Larger values like $\sigma/m \simeq 0.5\ \cmg$ are ruled out, even allowing for a factor of 2 uncertainty in the age of the halo. These conclusions are in substantial agreement with the detailed analysis of seven clusters \citep{Newman13b} by \citet{Kaplinghat15b}, which found a average value of $\sigma/m \simeq 0.1\ \cmg$ on cluster velocity scales for an assumed age of 5 Gyr. \end{itemize} Based on these results, an ideal scale to investigate possible DM self-interactions appears to be the dwarf galaxy scale with halo masses $10^{10-11}\ \msun$ scale, as they will have the largest interaction cross sections and the least contracted halos. However, these are precisely the halos expected to be most vulnerable to stellar feedback \citep{Pontzen12,Governato12,DiCintio14,Onorbe15}. Ongoing work \citep{Vogelsberger14,Fry15,Roblesinprep} is investigating the effects of feedback on the SIDM halos and their results suggest that dwarfs with stellar masses $M_\star \lesssim 10^6 \msun$ will have density profiles indistinguishable from the predictions of the dark matter-only simulations. This suggests that the faintest dwarf spheroidals provide excellent laboratories constraining SIDM models. For halo masses much larger than $10^{11}\ \msun$ that host a large stellar disk or bulge, as the inner halo becomes isothermal the SIDM halo retains the high densities created by adiabatic contraction following the formation of the disk. Thus, in Milky Way sized halos the CDM and SIDM halos densities are very similar beyond about a kpc, in marked contrast to the dark-matter-only predictions. As predicted by \citet{Kaplinghat14b}, the self-interactions also force the SIDM halo to be more compact along the stellar disk axis. We find that the SIDM halo in the inner region of Milky Way analogs is more compact along the galactic disk axis than the CDM halo. Thus, it may be possible to use the shape of the dark matter halo in the inner regions of large spiral galaxies to provide a sharp test of the SIDM paradigm. The predictive cross-talk between the dark matter and baryons in the SIDM models leads to a large diversity of halo profiles and halo shapes. This cross-talk is purely gravitational and the result of the dark matter becoming isothermal in the inner parts of the halos and they are fully explained by simple equilibrium models. The prospects for using these concrete predictions of the SIDM paradigm to rule in or rule out SIDM in the near future are excellent. | 16 | 9 | 1609.08626 |
1609 | 1609.05704_arXiv.txt | We investigate the physical properties of \textit{Fermi} TeV BL Lac objects jets by modeling the quasi-simultaneous spectral energy distribution of 29 \textit{Fermi} TeV BL Lacs in the frame of a one-zone leptonic synchrotron self-Compton model. Our main results are the following: (i) There is a negative correlation between $B$ and $\delta$ in our sample, which suggests that $B$ and $\delta$ are dependent on each other mainly in Thomson regime. (ii) There are negative correlations between $\nu_{\text{sy}}$ and $r$, the $\nu_{\text{IC}}$ and $r$, which is a signature of the energy-dependence statistical acceleration or the stochastic acceleration. There is a significant correlation between $r$ and $s$, which suggests that the curvature of the electron energy distribution is attributed to the energy-dependence statistical acceleration mechanism. (iii) By assuming one proton per relativistic electron, we estimate the jet power and radiative power. A size relation $P_{\text{e}} \sim P_{\text{p}} > P_{\text{r}} \gtrsim P_{\text{B}}$ is found in our sample. The $P_{\text{e}}>P_{\text{B}}$ suggests that the jets are particle dominated, and the $P_{\text{e}}\sim P_{\text{p}}$ means that the mean energy of relativistic electrons approaches $m_{\text{p}}/m_{\text{e}}$. There are not significant correlations between $P_{\text{jet}}$ and black hole mass in high or low state with a sub-sample of 18 sources, which suggests that the jet power weakly depends on the black hole mass. (iv) There is a correlation between the changes in the flux density at 1 TeV and the changes in the $\gamma_{\text{peak}}$, which suggests the change/evolution of electron energy distribution may be mainly responsible for the flux variation. | Blazars are the most extreme active galactic nuclei (AGNs) pointing their jets in the direction of the observer (Urry \& Padovani 1995). They have high luminosity, large amplitude and rapid variability, high and variable polarization, radio core dominance, and apparent super-luminal speeds (Urry \& Padovani 1995; Massaro et al. 2016). Generally, Blazars are divided into subcategories of BL Lacs objects (BL Lacs), characterized by almost completely lacking of emission lines or only showing weak emission lines (EW $\leq$ 5 $\mathring{\text{A}}$), and highly polarized quasars or flat spectral radio quasars (FSRQs), showing broad strong emission lines (Falomo et al. 2014; Massaro et al. 2014). The broadband spectral energy distributions (SEDs) of blazars are double peaked. The bump at the IR-optical-UV band is explained with the synchrotron emission of relativistic electrons, and the bump at the GeV-TeV gamma-ray band is due to the inverse Compton (IC) scattering (e.g., Dermer et al. 1995; Dermer et al. 2002; B\"{o}ttcher 2007). The seed photons for IC process could be from the local synchrotron radiation on the same relativistic electrons (i.e. synchrotron self-Compton (SSC); e.g., Tavecchio et al. 1998), or from the external photon fields (EC; e.g., Dermer et al. 2009), such as those from accretion disk (e.g., Dermer \& Schlickeiser 1993) and broad-line region (e.g., Sikora et al. 1994). The hadronic model is an alternative explanation for the high energy emissions from blazars (e.g., Dermer et al. 2012). The modeling of SED with a given radiation mechanism allow us to investigate the intrinsic physical properties of emitting region and the physical conditions of jet (e.g., Ghisellini \& Tavecchio 2008; Celotti \& Ghisellini 2008; Ghisellini et al. 2009; Ghisellini et al. 2010; Ghisellini et al. 2011; Zhang et al. 2012; Yan et al. 2014). Celotti \& Ghisellini (2008) estimated the powers of blazars jets based on EGRET observations and they found that the typical jet should comprise an energetically dominant proton component and only a small fraction of the jet power is radiated if there is one proton per relativistic electron. In addition, in their work, the TeV BL Lacs shows some special jet properties, such as $P_{e} \sim P_{p}$ and relatively high radiation efficiency. Ghisellini et al. (2009, 2010, 2011) mainly concerned the relation between the jet power and the accretion disk luminosity in \textit{Fermi} blazars and they found that there is a positive correlation between the jet power and the accretion disk luminosity for \textit{Fermi} broad-line blazars. The $\gamma$-ray extragalactic sky at high (> 100 MeV) and very high (> 100 GeV) energies is dominated by blazars. About 50 blazars have been detected in the TeV gamma-ray band\footnote[1]{http://tevcat.uchicago.edu/.}, and most of them are the BL Lacs. The SEDs of TeV BL Lacs suffer less contamination of the emission from the accretion disk and EC process, so it can be explained well by the one-zone SSC model (Paggi et al. 2009a; Dermer et al. 2015). Since the launch of the \textit{Fermi} satellite, we have entered in a new era of blazar research (Abdo et al. 2009; Abdo et al. 2010a). The abundant data observed by \textit{Fermi}/LAT in the MeV-GeV band, together with the multi-wavelength campaigns at the radio, optical, X-ray bands and the ground-based observations at the TeV gamma-ray band, now provide an excellent opportunity to study the TeV blazars (e.g., Massaro et al. 2011a, Giommi et al. 2012; Massaro et al. 2013). In this paper, we have collected the quasi-simultaneous broadband SEDs of 29 \textit{Fermi} TeV BL Lacs from the literatures. We used the one-zone leptonic synchrotron self-Compton model with the log-parabolic electron energy distribution to fit SEDs. And we used the Markov Chain Monte Carlo sampling method instead of the "eyeball" fitting to obtain the best-fit model parameters. Then, based on the model parameters, we systematically investigated the physical properties of \textit{Fermi} TeV BL Lacs jets through statistical analysis. This paper is organized as follows: In Sect.2, we present the sample, the model and the fitting strategy are presented in Sect.3. Then, results and discussions are showed in Sect.4. Finally, we end with a conclusion of the findings in Sect.5. The cosmological parameters $H_{0}=70~$Km~s$^{-1}$Mpc$^{-1}$, $\Omega_{m}=0.3$, and $\Omega_{\Lambda}=0.7$ are adopted in this work. | In this work, we have modeled the quasi-simultaneous broadband SEDs of 29 \textit{Fermi} TeV BL Lac objects by using a one-zone leptonic SSC model with the log-parabolic electron energy distribution. We obtain the best-fit model parameters by MCMC sampling method. Then, we systematically investigate the physical properties of \textit{Fermi} TeV BL Lac jets. Our main results are the following: (i) There is a negative correlation between $B$ and $\delta$ for our source sample, which suggests that $B$ and $\delta$ are dependent on each other mainly in Thomson regime. (ii) There are negative correlations between $\nu_{\text{sy}}$ and $r$, $\nu_{\text{IC}}$ and $r$ for our source sample, which confirms the consistency of the model and is a signature of the energy-dependence statistical acceleration or the stochastic acceleration. We check the correlation between $r$ and $s$, and find that they are significantly correlated. The result suggests that the curvature of log-parabolic electron energy distribution is attributed to the energy-dependence statistical acceleration mechanism working on the emitting electrons in \textit{Fermi} TeV BL Lacs. We do not find any correlations between $B$ and $r$, which may be because the curvature of the electron energy distribution is related to the particle acceleration rather than to the radiative cooling process. However, it should be pointed out that because we use a "static" SSC code and do not consider the evolution of the electron energy distribution, so the effects of radiation cooling on the electron energy distribution are not really addressed in here. (iii) By assuming one proton per relativistic electron and the jet power are be carried by relativistic electrons, cold protons, magnetic field (Celotti \& Ghisellini 2008). We calculate the jet power in different forms and the radiative power in the stationary frame of the host galaxy. We find a size relation $P_{\text{e}} \sim P_{\text{p}} > P_{\text{r}} \gtrsim P_{\text{B}}$. The $P_{\text{e}}>P_{\text{B}}$ suggests that the jets are particle dominated in \textit{Fermi} TeV BL Lacs. The $P_{\text{e}}\sim P_{\text{p}}$ is consistent with Celotti \& Ghisellini (2008) and means that the mean energy of relativistic electrons approaches $m_{\text{p}}/m_{\text{e}}$ in \textit{Fermi} TeV BL Lacs. Then, we explore the correlation between $P_{jet}$ and black hole masses with a sub-sample of 18 sources in our sample. There are not significant correlations whether in high or low state, which suggests that the jet power weakly depends on the black hole mass for \textit{Fermi} TeV BL Lac jets. And we find the $P_{\text{jet}} \propto \gamma_{\text{peak}}^{-1.02\pm0.19}$, which is consistent with the prediction of the blazar sequence. (iv) At last, we explore the cause of flux variation. We only find a correlation between the changes in the flux density at 1 TeV and the changes in the $\gamma_{\text{peak}}$, which suggests the change/evolution of electron energy distribution may be mainly responsible for the flux variation in our sample. | 16 | 9 | 1609.05704 |
1609 | 1609.02312_arXiv.txt | The early star-forming Universe is still poorly constrained, with the properties of high-redshift stars, the first heating sources, and reionization highly uncertain. This leaves observers planning 21-cm experiments with little theoretical guidance. In this work we explore the possible range of high-redshift parameters including the star formation efficiency and the minimal mass of star-forming halos; the efficiency, spectral energy distribution, and redshift evolution of the first X-ray sources; and the history of reionization. These parameters are only weakly constrained by available observations, mainly the optical depth to the cosmic microwave background. We use realistic semi-numerical simulations to produce the global 21-cm signal over the redshift range $z = 6-40$ for each of 193 different combinations of the astrophysical parameters spanning the allowed range. We show that the expected signal fills a large parameter space, but with a fixed general shape for the global 21-cm curve. Even with our wide selection of models we still find clear correlations between the key features of the global 21-cm signal and underlying astrophysical properties of the high redshift Universe, namely the Ly$\alpha$ intensity, the X-ray heating rate, and the production rate of ionizing photons. These correlations can be used to directly link future measurements of the global 21-cm signal to astrophysical quantities in a mostly model-independent way. We identify additional correlations that can be used as consistency checks. | \label{Sec:Intro} Some of the most exciting epochs in cosmic history, including the cosmic dark ages, the formation of the first radiative sources (cosmic dawn), and the onset of the epoch of reionization during which the entire Universe became ionized, are currently inaccessible observationally. Our theoretical understanding of galaxy formation gives us significant guidance, but this is limited by astrophysical uncertainties \citep{Barkana:2016}. A major focus are three cosmic events expected at early times \citep{Madau:1997}: cosmic reionization (known to have occurred given the highly ionized Universe at present \citep{GP}), cosmic heating (likely by X-rays), and \Lya coupling (an event specific to 21-cm cosmology). In the hierarchical picture of structure formation, halos grew gradually during the dark ages, assembling mass via gravitational interactions. Massive enough halos were able to retain gas which could radiatively cool, condense and form stars, with the first stellar objects forming at $z\sim 65$ \citep{Naoz:2006, Fialkov:2012}. The minimal mass of halos within which stars can form, M$_{\textrm{min}}$, depends on the chemical composition of the gas, and in the pristine conditions at high redshifts, two cooling channels dominate: (1) radiative cooling of molecular hydrogen happens in the smallest halos, with mass above $10^5\,$M$_{\odot}$ \citep[e.g.,][]{Tegmark:1997, Bromm:2002, Yoshida:2003}, and (2) radiative cooling of atomic hydrogen takes place in halos with mass above $10^7\,$M$_\odot$ \citep[e.g.,][]{Barkana:2001}. Star formation in small halos is a vulnerable process and is believed to be affected by several feedback mechanisms which can either boost or suppress the formation of the next generation of stars. One of the mechanisms discussed in the literature is the Lyman-Werner (LW) feedback. UV radiation in the LW band emitted by the first stars can dissociate hydrogen molecules \citep{Haiman:1997}, depleting the reservoirs of gas available for the formation of the future stars \citep[however, the efficiency of the LW feedback is poorly understood, e.g.,][]{Visbal:2014,Schauer:2015}. Because LW photons reach up to $\sim 100$ comoving Mpc, this feedback is not local and star formation activity at one site can potentially sterilize halos over a large cosmological volume. Another possible feedback mechanism is the stellar feedback from supernova explosions which can expel gas from the halo, effectively increasing the minimum cooling mass well above the atomic cooling threshold \citep{Wyithe:2013}. An additional feedback mechanism that can affect star formation is the photoheating feedback which becomes efficient once the intergalactic gas is photoheated above $10^4$ K by ionizing radiation emitted by stars; this gas stops accreting onto halos below $\sim 10^8-10^9$ M$_\odot$, thus quenching subsequent star formation within them. Because heavy halos are rare at high redshifts, LW, supernova and photoheating feedbacks can, when they are effective, delay major cosmological events such as the heating of the intergalactic gas and reionization. Finally, there is a possibility that light halos (below the atomic cooling mass) can continue to contribute to star formation even in the presence of LW radiation, via the metal-line cooling channel. Because metal-line cooling is more efficient than molecular cooling, this channel can dominate star formation in small halos once the gas is enriched by the first supernovae explosions. However, the possibility of star formation via metal cooling in the early universe and its contribution to the total star formation is highly uncertain \citep[e.g.,][]{Jeon:2014,Wise:2014,O'Shea:2015,Cohen:2016}. The fraction of gas that is converted into stars (the star formation efficiency, hereafter SFE) is another unknown and can be of order a few tens of percent or lower depending on the halo mass, redshift and dominant feedback mechanisms. Observations at low redshifts show that the star formation efficiency is a few percent in massive halos \citep{Tinker:2016}, while isolated dwarf galaxies show a very low SFE of order $\sim 0.1-0.01$ \% \citep{Read:2016}. Simulations of high-redshift stellar activity present a large scatter of values for SFE, especially for small halos which likely dominated the early Universe \citep[e.g.,][]{Wise:2014, O'Shea:2015, Xu:2016}. Matching the observed luminosity function to the expected number of halos at $z\gtrsim 6$ shows that the peak value of the SFE is 30\% for halos of M$_\textrm{h}\sim 10^{11}-10^{12}$ M$_\odot$, dropping to SFE$\sim 10\%$ at the low mass M$_\textrm{h}\sim 2\times 10^{10}$ M$_\odot$ and high mass M$_\textrm{h}\sim 3\times 10^{13}$ M$_\odot$ limits \citep{Behroozi:2015, Mirocha:2016, Mason:2015, Mashian:2016, Sun:2016}. As noted above, the formation of the first luminous objects had a dramatic effect on the Universe, completely changing the environment. The first astronomical objects emitted UV and X-ray radiation which heated and ionized the gas while supernova explosions enriched the primordial gas with metals leading to the formation of the next generation of stars. Stars are believed to have been the main origin of UV photons which reionized the neutral intergalactic medium (IGM), resulting in the total cosmic microwave background (CMB) optical depth of $\tau \sim 0.055\pm0.009$ \citep{Planck:2016b}. However, the origin of the first heating sources, which raised the temperature of the IGM above that of the CMB, is still debatable. The most plausible heating radiation is X-rays, which can travel far even in a neutral Universe. The X-ray efficiency of the sources as well as their spectral energy distribution (SED) remain very poorly constrained. Several different candidates have been proposed in the literature including X-ray binaries (XRBs) \citep{Mirabel:2011}, mini-quasars \citep{Madau:2004}, hot gas in the first galaxies, and hard X-rays produced via inverse Compton scattering of the CMB off electrons accelerated by supernovae \citep{Oh:2001}. Finally, there are more exotic possibilities such as dark matter annihilation \citep{Cirelli:2009}. Out of the plethora of candidates, XRBs (which have a hard SED which peaks around $1-3$~keV) are likely to be the dominant source of cosmic heating at $z\gtrsim 6$ \citep{Mirabel:2011,Fragos:2013}. The hard spectrum has a major effect on 21-cm cosmology, substantially delaying cosmic heating and decreasing the amplitude of 21-cm fluctuations from heating \citep{Fialkov:2014b}. Extrapolations of recent observations to high redshift continue to support such a scenario \citep{Madau:2016,Mirocha:2016}. However, direct observational constraints on the X-ray efficiency of the first sources are rather weak. Upper limits on the heating efficiency come from the soft unresolved cosmic X-ray background \citep{Fialkov:2016} and lower limits are given by the observed upper limits on the 21-cm power spectrum \citep{Ali:2015, Pober:2015, Fialkov:2016}. The most promising tool to explore the early universe is the redshifted 21-cm signal of neutral hydrogen. It is strongly affected by astrophysics and cosmology, and, thus, is believed to be an excellent probe of processes that took place at high redshifts. In particular, the first stars also are expected to have emitted Ly$\alpha$ photons (plus higher energy photons that redshifted down to Ly$\alpha$), which coupled the 21-cm line (in terms of the relative abundance of its ground and excited states) to the kinetic gas temperature, leading to a strong, potentially observable, 21-cm signal \citep{Madau:1997}, which otherwise would have faded away by $z \sim 30$. The currently unexplored parameter space of the early universe leaves a large window within which the 21-cm signal may fall, making it difficult to predict its shape and guide current and future radio telescopes. The signal has not been detected yet \footnote{\citet{Bernardi:2016} used a Bayesian method with a simplified Gaussian model for the absorption feature (see Section~\ref{Sec:Methods}) to constrain the global signal using early LEDA observations.} and only upper limits have been placed on its power spectrum at redshift $z<10$ \citep{Ali:2015,Pober:2015,Ewall-Wice:2016}; however, many current and future observations aim to detect and measure the signal out to $z\sim35$. Experiments such as the Experiment to Detect the Global EoR Step \citep[EDGES,][]{Bowman:2010}, the Shaped Antenna measurement of the background RAdio Spectrum \citep[SARAS,][]{Patra:2013}, the Large Aperture Experiment to Detect the Dark Age \citep[LEDA,][]{Bernardi:2015} and the Dark Ages Radio Explorer \citep[DARE,][]{Burns:2012}, are trying to measure the global signal, while the Low Frequency Array \citep[LOFAR,][]{van Haarlem:2013}, the Murchison Wide-field Array \citep[MWA,][]{Bowman:2013,Ewall-Wice:2016}, the Precision Array to Probe the Epoch of Reionization \citep[PAPER,][]{Ali:2015}, the Hydrogen Epoch of Reionization Array \citep[HERA,][]{Pober:2014,DeBoer:2016}, and the Square Kilometer Array \citep[SKA,][]{Koopmans:2015} are aiming to measure the power spectrum. Our goal in this paper is to explore the full parameter space of the global 21-cm signal resulting from the uncertainties in the astrophysical parameters of the high-redshift universe. Other recent work has focused on extrapolating low-redshift observations of galaxies to high redshift \citep{Madau:2016,Mirocha:2016}, but we adopt a more flexible approach. While it will be interesting to use observations to find out if such extrapolations are accurate, a~priori, this cannot be assumed. Compared to current observations (which are mostly at relatively low redshift), conditions are very different at redshift 20, e.g., in terms of the CMB temperature, the cosmic and virial halo densities (of both the dark matter and gas), the typical mass of galactic halos, and halo merger histories. Thus, the astrophysical properties of early galaxies could be quite different from those suggested by extrapolations of observed galaxies, and it is important to keep an open mind until direct observational evidence becomes available. In what follows, as we lay out the large parameter space possible for the global 21-cm signal, we try to characterize the properties of this signal and find relations between the shape of the global signal and the astrophysical parameters at high redshifts. \cite{Mirocha:2013} previously addressed parameter reconstruction using a physical model for the global signal. In this \citep[as well as the follow-up works by][where the authors study how well current and near-future experiments could constrain the four parameters of their model using the measurements of the signal's three key points and taking into account the foreground and the noise]{Mirocha:2015, Harker:2016}, the authors used analytical formulas or simple models that account only for the mean evolution of the Universe. In contrast, our more realistic simulations include spatial fluctuations in star formation and take into account the finite effective horizons of the radiative backgrounds, spatially inhomogeneous feedback processes, and time delay effects. We also capture a wider parameter space, as our code includes the possibility of having substantial star formation in halos below the atomic cooling threshold, in which case spatially-inhomogeneous processes such as the streaming velocity and LW feedback play a key role (and are included in our 21-cm code but not in others). This paper is organized as follows: In Section~\ref{Sec:Methods} we briefly discuss the general properties of the 21-cm signal as well as our numerical methods. We present and discuss our specific choice of the parameters and their ranges in Section~\ref{Sec:Param}, and show the resulting parameter space spun by the 21-cm signal in Section~\ref{Sec:results}. Finally, we summarize our results and discuss our conclusions in Section~\ref{Sec:sum}. | \label{Sec:sum} In this paper we have explored the allowed parameter space of the global 21-cm signal, varying the main high-redshift astrophysical parameters such as the minimal mass of star-forming halos, star formation efficiency, and heating and ionization rates, all of which are poorly constrained. The large uncertainty in high-redshift astrophysical processes results in weak limits on the predicted 21-cm signal. We used a realistic semi-numerical simulation to produce the 21-cm global signal in the redshift range $z=6-40$ for 193 different sets of astrophysical parameters in agreement with current observations (except that 21 are excluded by the Planck measurement of optical depth at 3~$\sigma$). We applied these data to establish universal patterns in the predicted global 21-cm curves. We found that the general shape of the signal can be predicted theoretically, but its features remain highly unconstrained. Still, there are clear correlations between the three key features of the global 21-cm signal (the high-$z$ maximum, the minimum and the low-$z$ maximum points) and underlying astrophysical parameters of the early universe. Our compilation of realistic models and fitting formulae for these correlations can be used to rule out portions of the parameter space as data from ongoing and future radio experiments becomes available. If and when the global signal is measured, our results can be used to reconstructed key aspects of the high-redshift population of sources including the first stars, X-ray binaries, and mini-quasars. The parameters that we varied can be divided into three categories. The first group consists of parameters related to primordial star formation, including the minimum mass of halos in which stars can form and the star formation efficiency. These parameters are the only ones that affect the shape of the global signal (through the \Lya intensity) from the formation of the first stars down to the redshift where X-ray sources turn on. The second group captures properties of the first heating sources, including their X-ray spectra, luminosity, and evolution with redshift (e.g., XRBs versus mini-quasars), which together with the properties of star formation affect the shape of the signal from the moment when X-ray sources turn on to the point when reionization becomes significant. Finally, the CMB optical depth is related to ionization properties of stars which drive the global signal at the low-redshift end. As anticipated, properties of the high-$z$ maximum point are relatively simple because it occurs at high redshift where significant \Lya coupling begins, which in all our models occurs before heating and ionizing sources complicate the evolution of the global signal. There is a close relation between the redshift of this turning point and the corresponding intensity of the 21-cm signal, providing a potential consistency check for observations. These observable quantities also correlate closely with the Ly$\alpha$ intensity and its derivative at that epoch, according to fitting formulae that we have obtained. Thus, measuring the global signal at this point would help determine the total star formation rate at this early epoch, thus constraining a combination of the minimum cooling mass of star forming halos and the star formation efficiency. The redshift and depth of the absorption trough show the largest scatter, since this minimum can occur under various physical conditions and is affected by many astrophysical parameters. It typically occurs when \Lya coupling approaches saturation and significant X-ray heating begins. In most models, the absorption trough is the strongest feature of the signal and its detection is one of the main goals of the global 21-cm experiments; the large predicted scatter in the location of this point should encourage observers to search for the signal in as wide a frequency range as possible. Measuring the redshift and depth of the absorption trough alone would not help us to strongly constrain any single parameter, but can be used to rule out some areas of the astrophysical parameter space. Despite this complexity, we have shown that the depth of the absorption trough of the 21-cm signal is strongly correlated with the ratio between the Ly$\alpha$ intensity and the X-ray heating rate, as given by a corresponding fitting formula. The 21-cm signal from the low-$z$ maximum is typically expected to be seen in emission, once heating approaches saturation (unless heating occurs very late). This maximum is affected by both cosmic heating and reionization, so in our models it is affected by both the total CMB optical depth and the properties of X-ray sources, with some scatter introduced by other parameters. For maxima that take place at late times the signal is weaker since reionization is then more advanced. We have fit simple functions to relations between (1) the redshift and the brightness temperature of this point, (2) the heating rate and the brightness temperature, and (3) the ratio of the heating rate to the ionization production and the brightness temperature. Therefore, measurement of the redshift and temperature at the emission peak would give a self-consistency check and allow us to estimate both the X-ray and ionizing intensity of sources. Taken together, the correlations in Eqs.~(\ref{eq:JA}), (\ref{eq:dJA}), (\ref{eq:minJdEPS}), (\ref{eq:lowzZ}), and (\ref{eq:lowzZratio}), can be used to directly link future measurements of the global 21-cm signal to astrophysical properties of the high redshift Universe, in a mostly model-independent way. Meanwhile, those in Eqs.~(\ref{eq:highz}) and (\ref{eq:lowzTZ}) can be used as consistency checks on the measurements (or on the theory, depending on one's point of view). Some caution is advisable, as models like ours do not capture the full possible complexity of high-redshift astrophysics. For example, $f_*$ and the other efficiency parameters could vary with redshift, with the local density, or show a large scatter among halos. We expect the main effect of this to be that the correlations that we identified at each key turning point will measure the astrophysical parameters as averaged spatially and over time (out to an earlier time than that corresponding to a given key feature). Also, the scatter could increase in some correlations, particularly those that depend on multiple redshifts as in Figure~\ref{fig:Slope}. We plan to explore how more elaborate models affect our results. Some of our conclusions are reminiscent of those found by \cite{Mirocha:2013}, now derived in the context of a wider array of astrophysical models and more realistic simulations; a conclusion that is particularly similar is that the high-$z$ maximum point reflects the Ly$\alpha$ intensity and its time derivative. We have tried to cover as large a parameter space as possible, in terms of astrophysical source formation, radiative efficiencies, feedback effects, and the mean free path of ionizing photons. The goal was to make our conclusions as robust as possible given current uncertainties about high-redshift astrophysics. However, in the results we have focused only on some of the parameters and a few correlations, namely those that were cleanest and thus most useful. We have explored many others that did not give a clearly useful result, and we plan to continue such studies. Current and future 21-cm observations, such as those mention in the Introduction, are expected to soon begin to exclude realistic possible realizations of the global 21-cm signal. We hope this is followed soon afterwards with detections, which will probe currently mysterious astrophysical processes at very high redshifts. | 16 | 9 | 1609.02312 |
1609 | 1609.02915_arXiv.txt | We present the first detailed chemical abundance study of the ultra-faint dwarf galaxy Tucana~II based on high-resolution Magellan/MIKE spectra of four red giant stars. The metallicity of these stars ranges from $\mbox{[Fe/H]} = -3.2$ to $-2.6$, and all stars are low in neutron-capture abundances ([Sr/Fe] and [Ba/Fe] $< -1$). However, a number of anomalous chemical signatures are present. Three stars are carbon-enhanced, including the most metal-rich star. This star ($\mbox{[Fe/H]}=-2.6$) shows [Na,$\alpha$,Sc/Fe] $< 0$, suggesting an extended star formation history with contributions from AGB stars and Type~Ia supernovae. The other carbon-enhanced stars have $\mbox{[Fe/H]} < -3$ and may be consistent with enrichment by faint supernovae, if such supernovae can produce neutron-capture elements. A fourth star with $\mbox{[Fe/H]} = -3$ is carbon-normal, and exhibits distinct light element abundance ratios from the carbon-enhanced stars. The carbon-normal star implies that at least two distinct nucleosynthesis sources, both possibly associated with Population~III stars, contributed to the early chemical enrichment of this galaxy. Despite its very low luminosity, Tucana~II shows a diversity of chemical signatures that preclude it from being a simple ``one-shot'' first galaxy, but still provide a window to star and galaxy formation in the early universe. | \label{s:intro} Ultra-faint dwarf galaxies (UFDs) are old, metal-poor galaxies with large mass-to-light ratios \citep{Simon07,Brown14}. These galaxies are ${>}30$\,kpc away, but detailed chemical abundances can be derived for the brightest stars in UFDs through high-resolution spectroscopy on 10\,m class telescopes. The abundances of these metal-poor stars likely trace the nucleosynthetic output of the first Population~III (Pop~III) stars that enriched their host galaxy. Since UFDs have relatively simple star formation histories, they are a particularly powerful probe for dwarf galaxy archaeology, as all their stars formed from the same galactic environment \citep[e.g.,][]{Frebel12,Karlsson13,Ji15}. This provides valuable constraints on the nature and site of the first nucleosynthesis events that cannot be derived for field stars from the chemical signatures alone \citep[e.g.,][]{Ji16b}. High-resolution spectroscopy has led to elemental abundance measurements of stars in ten different UFDs. The overarching message is that, in most respects, stars in UFDs are chemically similar to ordinary metal-poor halo stars. Considering the population of ten UFDs, the lowest metallicity stars tend to be carbon-enhanced, a likely signature of the first stars \citep[e.g.,][]{Cooke14, Placco14, Ji15, Yoon16}. Most UFDs show evidence for somewhat sustained star formation and chemical evolution, with [$\alpha$/Fe] ratios that decline over the range $\mbox{[Fe/H]} =-3$ to $-2$, with the notable exception of Segue~1 \citep{Vargas13,Frebel14}\footnote{$\mbox{[X/Y]} = \log_{10}(N_X/N_Y) - \log_{10}(N_X/N_Y)_\odot$ for elements X,Y}. The overall duration of star formation is expected to be very short \citep{Brown14,Webster15}, and these galaxies appear to completely lack stars with $\mbox{[Fe/H]} \gtrsim -1.5$. However, the heavy element abundances of UFD stars differ significantly from those of halo stars. Most UFDs display the by now typical extremely low neutron-capture element abundances \citep[e.g.,][]{Koch13,Frebel14,Ji16a}. But some UFDs deviate and contain distinctly different chemical signatures: Reticulum~II shows the clear signature of a prolific $r$-process event \citep{Ji16b,Ji16c,Roederer16b}; and Canes~Venatici~II contains a star that may have an abnormally high [Sr/Ba] ratio \citep{Francois16}. The diversity of neutron-capture element abundances in UFDs can be interpreted as resulting from highly stochastic production of neutron-capture elements at low [Fe/H] (e.g., \citealt{Lee13}). The UFD Tucana~II (henceforth Tuc~II) was recently discovered in the Dark Energy Survey \citep{Koposov15a,Bechtol15}. It was confirmed to be a galaxy by \citet{Walker16} since it displays a significant velocity dispersion ($8.6^{+4.4}_{-2.7}$\,{\kms}) and its stars span a range of up to 1\,dex in metallicity. The low luminosity ($M_V \sim -3.8$) and overall metallicity ($\left<\mbox{[Fe/H]}\right> \sim -2.2$) suggests that Tuc~II stars may contain clues to early nucleosynthesis and the nature of the first stars. | 16 | 9 | 1609.02915 |
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1609 | 1609.07647_arXiv.txt | In the standard model extended with a see-saw mass matrix, we study the production of sterile neutrinos from the decay of vector bosons at temperatures near the masses of the electroweak bosons. We derive a general quantum kinetic equation for the production of sterile neutrinos and their effective mixing angles, which is applicable over a wide range of temperature, to all orders in interactions of the standard model, and to leading order in a small mixing angle for the neutrinos. We emphasize the relation between the production rate and Landau damping at one loop order, and show that production rates and effective mixing angles depend sensitively upon the neutrino's helicity. Sterile neutrinos with positive helicity interact more weakly with the medium than those with negative helicity, and their effective mixing angle is not modified significantly. Negative helicity states couple more strongly to the vector bosons, but their mixing angle is strongly suppressed by the medium. Consequently, if the mass of the sterile neutrino is $\lesssim 8.35\,\mathrm{MeV}$, there are fewer states with negative helicity produced than those with positive helicity. There is an Mikheyev-–Smirnov-–Wolfenstein type resonance in the absence of lepton asymmetry, but due to screening by the damping rate, the production rate is not enhanced. Sterile neutrinos with negative helicity freeze-out at $T^-_f\simeq 5\,\mathrm{GeV}$ whereas positive helicity neutrinos freeze-out at $T^+_f \simeq 8\,\mathrm{GeV}$, with both distributions far from thermal. As the temperature decreases, due to competition between a decreasing production rate and an increasing mixing angle, the distribution function for states with negative helicity is broader in momentum and hotter than that for those with positive helicity. Sterile neutrinos produced via vector boson decay do not satisfy the abundance, lifetime and cosmological constraints to be the sole dark matter component in the universe. Massive sterile neutrinos produced via vector boson decay might solve the $^{7}Li$ problem, albeit at the very edge of the possible parameter space. A heavy sterile neutrino with a mass of a few MeV could decay into light sterile neutrinos, of a few keV in mass, that contribute to warm dark matter. We argue that heavy sterile neutrinos with lifetime $\leq 1/H_0$ reach local thermodynamic equilibrium. | \label{sec:intro} The paradigm of standard cosmology is inflation plus cold dark matter, $\Lambda CDM$. While it succeeds in describing the formation of structure at large scales, there are discrepancies at smaller scales, especially galactic and sub-galactic. There is the core-cusp problem: N-body simulations of cold dark matter produce dark matter profiles that generically feature cusps, yet observations suggest that the profile has a smooth core \cite{corecusp,dalcantonhogan}. There is also the missing satellites problem, where simulations also predict that typical galazies are surrounded by satellites dominated by dark matter, which is inconsistent with current observations \cite{toobig}. Both of these problems can be solved by allowing some fraction of the dark matter to be warm dark matter, WDM \cite{wdm1,wdm2,wdm3,wdm4,wdm5,wdm6}. A possible candidate for warm dark matter is a massive ``sterile" neutrino \cite{warmdm,dodwid,dodwid2,dolgovhansen,abazajian3,sterilesexperiment,kusenkorev,sterilereview}. Whether dark matter is hot or cold depends upon its free streaming length, $\lambda_{fs}$, which is the cut-off scale in the linear power spectrum of density perturbations. Cold dark matter with $\lambda_{fs} \lesssim$ pc produces cuspy profiles, while warm dark matter with $\lambda_{fs} \sim \mbox{few kpc}$ gives cored profiles. One important question is whether these disrepancies can be explained with the standard $\Lambda CDM$ model by including the effects of baryons in the simulations. Recent simulations of warm dark matter which include velocity dispersion suggest that cores do form, but do not yet reliably constrain the mass of the WDM candidate in a model independent manner \cite{padu}. In order to evade cosmological bounds the distribution functions of warm dark matter cannot be thermal \cite{planck}. How sterile neutrinos are produced in the early universe was studied originally in Ref. \cite{dolgovenqvist, dolgovreview}. In Refs. \cite{dodwid,shifuller,asaka,laine1,laine,laine2,abacir} it was argued that sterile neutrinos are a viable candidate for warm dark matter, which is produced out of local equilibrium either non-resonantly \cite{dolgovenqvist,dodwid,asaka,laine} or resonantly in the presence of a lepton asymmetry \cite{shifuller}. Models in which a scalar decays into a pair of sterile neutrinos at the electroweak scale (or higher) also yields a non-thermal distribution \cite{boyan1,kusenko1,petraki,merle,drewescalar}. Observations of the X-ray emission spectrum of the Andromeda galaxy with Chandra led to tight constraints on the non-resonant production of sterile neutrinos \cite{casey}. More recently, the report of observation of a 3.5 keV signal from the XMM Newton X-ray telescope has been argued to be due to a 7 keV sterile neutrino\cite{bulbul,boyarsky}, although this interpretation has been challenged \cite{jeltema,malysh,anderson,sekiya}. The prospect of a keV sterile dark matter candidate continues to motivate studies in both theory and observation \cite{merle,abazajian0,abazajian,abazajian2,kaplinghat2,kaplinghat3,kusenko1,lellolightsterile,added2,abacir,merleini}. Neutrino masses, mixing and oscillations are uncontroversial evidence of physics beyond the standard model. A robust experimental program has brought measurements of most of the parameters associated with light neutrino masses \cite{pdg,neutrinoexperiments} with several relevant questions poised to be answered in the near future \cite{upcomingexperiments}. Short baseline neutrino oscillation experiments such as LSND and MiniBooNE \cite{lsnd,miniboone} present a picture of the neutrino sector which may require an additional sterile neutrino species of mass $\sim 1 eV$ \cite{sterilesexperiment,giunti2,mirizzi} but there remains tension with other experiments \cite{nosterilesexperiment} and a definitive resolution of these anomalies requires further experiments \cite{lasserre,decay,lello1,shrock,barger,giunti}. An interpretation of short baseline experimental anomalies as a signal for sterile neutrinos leads to a relatively light mass $\sim eV$ which rules out this putative new sterile neutrino as a candidate for dark matter. Howevever, many well motivated extensions beyond the standard model posit the existence of heavy neutrinos. It has been argued that sterile neutrinos with a mass on the order of $\mathrm{MeV}$ or larger \cite{gninenko} could decay and explain the short baseline anomalies. Alternatively, heavy sterile neutrinos produced through rare decay channels could also explain the anomaly \cite{heavysterile}. Recent proposals make the case for a program to search for heavy neutrinos \cite{heavyoscil,boyancascade} in a wide range of experiments including hadron colliders \cite{goran,added1,added3,han,pilafsis,bonivento}. Furthermore, it has been argued that heavy sterile neutrinos in the mass range $100-500 \,\mathrm{MeV}$ can decay non thermally and so evade bounds from cosmology and accelerator experiments \cite{fullkuse}. Sterile neutrinos with mass $\simeq \mathrm{MeV}$ can be of cosmological relevance in models where the reheating temperature is low \cite{gelmini}. A heavy sterile neutrino with mass $\simeq 14 \,\mathrm{MeV}$, a mixing angle $\theta \simeq 10^{-3}$, and a lifetime $\tau_s \simeq 1.8\times 10^{5}\,s$ has been proposed \cite{ishida} as a novel solution to the ``Lithium-problem''. This is the nearly threefold discrepancy between the standard big-bang nucleosynthesis (BBN) prediction and observed abundance of $^{7}Li$ \cite{fields,serpico,pospelovbbn,ishida,poulin}. This solution relies on the energy injected by the decay of the sterile neutrino to destroy part of $^{7}Be$ prior to its conversion into $^{7}Li$ in the late stages of BBN \cite{poulin,ishida}. This mechanism has been recently re-analyzed and confirmed in Ref.\cite{salvati} with a sterile neutrino mass $\simeq 4.35 \,\mathrm{MeV}$, mixing angle $\theta \leq 10^{-5}-10^{-4}$ and lifetime $\simeq 1.8\,\times 10^{5}\, s$. An important bonus of this mechanism is that the decay of the heavy sterile neutrino, $\simeq \mathrm{MeV}$ in mass, yields an \emph{increase} in the effective number of relativistic species $\Delta N_{eff} \simeq 0.34$ at the $95\%\,\mathrm{CL}$\cite{salvati}. The energy injection from the decay of heavy sterile neutrinos with longer lifetime may also contribute to early ionization \cite{haiman}. Although there is no experimental evidence for such heavy sterile neutrinos, there are stringent accelerator and cosmological bounds on their possible masses and mixing angles with active neutrinos \cite{kusbounds,gelmini,oleg,aaron}. There is a hierarchy of masses for the light active neutrinos, with nearly two orders of magnitude between the mass squared differences for the explanation of solar and atmospheric neutrino mixing. Possible extensions beyond the standard model may also accommodate a hierarchy of \emph{heavy} neutrinos \cite{sterilereview,asaka2,drewes}. Current and future underground neutrino detectors may be able to probe dark matter candidates with $\simeq \mathrm{few}\,\mathrm{MeV}$ \cite{pospelov}. The possibility of a hierarchy of \emph{heavy} sterile neutrinos offers novel production mechanisms for warm (and hot) dark matter, from the cascade decays of heavy neutrinos to lighter ones. This possibility is similar to models of many dark matter components proposed recently \cite{dienes}, where the decay of a heavy field seeds the production of a light one. This leads to a scenario of \emph{mixed dark matter} described by several species of massive neutrinos with non-equilibrium distribution functions, and thereby evade Lyman-$\alpha$ constraints\cite{lymanboyar}. A recent article argued on various possible production mechanisms of sterile neutrino directly from standard model processes available throughout the thermal history of the universe and analyzed in detail the scenario of production of mixed dark matter (colder, warmer and hotter) from pion decay shortly after the QCD crossover\cite{mixedlou}. This analysis, along with previous work \cite{lellolightsterile}, also suggests that the decay of a heavy sterile neutrino into a light active neutrino increases the effective number of neutrinos, $N_{eff}$. This has been studied recently in Ref.\cite{salvati} in the context of energy injection post (BBN) from the decay of a heavy sterile neutrino with lifetime $\simeq 10^5\,\mathrm{secs}$. \vspace{5mm} \textbf{Motivation and Goals:} Sterile neutrinos with masses in the range $\mathrm{KeV}-\mathrm{few}\,\mathrm{MeV}$ may play an important role in cosmology. Most of the studies of their production and freeze-out have focused on the well motivated mass range of $\mathrm{few}\,\mathrm{KeV}$ as possible warm dark matter candidates. However, if the hierarchy of masses and mixing of light active neutrinos is of any guide in extensions beyond the standard model, a possible hierarchy of heavier sterile neutrinos that mix with the light active neutrinos, with very small mixing angles, may emerge. In this scenario, the possibility that heavier neutrinos yield a mixture of dark matter components, from cold, heavy species to warm, light ones, with important cosmological impact, and the possibility that $\simeq \mathrm{MeV}$ sterile neutrinos \emph{may} yield a solution to the $^{7}Li$ problem\cite{poulin,ishida,salvati} motivates our study of the production and freeze out of sterile neutrinos in a wider range of masses and temperatures. In this article we study the production of sterile neutrinos solely from \emph{standard model interactions}. Ref. \cite{mixedlou} identified several possible processes available throughout the thermal history of the Universe that \emph{may} lead to the production of a sterile species from its mixing with active neutrinos. Recently, Ref. \cite{laine2} studied the damping rate of $\mathrm{GeV}$ sterile neutrinos at high temperature within the context of the washout rate of leptonic densities. In contrast, we focus on sterile neutrinos with masses $ \lesssim \mathrm{few}\,\mathrm{MeV}$, which is appropriate both to dark matter and the possible solution of the $^7Li$ problem. We highlight the important role which different helicity channels play for the production rate and mixing angles, including cosmological expansion. We also compute the non-equilibrium distribution functions for different helicities, along with various cosmological constraints. In this article our goals are twofold: \begin{itemize} \item{ \textbf{i:)} Using quantum kinetics, we seek to provide a consistent description of both production and freeze-out, valid in a wide range of temperature, under a minimal set of assumptions. These are: \textbf{a)} except for the coupling between sterile and active neutrinos via a see-saw type mass matrix, we only consider the interactions of the standard model. \textbf{b)} consistent with bounds from accelerator experiments and cosmology, \cite{kusbounds,gelmini,casey,bulbul,boyarsky,kaplinghat2,oleg,aaron} we assume that the vacuum mixing angle, $\theta \ll 1$. Taken together, these bounds suggest that $\theta^2\lesssim 10^{-5}$ for a wide range of masses $M_s \lesssim 300 \,\mathrm{MeV}$. \textbf{c):} Interactions in the standard model can be treated perturbatively, and that the relevant degrees of freedom, including active neutrinos, are in local thermal equilibrium (LTE) during both production and freeze-out of sterile species. The latter is consistent with a small mixing angle. } \item{\textbf{ii:)} We work to leading order in the electro-weak coupling $\alpha_W$, without any assumption on the mass scales of the sterile neutrinos. To leading order in $\alpha_W$, the production of sterile neutrinos occurs from the decay of $W$ and $Z$ bosons in the thermal bath. We focus on the temperatures at the electroweak scale, $T \simeq M_{W}, M_Z$. This is sufficiently below the temperature for the electroweak phase transition, which is a crossover at $T_{ew} \simeq 160 \mathrm{GeV}$ \cite{laine}, so that the $W$ and $Z$ bosons are in local thermal equilibrium, with masses close to those at zero temperature. } \end{itemize} \vspace{2mm} \textbf{Brief summary of results:} For simplicity we consider a model with one active and one sterile neutrino. Our main results are: \vspace{2mm} \begin{itemize} \item{We obtain the mass eigenstates, effective mixing angles and damping rates directly from the equations of motion in the medium in terms of the full self-energy \emph{to all orders in weak interactions}. We give an expression for the effective mixing angles which is broadly valid for $\theta \ll 1$ and to all orders in perturbation theory in standard model couplings, at any temperature. The mixing angle in the medium depends strongly on helicity: negative helicity neutrinos (and positive helicity antineutrinos) feature mixing angles which are strongly suppressed at high temperature. In contrast, for positive helicity neutrinos (and negative helicity anti-neutrinos), the corrections to the mixing angle are subleading, so that the effective mixing angle is nearly the same as that in vacuum. This happens because the interaction of neutrinos with positive helicity is helicity suppressed. Damping rates are also helicity dependent and suppressed for those with positive helicity, however because the effective mixing angle is \emph{larger} than that of the negative helicity states, the resulting production rate is comparable in a wide range of masses. We obtain the general form of the quantum kinetic equation that describes the production and freeze out of sterile-like neutrinos The production rate is determined by the damping rate of sterile-like \emph{mass eigenstates} and the mixing angle in the medium. Although the production rate of positive helicity states is suppressed by helicity, it is comparable to the rate for those with negative helicity, because over a wide regime of masses it is compensated by the suppression of the mixing angle for neutrinos with negative helicity.} \item{For sterile-like masses $M_s \ll M_W$ we find a Mikheyev-Smirnov-Wolfenstein (MSW) \cite{msw} resonance in the \emph{absence of a leptonic asymmetry}. However, it is screened by the damping rate and so does \emph{not} lead to enhanced production. For this mass range of $M_s$ negative helicity states freeze out at $T^-_f \simeq 5 \,\mathrm{GeV}$ whereas positive helicity states freeze-out at $T^+_f \simeq \,8 \mathrm{GeV}$. Both feature highly non-thermal distribution functions, where for the negative helicity states the distribution function is broader and hotter than that for positive helicity. Paradoxically, this is a consequence of a \emph{longer freeze-out time} for the negative helicity states, despite the fact that their coupling to the environment is stronger. This is a surprising result, stemming from a competition between a diminishing damping rate and an \emph{increasing} effective mixing angle as the temperature decreases. We argue that this leading order production mechanism establishes a \emph{lower bound} for the abundance. We find however, that sterile-like neutrinos produced via vector boson decay do not satisfy the various bounds on lifetimes and mixing angles to be viable $\mathrm{keV}$ dark matter candidates. However they can be suitable as the $\mathrm{MeV}$ sterile neutrinos that are conjectured to provide a solution to the $^{7}Li$ problem. This is with the caveat that there seems to be tension among the various bounds available in the literature \cite{salvati,aaron}. Just as these heavier neutrinos may decay by injecting energy into the plasma as the solution to this problem, we also conjecture that they may also decay into lighter $\simeq \,\mathrm{KeV}$ sterile neutrinos, with a much smaller branching ratio, that could be suitable candidates for warm dark matter. } \end{itemize} To the best of our knowledge there has not yet been a systematic study of the \emph{full dynamics of production and freeze-out obtaining the non equilibrium distribution functions} of heavy sterile neutrinos with $M_s \lesssim \mathrm{few}\,\mathrm{MeV}$ at the scale $T \simeq 100 \mathrm{GeV}$ with cosmological expansion. Our analysis is motivated by the possible cosmological relevance of sterile neutrinos in a wide range of masses, and complements previous studies that focus on lower temperature regimes. | Our goals in this article are two-fold: i) to obtain the general form of the quantum kinetic equations and effective mixing angles in the medium to describe production and freeze-out of sterile-like (mass eigenstates) neutrinos in a broad range of temperature and under a minimal set of assumptions. Our study departs from previous ones (see the recent review\cite{revster}) in several important aspects: we focus on the different helicity contributions, and we systematically include the absorptive part of the self-energy in the in-medium modification of the mixing angle. ii) To apply the kinetic equations thus found to study the production to leading order in standard model couplings from vector boson decay at $T\simeq M_W$. We obtained the effective mixing angles in the medium directly from the equations of motion in the case of mixing of one sterile with one active neutrino via a see-saw mass matrix with standard model interactions for the active (flavor) neutrino valid when the vacuum mixing angle $\theta \ll 1$ but to \emph{all orders in standard model couplings}. Assuming that all standard model degrees of freedom are in (LTE) in the relevant temperature range we obtained the quantum kinetic equation that describes the production, evolution and freeze-out of sterile-like mass eigenstates. The mixing angles in the medium and the production rate are determined by the real and imaginary parts of the self-energy on the mass shell of the sterile-like mass eigenstate, and depend on helicity. The full quantum kinetic equation to leading order in $\theta \ll 1$ is \be \frac{d n^h_2(q;t)}{dt} = \Gamma^h_2(q) \Big[n_{LTE}(q) -n^h_2(q;t)\Big]\,, \nonumber \ee where $h=\pm$ correspond to helicity states and $\Gamma^\mp_2(q)$ are given by (\ref{gamasterfin}) with (\ref{gamneghel},\ref{gamneghel}), and $n_{LTE}$ is the Fermi-Dirac distribution function in (LTE). The full expression for the mixing angles in the medium, valid to all orders in standard model couplings and to leading order in $\theta \ll 1$ is given in the relativistic limit by \be \theta^h_{eff}(q) = \frac{\theta}{\Bigg[\Big( 1+\frac{\Delta^h(q)}{\xi}\Big)^2+\Big(\frac{\gamma^h(q)}{\xi} \Big)^2\Bigg]^{1/2}}\,, \nonumber \ee where $\Delta,\gamma,\xi$ are given by (\ref{gamneghel}-\ref{delposhel},\ref{chi}) respectively in terms of the real ($\Delta$) and imaginary ($\gamma$) part of the \emph{active neutrino self-energy on the mass shell of the sterile-like eigenstate}. We implemented the quantum kinetic equation to obtain the production of sterile-like neutrinos from vector boson decay at $T \simeq M_W$ including cosmological expansion. For negative helicity neutrinos (and positive helicity anti-neutrinos) the effective mixing angle is strongly suppressed by the medium, however for positive helicity neutrinos (and negative helicity anti-neutrinos) the medium corrections are negligible because the interaction with the medium is helicity suppressed. We find that there is a region of masses for which the production of both species is comparable. It is noteworthy that the mixing angle for negative helicity neutrinos features an MSW resonance \emph{in absence of lepton asymmetry}, which, however, is screened by the imaginary part of the self-energy. This is an important aspect that has not been previously addressed before: the absorptive (imaginary) part of the self-energy \emph{also} contributes to the mixing angle in the medium. Negative helicity neutrinos freeze-out at $T^-_f \simeq 5\,\mathrm{GeV}$ with a broader distribution as a consequence of a competition between a diminishing damping rate $\gamma$ and an \emph{increasing} effective mixing angle as temperature diminishes. Positive helicity neutrinos freeze-out temperature is $T^+_f \simeq 8\,\mathrm{GeV}$ with a distribution that peaks at much smaller momenta, describing a colder species. Accounting for both channels we find that the distribution function of sterile-like neutrinos of mass $M_2 \simeq M_s$ is given by \be n_2(y) = 3.6 ~\Bigg( \frac{\theta^2}{10^{-4}} \Bigg)\Big(\frac{M_s}{\mathrm{MeV}}\Big)^2\,f(M_s,y) \,,\nonumber \ee where $y = q/T$ and $y^2 f(M_s,y)$ is strongly non-thermal and is displayed in fig.(\ref{fig:hump}) revealing the competition between the colder (positive helicity) and hotter (negative helicity) components. The total abundance normalized to that of one relativistic degree of freedom in thermal equilibrium ($\mathcal{N_\nu}$) is \be \frac{\mathcal{N}_2}{\mathcal{N}_\nu} \simeq 2 ~\Bigg( \frac{\theta^2}{10^{-4}}\Bigg)~\Bigg(\frac{M_s}{\mathrm{MeV}} \Bigg)^2~\Bigg[1+ \Big(\frac{M_s}{8.35\,\mathrm{MeV}}\Big)^2\Bigg]\,. \label{totalabufin} \ee The first term in the bracket is the contribution from the positive helicity states and the second from the negative helicity, both become comparable for $M_s \simeq 8.35\,\mathrm{MeV}$. We argue that this expression is a \emph{lower bound} on the abundance of sterile-like neutrinos. The fractional abundance of dark matter contributed by both helicity components is given by (\ref{fracDM}). Constraints from X-ray data on masses and mixing angles suggest that sterile-like neutrinos produced by vector boson decay \emph{cannot} yield a substantial $M_s \simeq \mathrm{KeV}$ warm dark matter component. However, this production mechanism yield a substantial abundance of $M_s \simeq \,\mathrm{MeV}$ \emph{heavy} sterile-like neutrinos with $\theta^2 < 10^{-4}$ consistent with accelerator constraints. Therefore this production mechanism may yield the heavy neutrinos recently invoked to solve the $^{7}Li$ problem\cite{ishida,poulin,salvati}. However, the parameter range determined in \cite{salvati} also bounds $\mathcal{N}/\mathcal{N_\nu} \simeq 10^{-4}$ which is incompatible with the result (\ref{totalabufin}) for the range of mass and mixing angles reported in this reference, and is also in conflict with recent bounds reported in\cite{aaron}. The possibility that heavy $\simeq \mathrm{MeV}$ sterile-like neutrinos decaying after BBN injecting energy in the medium providing a solution of the $^{7}Li$ problem as suggested also in refs.\cite{ishida,poulin,salvati} merits a deeper study both of the production mechanism as well as the cosmological impact of this heavy neutrino species. \vspace{2mm} \textbf{Further questions:} The study of production at $T\gg M_{W,Z}$ requires a deeper understanding of the finite temperature corrections to the dispersion relations of the vector bosons near the electroweak crossover regime, a worthy study which is beyond the scope of this article. We have suggested several other processes that contribute to the production throughout the thermal history of the Universe, while these are higher order (two loops) processes they may comparable to the leading order processes or even dominate at temperature $T \ll M_{W}$. However the medium corrections to the mixing angles are completely determined by the one loop contribution and the effective mixing angle is given by eqn. (\ref{louTtetaf}). Further study of these processes is clearly warranted, they can be competitive near the freeze-out temperature of the leading one-loop contribution for heavy sterile neutrinos and, crucially contribute to the production of lighter mass eigenstates. We have also argued that sterile neutrinos with lifetimes shorter than the age of the Universe that are decaying today must necessarily have thermalized at some time in the past. Studying this thermalization process is of fundamental interest since most of the calculations of production neglect this possibility by neglecting the loss term in the kinetic equation, yet thermalization of heavy sterile neutrinos may have important cosmological consequence for the expansion history. In ref.\cite{hernandez} a study of low scale type I seesaw models suggest that heavy sterile neutrinos do thermalize through collision processes dominant at much lower temperatures motivating further studies of thermalization. We have also suggested that several compelling extensions beyond the standard model posit a hierarchy of sterile neutrino masses, and this opens the possibility that heavier sterile-states may be produced at high temperature, as analyzed here, and decay well after BBN or near the time of matter-radiation equality into lighter sterile states that may be suitable WDM candidates. This mechanism of cascade decay, which is fundamentally similar to that advocated for the solution of the $^{7}Li$ problem as an energy injection mechanism, is worthy of study. | 16 | 9 | 1609.07647 |
1609 | 1609.05197_arXiv.txt | {Particle production via parametric resonance in the early Universe, is a non-perturbative, non-linear and out-of-equilibrium phenomenon. Although it is a well studied topic, whenever a new scenario exhibits parametric resonance, a full re-analysis is normally required. To avoid this tedious task, many works present often only a simplified linear treatment of the problem. In order to surpass this circumstance in the future, we provide a fitting analysis of parametric resonance through all its relevant stages: initial linear growth, non-linear evolution, and relaxation towards equilibrium. Using lattice simulations in an expanding grid in $3+1$ dimensions, we parametrize the dynamics' outcome scanning over the relevant ingredients: role of the oscillatory field, particle coupling strength, initial conditions, and background expansion rate. We emphasize the inaccuracy of the linear calculation of the decay time of the oscillatory field, and propose a more appropriate definition of this scale based on the subsequent non-linear dynamics. We provide simple fits to the relevant time scales and particle energy fractions at each stage. Our fits can be applied to post-inflationary preheating scenarios, where the oscillatory field is the inflaton, or to spectator-field scenarios, where the oscillatory field can be e.g.~a curvaton, or the Standard Model Higgs.} \begin{document} | Compelling evidence supports the idea of an inflationary phase in the early Universe~\cite{Planck2015}. The specific particle physics realization of inflation is however uncertain, so the inflationary period is typically parametrized in terms of a scalar field, the inflaton, with a vacuum-like potential. After inflation, the reheating stage follows, converting all inflationary energy into different particle species, which represent all the matter and radiation in the Universe. Eventually, the created particles dominate the total energy budget and 'thermalize', signaling the onset of the 'hot Big Bang' thermal era. In this paper we consider inflaton potentials with simple monomial shapes, as this gives rise to one of the most important particle creation phenomena in the early universe: parametric resonance. This is the case of chaotic inflation models, where the inflaton rolls down a monomial potential during the whole inflationary period. Although these scenarios are under tension with cosmological data~\cite{Planck2015}, the simple addition of a small non-minimal gravitational coupling reconcile them with the observations~\cite{Tsujikawa:2013ila}. Some scenarios which fit perfectly well the observational data, e.g.~Higgs-Inflation~\cite{Bezrukov:2007ep,Bezrukov:2010jz} and Starobinsky inflation~\cite{Starobinsky:1980te}, also exhibit a monomial potential with a single minimum, but only during the stages following inflation. In all the scenarios we consider, soon after the end of inflation, the inflaton is in the form of a homogeneous condensate, and starts oscillating around the minimum of its potential. Each time the inflaton crosses zero, all particle species sufficiently strongly coupled to the inflaton, are created in energetic bursts. In the case of bosonic species, the production of particles is resonant, and the energy transferred grows exponentially within few oscillations of the inflaton~\cite{Traschen:1990sw,Kofman:1994rk,Shtanov:1994ce,Kaiser:1995fb,Kofman:1997yn,Greene:1997fu,Kaiser:1997mp,Kaiser:1997hg}. In the case of fermionic species, there is also a significant transfer of energy~\cite{Greene:1998nh,Greene:2000ew,Peloso:2000hy,Berges:2010zv}, but Pauli blocking prevents resonance from developing. The production of particles in this way, either of fermions or bosons, represents the archetypical example of what is meant by an initial 'preheating' stage of reheating. Inflationary preheating is however not the only case where parametric resonance takes place in the early Universe. If a light spectator field is present during inflation, this field forms a homogeneous condensate during the inflationary period, and oscillates around the minimum of its potential afterwards. This is the case e.g.~of the curvaton scenario~\cite{Enqvist:Curvaton,Lyth:Curvaton,Takahashi:Curvaton,Mazumdar:2010sa}. The curvaton may decay after inflation via parametric resonance, transferring abruptly all its energy to the particle species coupled to it~\cite{Enqvist:2008be, Enqvist:2012tc, Enqvist:2013qba, Enqvist:2013gwf}. Another example of a spectator field, naturally decaying through parametric resonance after inflation, is the Standard Model Higgs field. If the Higgs is weakly coupled to the inflationary sector, the Higgs is always excited either during inflation~\cite{Starobinsky:1994bd,Enqvist:2013kaa,DeSimone:2012qr}, or towards the end of it~\cite{Herranen:2015ima,Figueroa:2016dsc}. The Higgs is then 'forced' to decay into the rest of the SM species after inflation\footnote{Note that in the case of Higgs-Inflation~\cite{Bezrukov:2007ep,Bezrukov:2010jz}, the Higgs also decays after inflation via parametric resonance, into the rest of the SM fields~\cite{Bezrukov:2008ut,GarciaBellido:2008ab,Figueroa:2009jw}. In this scenario, the Higgs plays however the role of the inflaton. Therefore, the Higgs decay in the case of Higgs-inflation scenarios~\cite{Bezrukov:2008ut,GarciaBellido:2008ab,Figueroa:2009jw,Bezrukov:2014ipa}, should rather be categorized within the context of preheating scenarios.}, via parametric resonance~\cite{Enqvist:2013kaa,Enqvist:2014tta,Figueroa:2014aya,Kusenko:2014lra,Figueroa:2015hda,Enqvist:2015sua,Figueroa:2016dsc}. In this paper, independently of the context, we will often refer to the oscillatory field as the 'mother' field, and to the created species as the 'daughter' fields. Particle production of daughter fields via parametric resonance, corresponds to a non-perturbative effect, which cannot be captured by perturbative coupling expansions, not even if the couplings involved are small~\cite{Kofman:1997yn}. During the initial stage of parametric resonance, the system is linear, and analytical methods can be applied. As the particle production is exponential for bosonic species, the daughter field(s) eventually 'backreact' onto the mother field, making the system non-linear. In order to fully capture the non-linearities of the system, we need to study this phenomenon in the lattice. The approach of classical field theory real-time lattice simulations can be considered valid as long as the occupation number of the different species is much larger than one, and hence their quantum nature can be ignored~\cite{Khlebnikov:1996mc,Prokopec:1996rr}. Lattice simulations have been, in fact, successfully carried out for different preheating scenarios during the last years, see e.g.~\cite{Allahverdi:2010xz,Amin:2014eta} and references therein. However, each time a new scenario exhibits parametric resonance, a new re-analysis is often required. As lattice simulations are computationally expensive and time consuming, and not everybody has the expertise on the appropriate numerical packages~\cite{Latticeeasy-paper,Defrost-paper,Easther:2010qz,Huang:2011gf,GABElink}, many studies often resort to over-simplified analytical analysis, which capture only the initial linear stage. A systematic study of parametric resonance, fitting the dynamics through all the relevant stages, from the initial linear growth till the relaxation towards equilibrium, passing through an intermediate non-linear stage, is missing in the literature. In this work, we fill in this gap. We have used massively parallelized lattice simulations to charaterize the dynamics of parametric resonance through all its stages. We have parametrized the dynamics by scanning over the relevant circumstances and parameters: role of the oscillating field, particle coupling, initial conditions, and background rate of expansion. We have obtained in this way simple fits to the most significant quantities, like the characteristic time scales and energy fractions of the different particle species. Our fitted formulas can be applied to the study of parametric resonance in scenarios where the mother field dominates the energy budget of the universe (i.e.~preheating), or in scenarios where the mother field represents only a sub-dominant component (e.g.~inflationary spectator fields). As parametric resonance in the context of the early Universe has been well studied in the past, let us emphasize here that with our present work, we simply aim to aid in the analysis of future scenarios exhibiting parametric resonance. The advantage of using our fitted formulas will be twofold: on the one hand skipping the tedious task of running new simulations, and on the other hand preventing the use of over-simplified linear analysis of the problem. The structure of this work is as follows. In Section \ref{sec:AnalyticalParamRes} we describe general aspects of parametric resonance, while we derive an analytical estimate of the decay time of the mother field, based on a linear calculation. In Section~\ref{sec:Lattice} we preset the numerical results from our lattice simulations. We describe our results for preheating with a quartic potential in Section~\ref{sec:lphi4}, and for preheating with a quadratic potential in Section~\ref{sec:m2phi2}. We compare these results against the analytical estimations from Section~\ref{sec:AnalyticalParamRes}. In Section \ref{sec:specfields} we present the analogous numerical study for scenarios where the mother field represents only a sub-dominant energy component of the Universe. In Section \ref{sec:Summary} we list all fitted formulas together from all scenarios considered. In Section~\ref{sec:discussion} we discuss the context where our results can be useful. In the appendices we present details on the lattice formulation we have used, and discuss briefly the evolution of the field spectra in some of the scenarios considered. From now on we consider $\hbar = c = 1$ units, and represent the reduced Planck mass by {\small$m_p^2 = {1/{8\pi G}} \simeq 2.44\cdot10^{18}$ GeV}. We take a flat background with Friedman-Robertson-Walker (FRW) metric {\small$ds^2 = dt^2 - a^2 (t) dx^i dx^i$}, where {\small$a(t)$} is the scale factor, and {\small$t$} the cosmic time. | In this section we just collect together the fitted formulas from all the scenarios considered. In the case of a spectator-field with a quartic potential we just quote the results from~\cite{Figueroa:2015hda}. The interested reader can, in this manner, access rapid and easily to the key results from this paper (complemented with those from~\cite{Figueroa:2015hda}). For self-consistent reading of this Section, let us first summarize the dynamics of parametric resonance, and define the variables to which we provide fits. In parametric resonance with $q > 1$, as soon as the mother field $\phi$ starts oscillating, there is a fast transfer of energy into the daughter species. This occurs independently of whether the mother field dominates or not the energy budget of the universe. During few oscillations, the energy of the daughter fields $X$ remain orders of magnitude smaller than the energy of the mother field. Hence the $\phi$ field oscillates initially almost unaffected by the presence of its decay products. This corresponds to a linear regime, where the mode functions of the daughter field grow exponentially fast in some range of momenta. Due to this exponential excitation, there is always a time for any given resonance parameter, when the energy transferred becomes large enough so that the backreaction from the daughter species into the mother field cannot be further ignored. We refer to this moment as $z_{\rm br}$. From that moment onwards, the (conformal) amplitude of the mother field starts decreasing in a noticeable manner, see Fig.~\ref{fig:lphi4-init} for an example of this. The time scale $z_{\rm br}$ defines therefore the onset of the mother field decay, and not the time scale of the decay itself, as the linear calculation suggests. From then on, at $z \gtrsim z_{\rm br}$, the system becomes non-linear, and evolves towards a stationary state. The latter is characterized by the different energy fractions of the fields (kinetic, gradient and potential energies) evolving very slowly, while at the same time an equipartition distribution of energies is set. This regime is attained at a time $z_{\rm dec}$. We consider this moment as the truly decay time scale of the mother field: while during the non-linear regime $z_{\rm br} \leq z \leq z_{\rm dec}$ energy is significantly exchanged between the mother and the daughter fields, at $z \geq z_{\rm dec}$ the energy exchange ceases and the energy fractions evolve in a stationary regime. In the case of quadratic potentials, the system tends very slowly to restore, at long times $z \gg z_{\rm dec}$, the mother field energy dominance. Hence, we also provide the time scales $z_{0.80}$ and $z_{0.99}$ corresponding to moments when the mother field represents $\sim 80\%$ and $\sim 99\%$ of the energy budget of the mother-daughter system. In the following we summarize our fits for $z_{\rm br}$, $z_{\rm dec}$ (and $z_{0.8}$, $z_{0.99}$ when applicable) as a function of the resonance parameter, for all the scenarios we have analyzed. We take the coupling between the mother and the daughter field to be of the form $g^2 \phi^2 \chi^2$. In the case of preheating, where the mother field -- the inflaton -- dominates the energy budget of the universe, we also provide due to its interest, the stationary energy fractions. \begin{itemize} \item Preheating with inflationary potential $V_{\rm inf} (\phi) = \frac{1}{4} \lambda \phi^4$: \bea z_{\rm br} (q) \in [40, 250 ] &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:lame-zbtime} } \\ \nonumber\\ z_{\rm dec} (q) - z_{\rm br} (q) \simeq \left\lbrace \begin{array}{l} 51 q^{0.28} \hspace{0.3cm}\text{ if } q < 100 \vspace*{2mm}\\ 11 q^{0.56} \hspace{0.3cm}\text{ if } q \geq 100 \\ \end{array} \right. &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:lphi4-zdecay} }\label{eq:ze1} \eea where $q \equiv \frac{g^2}{\lambda}$.\vspace*{-2mm} \begin{center} Energy Fractions at $z \gtrsim z_{\rm dec}$:% \begin{eqnarray}\label{eq:EnergiesPhi4atZeV2} \begin{array}{c} {E_{K,\chi}\over E_t} \simeq (29.5 \pm 3.3) \%\,,\hspace*{0.2cm}{E_{K,\varphi}\over E_t} \simeq (22.6 \pm 3.4) \%\,,\hspace*{0.2cm} {E_{G,\chi}\over E_t} \simeq (26.2 \pm 3.4) \%\,,\hspace*{0.2cm}\vspace*{3mm}\\ {E_{G,\varphi}\over E_t} \simeq (17.7 \pm 3.0) \%\,,\hspace*{0.2cm} {E_{\rm int}\over E_t} \simeq (3.2 \pm 0.7) \%\,,\hspace*{0.2cm}{E_{V}\over E_t} \simeq (0.8 \pm 0.2) \% % \end{array} \end{eqnarray} \end{center} \vspace*{0.5cm} \item Preheating with inflationary potential $V_{\rm inf} (\phi) = \frac{1}{2} m^2 \phi^2$: \bea z_{\rm br} (q) \in [40, 135 ] &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:m2phi2-zbTime} } \\ z_{\rm dec} (q) \simeq 19.9 q_*^{0.27} &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:m2phi2-energy} } \\ z_{0.8} \simeq 0.26\,q_* &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:m2phi2-energy} }\\ z_{0.99} \sim 30 \,q_* &\ ;& \hspace{0.2cm} \text{ (extrapolated) } \label{eq:ze2} \eea where $q_* \equiv \frac{g^2 \phi_*^2}{4 m^2}$, with $\phi_*$ the initial value of the inflaton field.\vspace*{0.0cm}\\ \begin{center} Dominant Energy Fractions at $z \simeq z_{\rm dec}$ ($q_* \gtrsim 5\cdot 10^4$):\vspace*{-0mm} \begin{eqnarray}\label{eq:EnergiesM2phi2atZeV2} \begin{array}{c} {E_{K,\chi}\over E_t} \simeq (25.2 \pm 2.2) \%\,,\hspace*{0.2cm}{E_{K,\varphi}\over E_t} \simeq (26.0 \pm 2.3) \%\,,\hspace*{0.2cm} {E_{G,\chi}\over E_t} \simeq (22.9 \pm 2.5) \%\,,\hspace*{0.2cm} \end{array} \end{eqnarray} \vspace*{0.1cm}\\ Sub-dominant Energy Fractions at $z \simeq z_{\rm dec}$ ($q_* \gtrsim 7.5\cdot 10^3$):\vspace*{-0mm} \begin{eqnarray} \begin{array}{c} {E_{G,\varphi}\over E_t} \simeq {19\over(1+ 30000/q_*)^{1/2}} \% \,,\hspace*{0.2cm}{E_{V}\over E_t} \simeq {27\over(q_*/2000 -1)^{1/3}}\%\,, \hspace*{0.2cm} {E_{\rm int}\over E_t} \simeq (2.3 \pm 0.5) \% \end{array} \end{eqnarray} \end{center} \item Spectator-field with potential $V (\phi) = \frac{1}{2} m^2 \phi^2$ and RD expansion rate: \bea z_{\rm br} (g) \simeq 16.9 - 20.9 \log_{10} g &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:spec-zizetime} } \\ z_{\rm dec} (q) \simeq 27.3 q_*^{0.33} &\ ;& \hspace{0.2cm} \text{ See Fig.~\ref{fig:spec-zizetime2} }\\ z_{0.40} \simeq 0.18\,q_* &\ ;& \hspace{0.2cm} \text{ (measured) }\\ z_{0.99} \sim 8 \cdot 10^{-6} \,q_*^{3} &\ ;& \hspace{0.2cm} \text{ (extrapolated) } \label{eq:ze3} \eea where $q_* \equiv \frac{g^2 \phi_*^2}{4 m^2}$ with $\phi_*$ the initial value of the mother field.\vspace*{0.5cm} \item Spectator-field with potential $V (\phi) = \frac{1}{4} \lambda \phi^4$ (Standard Model Higgs, see \cite{Figueroa:2015hda}\footnote{In the notation of that reference, we use $z_i$ instead of $z_{\rm br}$, and $z_e$ instead of $z_{\rm dec}$.}): \bea z_{\rm br} (q) &\approx & \left\lbrace \begin{array}{l} 16 \beta^{\frac{-(1 + 3 \omega)}{3 (1 + \omega)}}\hspace{2.65cm} \text{ if } q \in \text{Resonance Band} \vspace*{2mm}\\ ( 86.9- 9.2 \log{q} ) \beta^{\frac{-(1 + 3 \omega)}{3 (1 + \omega)}} \hspace{0.31cm}\text{ if } q \notin \text{Resonance Band} \\ \end{array} \right. \\ z_{\rm dec} (q) & \approx & 50.7 \beta^{\frac{-(1 + 3 \omega)}{3 (1 + \omega)}} q^{0.44} \ ,\label{eq:ze4} \eea where $\beta \equiv \frac{\sqrt{\lambda} \phi_*}{H_*}$, $q \equiv \frac{g^2}{\lambda}$, and $\omega$ is the equation of state ($\omega=0, 1/3, 1$, for MD, RD, and KD respectively). \end{itemize} Note that the time scales reported here depend on our choice of å mother-daughter interaction $g^2 \phi^2 X^2$, representing this the only interaction the daughter field experiences. The time scales may change, for instance, in the presence of self-interactions of the $X$ field \cite{Podolsky:2005bw}. Thoughtful analysis of parametric resonance, including analytical calculations of the the Floquet index and analysis of the Floquet theorem, can be found e.g.~in~\cite{Kofman:1997yn,Greene:1997fu,Amin:2014eta}. In this work we rather concentrate in the study of parametric resonance using classical real time field theory lattice simulations. We have simulated an oscillating mother field $\phi$ coupled to a daughter field $X$, which is excited due to an interaction term $g^2\phi^2 X^2$. We have considered two main scenarios. First, when the mother field is the inflaton field, oscillating around the minimum of its potential after inflation. We have considered the case of chaotic inflation with $V \propto \phi^2$ and $V \propto \phi^4$ potentials. In a second type of scenarios, the oscillating field was just a spectator-field during inflation, playing no dynamical role on the expansion of the Universe. We have considered also $V \propto \phi^2$ and $V \propto \phi^4$ potentials, but analyzed only numerically the former, as the latter was already analyzed in~\cite{Figueroa:2015hda}. Our results show very clearly that the computation in the linear regime of the moment of efficient transfer of energy $z_{\rm eff}$, see Eqs.~(\ref{eq:EffEnergyTransferTimeScaleApproxPhi4}), (\ref{eq:EffEnergyTransferTimeScaleApproxPhi2}), does not represent a good estimation of the decay time scale $z_{\rm dec}$ of the mother field. Instead, $z_{\rm eff}$ indicates well (up to $\mathcal{O}(1)$ factors) the onset of the mother field decay at $z_{\rm br}$, when the back-reaction of the daughter field becomes noticeable. Despite the exponential transfer of energy into the daughter fields during the time $z < z_{\rm br}$, the daughter field fluctuations follow a linear equation, whilst the mother field amplitude remains almost unperturbed. At $z \gtrsim z_{\rm br}$, the presence of the excited daughter fields makes the amplitude and energy of the mother field to abruptly decrease. At $z \gtrsim z_{\rm br}$ the dynamics become non-linear, and there is a noticeable transfer of energy between the mother and the daughter fields. Eventually, at $z \gtrsim z_{\rm dec}$ the amplitude of the fields settle down to stationary values, with the energy equiparted among the different components. As for $z \geq z_{\rm dec}$ the dominant energy components do not evolve any more noticeably, we identify the onset of that stationary stage as the truly time scale of the decay of the mother field. In the case of a quadratic potential, at $z \gtrsim z_{\rm dec}$, in reality only the mother field kinetic and potential (conformal) terms remain almost constant, as the (conformal) energy components of the daughter fields decay slowly at long times. The linear calculation of $z_{\rm eff} \sim z_{\rm br}$ indicates that the stronger the coupling between mother and daughter field, the faster the system becomes non-linear. The dependence is however only logarithmic, so in practice $z_{\rm br}$ only changes by a factor $\mathcal{O}(1)$ when varying the strength of the coupling in more than 10 orders of magnitude, see e.g.~Figure~\ref{fig:spec-zizetime}. As the system becomes however non-linear after $z \gtrsim z_{\rm br}$, our numerical results show the rather counter-intuitive result, opposite to the intuition gained from the analytic estimations: the stronger the mother-daughter coupling, the longer the time decay $z_{\rm dec}$ scale is, with a typical power-law behavior with respect the resonance parameter, $z_{\rm dec} \propto q^r$, with $r \sim 1/4, 1/3$ or $1/2$, depending on the case, see Eqs.~(\ref{eq:ze1}), (\ref{eq:ze2}), (\ref{eq:ze3}), (\ref{eq:ze4}). Let us note that we have defined and obtained the decay time scale $z_{\rm dec}$ at the onset of the stationary regime, but we have not analyzed the evolution of the equation of state or the departure from thermal equilibrium. For a study of the subsequent evolution of the system at $z \gtrsim z_{\rm dec}$ towards thermalization, see~\cite{Micha:2002ey,Micha:2004bv,Podolsky:2005bw,Lozanov:2016hid}. We have found nonetheless a remarkable result: in the case of quadratic potentials, the energy components of the daughter field tends to decay at the very late times $z \gg z_{\rm dec}$, so that slowly but monotonically the mother field tends to dominate the energy budget of the mother-daughter system. Let us remark that in this work we have considered the decay products to be scalar fields. However, parametric resonance can also take place for all bosonic species, including gauge fields (either Abelian and non-Abelian). There are many scenarios where the decay products are gauge fields, see e.g.~\cite{GarciaBellido:1999sv,Rajantie:2000nj,GarciaBellido:2003wd,Bezrukov:2008ut,GarciaBellido:2008ab,Bezrukov:2014ipa,Dufaux:2010cf,Deskins:2013lfx,Adshead:2015pva,Figueroa:2015hda,Enqvist:2015sua,Lozanov:2016pac,Figueroa:2016dsc}, although not in all of them the driving particle production mechanism is parametric resonance. As we demonstrated in~\cite{Figueroa:2015hda}, the dynamics of parametric resonance into Abelian gauge fields (at least for a mother field with quartic potential), is only slightly modified in the linear regime, i.e.~$z_{\rm br}$ is only marginally changed. The late time non-linear dynamics remain however basically unchanged. Therefore, in principle, our fitted formulas can be applied equally to the case of parametric resonance of gauge bosons. % In the case of non-abelian gauge fields, the non-linearities in the gauge boson EOM due to the non-abelian structure, may block parametric resonance before reaching $z_{\rm br}$, if the resonance parameter is sufficiently large, see e.g.~\cite{Enqvist:2015sua}. It is well known that violent out-of-equilibrium phenomena like particle production via parametric resonance, can produce scalar metric perturbations~\cite{Bassett:1998wg,Bassett:1999mt,Bassett:1999ta,Finelli:2000ya,Chambers:2007se,Bond:2009xx} and a significant amount of gravitational waves (GW)~\cite{Khlebnikov:1997di,Easther:2006gt,Easther:2006vd,GarciaBellido:2007dg,GarciaBellido:2007af,Dufaux:2007pt,Dufaux:2008dn,Figueroa:2011ye,Bethke:2013aba,Bethke:2013vca,Figueroa:2016ojl,Antusch:2016con}. A natural extension of our present fitting analysis is to parametrize the production of GW from parametric resonance in the early Universe. Although GW production in preheating after chaotic inflation models has been widely considered in the literature, there is still lacking a systematic parametrization of the GW spectrum today as a function of the different couplings\footnote{A parameter-fitting analysis exists however for the GW production from Hybrid preheating, see~\cite{Dufaux:2008dn}.}. We plan to do this in a forthcoming publication. There are some scenarios of preheating where the daughter fields are scalar fields, but the mechanism responsible for the particle production is not parametric resonance, e.g.~hybrid preheating~\cite{Felder:2000hj,Felder:2001kt,Copeland:2002ku,GarciaBellido:2002aj,GarciaBellido:2007dg,GarciaBellido:2007af,Dufaux:2008dn}. Our fitted formulas, unfortunately, cannot be applied to these scenarios. The case of trilinear or non-renormalizable interactions between the mother and the daughter field(s)~\cite{Dufaux:2006ee,Croon:2015naa,Antusch:2015vna,Enqvist:2016mqj} are neither captured by our analysis. The case of oscillations of a multi-component field is neither captured well by our analysis\footnote{In the case of super-symmetric flat directions, it may well happen that the flat directions are never really excited in first place~\cite{Enqvist:2011pt}, and therefore it makes no sense to speak about oscillations after inflation.}, see e.g.~\cite{Tkachev:1998dc,Olive:2006uw,Gumrukcuoglu:2008fk,DeCross:2015uza,Ballesteros:2016euj}. We speculate nonetheless, that the non-linear dynamics after the initial excitation in all these scenarios, is probably very similar to the one after parametric resonance. However, only proper lattice simulations can prove this. As a final thought, it is interesting to note that, in the case of an inflaton with quartic potential, our results may challenge somehow the application of the standard perturbative calculation of the reheating temperature $T_{\rm RH} \sim 0.1\sqrt{\Gamma m_p}$, where $\Gamma$ is the inflaton decay width. It is often argued that, as preheating does not deplete completely the energy from the inflaton, % reheating will only be completed when the inflaton decays perturbatively into other matter fields. Our simulations for a potential $V \propto \phi^4$ show however, that at the onset of the stationary regime at $z \simeq z_{\rm dec}$, the energy fractions do not evolve significantly anymore, and the inflaton energy never represents more than $50$\% of the total energy budget. Therefore, even if eventually the inflaton decays perturbatively into some species, the originally produced daughter field from parametric resonance (which also represent $50\%$ of the energy budget), may have already thermalized and reheated the universe. As we expect equipartition in the different field components at the onset of the stationary regime, if there were several daughter fields experiencing parametric resonance (and not just one like in our simulations), in principle the fraction of energy stored in the inflaton at the stationary regime, should be approximately suppressed by the total number of species (i.e.~the number of daughter fields plus one inflaton). In that case, whether the inflaton decays perturbatively later or not, should be mostly irrelevant, since by then most of the energy is stored in the parametrically excited daughter fields, which may very well have thermalized before. In the case of a quadratic potential, our results tend however to reinforce the opposite circumstance, as the system approaches at long times $z \gg z_{\rm dec}$, slowly but monotonically, towards a complete energy dominance of the mother field over the daughter field(s). This reinforces the idea that for a quadratic potential, a perturbative decay (or further interactions besides $g^2\phi^2X^2$) are necessary ingredients in order for the mother field to decay at all. | 16 | 9 | 1609.05197 |
1609 | 1609.01850_arXiv.txt | Thermal Sunyaev-Zeldovich (tSZ) power spectrum is one of the most sensitive methods to constrain cosmological parameters, scaling as the amplitude $\sigma_8^8$. It is determined by the integral over the halo mass function multiplied by the total pressure content of clusters, and further convolved by the cluster gas pressure profile. It has been shown that various feedback effects can change significantly the pressure profile, possibly even pushing the gas out to the virial radius and beyond, strongly affecting the tSZ power spectrum at high $l$. Energetics arguments and SZ-halo mass scaling relations suggest feedback is unlikely to significantly change the total pressure content, making low $l$ tSZ power spectrum more robust against feedback effects. Furthermore, the separation between the cosmic infrared background (CIB) and tSZ is more reliable at low $l$. Low $l$ modes are however probing very small volumes, giving rise to very large non-gaussian sampling variance errors. By computing the trispectrum contribution we identify $90<l<350$ as the minimum variance scale where the combined error is minimized. We find constraints on $\sigma_8$ by marginalizing over the feedback nuisance parameter, obtaining $\sigma_8 =0.820^{+0.021}_{-0.009}\left(\Omega_m/0.31\right)^{0.4}$ when fixing other parameters to Planck cosmology values. Our results suggest it is possible to obtain competitive cosmological constraints from tSZ without cluster redshift information, and that the current tSZ power spectrum shows no evidence for a low amplitude of $\sigma_8$. | The thermal Sunyaev Zeldovich (tSZ) effect is a secondary anisotropy of the cosmic microwave background (CMB), where CMB photons inverse Compton scatter off of energetic electrons that lie along the line of site between us and the surface of last scattering \cite[]{sunyaev1970small}. Its amplitude $Y$ is determined by the projected gas pressure along the line of sight. The tSZ effect has been used for over a decade to study individual clusters ~\cite[]{reese2002determining}. The pressure can be expressed as a product of density and temperature, and in a virialized system the latter scales roughly as $M_{\rm vir}^{2/3}$, where $M_{\rm vir}$ is the halo virial mass. Integrating the tSZ signal across the cluster gives the scaling of $Y_{\rm vir} \propto M_{\rm vir}^{5/3}$. In recent years the multi wavelength, high angular resolution, large-array surveys has allowed the measurement of its power spectrum over a large range of scales ~\cite[]{aghanim2015planck}~\cite[]{story2013measurement}. The tSZ power spectrum has been advocated as a strong probe of cosmology, ranging from constraining $\Lambda$CDM cosmological parameters ~\cite[]{komatsu2002sunyaev}, to primordial non-gaussianities and massive neutrinos ~\cite[]{hill2013cosmology}. Its main advantage when compared to cluster abundance method is that one does not need to measure the cluster redshifts, or their virial halo mass. Instead, tSZ power spectrum probes an integral over the cluster halo abundance as a function of redshift and halo mass, multiplied by the total pressure content of clusters, and further convolved by the cluster gas pressure profile. tSZ power spectrum is sensitive to different halo masses and redshifts as a function of angular moment $l$. However, it is still a projection and as such it is difficult to disentangle the different redshifts and/or halo masses. The tSZ power spectrum is sensitive not only to cosmological parameters, but also to nonlinear gas physics found in the intra-cluster medium ~\cite[]{komatsu1999sunyaev, battaglia2010simulations}. The gas distribution depends on the dark matter gravitational potential well, stellar formation, AGN feedback, supernovae, and radiative cooling. In practice, simulated pressure profiles in different simulations often differ with one another, and the resulting power spectra differ as well, specially at high $l$ where the change of profile matters more \cite[]{mccarthy2014thermal}. Direct observation of the pressure profiles is a promising approach ~\cite[]{arnaud2010universal}, but is limited to observed massive halos where the effects of AGN and other feedback mechanism are less pronounced. In this paper we revisit tSZ power spectrum as a probe of cosmology. We take advantage of the fact that while the gas pressure profiles are very dependent on the detailed physical modeling inside the clusters, the total pressure content integrated over the cluster is a lot less model dependent. This is because while the feedback models can push the gas around, they cannot easily inject enough energy to change its total thermal content. This is confirmed in tSZ power spectrum simulations \cite[]{mccarthy2014thermal}, which are relatively unaffected at low $l$. Another argument are the scaling relations of Planck \cite[]{2013A&A...557A..52P} and by Greco et al.\cite[]{greco2015stacked}, where the simple $Y_{\rm vir} \propto M_{\rm vir}^{5/3}$ scaling holds over a large range of halo masses. A weak lensing calibration of this scaling relation has recently been provided in \cite[]{wang2016weak}. The scaling works because the Planck beam is very large and encompasses the entire cluster tSZ effect, but it also suggests that there is no significant change in the thermal gas content that would break the scaling. We can thus side-step the complex gas dynamics by focusing on large scales where the gas profile is less important, and only the total pressure cluster content contributes. In this paper we will use a simple one-parameter model for AGN and supernova feedback, similar to the model used by for CMB weak lensing statistics~\cite[]{mohammed2014baryonic}. Various simulations show that gas will be expelled in less massive halos ~\cite[]{read2005mass},~\cite[]{pontzen2012supernova}, and forced to outer reaches beyond the viral radius. This parameter has the effect of suppressing the high $l$ contribution from clusters below a given critical mass ($M_{crit}$). In effect, the galaxy is ``puffed up'' due to the feedback effects. However, as stated above, in our model we preserve the total gas pressure content, so at sufficiently large scales (low $l$) there is no effect. In section II we quickly overview the tSZ power spectrum calculations, its scaling relations, its covariance matrix model, and the feedback modeling. In section III we perform a likihood analysis over the $\sigma_8$, $M_\textrm{crit}$ parameter space and marginalize over $M_\textrm{crit}$ to find constraints on $\sigma_8$. In section IV we compare our constraints with other techniques and discuss the outlook for further method improvements. For all our analysis, we use the Planck TT,TE,EE+lowP+lensing+BAO fiducial cosmology for all our calculations; $\Omega_m = 0.3089$, $\Omega_\Lambda = 0.6911$, and $h = 67.74{\rm km/s/Mpc}$. We use for data the NILC - MILCA F/L cross-power spectrum after foreground subtraction as described in ~\cite{aghanim2015planck} and the ACT value at high $l$ from ~\cite{hill2014atacama}. \label{sec:intro} | \label{sec:discussion} In this paper we present a tSZ power spectrum cosmology analysis, where we extend previous work by accounting for feedback in clusters and including connected trispectrum in the analysis. We introduce a one-parameter model for feedback within galactic clusters which, when marginalized over, provides more realistic constraints on cosmological parameters. Our model is relatively simple, but also does not rely on any detailed understanding of gas dynamics or simulated results. A key assumption of our model is that while astrophysical processes within the cluster can push the gas around, possibly all the way to the virial radius, suppressing small scale clustering, its total thermal content does not change, guaranteeing that large scale clustering is unchanged. Maximizing likelihood with respect to this model we find we an updated constraint for $\sigma_8 =0.820^{+0.021}_{-0.009}\left(\Omega_m/0.31\right)^{0.4}$ when using our combined Planck/ACT data set, assuming Planck cosmology values for other parameters. Our results are consistent with Planck's overall normalization of $\sigma_8^{\textrm{Planck}}= 0.8159 \pm 0.0086$~\cite[]{planck2015cosmo}. Planck SZ analysis has been argued to be supporting low amplitude, $\sigma_8 = 0.78 \pm 0.02$ ~\cite[]{aghanim2015planck}, but we find no evidence of this in our tSZ power spectrum analysis. If anything, our results suggest a normalization that is even higher than that of Planck cosmology, which in itself is relatively high. Our analysis differs from that done in Planck in two substantial ways; we have a full treatment of the trispectrum which substantially reduces the weight of low-$l$ data which would otherwise support a lower value of $\sigma_8$ and the use of marginalization over feedback parameters which effectively reduces the weight of high-$l$ data. Our normalization is also somewhat higher than the values given by ACT-SZ analysis $\sigma_8^{\textrm{ACT-SZ}} = 0.793 \pm 0.018$~\cite[]{hill2014atacama}. Our results also differ from those found in \cite[]{hill2014detection} primarily due to a substantial amplitude shift between Planck 2013 and Planck 2015 tSZ Compton parameter map. It should be noted that we do not include data from the South Pole Telescope (SPT)\cite[]{george2015measurement}. There has been discussion in the literature about possible challenges facing this measurement due to the limited frequency channels compared to Planck, creating difficulty distinguishing tSZ signal from the primary CMB and other extra-galactic sources. This difficultly is further compounded by the fact that at $l=3000$ the tSZ signal is sub dominant to the foregrounds\cite[]{dolag2015sz}. While these same challenges face ACT's tSZ analysis, we use ACT's value due to their higher error bars which we view as more reflective of the current uncertainty in this difficult measurement. Like previous work, our model has more small scale power than predicted by SPT (or ACT) and it is likely that additional data at $l>2000$ will help resolve this tension. We have identified $100<l<350$ as the sweet spot where the current errors and modeling uncertainties are minimized; at $l<100$ the amplitude has significant trispectrum error and at $l>350$ the amplitude has significant error from AGN modeling and foregrounds. CIB contamination is one of the dominant sources of error, and if the foreground separation can be improved one may be able to push the analysis to higher $l$, where connected trispectrum term become less important. However, pushing to higher $l$ would also require improved modeling of feedback effects on the pressure profile, which is not necessary for low $l$ used in this paper. Similarly, one might be able to expand the "sweet spot" to lower $l$ by masking nearby clusters. This has been shown (\cite[]{shaw2009sharpening}, \cite[]{hill2013cosmology}) as a possible technique to reduce the trispectrum error, since large nearby clusters contribute significantly to the signal at low $l$, but contain little volume and hence are subject to significant sampling variance and Poisson fluctuations. Going forward it will be useful to further extend our model based on a more nuanced understanding of the gas physics involved. Our model is fairly simplistic in how we treat the expelled gas pressure distribution. It is possible that with new tSZ, kSZ, and X-ray studies understanding the pressure profile at the outer edges of clusters would be improved. Similarly, our flat prior for $\log{M_{crit}}$ is simple and better independent constraints of this parameter would allow similarly better constraints of $\sigma_8$. As new generations of CMB experiments start collecting data at higher spatial resolutions, one would like to be able to further improve cosmological constraints from the tSZ power spectrum. Our analysis suggests this may in principle be possible, but it will not be easy, and will require improvements in our understanding of the cluster pressure profiles on scales at and beyond the virial radius. | 16 | 9 | 1609.01850 |
1609 | 1609.03587_arXiv.txt | We present \INSERTNIGHTSSIMP \ nights of ground-based, near-infrared photometry of the variable L/T transition brown dwarf SIMP J013656.5+093347 and an additional \INSERTNIGHTSTVLM \ nights of ground-based photometry of the radio-active late M-dwarf TVLM 513-46546. Our TVLM 513-46546 photometry includes \INSERTNIGHTSSIMULTANEOUSTVLM \ nights of simultaneous, multiwavelength, ground-based photometry, in which we detect obvious J-band variability, but do not detect I-band variability of similar amplitude, confirming that the variability of TVLM 513-46546 most likely arises from clouds or aurorae, rather than starspots. Our photometry of SIMP J013656.5+093347 includes \INSERTNIGHTSSIMPJBAND \ nights of J-band photometry that allow us to observe how the variable light curve of this L/T transition brown dwarf evolves from rotation period to rotation period, night-to-night and week-to-week. We estimate the rotation period of SIMP J013656.5+093347 as \WeightedMeanSIMPPeriod \ $\pm$ \WeightedErrorSIMPPeriodScaled \ hours, and do not find evidence for obvious differential rotation. The peak-to-peak amplitude displayed by SIMP J013656.5+093347 in our light curves evolves from greater than \SIMPMaximumRot\% to less than \SIMPMinimumRot\% in a matter of days, and the typical timescale for significant evolution of the SIMP J013656.5+093347 light curve appears to be approximately $<$1 to 10 rotation periods. This suggests that those performing spectrophotometric observations of brown dwarfs should be cautious in their interpretations comparing the spectra between a variable brown dwarf's maximum flux and minimum flux from observations lasting only approximately a rotation period, as these comparisons may depict the spectral characteristics of a single, ephemeral snapshot, rather than the full range of characteristics. | Over the last number of years simultaneous, multiwavelength photometric and spectrophotometric observations of variable brown dwarfs at the L/T transition have enabled a number of detailed investigations attempting to elucidate the cause of the observed variability (e.g. \citealt{Artigau09,Radigan12,Buenzli12,Apai13,Buenzli14,Buenzli15,Yang15,Yang16}). The result of these investigations has been increased understanding of the role that fluctuations in cloud condensate opacity play in the observed variability. Upon the first detections of significant variability at the L/T transition (e.g. \citealt{Artigau09,Radigan12}) the belief was that the observed fluctuations were caused by gaps or holes in thick clouds exposing hotter regions underneath, rotating in and out of view \citep{AckermanMarley01,Burgasser02,Artigau09}. More recent simultaneous multiwavelength photometry, or spectrophotometry -- especially driven by the impressive, spectrophotometric capabilities of the WFC3 instrument \citep{MacKenty10} on the Hubble Space Telescope ({\it HST}) -- have revealed the possibility that these variations were more specifically caused by thick or thin clouds \citep{Radigan12,Buenzli15}, rather than outright cloud-free regions, or caused by varying haze layers \citep{Yang15}. However, a number of these detailed studies only observed these variable brown dwarfs for tens of minutes (e.g. \citealt{Buenzli14}), to hours at a time (e.g. \citealt{Apai13,Buenzli15,Yang15}). This is in spite of the fact that, since the first detections of large amplitude variability in L/T transition brown dwarfs, the photometric light curves of some of these objects have been shown to display significant evolution in their light curves (e.g. \citealt{Artigau09,Radigan12,Gillon13}). Thus, the conclusions these detailed, spectrophotometric studies reach are based inherently on an ephemeral snapshot, that might not be indicative of the usual behaviour of the L/T dwarf, or might not cover the extreme range of variability behaviour displayed by the ultra-cool dwarf. To determine whether these detailed, spectrophotometric observations capture the range of characteristics of ultra-cool dwarfs, longer-term light curves of these variable ultra-cool dwarfs are required. To date, a handful of nights of photometry have captured the evolution of the light curves of the L/T transition brown dwarf SIMP J013656.5+093347 (e.g. \citealt{Artigau09,Metchev13,Radigan14}), and the significant night-to-night evolving variability of the L/T transition Luhman-16 binary system \citep{Gillon13}. {\it Kepler} has now returned photometry of two early L dwarfs allowing the evolution, or lack thereof, of their optical light curves to be tracked for two years for an L0 dwarf (WISEP J190648.47+401106.8; \citealt{Gizis13,Gizis15}) and $\sim$36 days for an L8 dwarf (WISEP J060738.65+242953.4; \citealt{Gizis16}). Relatively constant variability was detected for the L0 dwarf WISEP J190648.47+401106.8 \citep{Gizis13,Gizis15}, while for the L8 dwarf WISEP J060738.65+242953.4 a lack of variability was detected, arguably due to a pole-on inclination of the ultra-cool dwarf \citep{Gizis16}. What is needed to advance this science is many more long term light curves of ultra-cool dwarfs that accurately track the evolution, or lack thereof, of the light curves of these objects from rotation period to rotation period, night-to-night, week-to-week, season-to-season and even year-to-year. In addition to L/T transition objects, the cause of variability of M/L transition objects has also attracted growing interest recently (e.g. \citealt{Gelino02,Littlefair08,Harding13,Metchev15,Hallinan15}). Starspots are the usual explanation for variability of ultra-cool dwarfs on the stellar side of the hydrogen-fusing limit. M-dwarfs are notoriously active -- nearly all very late M-dwarfs are active, with detections of H$\alpha$, a common activity marker, rising throughout the M-spectral class (\citealt{West04}; \citealt{Schmidt15}). Rotation periods for M-dwarfs, revealed by Doppler imaging techniques (\citealt{BarnesCollier01} ; \citealt{Barnes04}), and by photometry (\citealt{Rockenfeller06}; \citealt{Irwin11}), have indicated that starspots are pervasive on M-dwarfs. Recently, clouds have been found to play a role not only on L/T transition objects, but possibly throughout the whole L spectral class \citep{Metchev15}. Therefore, it is possible that clouds may even lead to some of the observed variability for very late M-dwarfs. The detection of anti-correlated light curves on the M9 dwarf TVLM 513-46546 was at first inferred to be due to the presence of clouds even on this late M-dwarf \citep{Littlefair08}. More recently, in addition to clouds, aurorae -- similar to the planetary aurorae in our solar system (e.g. \citealt{Clarke80}) -- have been suggested to be another possibility to explain the observed variability of ultra-cool dwarfs. Such auroral activity has been observed on an M8.5 dwarf at the end of the main sequence \citep{Hallinan15}. Multiwavelength photometry of this dwarf displays light curves that are anti-correlated in phase, and \citet{Hallinan15} speculate that the aforementioned anti-correlated light curves of TVLM 513-46546 \citep{Littlefair08}, and possibly other late M-dwarfs, may result from auroral activity. In addition, auroral activity may not simply be constrained to the M/L transition, as recent radio detections of polarized, pulsed emissions from a T2.5 dwarf \citep{Kao16} and a T6.5 dwarf \citep{RouteWolszczan12,WilliamsBerger15} indicate this phenomenon may extend even to the L/T transition. Therefore, an additional possibility is that the variability at the L/T transition is caused in part or wholly by auroral activity, rather than the commonly accepted fluctuations in cloud condensate opacity. Furthermore, it is also possible, or arguably likely, that the variability of a single ultra-cool dwarf may be driven by more than a single one of the previously mentioned astrophysical variability mechanisms. Magnetically driven cool or hot starspots may be periodically obscured by time evolving clouds (\citealt{Lane07}; \citealt{Heinze13}; \citealt{Metchev15}), or the variability of predominantly cloudy brown dwarfs may be affected by the presence of occasional aurorae \citep{Hallinan15}. Longer term light curves of ultra-cool dwarfs can also inform our understanding of whether one or more of these physical mechanisms play a role in the observed variability on a single ultra-cool dwarf. To illuminate these various questions here we attempt to observe how the variability of two well known variable ultra-cool dwarfs evolves from rotation period to rotation period, night-to-night and week-to-week. In Section \ref{SecUCDTargets} we provide an overview of our two targets: the variable T2.5 brown dwarf SIMP J013656.5+093347, and the variable M9 dwarf TVLM 513-46546. In Section \ref{SecObs} we present \INSERTNIGHTSSIMP \ nights of photometry of SIMP J013656.5+093347, and \INSERTNIGHTSTVLM \ nights of photometry of TVLM 513-46546, including \INSERTNIGHTSSIMULTANEOUSTVLM \ nights of simultaneous multiwavelength photometry of TVLM 513-46546. In Section \ref{SecAnalysisSIMP} we analyze the rapidly evolving light curves of SIMP J013656.5+093347; for SIMP J013656.5+093347 the peak-to-peak amplitude of variability evolves from a minimum of less than \SIMPMinimumRot\% to a maximum of greater than \SIMPMaximumRot\% in just a handful of nights. This suggests that those performing detailed comparisons of the spectra of L/T transition brown dwarfs between the maximum flux and minimum flux displayed in a observation lasting only a single rotation period, should be aware that this comparison may reveal very different conclusions if the variability amplitude is small (less than $\sim$\SIMPMinimumRot\% peak-to-peak amplitudes), compared to when the variability is considerably greater. In Section \ref{SecAnalysisTVLM} we show that the fact we detect near-infrared variability without detectable accompanying optical variability from our multiwavelength photometry confirms that the variability of TVLM 513-46546 most likely arises from clouds or aurorae. | We have returned long-term, ground-based photometry of two ultra-cool dwarfs: the T2.5 L/T transition brown dwarf SIMP 0136 and the M9 radio-active, ultra-cool dwarf TVLM 513. Our \INSERTNIGHTSSIMP \ nights of ground-based, near-infrared photometry of SIMP 0136, including \INSERTNIGHTSSIMPJBAND \ nights of J-band photometric monitoring spread out over \INSERTSPREADDAYSSIMPJBAND\ days, allow us to observe how the variability of SIMP 0136 evolves from rotation period to rotation period, night-to-night, \& week to week. We estimate the rotation period of SIMP 0136 as \WeightedMeanSIMPPeriod \ $\pm$ \WeightedErrorSIMPPeriodScaled \ hours. We do not detect any significant deviations in the photometric period from night-to-night, and therefore constrain the frequency and magnitude of obvious differential rotation; if differential rotation is present on SIMP 0136 and drives the variability with a different period for an entire night during our observations, then the period must deviate by less than $\sim$\DiffRotationPercentage\% from our estimate of the rotation period of SIMP 0136. The peak-to-peak amplitude displayed by SIMP 0136 in our J-band light curves evolves from greater than \SIMPMaximumRot\% to less than \SIMPMinimumRot\% in the space of just a few days. Longer-term spectrophotomeric studies, and/or Doppler imaging studies, of individual L/T transition brown dwarfs are crucial in determining whether the rapid evolution of the variability that we observe from rotation period to rotation period is due to the size of thinner-cloud (or cloud-free) regions growing or shrinking, or that the characteristics of the clouds (e.g. the thickness, height, or type of clouds) in thinner-cloud regions of approximately constant size are changing. Although spectrophotomeric studies to date have suggested that that the variability observed in L/T transition brown dwarfs over approximately a single rotation period are well modeled by only two surfaces consisting of thick clouds, and warmer, thinner clouds \citep{Apai13,Buenzli15}, our observations indicate that this comparison may only be representative of the spectral characteristics at a single ephemeral snapshot in time, and may not represent the full continuum of spectral characteristics. The timescale for significant evolution of the SIMP 0136 light curve appears to be as short as a rotation period, and as long as approximately a day. This is arguably the first time that the timescale for the evolution of the light curve of a L/T transition brown dwarf has been robustly measured. With the possible exception of Luhman-16, for other highly variable L/T transition brown dwarfs there have yet to be published a sufficient number of multi-night light curves in a single observing band for the timescale of evolution to be robustly measured (e.g. \citealt{Radigan12} for 2MASS J21392676+0220226). As for Luhman-16, \citet{Gillon13} published observations suggesting strong night-to-night variations, but it is less clear if the binary exhibits significant rotation period to rotation period variations. Nonetheless, determining a precise timescale for the light curve evolution of a single member of the Luhman-16 binary brown dwarf system is challenging due to the fact that both L/T transition binary members are variable \citep{Biller13,Buenzli15b}. However, several studies have indicated that the majority of the variability from the Luhman 16 system originates from Luhman 16B \citep{Gillon13,Burgasser14,Buenzli15}; if this is the case, then the significant correlation that we observe of the peak-to-peak amplitudes from one rotation to the next for SIMP 0136 suggests that the timescale for significant evolution of the light curve might be longer for SIMP 0136 than for Luhman-16B. For the radio-active, ultra-cool dwarf TVLM 513, our \INSERTNIGHTSSIMULTANEOUSTVLM \ nights of simultaneous, ground-based photometry display obvious J-band variability, without accompanying obvious I-band variability of similar amplitude. This confirms that the variability of TVLM 513 likely arises from clouds or aurorae, rather than starspots. We encourage further monitoring of these two intriguing ultra-cool dwarfs, and other ultra-cool dwarfs, to better constrain the long-term evolution, or lack thereof, of their variability. | 16 | 9 | 1609.03587 |
1609 | 1609.03112.txt | The interstellar medium is characterized by a rich and diverse chemistry. Many of its complex organic molecules are proposed to form through radical chemistry in icy grain mantles. Radicals form readily when interstellar ices (composed of water and other volatiles) are exposed to UV photons and other sources of dissociative radiation, and if sufficiently mobile the radicals can react to form larger, more complex molecules. The resulting complex organic molecules (COMs) accompany star and planet formation, and may eventually seed the origins of life on nascent planets. Experiments of increasing sophistication have demonstrated that known interstellar COMs as well as the prebiotically interesting amino acids can form through ice photochemistry. We review these experiments and discuss the qualitative and quantitative kinetic and mechanistic constraints they have provided. We finally compare the effects of UV radiation with those of three other potential sources of radical production and chemistry in interstellar ices: electrons, ions and X-rays. | The interstellar medium (ISM), the medium between stars, is host to a rich, organic chemistry. The presence of large molecules may be a surprise considering that the ISM is permeated by high-energy radiation fields including UV radiation. The interstellar UV field has been known or inferred to be present since the dawn of modern astronomy. The potentially destructive power of UV radiation on molecules was a major argument against the presence of an interstellar chemistry. Today we know that some parts of the interstellar medium, so called molecular clouds, are shielded from the onslaught of external UV irradiation by interstellar dust particles that efficiently absorb UV and visible radiation. These cold ($\sim$10~K) clouds are rich in molecules, most of which can be explained by gas-phase chemical reactions between ions and molecules. This gas-phase chemistry cannot explain the presence of some of the most complex organic molecules found in the ISM, however, i.e. molecules such as ethylene glycol and ethanol. These highly saturated organics are instead proposed to originate from radical chemistry in icy mantles embedding the interstellar grains that reside in clouds. Radicals form in ices through either successive atom-addition reactions or through dissociation of pre-existing molecules. Radical reactions generally lack energy barriers and can thus proceed at the low temperatures characteristic of interstellar ices (10--100~K), as long as the radical reactants either form on neighboring sites or can diffuse through the ice. Despite the absence of external UV photons, there are several potential sources of dissociative radiation in clouds. Cosmic ray interactions with molecular hydrogen and grain mantles cause secondary UV photons and electrons. Close to young stars there may also be substantial X-ray radiation. The same kind of radiation that precludes molecular survival exterior to clouds is thus in small doses a source of chemical complexity in cloud interiors. The relative importance of different radiation sources for ice chemistry is contested and may vary between different interstellar environments. The primary aim of this review is to describe and evaluate the constraints that laboratory experiments have provided on ice photochemistry in astrophysical environments. In the final section these results are then compared with ice experiments that use electrons, ions and X-rays as an energy source. \subsection{Astrochemistry in the Context of Star and Planet Formation} \begin{figure}[htbp] \begin{center} \vspace{-5mm} \includegraphics[width=5in]{fig1.pdf} \vspace{-5mm} \caption{Solar-type star formation begins with collapse of a core in a dense molecular cloud\cite{Shu77}. During collapse a disk and outflows form to mediate transport of angular momentum. Once collapse has proceeded far enough for the core to heat up a protostar is formed. The remaining cloud envelope accretes onto the star or disperses leaving a pre-main sequence star with a disk. The disk is the formation site of planets, which must form on time scales of 1-10~Myrs, the life time of such disks. } \label{fig:sf} \end{center} \end{figure} The context of this review is the interstellar and circumstellar media from which stars and planets form. This medium consist of $\sim$99\% gas and 1\% grains, by mass. The gas is mainly atomic or molecular hydrogen with trace amounts of heavier elements. Figure \ref{fig:sf} shows the different phases of star and planet formation for an isolated Solar-type star. In the Galaxy, most of the space between stars is characterized by a diffuse mixture of gas (mainly atomic and molecular hydrogen) and small, sub-micron-sized dust grains. The gas and grains are almost completely exposed to the interstellar radiation field (ISRF), which has a substantial far-ultra-violet (FUV) or vacuum-ultra-violet (VUV) component that readily dissociates most molecules (both FUV and VUV emission is defined as UV emission with shorter wavelengths than 200 nm). The diffuse ISM is therefore characterized by atoms and ions, though low abundances of molecules survive\cite{Snow06}. Over-densities in the ISM are referred to as clouds and the beginning of star formation is the assembly of a dense molecular cloud with even denser cloud cores. The edges of the cloud are still exposed to the ISRF as well as to radiation of any nearby stars, which results in so called photon dominated regions (PDRs) where the combination of high UV fields and dense, relatively cold gas and dust result in a characteristic chemistry\cite{Hollenbach97}. The very edges of such PDRs present, similarly to the diffuse ISM, mainly atoms, ions and radicals, but gas-phase chemistry quickly reforms simple molecules behind the UV front. Some PDRs display a rich molecular chemistry, which may be the result of UV-mediated destruction of macromolecules and volatile grain mantles\cite{Guzman14,Guzman15}. Molecular cloud interiors are protected from external UV radiation, but not from cosmic rays. Cosmic rays interact with hydrogen to produce secondary UV radiation and electrons, and are responsible for partially ionizing the cloud gas. The energy input from the cosmic rays is relatively low, which in combination with efficient molecular cooling results in low ($\sim$10~K) temperatures. Densities are high compared to the diffuse interstellar medium (the hydrogen density n$_{\rm H}\sim$10$^4$ cm$^{-3}$), but very low compared to terrestrial conditions. Under these conditions, ion-molecule reactions in the gas-phase drive an efficient chemistry resulting in e.g. formation of large carbon chains, as well as many smaller molecules\cite{Bergin07}. Most gas-phase molecules observed in clouds can be explained by this ion-molecule chemistry\cite{Herbst73}. Dense molecular clouds are also characterized by large volatile depletions from the gas-phase due condensation or freeze-out of molecules onto interstellar dust grains. At grain temperatures of $\sim$10~K all molecules but H$_2$ and He readily stick onto grains upon collision and the freeze-out time scales are set by collision time scales. Freeze-out combined with atom reactions on the grain surfaces produce large quantities of small hydrogenated species such as water, which remain on the grains forming an icy grain mantle. If a core within the dense cloud becomes sufficiently massive it begins to collapse due to self-gravity\cite{Shu77}. Initially the core is kept cool by molecules radiating away the heat produced by the collapse. As the collapse proceeds the core becomes optically thick, however, and begins to warm up. Eventually it becomes warm enough for deuterium fusion and a protostar has formed. Protostars remain enveloped in molecular material for some time and are also characterized by the presence of a circumstellar disk and/or outflows. The outflows help transport away angular momenta. The disk serves a similar purpose, but also aids in accretion of additional material onto the star. Chemically this stage is characterized by sublimation of ice mantles, and ensuing gas-phase chemistry, as grains flow toward the protostar\cite{Herbst09}. Some of the cloud chemistry products probably survives incorporation into the accretion disk\cite{Visser09}. Within a few hundred thousand years, the initial envelope of cloud material either accretes or disperses and what remains is a pre-main sequence star and a circumstellar disk. The disk is the formation site of planets and is generally referred to as a protoplanetary disk. The disk structure and composition regulates the formation of planets and both are topics of very active research (see Williams \& Cieza (2011)\cite{Williams11} for a review). The chemical structure of protoplanetary disks is less well known compared to protostellar envelopes and clouds because of small angular sizes and intrinsically weak emission. The remnant of our own protoplanetary disk, comets, display a rich organic chemistry\cite{Mumma11,LeRoy15,Wright15,Goesmann15} and recent observations suggest that protoplanetary disks can host a complex organic chemistry as well\cite{Oberg15}. This is important when considering the connection between astrochemistry and prebiotic chemistry, since it is the disk material that eventually becomes incorporated into planets and comets, and thus set the prebiotic potential of nascent planets. In following sub-sections we review in some further depth the aspects of astrochemistry that directly relates to the photochemical formation of complex organics. \subsection{Ice Formation and Demographics in the ISM} Icy grain mantles are the starting point of complex organic molecule formation in the interstellar medium and therefore worth considering in some detail. In clouds, protostellar envelopes and protoplanetary disks, such icy mantles constitute the main reservoir of volatiles. Ices form by a condensation of atoms and molecules from the gas-phase and subsequent grain surface chemistry\cite{Tielens82}.. The latter mostly involves hydrogen atoms, because at grain temperatures of 10~K (typical of molecular clouds), H atoms are much more mobile than any other ice constituent. While the dense cloud is forming out of the diffuse media, atoms dominate the gas-phase. The most abundant atoms are H, He (which does not partake in chemistry), O, C and N. O, C, N and H accrete onto grains and H$_2$O, CH$_4$ and NH$_3$ form by hydrogenation. Some C and O are converted to CO in the gas-phase relatively quickly and then condense out. As long as O condensation is also taking place, OH will be abundant on the grain surface and condensed CO can react with OH to form CO$_2$\cite{Ioppolo11}. As the gas-phase becomes dominated by molecules, CO becomes the most abundant species to accrete onto grains. Some CO will still react to form CO$_2$, while other CO molecules will react with H to form H$_2$CO and CH$_3$OH. \begin{figure}[htbp] \begin{center} \includegraphics[width=5in]{fig2.pdf} \caption{The median composition of interstellar icy grain mantles (bars) normalized to the most abundant ice species, water\cite{Boogert15}. The superimposed error bars indicate the minimum and maximum abundance detected of different species, relative to water ice in that line of sight.} \label{fig:ice} \end{center} \end{figure} Interstellar ice compositions are primarily studied through absorption infrared spectroscopy toward a background source such as a protostar\cite{Gibb04,Oberg11c,Boogert15}. Based on astronomical observations, ice mantle compositions can vary substantially between different clouds, and also between different protostellar envelopes within the same dense cloud. There are some general trends, however. First, water ice is the most abundant ice constituent, followed by CO and CO$_2$ at $\sim$20--30\% each with respect to water. CH$_4$, NH$_3$ and CH$_3$OH are also detected in many lines of sight at typical abundances of $\sim$5\% with respect to water ice. Figure \ref{fig:ice} shows the median ice abundances toward Solar-type protostars, as well as the minimum and maximum abundance with respect to water ice observed for each of the other five species. Analysis of observed ice spectral profile reveal that interstellar ices are not perfectly mixed, but rather present in at least two different phases, a water-rich phase and water-poor-phase\cite{Pontoppidan03,Pontoppidan04}. The water-rich phase is observed to form first and contains most of the water and CO$_2$, and probably most NH$_3$ and CH$_4$ ice as well. A second CO-rich ice contains the remaining CO$_2$ ice and probably most of the CH$_3$OH (Fig. \ref{fig:ice})\cite{Penteado15}. Theoretically these two separate phases form because ice formation is dominated by hydrogenation of atoms at early times, and by reactions involving CO at late times\cite{Garrod11}. It is important to note that none of the ice experiments reviewed below mimic the observed average ice composition and morphology perfectly. The reasons for differences between experiment and observations vary. Early experiments were conducted before the large surveys of ices toward Solar-type protostars in the 2000s, and thus relied on more anecdotal ice observations. Most of the differences between astrophysical and laboratory ice compositions are instead motivated by experimental concerns however. Simple ice mixtures or even pure ices are often used to isolate the kinetics and mechanisms of particular ice formation pathways. In other experiments the concentrations of the most reactive ice species are increased to increase the yield of prebiotically interesting products. In astrophysical environments it is difficult to directly detect organics larger than CH$_3$OH in ices due to overlapping ice bands from more complex organic ices. A mixture of moderately complex organics, including CH$_3$CH$_2$OH and HCOOH, have been proposed to contribute to observed ice bands between 5 and 7$\mu$m toward protostars\cite{Gibb04,Oberg11c,Boogert15}, but unique identifications have been controversial. Similarly, families of very large COMs may be possible to detect with the next generation of infrared telescopes using features at $\sim$3.4$\mu$m, but only if the original ice mantles has first sublimated\cite{MunozCaro09,MunozCaro13}. Detections of specific molecules in both ice and ice residues will likely remain elusive even with large increases in sensitivity, however, due to the intrinsic overlap between infrared ice features of related molecules. \subsection{Interstellar UV Fields} \begin{figure}[htbp] \begin{center} \includegraphics[width=5in]{fig3.pdf} \caption{Ices can be exposed to a range of UV spectra in the ISM, including black-body radiation from nearby massive stars\cite{Mathis83} (top), line emission from a young Solar-type star\cite{Herczeg02,Valenti03,Johns-Krull07} (middle), and line emission from cosmic ray interactions with molecular hydrogen\cite{Gredel87} (bottom; for a high resolution spectrum see the original publication). All spectra have been normalized to have a peak flux of unity to enable easy comparison of the different spectral energy distributions.} \label{fig:uv} \end{center} \end{figure} In the context of this review it is important to consider what kind of UV fields icy grain mantles are exposed to from formation to sublimation (Fig. \ref{fig:uv}), and how the energy deposition from UV radiation compares to the energy deposition from other dissociative radiation and particles. Close to the edges of clouds, UV fields are dominated by radiation from massive stars, which can be approximated as a black body with a temperature of 10,000--30,000~K. If there is a massive star nearby, this field will be very intense. If no such star is present close to the clouds, the grain will still be exposed to the interstellar radiation field (ISRF), which is similar in shape, but lower in flux. Models of the interstellar radiation field predict integrated VUV (912--2050 \AA) fluxes of $2.67\times10^{-3}$ erg cm$^{-2}$ s$^{-1}$\cite{Draine78,vanDishoeck06}, corresponding to $\sim$10$^8$ photons cm$^{-2}$ s$^{-1}$. No photons with shorter wavelengths than 912 {\AA} are present in the diffuse ISM due to absorptions by atomic hydrogen. In diffuse environments the timescales of ice photodesorption are short. Build-up of icy grain mantles therefore take place only in dense molecular clouds, where the grains are shielded by dust absorption from any external sources of UV radiation. They are instead exposed to a low-intensity UV field, which is due to cosmic ray interactions with H$_2$. The resulting VUV flux is $\sim$10$^{4}$ photons cm$^{-2}$\cite{Shen04}. The VUV spectra that emerges from this process has been calculated (Fig. \ref{fig:uv})\cite{Gredel87}, and the calculation demonstrates that a substantial portion of the VUV flux comes in the form of Lyman-$\alpha$. At the later stages of star formation, icy mantles can be directly exposed to UV radiation from the central protostar or pre-main sequence star. Solar-type pre-main sequence stars can have large VUV excesses in the form of Lyman-$\alpha$ radiation\cite{Herczeg02,Valenti03,Johns-Krull07,Bergin03}, which may induce additional UV processing in ices residing in protoplanetary disks. In addition to UV radiation, interstellar ices are exposed to higher energy photons (especially X-rays), cosmic rays (i.e. high energy atomic nuclei) and electrons. Which kind of radiation that is the most important for ice chemistry in molecular clouds, protostellar envelopes and in protoplanetary disks is debated. Calculations relevant to molecular clouds suggest that an order of magnitude more energy is deposited into grain mantles from secondary UV photons compared to direct interactions with cosmic rays\cite{Shen04}. These estimates may underestimate, however, the importance of secondary electrons created by a cosmic ray as it is passing through the grain and ice mantle\cite{Bennett05}. That is, cosmic rays passing through an icy mantle can generate millions of secondary electrons that can induce dissociations similarly to UV radiation. The interactions of electrons with ices was recently reviewed\cite{Arumainayagam10}, and there are many studies on the generation of secondary electrons by cosmic ray interactions with ice, and their effects on ice chemistry\cite{Moore07,Kaiser13}. In the protoplanetary disk stage, cosmic rays can be efficiently excluded from the disk itself due to shielding by strong T Tauri (pre-main sequence Solar-type stars) stellar winds\cite{Cleeves15}. In this environment, X-rays may be the most important source of molecular dissociation in ices; young stars have X-ray excesses, and X-rays can penetrate deeper into the disk than VUV photons. \subsection{Complex Organic Molecules during Star Formation} \begin{figure}[htbp] \begin{center} \includegraphics[width=6in]{fig5.pdf} \caption{Spectra between 239 and 243 GHz toward the massive protostar NGC 7538 IRS1\cite{Oberg14}.} \label{fig:irs1} \end{center} \end{figure} COMs were first discovered in massive-star forming regions, where clusters of stars, some of which are many times more massive than the Sun, are currently forming and heating up their surrounding dense media and sublimating icy grain mantles\cite{Ball70,Solomon71}. In these, and other interstellar environments, COMs are observed using rotational spectroscopy. Surveys toward large numbers of high-mass star forming regions have revealed that cm and mm emission from COMs are common attributes of high-mass star formation (Fig. \ref{fig:irs1} shows one example) and the presence of lines from acetonitrile and other COMs toward a star forming region is now used as a sign post for ongoing high-mass formation \cite{Bisschop07,Rosero13}. The emitted molecules generally present high excitation temperatures (100--200~K). The high excitation temperatures as well as spatial resolved observations show that the molecular emission originates in a hot core close to the protostar, where ices have sublimated. In early models, COMs were proposed to form in the hot gas close to the protostar where they are observed. Sublimation of CH$_3$OH ice and other volatiles close to the protostar were supposed to initiate an efficient gas-phase ion-molecule chemistry resulting in the rapid production of HCOOCH$_3$, CH$_3$OCH$_3$ and other complex organic molecules\cite{Charnley95}. A key step in this ion-molecule chemistry is the formation of protonated form of the final product (e.g. protonated HCOOCH$_3$) followed by a recombination with an electron to produce the neutral molecule. The viability of this scenario was challenged, however, by experimental measurements revealing that the recombination step had a very low yield for most molecules and instead preferentially resulted in dissociation into smaller fragments\cite{Geppert06}. Gas-phase chemistry following ice sublimation is therefore no longer considered a plausible scenario for the formation of the majority of the observed COMs toward protostars, but can still contribute to the formation of a subset of COMs\cite{Garrod08}. \begin{figure}[htbp] \begin{center} \includegraphics[width=5in]{fig4.pdf} \caption{Cartoon of complex organic molecule formation through ice photochemistry during star formation. Simple ices form from atom addition reactions at low temperatures in dense molecular clouds. When these ices are exposed to UV radiation, some portion of the original ice is dissociated into radicals. As the icy grain heats up during star formation, the radicals become mobile and can combine to form more complex species. Some radical chemistry may also be possible in the coldest regions dependent on radical concentration and the efficiency of non-thermal diffusion.} \label{fig:com_form} \end{center} \end{figure} The current understanding is that COMs mainly form through ice chemistry on grains. Figure \ref{fig:com_form} illustrates how COMs are proposed to form during star formation through ice photochemistry. As outlined above and below this model has several caveats, including few constraints on the dominant radiation source, and on the relative importance of thermal and non-thermal diffusion of reactants. There are also alternative COM formation pathways that have been proposed\cite{Fedoseev15,Balucani15}. Nevertheless, this is the formation scheme employed in most contemporary astrochemistry models, and it thus provides a useful starting point\cite{Herbst09}. In this scheme, the simple ices formed during the molecular cloud phase partially dissociate into radicals when exposed to UV radiation from e.g. cosmic ray interactions with H$_2$. For example CH$_3$OH can be dissociated into CH$_3$+OH, CH$_2$OH+H and other species. When a protostar forms and begins to accrete nebular material it accretes dust grains together with gas. As the dust grain approaches the protostar, the icy mantle becomes warm enough for radicals to become mobile in the ice. The radicals combine to form new, often more complex species, such as CH$_3$CH$_2$OH from CH$_3$ and CH$_2$OH\cite{Garrod06,Garrod08,Garrod13}. The icy grain continues to flow toward the central protostar and eventually becomes warm enough for the ices to sublimate (T$\sim$100--150~K) populating the gas-phase with a mixture of simple molecules from the original ice mantle and the newly formed more complex organics. This model explains observed COM abundances toward massive protostars very well\cite{Herbst09}. During the past couple of decades it has become clear that interstellar COMs are not limited to hot cores in high-mass star forming regions. Large numbers of complex organics have been detected toward the low-mass Solar-type protostar IRAS 16293-2422 \cite{Cazaux03,Jorgensen11}. In this protostar the excitation pattern and spatial profiles \cite{Bisschop08} of the detected molecules indicate that the molecules are emitting from a hot and compact region that is associated with ice sublimation close to the protostar, analogously to the hot cores found around massive protostars. The chemical composition is also similar to what is found toward more massive protostars, and the chemistry has been successfully modeled using adaptations of the scenario presented in Fig. \ref{fig:com_form} \cite{Herbst09}. More recently several COMs have been observed in a number of sources that are cold compared to the hot cores associated with the innermost regions of low-mass and high-mass protostars. These observations cannot be explained by thermal sublimation of complex ices, but may still be consistent with an ice photochemistry origin. In fact some such observations, including one of COMs in a protostellar outflow, provide some of the strongest support of an icy origin of gas-phase interstellar COMs\cite{Arce08}. Protostellar outflows are energetic outflows of gas and dust that appear during the early stages of star formation. The outflows cause shocks, which induce ice mantle sublimation. Compared to protostars, the timescales in these outflows are short, too short for any significant gas-phase formation of COMs. Outflows thus provide a good opportunity to observe ice chemistry products in isolation from gas-phase COM chemistry. The presence of COMs in such an outflow is thus consistent with an ice formation scenario, but not with gas-phase formation of COMs. COMs have also been detected in the lukewarm envelopes of several low-mass and high-mass protostars. The COM excitation temperatures in these sources are well below the expected ice sublimation temperature of $\sim$100~K\cite{Bottinelli07,Oberg11b,Oberg13,Oberg14,Fayolle15}. COMs are thus present in protostellar gas prior to the onset of thermal ice sublimation. The protostellar envelopes where these COMs are detected are generally warm enough for diffusion-limited COM formation to be efficient in ices. The detections do not then pose a fundamental challenge to the ice photochemistry model, as long as non-thermal desorption, e.g. UV photodesorption, chemical desorption and sputtering from cosmic rays and energetic electrons, is sufficiently efficient\cite{Oberg08,Garrod07,Dartois15}. These observations do, however, complicate the scenario presented in Fig. \ref{fig:com_form}. On a more positive note, they also provide a potential window into COM formation at different ice temperatures. Finally, COMs have been detected in few very cold, molecular cloud core environments\cite{Oberg10a,Bacmann12,Cernicharo12}. These observations are more challenging to explain by an ice radical combination chemistry, since radical diffusion is expected to be inhibited in ice at such low temperatures\cite{Garrod08}. Non-thermal or hot-atom diffusion of radicals following dissociation of the parent molecule may explain these observations, since it would allow radicals forming in the ice at some distance from one another to meet and react, but there is a lack of data or theory on the efficiency of non-thermal diffusion in ices. Another possible explanation is cosmic ray driven chemistry, since a single cosmic ray interaction with the ice will result in a large number of secondary electrons, which may locally heat up the ice and thus enable radical diffusion for a short period of time\cite{Bennett05}. \subsection{Review Outline} In summary, complex organic molecules are observed in a range of interstellar sources and ice photochemistry provides an attractive explanation for the their existence. The idea that UV irradiation of interstellar ices is a source of chemical complexity has its roots in a series of laboratory experiments carried out in the 1970s and 1980s\cite{Greenberg76}, that found that UV irradiation of interstellar ice analogs resulted in the production of large abundances of large and complex organics. This kind of experiment revealed that ice photochemistry is possible pathway to chemical complexity in space. To determine that ice photochemistry is the dominant pathway requires additional information, however, on the kinetics and mechanisms of ice photochemistry reactions, and on competing formation pathways that are regulated by electrons, cosmic rays and other sources of energy. Some of this information has been acquired over the past decades following dedicated efforts to unravel ice photochemical systems and their electron and cosmic ray induced counterparts. The information is far from complete, however. The aim of this review is therefore not to give a final answer to how much ice photochemistry contributes to chemical complexity in space, but rather to present and evaluate the experiments that have taken us to the current level of understanding of how UV photons interact with interstellar ice analogs to form new organic molecules. Prior to the experimental work, we briefly review the theoretical models that have been developed to interpret experiments and astronomical observations of ice photochemistry (\S2). \S3 reviews experimental techniques employed in laboratory studies of ice photochemistry. In \S4--7 we then review ice photochemistry experiments and their contribution to our understanding of ice photochemistry. \S4 evaluates existing data on the initiation of photochemistry in ices, i.e. ice photodissociation, and the constraints this data provides on the photodissociation process in different ice set-ups. \S5 reviews photochemistry experiments on pure ices, with a focus on the constraints they have provided on photodissociation branching ratios (i.e. what radicals ices dissociate into and how this differ from the gas-phase) and radical diffusion. Photochemistry in simple ice mixtures is treated in \S6. In \S7 we then present and discuss photochemistry experiments that employ more complex ice mixtures as the starting material. The experiments are typically motivated by the question of what prebiotically interesting molecules could form in analogs to interstellar ices, but sometime provide mechanistic constraints as well. \S8 reviews how UV photochemistry compares with the chemistry induced by other types of dissociative radiation, most notably X-rays, energetic protons, and electrons, before a summary and concluding remarks in \S9. | Interstellar ice photochemistry is an efficient pathway to chemical complexity in space. It is a source of prebiotic amino acids and sugars, and may be the original source of enantiomeric excess on the nascent Earth. How much and under which conditions such complex organics can form during star and planet formation is still largely unknown, however, due to largely unconstrained ice photochemistry kinetics and mechanisms. The quantitative time-resolved studies that do exist on ice photochemistry have demonstrated that all steps of the physiochemical process, from initial photodissociation to radical reactions to form new molecules, are difficult to predict from theory or analogy with gas-phase photochemical reactions. The ice environment regulates the effective photodissociation cross section, the radical production branching ratio, radical diffusion rates and therefore radical combination rates, as well as the relative importance of reaction barriers. Experiments have provided some constraints on these processes for specific ice constituents and ice matrices, but much work remains to obtain a comprehensive and quantitative understanding of interstellar ice photochemistry and the role it plays in seeding nascent planets with prebiotic molecules and precursors. Below we summarize the findings in this review as well as provide some final recommendations. \begin{enumerate} \item Measured VUV photodissociation rates of pure and mixed ices are low compared to equivalent measurements for gas-phase species, which can be explained by a combination of ice opacity (most measurements are done on thick ices), and the cage effect, which results in immediate recombinations of some of the photo produced radicals. Because of the importance of the cage effect, gas phase photodissociation cross sections cannot be applied to ice chemistry models. \item Photodissociation branching ratios into different radicals have been estimated for CH$_4$ and CH$_3$OH ice and the results are substantially different from gas-phase measurements. Branching ratios cannot then be taken from gas-phase measurements when modeling ice photochemistry. \item Diffusion and/or reorientation of radicals often, but not always, regulate photochemistry kinetics. In ices that are too cold for thermal diffusion to be active, complex molecule formation can be explained by a combination of non-thermal diffusion and reactions between radicals formed at neighboring sites. It is important to note that both processes may be less important in astrophysical environments where radical production rates are lower. Measured product yields and kinetics in laboratory ices cannot be extrapolated directly to interstellar ices without taking into account the different radical production time scales in the two settings. \item Photochemistry kinetics in ice mixtures often differ substantially from the photochemistry kinetics observed in pure ices due to a combination of new radical formation pathways, larger separation of radicals (if the matrix is inert), and different diffusion environments. These aspects must be taken into account when using data from pure ice experiments to model interstellar ice chemistry. \item Many prebiotically interesting molecules, including amino acids, form in ice photochemistry experiments with ice compositions inspired by observed interstellar ices. It is important to realize, however, that the experiments are often designed to maximize yields, and that most of the individual amino acids seem to form during residual hydrolysis rather than through pure ice chemistry. A more detailed understanding of ice chemistry kinetics is required to predict the typical amino acid concentration in interstellar ices, and thus the abundance delivered to comets and further to nascent planets. \item Electrons, Ions and X-rays can induce an ice chemistry similar to than of VUV photons. Initial radical formation branching ratios may be different for different sources of energy, but this has yet to be firmly established. The subsequent initial non-thermal diffusion step may also depend somewhat on the details of the dissociation process. The diffusion barriers themselves, and thus the thermal diffusion rate, should by contrast not depend on the details of the radical production step. Based on experiments the similarities of photolysis and radiolysis experiments are more important than the differences. That is, product compositions in photolysis and radiolysis ice experiments are often remarkably similar. This suggests that the overall chemical evolution of interstellar ices will mainly depend on the amount of energy deposited into the ice and not on in what kind it is delivered. \end{enumerate} Ice photochemistry is clearly a plausible pathway to chemical complexity in space. It is not the only one however. To evaluate its importance compared to other sources of chemical complexity it is key to develop experimental constraints on the mechanisms and kinetics that drive ice photochemistry and related chemistries in a range of plausible ice morphologies. Existing experiments have demonstrated the difficulty in extrapolating such kinetics from gas to ice, and between different ice systems. As more experiments are carried out, more patterns may emerge however. Energy-resolved experiments seems to be an especially promising avenue for progress, since they enable direct comparison between the initial excitation step and the chemistry. Perhaps the most important missing information is how different radicals diffuse through realistic ices, and how thermal diffusion compares with non-thermal diffusion at different ice temperatures. Constraining these processes is absolutely crucial to extrapolate laboratory results on any kind of ice chemistry to interstellar settings where radical production rates are slow and time scales long. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% The "Acknowledgement" section can be given in all manuscript %% classes. This should be given within the "acknowledgement" %% environment, which will make the correct section or running title. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \begin{acknowledgement} The author thanks Edith Fayolle and Jennifer Bergner, and four anonymous referees for valuable feedback on the manuscript. The author also acknowledges funding from the Simons Collaboration on the Origins of Life (SCOL), award number 321183, an Alfred P. Sloan fellowship, and a Packard fellowship. \end{acknowledgement} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% The appropriate | 16 | 9 | 1609.03112 |
1609 | 1609.02039_arXiv.txt | A peculiar source in the Galactic center known as the Dusty S-cluster Object (DSO/G2) moves on a highly eccentric orbit around the supermassive black hole with the pericenter passage in the spring of 2014. Its nature has been uncertain mainly because of the lack of any information about its intrinsic geometry. For the first time, we use near-infrared polarimetric imaging data to obtain constraints about the geometrical properties of the DSO. We find out that DSO is an intrinsically polarized source, based on the significance analysis of polarization parameters, with the degree of the polarization of $\sim 30\%$ and an alternating polarization angle as it approaches the position of Sgr~A*. Since the DSO exhibits a near-infrared excess of $K_{\rm s}-L'>3$ and remains rather compact in emission-line maps, its main characteristics may be explained with the model of a pre-main-sequence star embedded in a non-spherical dusty envelope. | 16 | 9 | 1609.02039 |
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1609 | 1609.00499_arXiv.txt | Radio astronomy has changed. For years it studied relatively rare sources, which emit mostly non-thermal radiation across the entire electromagnetic spectrum, i.e. radio quasars and radio galaxies. Now it is reaching such faint flux densities that it detects mainly star-forming galaxies and the more common radio-quiet active galactic nuclei. These sources make up the bulk of the extragalactic sky, which has been studied for decades in the infrared, optical, and X-ray bands. I follow the transformation of radio astronomy by reviewing the main components of the radio sky at the bright and faint ends, the issue of their proper classification, their number counts, luminosity functions, and evolution. The overall ``big picture'' astrophysical implications of these results, and their relevance for a number of hot topics in extragalactic astronomy, are also discussed. The future prospects of the faint radio sky are very bright, as we will soon be flooded with survey data. This review should be useful to all extragalactic astronomers, irrespective of their favourite electromagnetic band(s), and even stellar astronomers might find it somewhat gratifying. | \begin{figure} \centering \includegraphics[width=8.4cm]{Fig1.eps} \caption{The radio sky at 4.85 GHz above an optical photograph of the NRAO site in Green Bank, West Virginia (USA). The former 300-foot telescope made this image, which is about 45$^{\circ}$ across. Increasing radio brightness is indicated by lighter shades to indicate how the sky would appear to someone with a "radio eye" 300 feet ($\sim 91$ metres) in diameter. The flux density limit is $\sim 25$ mJy \citep{Gregory_1996}. Copyright NRAO.} \label{fig:radionightsky} \end{figure} The radio sky is very different from the optical sky. When we look at the sky with the naked eye, we practically see only stars. These approximate blackbody radiators and therefore their emission covers a relatively narrow range of frequencies, centred at values ranging from the ultraviolet (UV) to the infrared (IR), depending on the star temperature. Therefore, most bright stars are extremely faint at radio frequencies. Figure \ref{fig:radionightsky} displays the radio sky as it would appear to someone with a ``radio eye'' with a diameter equal to that of the former National Radio Astronomy Observatory (NRAO) Green Bank 300-foot radio telescope. Most of the ``dots'' in the figure, which are unresolved radio sources, are actually distant \cite[$z \approx 0.8$;][]{Condon_1989} luminous radio galaxies (RGs) and quasars. The very few extended sources are mostly Galactic supernova remnants. Radio emission from RGs and quasars is due to ultra-relativistic ($E \gg m_{e} c^{2}$) electrons moving in a magnetic field and thereby emitting synchrotron radiation, which, unlike blackbody emission, can cover a very large range in frequency, reaching $\sim 10$ decades in some sources. \begin{figure} \centering \subfloat[3C 31]{{\includegraphics[width=5.0cm]{Fig2a.eps} }}% \qquad \subfloat[Fornax A]{{\includegraphics[width=5.0cm]{Fig2b.eps} }}% \caption{{\it (a):} large scale radio map at 1.4 GHz of 3C 31, showing filamentary plumes extending over 400 kpc from the galaxy. Copyright NRAO 1996. Image from http://www.cv.nrao.edu/$\sim$abridle/images.htm. {\it (b):} Radio emission (orange) associated with the giant elliptical galaxy NGC1316 (centre of the image), consisting of two large radio lobes, each extending over $\sim 180$ kpc. Image courtesy of NRAO/AUI and J. M. Uson.} \label{fig:3C_31_Fornax_A} \end{figure} As one goes fainter by using telescopes, galaxies take over even in the optical sky: for AB $\gtrsim 20 - 22$ mag, depending on the filter, galaxies outnumber stars by a large margin \citep{Windhorst_2011}. The Hubble Ultra Deep Field (HUDF), which covers 11 arcmin$^2$ in four filters ($B$ to $z$) down to approximately uniform limiting magnitudes AB $\sim 29$ for point sources, contains at least 10,000 objects, almost all of them galaxies \citep{Beckwith_2006}. Nevertheless, these galaxies are very different from those seen in the radio ``bright''\footnote{The standard flux density unit in radio astronomy is the Jansky (Jy), which is equivalent to $10^{-23}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$. By today's standards, strong radio sources have $S_{\rm r} \gtrsim 1$ Jy, intermediate ones have 1 mJy $\lesssim S_{\rm r} \lesssim 1$ Jy, while weak radio sources are below the mJy (soon $\mu$Jy) level.} sky. Radio quasars and RGs, in fact, are somewhat rare, atypical, mostly non-thermal sources across the entire electromagnetic spectrum, in which a large fraction of the total emission comes from relativistic jets, that is streams of plasma with speeds getting close to the speed of light, and associated lobes (Fig. \ref{fig:3C_31_Fornax_A}). Most of the galaxies detected in the HUDF, on the other hand, are undergoing episodes of star formation (SF) and therefore are strong thermal emitters. This review\footnote{There was obviously no way I could mention {\it all} papers dealing with the many topics related to this review, which have appeared in the literature. I have therefore had to make choices and often resorted to the sentence ``and references therein''. Moreover, I here deal exclusively with extragalactic sources; see, e.g., Sect. 3.8 of \cite{Norris_2013} for a discussion of radio surveys of the Galactic plane.} aims to discuss very recent developments in our understanding of the faint radio sky and how our radio view of the Universe has changed and got much similar to the optical one. By going radio faint one is in fact detecting the bulk of the active galactic nuclei (AGN) population, and not only the small minority of radio quasars and RGs, and also plenty of star-forming galaxies (SFGs). These developments are having (or should have) a strong effect on radio astronomy, but should also change the perception that astronomers working in other bands have of it. One of the main messages of this review, in fact, is that radio astronomy is not a ``niche'' activity but is extremely relevant also to more classical aspects of extragalactic astronomy, such as star formation and galaxy evolution. I wrote this review primarily with non radio-astronomers in mind but I believe that some radio astronomers might still be not fully aware of their changing landscape and especially of what is in store for them. Therefore, this review should be useful to all extragalactic astronomers irrespective of their preferred electromagnetic band(s). Stellar astronomers might still find some satisfaction in learning that the faint radio sky is dominated by SF related processes (which is another ``take home'' message)! The structure of this review is as follows: Sect. 2 discusses the main astronomical components of the radio sky, while Sect. 3 deals with the radio number counts. In Sect. 4 and 5 I describe the bright and faint radio sky populations respectively, expanding in the latter on source classification, radio number counts by class, luminosity functions (LFs), and evolution. Sect. 6 puts these results into the bigger picture by discussing some astrophysical implications, while Sect. 7 dwells on future prospects by discussing upcoming radio facilities and their astrophysical impact, predictions for deeper radio surveys, and the issue of source classification of very faint ($\lesssim 1~\mu$Jy) radio sources. Finally, Sect. 8 gives my conclusions and some messages. Throughout this paper, spectral indices are defined by $S_{\nu} \propto \nu^{-\alpha}$, magnitudes are in the AB system, and the values $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\rm \Omega_{\rm m} = 0.27$, and $\rm \Omega_{\rm \Lambda} = 0.73$ have been used. The state of the art of this field until early 2009 is reviewed by \cite{deZotti_2010}. | At this point, it should be clear to all astronomers that the faint radio sky plays a vibrant role in a variety of astrophysical topics, including the cosmic star formation history, galaxy evolution, the existence of powerful jets, and radio emission in RQ AGN. This role will grow even further in the near future. Radio observations are also unaffected by absorption, which means, for example, that they are sensitive to all types of AGN, irrespective of obscuration and orientation (i.e., Type 1s and Type 2s). I conclude this review by sending the following: \subsection{Messages to all astronomers} \begin{enumerate} \item Do not assume that {\it radio-detected} means {\it radio-loud}. While this was almost always true when radio surveys only reached the $\approx$ Jy level, this is no longer the case, quite the opposite: a sub-mJy AGN is more likely to be RQ than RL! \item The ``radio-quiet AGN'' label is obsolete, misleading, and wrong. The major difference between RL and RQ AGN is, based on the available evidence, the presence or lack of {\it strong} (relativistic) jets, which in practice translates into them being off or on the FIR -- radio correlation. I therefore propose that we start using "jetted AGN" and "non-jetted" AGN. This name has been used already (albeit very sparsely) in the literature\footnote{At the time of writing (mid-2016) I have found 16 refereed papers with the words "jetted AGN" in their abstract.}. I think it is high time it becomes the norm. \item Do not look for a bimodality in $R$ or $P_{\rm r}$ in quasars, as we already know that there are two main classes of AGN: jetted and non-jetted. \item Most importantly: radio astronomy is not a ``niche'' activity but is extremely relevant to a whole range of extragalactic studies related, for example, to star formation and galaxy evolution. Take advantage of that and use radio data! \end{enumerate} \subsection{Messages to radio astronomers} \begin{enumerate} \item The flattening of deep normalized radio counts is not an open issue: we have sorted out the source population of the sub-mJy GHz radio sky and learnt that below $\approx 0.1$ mJy the radio sky is dominated by SF related processes. This process took more than thirty years because the multi-wavelength data necessary to properly classify sources were not available and radio astronomy was ahead of the other bands. History does not have to repeat itself, although there is some risk that this might happen. This can be avoided if radio astronomers understand that ... \item ... radio astronomy is now a fully multi-wavelength enterprise. This is also evinced by the fact that about one third of the papers referenced in this review are {\it not} radio papers. The full, proper exploitation of data from the SKA and its precursors will {\it require} (this is not an option!) synergy with other contemporaneous astronomical facilities. These will include, among others, {\it Athena}, WFIRST, the LSST, the ELTs, JWST, and SPICA. \end{enumerate} | 16 | 9 | 1609.00499 |
1609 | 1609.05796_arXiv.txt | Understanding the nature of dark energy, the mysterious force driving the accelerated expansion of the Universe, is a major challenge of modern cosmology. The next generation of cosmological surveys, specifically designed to address this issue, rely on accurate measurements of the apparent shapes of distant galaxies. However, shape measurement methods suffer from various unavoidable biases and therefore will rely on a precise calibration to meet the accuracy requirements of the science analysis. This calibration process remains an open challenge as it requires large sets of high quality galaxy images. To this end, we study the application of deep conditional generative models in generating realistic galaxy images. In particular we consider variations on conditional variational autoencoder and introduce a new adversarial objective for training of conditional generative networks. Our results suggest a reliable alternative to the acquisition of expensive high quality observations for generating the calibration data needed by the next generation of cosmological surveys. | \label{sec:calibration} In the weak regime of gravitational lensing, the distortion of background galaxy images can be modeled by an anisotropic shear, noted $\gamma$, whose amplitude and orientation depend on the matter distribution between the observer and these distant galaxies. This shear affects in particular the apparent ellipticity of galaxies, denoted $e$. % Measuring this weak lensing effect is made possible under the assumption that background galaxies are randomly oriented, so that the ensemble average of the shapes would average to zero in the absence of lensing. Their apparent ellipticity $e$ can then be used as a noisy but unbiased estimator of the shear field $\gamma$: $\Ex [ e ] = \gamma$. The cosmological analysis then involves computing auto- and cross-correlations of the measured ellipticities for galaxies at different distances. These correlation functions are compared to theoretical predictions in order to constrain cosmological models and shed light on the nature of dark energy. However, measuring galaxy ellipticities such that their ensemble average (used for the cosmological analysis) is unbiased is an extremely challenging task. \cref{fig:great} illustrates the main steps involved in the acquisition of the science images. The weakly sheared galaxy images undergo additional distortions (essentially blurring) as they go through the atmosphere and telescope optics, before being acquired by the imaging sensor which pixelates the noisy image. As this figure illustrates, the cosmological shear is clearly a subdominant effect in the final image and needs to be disentangled from subsequent blurring by the atmosphere and telescope options. This blurring, or Point Spread Function (PSF), can be directly measured by using stars as point sources, as shown at the top of~\cref{fig:great}. \begin{figure*} \centering \includegraphics[width=1\linewidth]{figures/cosmos.pdf} \caption{{\small Samples from the \COSMOS dataset and generated samples using the conditional variational autoencoder (C-VAE, scheme I) and our variation on conditional generative adversarial network (C-GAN). Each column image shows three $64 \times 64$ images (here inverted) produced by conditioning on the same set of features $y \in \Re^{3}$ in the test-set. Due to its high dynamic range, most figures are very faint. In the bottom three rows, each image is individually normalized.}} \label{fig:cosmos_images} \end{figure*} Once the image is acquired, shape measurement algorithms are used to estimate the ellipticity of the galaxy while correcting for the PSF. However, despite the best efforts of the weak lensing community for nearly two decades, all current state-of-the-art shape measurement algorithms are still susceptible to biases in the inferred shears. These measurement biases are commonly modeled in terms of additive and multiplicative bias parameters $c$ and $m$ defined as: \begin{equation} \Ex[e] = (1 + m) \ \gamma + c \label{eq:shear_bias} \end{equation} where $\gamma$ is the true shear. Depending on the shape measurement method being used, $m$ and $c$ can depend on factors such as the PSF size/shape, the level of noise in the images or, more generally, intrinsic properties of the galaxy population (like their size and ellipticity distributions, \etc). Calibration of these biases can be achieved using image simulations, closely mimicking real observations for a given survey but using galaxy images distorted with a known shear, thus allowing the measurement of the bias parameters in \cref{eq:shear_bias}. \textbf{Image simulation pipelines}, such as the \textit{GalSim} package \cite{Rowe2015}, use a forward modeling of the observations, reproducing all the steps of the image acquisition process in \cref{fig:great}, and therefore require as a starting point galaxy images with high resolution and S/N. The main difficulty in these image simulations is therefore the need for a calibration sample of high quality galaxy images representative of the galaxy population of the survey being simulated. Our aim in this work is to train a deep generative model which can be used to cheaply synthesize such data sets for specific galaxy populations, by conditioning the samples on measurable quantities. \subsection{Data set}\label{sec:dataset} As our main dataset, we use the \COSMOS survey to build a training and validation set of galaxy images and extract from the corresponding catalog a condition vector $y$ with three features: half-light radius (measure of size), magnitude (measure of brightness) and redshift (cosmological measure of distance). To facilitate the training, we align all galaxies along their major axis and produce 85,000 instances of 64x64 image stamps using the GalSim package. We also use the \GALAXYZOO dataset~\cite{willett2013galaxy} to demonstrate the abilities of our alternative conditional adversarial objective. Each of the 61,000 galaxy images in this dataset is accompanied by $\yv \in [0,1]^{37}$ features produced using a crowd-sourced set of questions that form a decision tree. We cropped the central $50\%$ of these images and resized them to $128\times128$ pixels. We augmented both datasets by flipping the images along the vertical and horizontal axes. | 16 | 9 | 1609.05796 |
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1609 | 1609.03386_arXiv.txt | {We present cosmic microwave background (CMB) power spectra from recent numerical simulations of cosmic strings in the Abelian Higgs model and compare them to CMB power spectra measured by \Planck. We obtain revised constraints on the cosmic string tension parameter $G\mu$. For example, in the \LCDM\ model with the addition of strings and no primordial tensor perturbations, we find $G\mu < 2.0 \times 10^{-7}$ at 95\% confidence, about 20\% lower than the value obtained from previous simulations, which had 1/64 of the spatial volume. The increased computational volume also makes it possible to simulate fully the physical equations of motion, in which the string cores shrink in comoving coordinates. We find however that this, and the larger dynamic range, changes the amplitude of the power spectra by only about 10\%. The main cause of the stronger constraints on $G\mu$ is instead an improved treatment of the string evolution across the radiation-matter transition.} | \label{sec:intro} The Cosmic Microwave Background (CMB) has demonstrated many times its power as a clean and reliable probe for early and late-time cosmological physics. The latest results from the \Planck\ collaboration \cite{Ade:2015xua} have reached percent-level precision on nearly all standard parameters, in some cases doing even better. The CMB is also an important probe for cosmic defects \cite{Kibble:1976sj} as it is mainly sensitive to the large-scale properties of the defect distribution which is relatively well understood \cite{Pen:1997ae,Pogosian:1999np,Martins:2000cs,Durrer:2001cg,Martins:2003vd,Bevis:2010gj,BlancoPillado:2011dq}. The \Planck\ data currently limits the contribution of strings and other defects to the temperature anisotropies on large scales to be of the order of a few percent or less \cite{Ade:2013xla,Lizarraga:2014xza,Lazanu:2014eya,Ade:2015xua,Charnock:2016nzm}. However, the limits can be only as accurate and precise as the theoretical forecasts that are compared to the data. In this paper we use the recently updated unequal-time correlation functions (UETC) from \cite{Daverio:2015nva} to compute new CMB power spectra for cosmic strings in the Abelian Higgs model, the prototypical field theory with such topological defects. \footnote{For reviews of cosmic strings in cosmology and a discussion of the differences between calculations based on field theory and modelling with ideal Nambu-Goto strings see Refs.~\cite{Copeland:2011dx,Hindmarsh:2011qj}.} The new UETCs are based on simulations that are four times larger (64 times larger in volume) than those used previously \cite{Bevis:2010gj}, \ie\ they have a four times larger dynamical range in both space and time. As usual, the Abelian Higgs (AH) strings were simulated with couplings at the TypeI/TypeII boundary. Thanks to the much larger simulation volume, we were able to simulate directly scales that previously we could only infer. It was also possible to keep the string cores at a constant physical width, instead of letting them grow with the expansion of the universe as was done previously. In fact, in previous works with smaller simulation volumes, the string width was allowed to grow artificially in order to keep the core of the string resolved in a grid in comoving coordinates \cite{Press:1989yh,Bevis:2010gj}. The growth of the string was controlled by a parameter called $s$, which ranges from 0 to 1: $s=0$ corresponds to a string whose width stays constant in comoving coordinates, and $s=1$ corresponds to strings without artificial core-growth, \ie\ with a constant physical width throughout the simulation. In the new simulations presented in \cite{Daverio:2015nva}, we were able to simulate strings with $s=1$. In addition, the modelling of the evolution of the string network across the cosmological radiation-matter and matter-$\Lambda$ transitions has been treated much more carefully than in previous works. The differences in observables between between $s=1$ and $s=0$ were not great, and in particular the equal time correlators of the new simulations were consistent with the old ones (see Fig. 6 of ref.~ \cite{Daverio:2015nva}). The consistency meant that we were able to merge the UETCs from $s=1$ simulations with ones computed at $s=0$, thereby increasing the accessible range of time differences. A more detailed comparison of $s=1$ and $s=0$, and tests of the scaling assumption, will be the subject of a future publication. In the next section we describe briefly how we compute the power spectra of temperature and polarisation anisotropies in the CMB from the new UETCs obtained in \cite{Daverio:2015nva}. In section \ref{sec_cmbLAH} we compare how the different improvements in the simulations affect the $C_\ell$ and what the main causes of the differences with the previous spectra are. In particular we demonstrate that the new simulations change the $C_\ell$ by only about 10\%, when they are computed with the same methods as in \cite{Bevis:2010gj}. We then use the new $C_\ell$ to place constraints on the abundance of cosmic strings in section \ref{sec:fits} before concluding. | We have calculated the CMB power spectra from the energy-momentum correlations computed in recent large-scale numerical simulations of a network of cosmic strings in the Abelian Higgs model \cite{Daverio:2015nva}, and compared them to \Planck\ CMB power spectra. We obtain a revised constraint on the cosmic string tension parameter $G\mu$, and investigate the source of the difference. The new numerical simulations represent a significant improvement over those \cite{Bevis:2010gj} used in previous comparisons to \Planck\ data \cite{Ade:2015xua}, with a factor 64 increase in volume and sufficient resources to properly solve the equations in an expanding universe, and to investigate the radiation-matter and matter-$\La$ transitions, both for the first time. The larger simulations confirmed the shape and normalisation of the unequal time correlators (UETCs) computed in \cite{Bevis:2010gj}, within our original uncertainty estimates. The biggest change comes from the improved method of treating the radiation-matter transition \cite{Daverio:2015nva}, which demonstrates that the larger UETCs of the radiation era are preserved for longer than previously thought. The consequence is that strings produce 30\% higher CMB power spectra for a given $G\mu$. We trace the effect of the changes in Fig.~\ref{figure_cl5}. It can be seen that the TT power spectrum from the new simulations computed by the 2010 methods are very close to the 2010 power spectra, and that the major change comes from improvements in the treatment of the transition in the computation. We investigate two new methods for CMB power spectrum computation from defects, one proposed in \cite{Fenu:2013tea}, and the other our own \cite{Daverio:2015nva}, which is conceptually simpler and easier to implement. They differ little: the main effect comes from the increase in amplitude of the UETCs at the time of decoupling, when the CMB perturbations are generated. Given the agreement between the old and new simulations and the new treatment of the cosmological transitions, we expect that the main theoretical uncertainty in the resulting $C_\ell$ is now due to the old Einstein-Boltzmann solver that we use to compute the power spectrum from the UETC. That code is based on an old version of CMBEasy \cite{Doran:2003sy} with a by-now outdated recombination history and other legacy precision issues. In the future we plan to move to CLASS \cite{Lesgourgues:2011re}, a project that is currently under way. Our final constraints are given in Table \ref{t:gmulimits}, quoted as limits on the fraction of the power spectrum due to cosmic strings at multipole $\ell=10$ $\fd$, and the square of the string tension parameter $(G\mu)^2$. We see that whether or not we allow for primordial tensor fluctuations in the model, no more than about 1\% of the CMB power spectrum can be due to cosmic strings. With our improved calculation of the power spectrum, the limit on $G\mu$ (assuming no primordial tensor fluctuations) is approximately $2.0\times 10^{-7}$, 17\% lower than the value $2.4\times10^{-7}$ obtained with the old method \cite{Ade:2015xua}. The new constraint corresponds to a symmetry-breaking scale of $2.2 \times 10^{15}$ GeV in the Abelian Higgs model at critical coupling. \vspace{1cm} | 16 | 9 | 1609.03386 |
1609 | 1609.04273_arXiv.txt | { The first data release from the Gaia mission contains accurate positions and magnitudes for more than a billion sources, and proper motions and parallaxes for the majority of the 2.5~million Hipparcos and Tycho-2 stars. }{ We describe three essential elements of the initial data treatment leading to this catalogue: the image analysis, the construction of a source list, and the near real-time monitoring of the payload health. We also discuss some weak points that set limitations for the attainable precision at the present stage of the mission. }{ Image parameters for point sources are derived from one-dimensional scans, using a maximum likelihood method, under the assumption of a line spread function constant in time, and a complete modelling of bias and background. These conditions are, however, not completely fulfilled. The Gaia source list is built starting from a large ground-based catalogue, but even so a significant number of new entries have been added, and a large number have been removed. The autonomous onboard star image detection will pick up many spurious images, especially around bright sources, and such unwanted detections must be identified. Another key step of the source list creation consists in arranging the more than $10^{10}$ individual detections in spatially isolated groups that can be analysed individually. }{ Complete software systems have been built for the Gaia initial data treatment, that manage approximately 50~million focal plane transits daily, giving transit times and fluxes for 500~million individual CCD images to the astrometric and photometric processing chains. The software also carries out a successful and detailed daily monitoring of Gaia health. }{} | 16 | 9 | 1609.04273 |
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1609 | 1609.03453_arXiv.txt | We report the analysis of 22 B-band light curves of the dwarf nova V4140~Sgr obtained with SOI/SOAR during two nights along the decline of a superoutburst in 2006 Sep 12-24 and in quiescence over 50 days following the superoutburst. Three-dimensional eclipse mapping of the outburst light curves indicates that the accretion disc is elliptical (eccentricity $e=0.13$) and that superhump maximum occurs when the mass donor star is aligned with the bulge of the elliptical disc. The accretion disc is geometrically thin both in outburst and in quiescence; it fills the primary Roche lobe in outburst and shrinks to about half this size in quiescence. The stability of the eclipse shape, width and depth along quiescence and the derived disc surface brightness distribution indicate that the quiescent accretion disc is in a high-viscosity, steady-state. Flickering mapping of the quiescent data reveal that the low-frequency flickering arises from an azimuthally-extended stream-disc impact region at disc rim and from the innermost disc region, whereas the high-frequency flickering originates in the accretion disc. Assuming the disc-related flickering to be caused by fluctuations in the energy dissipation rate induced by magneto-hydrodynamic turbulence \citep{ga}, we find that the quiescent disc viscosity parameter is large $\alpha\simeq 0.2-0.4$ at all radii. The high-viscosity quiescent disc and the inferred low disc temperatures in superoutburst are inconsistent with expectations of the disc-instability model, and lead to the conclusion that the outbursts of V4140~Sgr are powered by mass transfer bursts from its donor star. | In dwarf novae, mass is transferred from a late-type star (the secondary) to a companion white dwarf (WD) via an accretion disc. They show recurrent outbursts on timescales of days-months, in which the disc brightens by factors 20-100 for $\simeq 1-10$ days. Outbursts are explained in terms of either a thermal-viscous disc-instability model \cite[DIM, e.g.,][]{c93,lasota} or a mass-transfer instability model \cite[MTIM, e.g.,][]{bath}. DIM predicts matter accumulates in a low viscosity \footnote{here we adopt the prescription of \cite{ss} for the accretion disc viscosity, $\nu = \alpha c_s H$, where $\alpha$ is the non-dimensional viscosity parameter, $c_s$ is the local sound speed and $H$ is the disc scaleheight.} disc during quiescence ($\alpha_{\rm quies}\sim 10^{-2}$) which switches to a high-viscosity regime during outbursts, whereas in MTIM the disc viscosity is always the same \cite[$\alpha\sim 0.1-1$, from the decline timescale of outbursting dwarf novae, e.g.,][]{w95}. Therefore, estimating $\alpha$ of a quiescent disc is key to gauge which model is at work in a given dwarf nova. Aside of the normal outbursts, the short-period dwarf novae of the SU\,UMa type show longer, slightly brighter and more regular superoutbursts, during which a hump-shaped brightness modulation (named {\em superhump}) with period slightly longer than the orbital period $P_\mathrm{orb}$ is seen in their light curves \citep{w95,hellier}. The most promising explanation for superhumps is given by the tidal resonance instability model \citep{whitehurst88,ho90,lubow94}: during a superoutburst, the accretion disc expands beyond the 3:1 resonance radius, $R_{31}$, and a tidal instability sets in, giving rise to an elliptical, slowly precessing disc. Superhump modulation then arises from the periodic tidal interaction between the outer elliptical disc and the mass-donor star (normal superhumps) or by the varying release of gravitational energy at the point where the gas stream hits the precessing elliptical disc outer edge (late superhumps), at the beat period between the orbital and the disc precession period, $1/P_\mathrm{sh}= 1/P_\mathrm{orb} - 1/P_\mathrm{prec}$. Eclipse mapping \citep{horne85} is a powerful tool to test for the presence of elliptical discs in outbursting dwarf novae as well as to check whether the orientation of the ellipse is according to the expectation of the tidal instability model \cite[e.g.,][]{odonoghue,rolfe00}. Flickering is the intrinsic brightness fluctuation of 0.01-1 mag on timescales of seconds to dozens of minutes seen in dwarf novae light curves. Optical studies \citep[e.g.,][]{bruch92,bruch96,bruch00,bb04} suggest there may be three different sources of flickering in dwarf novae and novalike systems, the relative importance of which varies from system to system: (i) the stream-disc impact region \cite[possibly because of unsteady mass transfer or post-shock turbulence,][]{wn,shu76}, (ii) the innermost disc regions around the WD \cite[possibly powered by unsteady WD accretion or post-shock turbulence in the boundary layer between disc and WD,] []{ej82,bruch92}, and (iii) the accretion disc itself \cite[probably as a consequence of magneto-hydrodynamic (MHD) turbulence or events of magnetic reconnection at the disc atmosphere,][]{ga,kawa,bb04}. The power density spectrum (PDS) of the flickering is characterized by a continuum power-law, $P(f) \propto f^{-n}$, with $n\simeq 1-3$ \citep{bruch92}, which flattens below a given cut-off frequency $f_c$. Currently, the most promising explanation for the anomalously large viscosity of accretion discs is related to MHD turbulence in the differentially rotating disc gas \citep[driven by the magnetorotational instability, MRI, see][]{bh91,hb91}. Most of the studies on this subject over the last two decades focused on confirming that MRI leads to self-sustained, turbulent and efficient outward flow of angular momentum and inward flow of disc matter \citep[e.g.,][and references therein]{bh98,beckwith11}, and on the comparison of the numerically derived values of $\alpha$ with those inferred by the decline timescale of outbursting dwarf novae \citep[e.g.,][]{king07,hirose14}. On the other hand, the study of \cite{ga} focused on the time variability of the viscous energy release in a MHD turbulent disc. They found that MHD turbulence leads to large fluctuations in the energy dissipation rate per unit area at the disc surface, $D(R)$, which they suggested could be a source of flickering in mass-exchanging binaries. Indeed, the PDS of the fluctuations in their study resemble those of flickering sources, with a power-law dependency of similar index range and a flat slope at low-frequencies. The interpretation of the fluctuations in $D(R)$ as the stochastic and statistically independent release of energy from a large number of turbulent eddies leads to a direct relation between the relative amplitude of the energy fluctuations $\sigma_\mathrm{D}/\langle D \rangle$ and the disc viscosity parameter, $\sigma_\mathrm{D}/\langle D \rangle \propto \alpha^{1/2}\,(H/R)^{1/2}$, providing an interesting observational way to estimate the local accretion disc viscosity parameter \cite[see, e.g.,][]{bb04}. For a thin accretion disc ($H\simeq 0.01\,R$), the above relation predicts that low-viscosity discs ($\alpha \sim 10^{-2}$) should show low-amplitude disc flickering (hardly detectable at a level $\leq 1$ per cent), whereas high-viscosity discs ($\alpha=$ 0.1-1) should display detectable flickering with relative amplitudes in the range 2.5-7.5 per cent. There is observational support for this prediction: e.g., the accretion disc seems the dominant source of flickering in the high-viscosity discs of the novalike systems RW\,Tri \citep{hs85} and UX\,UMa \citep{bruch00}, whereas there is no evidence of disc-related flickering in the low-viscosity accretion discs of the dwarf novae U\,Gem \citep{wn} and IP\,Peg \citep{bruch00} in quiescence. V4140 Sgr is an 88-min period eclipsing SU~UMa type dwarf nova \citep{js,bjs89} showing low-amplitude ($\sim 1$ mag), 5-10\,d long outbursts recurring every 80-90\,d and longer, brighter superoutbursts where superhumps appear in its light curve \citep{bb05}. Here we report the analysis of a sample of light curves of V4140~Sgr with eclipse mapping techniques to trace the evolution of the surface brightness of its accretion disc during decline from a superoutbursts, to locate the sources of flickering in the binary and to estimate the radial run of the quiescent disc viscosity parameter. Section\,\ref{observa} reports the observations and data reduction procedures. Data analysis and results are presented in section\,\ref{analysis}, discussed in section\,\ref{discuss}, and summarized in section\,\ref{conclusions}. | \label{conclusions} We applied 3D eclipse mapping techniques to follow the evolution of the surface brightness of the accretion disc of V4140~Sgr in a superoutburst. We find that the disc is elliptical in outburst and decline, with an eccentricity $e=0.13$. In both outburst stages, the disc orientation is such that superhump maximum occurs when the secondary star is aligned with the bulge of the elliptical disc. This lends observational support for the tidal resonance instability model of superhumps. The accretion disc fills the primary Roche lobe at outburst ($R_d=0.85\,R_\mathrm{L1}$), shrinks to $0.6\,R_\mathrm{L1}$ during decline and reaches a lower value of $0.45\,R_\mathrm{L1}$ in quiescence. There is marginal evidence that the disc half-opening angle is larger in outburst ($\beta= 1.0^o\pm 0.5^o$) than in quiescence ($\beta= 0.5^o\pm 0.5^o$), but the disc is geometrically thin in all three cases. The superoutburst occurs at disc temperatures too low to be accounted for by the disc-instability model even at the upper $4\,\sigma$ limit on the inferred distance to the binary. The stability of the eclipse shape, width and depth along 50\,d in the quiescent period following the superoutburst and the derived disc surface brightness distribution (outshining any contribution from the white dwarf) indicate that the quiescent accretion disc of V4140~Sgr is in a high-viscosity, steady-state regime. Flickering mapping of the quiescent data reveal three different sources of flickering in V4140~Sgr: an azimuthally-extended stream-disc impact region at disc rim (statistically significant at the $3\,\sigma$ level) and the innermost disc region (statistically significant at the $2.5\,\sigma$ level), responsible for the low-frequency flickering, and an extended, high-frequency disc component responsible for 2/3 of the total optical flickering in V4140\,Sgr. Assuming that the disc-related flickering is caused by fluctuations in the energy dissipation rate induced by MHD turbulence according to the model of \cite{ga}, we find that the quiescent disc viscosity parameter in V4140~Sgr is large ($\alpha \simeq 0.2-0.4$) at all disc radii, in agreement with MTIM predictions and in marked contrast with DIM predictions. The discrepancy between the inferred $\alpha$ values and DIM predictions is statistically significant at the $2\,\sigma$ confidence level outside the region $0.2\leq R/R_\mathrm{L1}\leq 0.3$. The high-viscosity, steady-state quiescent disc of V4140~Sgr and the inferred low disc temperatures in superoutburst are inconsistent with expectations of the disc-instability model, and lead to the conclusion that the outbursts of V4140~Sgr are powered by the only other mechanism considered so far, namely, bursts of enhanced mass transfer rate from its donor star. The dominant source of uncertainty in the computed $\alpha(R)$ distribution are the errors in the relative amplitude of the high-frequency, disc-related flickering component. These errors can be reduced by increasing the $S/N$ of the flickering curves, i.e., by increasing the statistics and/or the $S/N$ of the individual light curves included in the data sample. The benefits of this observational effort are double: it will allow a more meaningful and statistically robust estimate of the magnitude of the disc viscosity parameter and it will open the possibility to properly probe changes in $\alpha$ with radius. | 16 | 9 | 1609.03453 |
1609 | 1609.03179_arXiv.txt | {About 20 percent of all nearby early-type galaxies ($M_{\star} \gtrsim~6~\times 10^{9}~\mathrm{M}_{\odot}$ ) outside the Virgo cluster are surrounded by a disc or ring of low-column-density neutral hydrogen (\hi) gas with typical radii of tens of kpc, much larger than the stellar body. In order to understand the impact of these gas reservoirs on the host galaxies, we analyse the distribution of star formation out to large radii as a function of \hi \ properties using GALEX UV and SDSS optical images. Our sample consists of 18 \hi-rich galaxies as well as 55 control galaxies where no \hi \ has been detected. In half of the \hi-rich galaxies the radial UV profile changes slope at the position of the \hi \ radial profile peak. To study the stellar populations, we calculate the FUV-NUV and UV-optical colours in two apertures, 1-3 and 3-10 $R_\mathrm{eff}$. We find that \hi-rich galaxies are on average 0.5 and 0.8 mag bluer than the \hi-poor ones, respectively. This indicates that a significant fraction of the UV emission traces recent star formation and is associated with the \hi \ gas. Using FUV emission as a proxy for star formation, we estimate the integrated star formation rate in the outer regions ($R > 1R_\mathrm{eff}$) to be on average $\sim$ 6$\times$10$^{-3}$ M$_{\odot}~\mathrm{yr}^{-1}$ for the \hi-rich galaxies. This rate is too low to build a substantial stellar disc and, therefore, change the morphology of the host. We find that the star formation efficiency and the gas depletion time are similar to those at the outskirts of spirals. } | \label{sec:introduction} Since a long time evidence has appeared in the literature that early-type galaxies (ellipticals and lenticulars, hereafter ETGs) harbour central discs \citep{1951ApJ...113..413S, 1970ApJ...160..831S, 1976ApJ...206..883V}. In recent years, this has been quantified by several optical studies, based on broad band imaging and integral-field spectroscopy \citep[e.g.,][]{2011MNRAS.413..813C, 2011MNRAS.418.1452L, 2012ApJS..198....2K, 2013MNRAS.432.1768K,2014MNRAS.444.3340W}. Since they are relatively shallow and have a small field of view, these studies are limited to the inner regions of nearby ETGs. However, deep optical imaging, together with the kinematics of planetary nebulae and of globular clusters, have revealed that these discs can extend much further out \citep{2003Sci...301.1696R,2013MNRAS.428..389P,2014MNRAS.440.1458D,2015MNRAS.446..120D}. Rotating discs are often detected in the neutral hydrogen (\hi) gas phase, too \citep[e.g.,][]{2006MNRAS.371..157M, 2007A&A...465..787O, 2010MNRAS.409..500O}. The most complete census of \hi \ discs in ETGs to date was obtained by \citet[][hereafter S12]{2012MNRAS.422.1835S} as part of the \atlas \ survey\footnote{http://www-astro.physics.ox.ac.uk/atlas3d/}. Their results show that $\sim$20 percent of all nearby ETGs outside the Virgo cluster are surrounded by a low-column-density \hi \ disc with typical size of many tens of kpc, much larger than the stellar body. Some of these \hi \ discs are thought to originate from the accretion of gas-rich satellites as indicated by the misalignment of their angular momentum vectors with respect to those of the stellar discs \citep{2014MNRAS.444.3388S}. Some of the kinematically aligned gas discs may originate from the cooling of hot gas in the halo \citep[][]{2015MNRAS.451.1212N}. Whatever their origin, this paper is concerned with the impact of these large \hi \ discs on the properties of the host ETGs: do these \hi \ discs host any star formation (SF), and if so what impact does this have on the morphology and stellar populations of the host galaxy? \begin{table*} \caption{General properties of the \hi-rich sample} \begin{center} \begin{tabular}{lccccccccccc} \hline \hline No &Name & $D$ &$V_{hel}$ &$M_{K}$ &log$_{10}$($R_{eff}$) &$P.A.$ &$\varepsilon$ &$E(B-V)$ &log$_{10}M$(\hi) &log$_{10}M_{\star, r}$\\ & &[Mpc] &[km/s] &[mag] &[arcsec] &[degrees] & [ ] &[mag] &[M$_{\odot}$] &[M$_{\odot}$]\\ & (1) &(2) &(3) &(4) &(5) &(6) &(7)& (8) &(9) &(10)\\ \hline 1 &NGC~2594 &35.1 &2362 &-22.36 &0.82 &306 &0.122 &0.050 &8.91 &10.47\\ 2 &NGC~2685 &16.7 &875 &-22.78 &1.41 &37 &0.402 &0.051 &9.33 &10.31\\ 3 &NGC~2764 &39.6 &2706 &-23.19 &1.09 &173 &0.218 &0.034 &9.28 &10.64\\ 4 &NGC~2859 &27.0 &1690 &-24.13 &1.43 &269 &0.272 &0.017 &8.46 &10.97\\ 5 &NGC~3414 &24.5 &1470 &-23.98 &1.38 &135 &0.079 &0.021 &8.28 &11.11\\ 6 &NGC~3522 &25.5 &1228 &-21.67 &1.01 &241 &0.504 &0.020 &8.47 &10.31\\ 7 &NGC~3619 &26.8 &1560 &-23.57 &1.42 &72 &0.020 &0.013 &9.00 &10.91\\ 8 &NGC~3941 &11.9 &930 &-23.06 &1.40 &189 &0.338 &0.018 &8.73 &10.34\\ 9 &NGC~3945 &23.2 &1281 &-24.31 &1.45 &159 &0.435 &0.023 &8.85 &11.02\\ 10 &NGC~4036 &24.6 &1385 &-24.40 &1.46 &272 &0.507 &0.019 &8.41 &11.16\\ 11 &NGC~4203 &14.7 &1087 &-23.44 &1.47 &206 &0.153 &0.012 &9.15 &10.60\\ 12 &NGC~4262 &15.4 &1375 &-22.60 &1.10 &216 &0.427 &0.031 &8.69 &10.48\\ 13 &NGC~4278 &15.6 &620 &-23.80 &1.50 &50 &0.230 &0.023 &8.80 &11.08\\ 14 &NGC~5103 &23.4 &1273 &-22.36 &1.02 &108 &0.278 &0.015 &8.57 &10.29\\ 15 &NGC~5173 &38.4 &2424 &-22.88 &1.01 &261 &0.214 &0.024 &9.33 &10.42\\ 16 &NGC~5631 &27.0 &1944 &-23.70 &1.32 &34 &0.198 &0.017 &8.89 &10.89\\ 17 &UGC~03960 &33.2 &2255 &-21.89 &1.24 &108 &0.645 &0.040 &7.79 &10.39\\ 18 &UGC~09519 &27.6 &1631 &-21.98 &0.87 &198 &0.402 &0.018 &9.27 &10.06\\ \hline \hline \end{tabular} \label{table:HIrich} \end{center} \begin{tablenotes}[para,flushleft]\footnotesize Note.$-$ Column (1): The name is the principal designation from LEDA \citep{2003A&A...412...45P}. Column (2): distance in Mpc \citet{2011MNRAS.413..813C}. Column (3): heliocentric velocity \citet{2011MNRAS.413..813C}. Column (4): total galaxy absolute magnitude derived from the apparent magnitude in K~band \citet{2011MNRAS.413..813C}. Column (5): projected half-light effective radius \citet{2011MNRAS.413..813C}. Column (6): position angle of the galaxy calculated from the \hi \ image \citet{2014MNRAS.444.3388S}. Column (7): ellipticity of the galaxy calculated from the \hi \ image \citet{2014MNRAS.444.3388S}. Column (8): estimates of Galactic dust extinction \citet{2011ApJ...737..103S}. Column (9): total \hi \ mass calculated assuming galaxy distances given in column~(2) \citet{2012MNRAS.422.1835S}. Column (10): stellar mass calculated by using a mass-to-light ratio and a total luminosity in the r-band \citep{2013MNRAS.432.1709C}. \end{tablenotes} \end{table*} The relation between star formation (SF), \hi \ and colour for different types of galaxies has been known for decades \citep[see, e.g.,][and references in]{1994ARA&A..32..115R}. In late-type galaxies we know that the \hi \ content is an important driving factor of the star formation history. For example, \citet{2009MNRAS.396L..41H} show that \hi \ deficient spiral galaxies are at least 1 mag redder in UV-infrared colours than gas-rich ones, suggesting that quenching of SF and depletion of cold gas happen together. Furthermore, \citet{2011MNRAS.412.1081W} find that late-type galaxies with an \hi \ content larger than average have bluer outer regions. Similarly, \citet{2014ApJ...793...40H} report that most of the massive \hi \ galaxies $-$ mostly blue spirals$-$ detected in the ALFALFA survey have strong colour gradients, being bluer in the outer regions. These blue colours indicate existence of young stellar populations, and therefore, a link between the presence of \hi \ and relatively recent SF. Such link is known to exist in the \hi-dominated outer regions of spirals and in dwarf galaxies \citep[e.g.,][]{2008AJ....136.2846B, 2010AJ....140.1194B,2011AJ....142...37S,2012A&A...545A.142B,2015MNRAS.449.3700R}. Similar to late-type galaxies, evidence of SF is found in the outer regions of ETGs \citep{2009AJ....137.5037D, 2010ApJ...714L.290S, 2010ApJ...714L.171T, 2011ApJ...733...74L, 2012ApJ...755..105S}. For example, \citet{2012ApJ...755..105S} find extended UV discs in 76 percent of their ETGs (mostly UV rings with diameters of tens of kpc). In addition, \citet{2012ApJ...745...34M} find that $\sim$~42 percent of their ETGs with stellar mass between $10^{8}~\mathrm{and}~4\times10^{10}$ M$_{\odot}$ are UV-bright. These objects contain \hi \ and have blue UV-infrared colours together with enhanced SF outside one effective radius. Their results also support the idea that galaxy evolution can proceed from early- to late-type as discussed by, e.g., \citet{2009MNRAS.400.1225C}, \citet{2012ApJ...761...23F} and \citet{2012ApJ...755..105S}. In this paper, we aim to further explore the spatially-resolved link between \hi \ and SF in the outer regions of ETGs by analysing the UV, optical and \hi \ images of galaxies in the S12 sample as well as a control sample of \hi-poor control ETGs. For this purpose we use the \hi \ images published by S12, UV imaging from the Galaxy Evolution Explorer (GALEX) and optical imaging from the Sloan Digital Sky Survey (SDSS). In Section 2, we describe the \hi-rich and -poor sample as well as their selection criteria. In section 3, we describe the data (\hi, FUV, and optical imaging) and the methodology used for our work. In particular, we study the effect of the GALEX point spread function (PSF) on our results. In Section 4, we discuss the UV and optical properties of ETGs in our sample as a function of \hi \ content. In section 5, we compare star formation rate (SFR) with the studies in the literature. We also compare the integrated SF and SF efficiency in the \hi \ disc with that of late-type and dwarf galaxies. In Section 6, we present the conclusions of our work. | \label{sec:conclusions} In this paper we have investigated the relation between \hi, star formation and colours of the outer regions of ETGs by comparing an \hi-rich sample to an \hi-poor one. We have used spatially resolved \hi \ images together with the GALEX UV and SDSS $g,r$ band images. We have used two apertures to study the outer regions: 1-3 and 3-10 R$_{eff}$. In addition to the fixed apertures, we also study the SF as a function of \hi \ column density. Our main conclusions are given below. 1. Beyond 1~R$_{eff}$, \hi-rich ETGs are bluer (in UV-optical colours) than \hi-poor control ETGs. This holds also at fixed stellar mass for $M_{\star}<6 \times 10^{10}~\mathrm{M}_{\odot}$. In some extreme ETGs the outer colour is comparable to that of late-type galaxies.\\ \indent2. The \hi-rich galaxies have much stronger colour gradients --between the first and second apertures-- than the \hi-poor control galaxies. This means that the presence of the \hi \ has led to star formation in the outer parts of ETGs.\\ \indent3. In $\sim$ 89 percent of the \hi-rich ETGs, the \hi \ mass is higher in the second aperture than in the first. On average the \hi \ mass in the second aperture is 5 times higher than in the first aperture.\\ \indent4. Ten \hi-rich galaxies show increasing \hi \ profiles. In relation to that the peak \hi \ column density for 9 of these \hi-rich ETGs is beyond the SF threshold radius. More importantly, in five of the cases, at the same position, the FUV or NUV profile shows a change in their slope.\\ \indent5. The SFR surface density in the in the first aperture is almost 1 dex higher than in the second aperture: 9.1 and $1.3 \times 10^{-5}~\mathrm{M}_{\odot}~\mathrm{yr}^{-1}~\mathrm{kpc}^{-2}$, respectively.\\ \indent6. We have found that outermost regions of 8 \hi-rich ETGs reside between the blue-cloud and red-sequence. This situation is likely to persist for a long time due to their low efficient SF: the average SF efficiency of the second aperture is $1\times10^{-11}$~yr$^{-1}$. Although ETGs are quite different from late-type or dwarf galaxies in their central regions, they are very similar when considering the \hi \ dominated outer regions. Early-types, spirals and dwarfs all show a similar $\Sigma$~SFR vs $\Sigma$~\hi \ relation, which shows that SFR is increasing with increasing \hi \ surface density. Additionally, they show similarly low SF efficiency. The gas depletion time for the outermost regions of ETGs and late-type galaxies is almost the same: $\sim10^{11}$~yr. Another meaning of these results could be that forming stars in a low column density region does not depend on the host galaxy type. | 16 | 9 | 1609.03179 |
1609 | 1609.03723_arXiv.txt | {} {Although the temporal evolution of active regions (ARs) is relatively well understood, the processes involved continue to be the subject of investigation. We study how the magnetic field of a series of ARs evolves with time to better characterise how ARs emerge and disperse.} {We examine the temporal variation in the magnetic field distribution of 37 emerging ARs. A kernel density estimation plot of the field distribution was created on a log-log scale for each AR at each time step. We found that the central portion of the distribution is typically linear and its slope was used to characterise the evolution of the magnetic field. } {The slopes were seen to evolve with time, becoming less steep as the fragmented emerging flux coalesces. The slopes reached a maximum value of $\sim -1.5$ just before the time of maximum flux before becoming steeper during the decay phase towards the quiet Sun value of $\sim -3$. This behaviour differs significantly from a classical diffusion model, which produces a slope of $-1$. These results suggest that simple classical diffusion is not responsible for the observed changes in field distribution, but that other processes play a significant role in flux dispersion.} {We propose that the steep negative slope seen during the late decay phase is due to magnetic flux reprocessing by (super)granular convective cells. } | \label{sect_Introduction} \idea{Evolution in the convective zone} The evolution of active regions (ARs) in time is a well studied topic in solar physics \citep[see][]{vanDriel15}, yet the processes are still not fully understood. ARs result from the emergence of buoyant magnetic flux through the photosphere. This magnetic flux originates at the tachocline and rises through the convection zone in the form of a magnetic flux tube. During the rise, plasma drains into the flux tube legs, with conservation of angular momentum causing more plasma to drain into the following leg. This distorts the shape of the flux tube and causes an increase (decrease) in magnetic pressure of the leading (following) leg, as a result of the total pressure balance. \idea{Crossing the photosphere} When the flux tube reaches the base of the photosphere, its environment changes dramatically; % in particular it is no longer buoyant. Magnetic field accumulates here until the undulatory instability or convective upward motions allow fragments of the field to rise and break through the photosphere in a series of small magnetic loops \citep[\eg][ and references therein]{Pariat04}. The opposite polarities of these loops diverge and the like polarities of many of these small loops coalesce to form strong concentrated spots. The higher magnetic pressure of the leading flux tube leg means that the leading polarity forms a stronger, more compact spot(s) than the following polarity. This process of fragmented emergence followed by coalescence has been well observed \citep[\eg][]{Zwaan78, Strous96} and also modelled \citep[\eg][]{Cheung10}. \idea{Longterm evolution} During the emergence phase, the two polarity centres diverge, with their separation reaching a plateau around the time the AR achieves its peak flux. This indicates that the flux tube is no longer emerging. In addition, convective motions of supergranular cells advect the region's field, breaking it apart. Other processes may also play a role in the dispersion of the AR. Moving magnetic features are magnetic flux fragments that are observed to move radially outward from sunspots advected by the moat flow \citep{Harvey73}. They may also contribute to the removal of flux from the spot, although there is not yet conclusive evidence for this \citep{vanDriel15}. Cancellation of AR flux with the background field also contributes to the removal of magnetic flux from the photosphere, as well as cancellation between the two opposite polarities along the internal polarity inversion line of the AR. The decay phase of ARs is much longer than the emergence phase and can last for several weeks \citep[\eg][]{Hathaway08} or even months \citep[\eg][]{vanDriel99}, with the weaker following spot decaying much faster than the leading spot. \idea{Roadmap} In this study, we analyse the distribution (probability density) of the vertical component of the photospheric magnetic field (or flux density), and below we refer to it simply as the field distribution. This is a new method of characterising the evolution of ARs, by looking at changes in the field distribution as the regions evolve. Our analysis is different from the previous studies which analyse the magnetic flux distributions of photospheric magnetic clusters \citep[\eg][]{Parnell2009, Gosic2014}, % as we do not cluster the photospheric magnetic field in magnetic entities. \Sect{theory} describes the theoretical background of emergence, clustering (merging) and diffusion with a focus on the magnetic field distribution expected with these physical processes. The data used for the observational study are described in \sect{data_methods}, along with the AR area selection code, which defines the pixels used to calculate the field distribution. The field distribution plots and their characterisation are explained in \sect{KDE_plots}. Sections \ref{sect_Obs_Temporal} and \ref{sect_Obs_Flux} show the statistical results of the characterisation of 37 ARs. The characterisation reflects the different evolutionary stages; fragmented emergence, coalescence to form strong sunspots and gradual dispersion of the AR. Then, in Sections \ref{sect_Obs_Decayed} and \ref{sect_Obs_QS} we explore the dispersing phase of ARs as well as the quiet Sun. We next investigate possible issues present for the derived field distributions in \sect{Obs_Issues}. Finally, the observational and theoretical results are discussed and compared, allowing conclusions to be drawn in \sect{Conclusions}. | \label{sect_Conclusions} \idea{Summary of the results} In this study of 37 emerging ARs, we have shown that there is a relationship between the slope of the vertical component of the photospheric field distribution and the age of an AR. This is summarised in \Fig{prob_evolution}. At the beginning of a region's emergence, the slope is steep and negative. The slope becomes less steep, which indicates the coalescence of the fragmented flux that emerges. Later, the slope reaches a maximum just before the region achieves its peak flux value (at $\sim 0.75-0.8$ peak flux), before the decay processes become dominant. The slope becomes more negative as the region disperses and this decreasing trend continues towards the quiet Sun slope value of $\sim -3$ (\Fig{decaying_region}). \idea{Model the strong field} A comparison between the observational and theoretical results shows that a simple model of magnetic concentrations can describe the field (flux density) distribution in emerging ARs during the coalescence phase when smaller flux concentrations merge to form larger ones, leading to sunspot formation. The model predicts a slope of $\approx -1.67$ for $n=3$, in good agreement with the slope values found in observations of the coalescence phase (\Figs{gradvstime}{gradvsf(flux)}). \idea{Difficulties with the decay} However, later on there is a major deviation from the classical-diffusion model in the decay phase, indicating that AR magnetic fields do not disperse by simple diffusion. The latter predicts that after reaching peak flux, the field strength distribution should be characterised by a slope which is evolving from the range $[-1.6,-1.4]$ towards the diffusion exponent value of $-1$. However, as \Figs{gradvsf(flux)}{gradnearcm} clearly demonstrate, once ARs pass their peak flux and start decaying, their field strength distribution slopes evolve quite differently from these expectations: they start to attain higher negative values. Furthermore, ARs measured in the later decay phase display slopes in the range of $[-2.3,-1.6]$, as shown in \Fig{slope_vs_days}, while the quiet Sun, which can be regarded as the end-product of AR decay, shows a slope $\approx -3$. How can we understand this behaviour, which is so clearly opposite to the classical diffusion scenario? \idea{Effect of the convection: proposed scenario} We suggest that magnetic flux reprocessing by convective cells is responsible for the observed evolution of field distributions. Magnetic flux is being gnawed away by granular and supergranular convective cells, in agreement with the turbulent diffusion model \citep[\eg][]{Petrovay97}, which carry away flux concentrations from the strong-field area of ARs. \citet{Petrovay97_obs} analysed various theoretical models of sunspot decay using observational data of umbral areas and found that the turbulent diffusion model was the only model supported by the data. The turbulent diffusion model is also consistent with the removal of active region magnetic field by moving magnetic features \citep[\eg][]{Kubo08}, the movement of which is a result of convective flows. The advected field fragments become concentrated along the boundaries of supergranular cells, where they occasionally meet and cancel with opposite polarity field. The cancelled flux submerges and is re-processed by convection. Part of the reprocessed flux emerges in the centre of supergranular cells as weak intranetwork flux. This process breaks up strong-field flux tubes and makes them emerge as weaker field, effectively transferring strong field to weak field in the field distribution of a decaying AR, making the slope of the field distribution steeper (more negative). The weak intranetwork field is carried to the boundary of the supergranular cells and becomes more concentrated there, which is a counter-mechanism to the former scenario. We propose that the $-3$ slope found in the quiet-Sun field is the result of the combination of these two mechanisms. \idea{Compatibility for global dispersion of ARs?} Although the evolution of the magnetic field in a decaying AR is very different from what is expected in classical diffusion (slope of $-1$), % the AR area has previously been found to increase approximately linearly with time \citep[\eg][]{vanDriel-Gesztelyi03}, in accordance with the classical diffusion model. Is there a contradiction here? We propose that the increase of AR area is due to the non-stationarity of the supergranular convective cells, which have a lifetime of about one day. This non-stationarity creates a random walk of the flux tubes, which is the underlying physics of classical diffusion. Therefore the AR evolution can be seen as classical diffusion regarding area increase, while the magnetic field distribution is governed by magneto-convection. These ideas have to be tested in simulations of emerging ARs that include (or not) magneto-convection. Other studies \citep[\eg][]{Abramenko10} have characterised the active region magnetic field by its degree of intermittency. It could be interesting to compare this with changes in the field distribution slopes as the resolution of the magnetogram is decreased and particularly for resolutions which do not resolve the highly intermittent small scale structure. | 16 | 9 | 1609.03723 |
1609 | 1609.03209_arXiv.txt | \label{introduction} This contribution constitutes but a small part of the two-day splinter session entitled ``Star clusters from space, from the ground, and over time'' which took place at the 19th Cambridge Workshop on Cool Stars, Stellar Systems, and the Sun in Uppsala, Sweden in June 2016. Given the nature of this contribution it will be far from comprehensive, both in terms of the range of fundamental stellar parameters covered, but also in terms of the various methods employed to estimate such parameters (for more information on these the reader is referred to other Cool Stars 19 proceedings contributions). A recent and more comprehensive discussion on star clusters which not only covers global properties of both young and old clusters, but also our current understanding of the dynamical evolution of clusters, can be found in the proceedings of the Ecole Evry Schatzman 2015 (EES2015) school ``Stellar clusters: benchmarks of stellar physics and galactic evolution''. Clusters have long represented benchmarks with regard to the determination of fundamental stellar parameters, in large part due to the underlying assumption that members within such ensembles share several common properties; namely they are coeval, have the same chemical composition and are located at roughly the same distance. In addition to these shared characteristics, it is also observationally advantageous to focus on clusters as they have a significantly higher stellar number density (per unit area on the sky) compared to either young associations/moving groups or field stars and so for a given allocation of telescope time one can thus maximise the number of stars in ones sample. Studies of clusters have also been instrumental in driving our understanding of the formation and evolution of stars. By studying a given cluster we can infer the mass dependence of astrophysical phenomena at a given epoch and by studying several clusters spanning a range of ages we can track how such phenomena evolve with time, as well as investigate second-order effects such as the local environment. | \label{concluding_thoughts} Below I briefly reiterate the main conclusions from this contribution. \begin{enumerate} \item The colour-magnitude diagram of a given cluster can provide several global parameters shared by constituent members (including age, distance and the presence/uniformity of interstellar extinction), although one should be aware of the underlying uncertainties as regards the use of stellar evolutionary models and carefully assess the pedigree of adopted empirical relations. \item Model-dependent estimates of low-mass pre-MS stellar parameters are unreliable. Specifically, below 1\,\msun\ the models tend to underestimate the mass of a given star based on its position in the H-R diagram as well as predict radii which are too small based on its position in the mass-radius diagram. \item There is tentative evidence that the introduction of magnetic field-related phenomena (such as star spots and/or the inhibition of convective flows) may help to resolve the discrepancy between dynamically-determined parameters and those predicted by stellar evolutionary models. \item Recent spectroscopic survey results have demonstrated that clusters are more complex entities than previously thought (e.g. age spreads and kinematic substructure) and the continuing \emph{Gaia}-ESO and APOGEE/IN-SYNC surveys will only highlight further examples of this and hence continue to shape our understanding of the formation and early evolutionary stages of young clusters. \end{enumerate} | 16 | 9 | 1609.03209 |
|
1609 | 1609.00931_arXiv.txt | We construct a scenario where the outburst of the young-stellar-object ASASSN-15qi is an intermediate luminosity optical transient (ILOT). In this scenario a sub-Jupiter young planet was tidally destructed on to a young main-sequence star. The system is young, therefore the radius of the planet is larger than its final value, and consequently its density is smaller. The lower density allows the tidal destruction of the young Saturn-like planet on to the main-sequence star of mass $\approx 2.4 \rmModot$, resulting in a formation of a disc and a gravitationally-powered ILOT. Unlike the case of the more energetic ILOT V838~Mon, the mass of the destructed planet is too low to inflate a giant envelope, and hence the merger remnant stays hot. If our suggested model holds, this ILOT possesses two interesting properties: (1) its luminosity and total energy are below those of novae, and (2) it is not as red as other ILOTs. The unusual outburst of ASASSN-15qi, if indeed is an ILOT, further increases the diversity of the already heterogeneous group of ILOTs. We mark the region on the energy-time diagram occupied by such young ILOTs. | \label{sec:intro} \subsection{General} \label{subsec:general} As in recent years the quality and quantity of sky surveys increase, more attention is given to rare explosions and outbursts in the energy gap between novae and supernovae (e.g. \citealt{Mouldetal1990, Rauetal2007, Ofeketal2008, Ofeketal2016, Prietoetal2009, Botticella2009, Smithetal2009, Berger2009a, Berger2009b, KulkarniKasliwal2009, Mason2010, Pastorello2010, Kasliwaletal2011, Tylendaetal2013, Kasliwaletal2011, Kurtenkovetal2015, Tartagliaetal2016, Villaretal2016, Blagorodnovaetal2017}). Those outbursts, known as intermediate luminosity optical transients (ILOTs) form an extended family that has a number of subgroups (see \citealt{KashiSoker2016} for a detailed nomenclature: Intermediate-Luminous Red Transients, LBV giant eruptions and SN Impostors, and Luminous Red Novae or Red Transients or Merger-bursts). Researchers have been modeling ILOTs, or subgroups of ILOTs, either as single-star phenomena (e.g., \citealt{Thompsonetal2009, Kochanek2011} for eruptive red giants and \citealt{Ofeketal2013} for a SN impostors), or as interacting binary systems (\citealt{Kashietal2010, KashiSoker2010b, SokerKashi2011, SokerKashi2012, SokerKashi2013, McleySoker2014, Nandezetal2014, Goranskijetal2016, Pejchaetal2016b, Soker2016}), including a common envelope evolution \citep{RetterMarom2003, Retteretal2006, Tylendaetal2011, Ivanovaetal2013, IvanovaNandez2016, Tylendaetal2013, Nandezetal2014, Kaminskietal2015b, Soker2015, MacLeodetal2016, Blagorodnovaetal2017}. There are two main diagrams to characterize ILOTs. One is the peak luminosity versus eruption duration (time scale; \citealt{Rauetal2009,Kasliwal2013}), and the second diagram is that of the total eruption energy versus eruption time. The latter is called the energy-time diagram\footnote{An updated version of the energy time diagram is available at \url{http://phsites.technion.ac.il/soker/ilot-club/}}, and we present it in Fig. \ref{fig:etd}. \begin{figure*} \centering \includegraphics[trim= 0.8cm 0.1cm 1.5cm 0.5cm,clip=true,width=1.0\textwidth]{etd.eps} \caption{ Observed transient events on the energy time diagram. Blue empty circles represent the total (radiated plus kinetic) energy of the observed transients as a function of the duration of their eruptions, i.e., usually the time for the visible luminosity to decrease by 3 magnitudes. The Optical Transient Stripe is populated by ILOT events that we suggest are powered by gravitational energy of complete merger events or vigorous mass transfer events (\citealt{KashiSoker2010b, KashiSoker2016, SokerKashi2016}). For ILOTs that had sufficient data to create a model to calculate their total available energy, we mark it by a black asterisk above, or overlapping with, the blue circle. The total energy does not include the energy that goes to lifting the envelope and does not escape from the star. Novae models are marked with a green line \citep{dellaValleLivio1995}, with red crosses \citep{Yaronetal2005}, or with diamonds \citep{Sharaetal2010}. The four horizontal lines represent planetary nebulae (PNe) and pre-PNe that might have been formed by ILOT events \citep{SokerKashi2012}. Merger models of a planet with a planet/BD/star \citep{Bearetal2011} are shown on the left hand side, together with models we added for smaller merging planet with mass $0.1 \rmMJ$. The lower-left part (hatched in green) is our \it{new extension for younger objects}, including ASASSN-15qi (red square), where the planets are of lower density and can more easily undergo tidal destruction. } \label{fig:etd} \end{figure*} In the present study we deal with a subgroup of ILOTs that are powered by gravitational energy which is released from a complete merger process of two stars (termed Luminous Red Novae, or Red Transients, or Merger-bursts), such as V838~Mon \citep{SokerTylenda2003, TylendaSoker2006} and V1309~Sco \citep{Tylendaetal2011, Nandezetal2014, Kaminskietal2015a}. Differently from the objects above, in this study we discuss the destruction of a planet on a star, rather than a star on a star. \subsection{Star-Planet Mergers} \label{subsec:merger} \cite{Bearetal2011} propose that a V838~Mon like merger-burst can happen on smaller scales and low energies -- between a planet and a low mass main-sequence (MS) star, between a planet and a brown dwarf (BD) or between two planets. In this process the planet is tidally shredded into a disc, and the accretion of the gas in the disc onto the star, onto a brown dwarf, or onto another planet, leads to an outburst. According to this model, the destruction of the planet occurs before it touches the more massive object, because the density of the planet is lower than the density of the more massive object. For a typical mass of the destroyed planet of $\approx 1 \rmMJ$, these outbursts populate the lower left part of the optical transient stripe, with timescales of a few days and total energies of $10^{44.3}$--$10^{46.3} \erg$ \citep{Bearetal2011}. We note that in the triple-planet scenario proposed by \cite{RetterMarom2003} and \cite{Retteretal2006} for the outburst of V838~Mon, the planets enter the envelope of the star intact, and hence their scenario is different than the scenario we propose here. The thorough study by \cite{Metzgeretal2012} further established the star-planet merger process as member of the ILOT heterogeneous group. They study the interaction between a Sun-like star and planets of masses of $1 \textrm{--} 10 \rmMJ$, and find that the ratio of the mean densities of the planet and the star determines the outcome. For low enough mean density of the planet, the interaction can lead to tidal-dissipation event where the planet transfers mass at about steady rate to the star. For density ratio in the range $\approx 1$--$5$, \cite{Metzgeretal2012} find that the mass transfer is unsteady, resulting in a dynamical disruption of the planet into an accretion disc around the star. \cite{Metzgeretal2012} further calculate the luminosity of such a transient to be in the order of $\approx 10^{37} \erg$ and the time scale to be in the order of a few weeks. \subsection{The transient ASASSN-15qi} \label{subsec:ASASSN} The All-Sky Automated Survey for Supernovae (ASAS-SN) variability survey \citep{Shappeeetal2014} discovered an outburst designated ASASSN-15qi (also referred to as 2MASS J22560882+5831040) on JD 2,457,298 (2 October 2015). \cite{Herczegetal2016} report the observational properties of ASASSN-15qi in detail. We present the light curve in V in Fig. \ref{fig:lightcurve}. The following make the ASASSN-15qi outburst interesting. (1) Its location among young objects, suggesting it is associated with a young stellar object (YSO). (2) It showed a fast brightening of 3.5 mag in the optical bands in less than 23 hours. (3) It blew a fast wind that faded as the outburst decayed over 4-5 months. \begin{figure*} \centering \includegraphics[trim= 0.0cm 0.1cm 0.5cm 0.0cm,clip=true,width=1.0\textwidth]{ASASSN-15qi_lightcurve_3.eps} \caption{ The light curve of ASASSN-15qi from the observations in Herczeg et al. (2016). A change from a sharp decline to a moderate slope occurs about 5--6 days after the peak. The red arrows indicate times at the light curve where a break in the slope of the light curve appears. } \label{fig:lightcurve} \end{figure*} \cite{Herczegetal2016} calculate the outburst radiative energy to be $E_{\rm rad} \approx 7 \times 10^{42} \erg$ over a duration of 6 months. The kinetic energy might be much larger than the radiated energy, and it might accounts for most of the energy of the outburst. \cite{Herczegetal2016} mention that other outbursts from young stellar objects, such as V899~Mon and Z~CMa, share some similar spectroscopic features with ASASSN-15qi. They speculated that ASASSN-15qi might be either a mass transfer event, connected to interactions between a star and a planet on an eccentric orbit, formation of an excretion disc, or be some kind of a magnetic reconnection and outflow event. The energy and timescale of ASASSN-15qi place this event on the energy-time diagram just below the region where \cite{Bearetal2011} predict the location of events where a merger of a planet with a low-mass main-sequence star take place. This close location motivates us to propose a planet-destruction model for the intermediate luminosity optical transient ASASSN-15qi. In addition to the event in 2015, \cite{Herczegetal2016} mention an earlier outburst in 1976. We here concentrate on the event of 2015. If our scenario holds, then two planets have been destructed on the star, one in 1976 and one in 2015. Surely two such events make our proposed scenario much rarer. Yet, there are now many known planetary systems with a large number of planets and with rich variety of properties. Most pronounced is the detection of 7 earth-like planets around a very low mass star \citep{Gillonetal2017}. Our proposed scenario requires a planetary system with several planets, as one or more planets should perturb the orbit of the planet to be destroyed. We note that a scenario for an ILOT as a result of triple-planet collisions have been proposed by \cite{Retteretal2006} for the outburst of V838~Mon. | \label{sec:summary} We examined the possibility that the unusual outburst of the YSO ASASSN-15qi is an ILOT event, similar in many respects to V838~Mon, but much fainter and of lower total energy. As in the model of V838~Mon, the erupting system was young, but unlike the model for V838~Mon, we here suggested that the secondary object that was tidally destroyed onto the primary main-sequence star was a Saturn-like planet rather than another low mass main-sequence star. The young age in both systems is crucial. The reason is that along most of the main-sequence the density decreases with increasing stellar mass. Therefore, low mass stars on the main-sequence will not suffer a tidal destruction by more massive primary main-sequence stars. This holds for sub-stellar secondary objects as well. In young systems, the time scale on which the object shrinks to its final radius is longer for smaller mass. Therefore, young low-mass objects have lower density than their final density, while the much more massive star can already be very close to its zero-age main sequence density. The density of the low mass secondary object can be low enough to allow tidal destruction if it comes close enough to the primary star. Very low-mass main-sequence stars and brown dwarfs in old systems have higher densities than more massive main sequence stars, and they can tidally destroy old planets, as well as young ones. This is indicated by the three lines in the lower left of Fig. \ref{fig:etd}. The destruction of young planets on stars results in weak ILOTs, as we suggest here for ASASSN-15qi. Such events occupy the lower left part of the energy time diagram (hatched-green region in Fig. \ref{fig:etd}). The typical luminosity of these events is not necessarily above novae, and they do not have a red atmosphere. So they are not really intermediate between the luminosities of novae and supernovae, and they are not internally red. The difference in the mass of the destroyed object between V838~Mon and ASASSN-15qi brings another significant difference in the properties of the merger product. The luminosity of the outburst in both cases is in the range of giant stars. In the case of V838~Mon, part of the mass of the destroyed secondary star inflated a huge envelope. In our proposed scenario for the 2015 outburst of ASASSN-15qi, the mass that is available to inflate an envelope is $<0.001 \rmModot$. As is known from the evolution of post-AGB stars, this mass is too low to build a large envelope. This is the reason that ASASSN-15qi had a relative small radius during the event, and hence it was relatively hot and not red. ILOTs constitute an heterogeneous group of objects that covers a large area in the energy-time diagram or the luminosity-time diagram. We here suggested that this group, that is powered by gravitational energy in binary systems, is more heterogeneous and covers a larger area in these diagrams than the conventional view until now was. | 16 | 9 | 1609.00931 |
1609 | 1609.05808_arXiv.txt | {}% {High Synchrotron Peaked blazars (HSPs) dominate the \gr\ sky at energies larger than a few GeV; however, only a few hundred blazars of this type have been catalogued so far. In this paper we present the 2WHSP sample, the largest and most complete list of HSP blazars available to date, which is an expansion of the 1WHSP catalog of $\gamma$-ray source candidates off the Galactic plane.} {We cross-matched a number of multi-wavelength surveys (in the radio, infrared and X-ray bands) and applied selection criteria based on the radio to IR and IR to X-ray spectral slopes. To ensure the selection of genuine HSPs we examined the SED of each candidate and estimated the peak frequency of its synchrotron emission (\nupeak) using the ASDC SED tool, including only sources with \nupeak~$> 10^{15}$~Hz (equivalent to \nupeak~$> 4$~eV).} {We have assembled the largest and most complete catalog of HSP blazars to date, which includes 1691 sources. A number of population properties, such as infrared colours, synchrotron peak, redshift distributions, and $\gamma$-ray spectral properties, have been used to characterise the sample and maximize completeness. We also derived the radio $\log$N-$\log$S distribution. This catalog has already been used to provide seeds to discover new very high energy objects within {\it Fermi}-LAT data and to look for the counterparts of neutrino and ultra high energy cosmic ray sources, showing its potential for the identification of promising high-energy \gr\, sources and multi-messenger targets.} {} | Blazars are a class of radio-loud Active Galactic Nuclei (AGN) hosting a jet oriented at a small angle with respect to the line of sight \citep{Blandford1978, Antonucci1993, Urry1995}. The emission of these objects is non-thermal over most or the entire electromagnetic spectrum, from radio frequencies to hard $\gamma$-rays. The observed radiation shows extreme properties, mostly due to relativistic amplification effects. The observed Spectral Energy Distribution (SED) presents a general shape composed of two bumps, one typically located in the infrared (IR) and sometimes extending to the X-ray band and the other one in the hard X-ray to $\gamma$-rays. If the peak frequency of the synchrotron bump (\nupeak) in $\nu$ - $\nu$F$_{\nu}$ space is larger than $10^{15}$~Hz (corresponding to $\sim$ 4~eV), a source is usually called High Synchrotron Peaked (HSP) blazars\citep{Padovani1995,Abdo2010}. HSP blazars are also considered to be extreme sources since the Lorentz factor of the electrons radiating at the peak of the synchrotron bump $\gamma_{peak}$ are the highest within the blazar population, and likely of any other type of steady cosmic sources. Considering a simple SSC model where $\nu_{\rm peak}=3.2 \times 10^6 \gamma^{2}_{\rm peak} B \delta $ \citep[e.g.][]{Giommi2012a}, assuming $B=0.1$ Gauss and Doppler factor $\delta =10$, HSPs characterized by $\nu_{\rm peak} $ ranging between $ 10^{15}$ and $\gsim 10^{18}$~Hz demand $\gamma_{\rm peak} \approx 10^4-10^6$. The typical two-bump SED of blazars and the high energies that characterize HSPs imply that these objects occupy a distinct position in the optical to X-ray spectral index ($\alpha_{\rm ox}$) versus the radio to optical spectral index ($\alpha_{\rm ro}$) colour-colour diagram \citep{Stocke1991}. Considering the distinct spectral properties of blazars over the whole electromagnetic spectrum, selection methods based on $\alpha_{\rm ox}$ and $\alpha_{\rm ro}$ have long been used to search for new blazars. For example, \citet{Schachter1993} discovered 10 new BL Lacs via a multi-frequency approach with radio, optical and X-ray data, and their BL Lac nature with optical spectra. HSP blazars play a crucial role in very high energy (VHE) astronomy. Observations have shown that HSPs are bright and variable sources of high energy \gr\, photons (TeVCat)\footnote{http://tevcat.uchicago.edu} and that they are likely the dominant component of the extragalactic VHE background \citep{Padovani1993,Giommi2006,DiMauro2014,Giommi2015,Ajello2015}. In fact, most of the extragalactic objects detected so far above a few GeV are HSPs \citep[][see also TeVCat]{Giommi2009,Padovani2015a,Arsioli2015a,Fermi2fhl2015}. However, it is known that only a few hundred HSP blazars are above the sensitivity limits of currently available $\gamma$-ray surveys. For example, the 1WHSP catalog \citep[][hereafter Paper I]{Arsioli2015a}, which was the largest sample of HSP blazars when it was published, shows that out of the 992 objects in the sample, 299 have an associated \gr\, counterpart in the {\it Fermi} 1/2/3FGL catalogs. Nevertheless there is a considerable number of relatively bright HSPs which still lack a \gr\, counterpart. These are likely faint point-like sources at or below the {\it Fermi}-LAT, detectability threshold and were not found by the automated searches carried out so far. Indeed, \citet{Arsioli2016} have detected $\approx\,150$ new \gr\, blazars based on a specific search around bright WHSP sources, using over 7 years of {\it Fermi}-LAT Pass 8 data. In the most energetic part of the \gr\, band photons from high redshift sources are absorbed by the extragalactic background light (EBL) emitted by galaxies and quasars \citep{Dermer2011,Pfrommer2013,Bonnoli2015}. Therefore, the TeV flux can drop by a very large factor compared to GeV fluxes, making distant TeV sources much more difficult to detect. Paper I has shown that with the help of multi-wavelength analysis, HSP catalogs can provide many good candidates for VHE detection. The currently known HSP blazars are listed in catalogs such as the 5th {\it Roma-BZCAT} \citep[][hereafter 5BZCat]{Massaro2015}, the Sedentary Survey \citep{Giommi1999,Giommi2005,Piranomonte2007}, \citet{Kapanadze2013}, and Paper I. However, the number of known HSPs is still relatively small with less than $\approx 1000$ cataloged HSPs till now. Significantly enlarging the number of high energy blazars is important to better understand their role within the AGN phenomenon, and should shed light on the cosmological evolution of blazars, which is still a matter of debate. The 5BZCat is the largest compilation of confirmed blazars, containing 3561 sources, around 500 of which are of the HSP type. It includes blazars discovered in surveys carried out in all parts of the electromagnetic spectrum and is also based on an extensive review of the literature and optical spectra. The Sedentary survey comprises 150 extremely high X-ray to radio flux ratio $(\log f_{\rm x}/f_{\rm r}\geq 3\times10^{-10}~{\rm erg}~{\rm cm}^{-2}~{\rm s}^{-1}~{\rm Jy}^{-1})$ HSP BL Lacs. The sample was obtained by cross-matching the RASS catalog of bright X-ray sources \citep{Voges1999} and the NVSS 1.4~GHz radio catalog \citep{Condon1998}. \citet{Kapanadze2013} built a catalog of 312 HSPs with flux ratio $(f_{\rm x}/f_{\rm r}\geq 10^{-11}~{\rm erg}~{\rm cm}^{-2}~{\rm s}^{-1}~{\rm Jy}^{-1})$ selected from various X-ray catalogs, the NVSS catalog of radio sources, and the first edition of the $Roma-BZCAT$ catalog \citep{Massaro2009}. The 1WHSP sample relied on a pre-selection based on Wide-field Infrared Survey Explorer (WISE) IR colours, SED slope criteria, and \nupeak~$>10^{15}$ Hz. It includes 992 known, newly-identified, and candidate high galactic latitude ($b>|20^\circ|$) HSPs. In a series of papers \cite{Massaro2011,DAbrusco2012,Massaro2012} showed that most blazars occupy a specific region of the IR colour-colour diagram, which they termed the blazar strip. In Paper I we extended the blazar strip in the WISE colour-colour diagram to include all the Sedentary Survey blazars and called it the {\it Sedentary WISE colour domain} (SWCD). The SWCD is wider than the WISE blazar strip since it contains some blazars whose host galaxy is very bright, such as MKN421 (2WHSP J110427.3+381230) and MKN 501 (2WHSP J165353.2+394536). We understood from previous work that many low-luminosity HSP blazars have the IR colours dominated by the thermal component of the host giant elliptical galaxy. Therefore, a selection scheme adopting IR colour restrictions may work effectively for selecting cases where the non-thermal jet component dominates the IR band but is less efficient for selecting galaxy dominated sources (since they are spread over a larger area in the IR colour-colour plot). In the present paper we extend the previous 1WHSP catalog to lower Galactic latitudes ($b>|10^\circ|$) building the larger and more complete 2WHSP catalog including over 1600 blazars expected to emit at VHE energies by means of multi-frequency data. | \label{conclusion} We have assembled the 2WHSP catalog, currently the largest and most complete existing catalog of HSP blazars, by using a multi-frequency method and a detailed comparison with existing lists of \gr\, emitting blazars. 2WHSP extends the previous 1WHSP catalog \citep{Arsioli2015a} down to lower Galactic latitudes ($ |b|>10^{\circ} $) and to fainter IR fluxes. In addition, it includes all the bright known HSP blazars close to the Galactic plane. The 2WHSP sample includes 1,693 confirmed or candidates HSP blazars and was also put together to provide a large list of potential targets for VHE and multi-messenger observations. The average \nupeak\, for our catalog is $ \langle\log\nu_{\rm peak}\rangle = 16.22 \pm 0.02$ Hz and the average redshift is $\langle z \rangle=0.331 \pm 0.008$. We have shown that the SWCD region needs to be extended to include HSPs in which the host galaxy is dominant. Our radio $\log$N-$\log$S shows that the number of HSP blazars over the whole sky is $> 2,000$ and that HBL make up $\sim 10\%$ of all BL Lacs. Finally, we note that this catalog has already been used to provide seeds for the identification of new {\it Fermi}-LAT objects and to look for astrophysical counterparts to neutrino and UHECR sources \citep{Padovani2016,Resconi2016}, which proves the relevance of having a large HSP catalog for multi-messenger astronomy. | 16 | 9 | 1609.05808 |
1609 | 1609.02619_arXiv.txt | We have collected high-dispersion echelle spectra of red giant members in the twelve open clusters (OCs) and derived stellar parameters and chemical abundances for 26 species by either line equivalent widths or synthetic spectrum analyses. We confirm the lack of an age$-$metallicity relation for OCs but argue that such a lack of trend for OCs arise from the limited coverage in metallicity compared to that of field stars which span a wide range in metallicity and age. We confirm that the radial metallicity gradient of OCs is steeper (flatter) for R$_{\rm gc}<\,$12 kpc ($>$12 kpc). We demonstrate that the sample of clusters constituting a steep radial metallicity gradient of slope $-$0.052$\pm$0.011 dex kpc$^{-1}$ at R$_{\rm gc}<$ 12 kpc are younger than 1.5 Gyr and located close to the Galactic midplane ($\lvert\,z\rvert<\,$0.5 kpc) with kinematics typical of the thin disc. Whereas the clusters describing a shallow slope of $-$0.015$\pm$0.007 dex kpc$^{-1}$ at R$_{\rm gc}>$ 12 kpc are relatively old, thick disc members with a striking spread in age and height above the midplane (0.5$\,<\lvert\,z\rvert<\,$2.5 kpc). Our investigation reveals that the OCs and field stars yield consistent radial metallicity gradients if the comparison is limited to samples drawn from the similar vertical heights. We argue via the computation of Galactic orbits that all the outer disc clusters were actually born inward of 12 kpc but the orbital eccentricity has taken them to present locations very far from their birthplaces. | The formation and evolution of galaxies is one of the major puzzles of astrophysics. Large-scale cosmological simulations utilizing the currently favoured Lambda Cold Dark Matter ($\Lambda$CDM) paradigm of cosmology predicts the large-scale distribution of galaxies in space as observed today, but we have not yet met with equal success on small scales due to incomplete knowledge of how the baryons are distributed (Vogelsberger et al. 2014). As these baryons mostly reside in stellar discs in galaxies, understanding the formation and evolution of galactic discs is the heart of the galaxy formation theory (Feltzing 2015 and references therein). In this regard, the Milky Way Galaxy provides an excellent testing ground for models of galaxy formation and evolution thanks to the ability to resolve individual stars of its stellar populations and analyse them in exquisite detail. As the disc formed dissipatively and evolved dynamically, much of the dynamical information is lost. However, the chemical content of the disc stars may have preserved the dissipative history, the key to unraveling the formation history of the Milky Way (De Silva et al. 2009; Freeman \& Bland-Hawthorn 2002). Since the surface density of gas and star formation rate (SFR) in the Galactic disc have varied, as seen in many galaxies, the measured chemical abundances are a function of position as well. Therefore, measurement of the chemical abundance distribution in the disc (i.e. the radial abundance gradient) and the gradient's temporal variation over the disc's lifetime present strong observational constraints on the theoretical models of Galactic chemical evolution (GCE; Magrini et al. 2009; Minchev et al. 2013; Kubryk et al. 2015). Abundance gradients in the Milky Way can be measured using a wide variety of tracers including the young population of H\,{\sc ii} regions, OB stars, and Cepheid variables, and planetary nebulae of different ages, cool unevolved stars and red giants in the field (see for example, Magrini et al. 2010; Cheng et al. 2012; Hayden et al. 2014) and in open clusters (OCs; Friel et al. 2002, 2010; Magrini et al. 2010; Pancino et al. 2010; Yong et al. 2012; Frinchaboy et al. 2013). The cooler main sequence stars or red giants which are members of OCs provide not only abundance estimates for elements sampling all the major processes of stellar nucleosynthesis -- but a collection of stars whose age, distance, kinematics and metallicity can be measured with greater certainty than for field stars. Moreover, it is possible from an OC's space motion and a model of the Galactic gravitational potential to study the dynamics and estimate a cluster's birthplace (Wu et al. 2009; Vande Putte et al. 2010). In recent years, modern high-resolution spectrographs and large reflectors have provided high-quality spectra of stars in OCs for secure measures of abundances, and extended abundance estimates to distant OCs in the direction of Galactic anti-centre (Carraro et al. 2007; Yong et al. 2005, 2012; Sestito et al. 2006, 2008). However, not all the investigators have arrived at similar measures of the radial metallicity gradient which may be partly related to the different methods adopted (see for example, Magrini et al. 2009; Heiter et al. 2014). Our current understanding suggests a gradient of about $-$0.06 dex kpc$^{-1}$ (Friel et al. 2002; Pancino et al. 2010) to $-$0.20 dex kpc$^{-1}$ (Frinchaboy et al. 2013) in the radial range 5 to 10 kpc with a nearly flat trend ($-$0.02 dex kpc$^{-1}$) beyond 13 kpc (Carraro et al. 2007; Sestito et al. 2008; Pancino et al. 2010; Yong et al. 2012; Frinchaboy et al. 2013; Cantat-Gaudin et al. 2016). It is especially puzzling that the flattening of the metallicity distribution of OCs at large radii is not shown by the Galactic field stars (Cheng et al. 2012; Hayden et al. 2014) including Cepheids (Luck \& Lambert 2011; Genovali et al. 2014) where field stars suggest a constant steep decline of metallicity out to R$_{\rm gc}$ of 18 kpc. Another fascinating result, in common for both the field stars and OCs at large radii, is the evidence for enhanced [$\alpha$/Fe] ratios in the outer Galactic disc (Yong et al. 2005; Bensby et al. 2011; Luck \& Lambert 2011). The fact that the $\alpha$-enriched clusters older than 4$-$5 Gyr, reminiscent of rapid star formation history, have populated the Galactic disc beyond 12 kpc, where the star formation is slow due to low surface density of gas, is a puzzle in conflict with the inside-out formation scenarios of the Milky Way (Bovy et al. 2012; Ro\v{s}kar et al. 2013). A significant flattening of the metallicity gradient observed in the outskirts of other spiral galaxies (Scarano \& L\'{e}pine 2013) calls for a thorough understanding of many processes regulating the evolution of Milky Way. This paper extends our homogeneous abundance analysis of OCs to a new sample of OCs to aid in exploring further the evolution of the Galactic disc. This is our fourth paper reporting a comprehensive abundance measurements for red giants in twelve OCs in the radial 7.7 to 9.7 kpc with ages 71 to 730 Myr lacking detailed information on their chemical composition. The clusters selected for abundance analysis included: NGC 1647, NGC 1664, NGC 2099, NGC 2281, NGC 2287, NGC 2345, NGC 2437, NGC 2548, NGC 2632, NGC 6633, NGC 6940, and NGC 7209. All these clusters save NGC 2099, NGC 2287, NGC 2632 and NGC 6633 are analysed for the first time. The chemical content of NGC 2099 was previously measured for the elements Na, Al, $\alpha$-elements (Mg, Si, Ca, and Ti), iron-peak elements (Sc, V, Cr, Fe, Co, and Ni), and the $s$-process elements (Y, Ba, La, and Nd) by Pancino et al. (2010). The chemical composition of NGC 2632 was estimated previously for the elements Na, Al, Si, Ca, Ti, Cr, Fe and Ni by Pace, Pasquini \& Fran{\c c}ois (2008), and for 17 elements in the range from O to Ba by Yang, Chen \& Zhao (2015). The chemical abundances of none of the elements but for iron was explored previously for red giants in NGC 2287 and NGC 6633 by Santos et al. (2012) and Santos et al. (2009), respectively. For NGC 2099 and NGC 2632, we either provide a more extensive analysis of stars previously analysed or the first analysis of a member star while NGC 2287 and NGC 6633 are analysed comprehensively for the first time for many elements. We have presented in our previous papers (Reddy et al. 2012, 2013, 2015) abundance measurements for 16 OCs with ages from 130 Myr to 4.3 Gyr and R$_{\rm gc}$s from 8.3 and 11.3 kpc. We combine our sample of twenty-eight OCs (12 from this study and 16 from our previous papers) with a sample of 51 OCs drawn from the literature for which we have remeasured the chemical content using their published EWs, our models, linelists and reference solar abundances to establish a common abundance scale. Useful data on chemical composition of the literature sample is presented previously in Table 13 from Reddy et al. (2015). This paper is organized as follows: In Section 2 we describe observations, data reduction and radial velocity measurements. Section 3 is devoted to the abundance analysis and Section 4 to discussing the age$-$abundance relations. We present in Section 5 the radial abundance distribution of OCs and field stars and discuss the results in comparison with theoretical chemodynamical models of GCE. In Section 6 we compute the birthplace of OCs with other relevant orbital parameters and focus on discussing the origin of older OCs in the outer Galactic disc. Finally, Section 7 provides the conclusions. \begin{table*} \centering \caption{The journal of the observations for each of the cluster members analysed in this paper.} \label{log_observations} \begin{tabular}{lccccccccclc} \hline \multicolumn{1}{l}{Cluster}& \multicolumn{1}{c}{Star} & \multicolumn{1}{c}{$\alpha(2000.0)$}& \multicolumn{1}{c}{$\delta(2000.0)$}& \multicolumn{1}{c}{V}& \multicolumn{1}{c}{B-V} & \multicolumn{1}{c}{V-K$_{\rm s}$} & \multicolumn{1}{c}{J-K$_{\rm s}$} & \multicolumn{1}{c}{$RV_{\rm helio}$} & \multicolumn{1}{l}{S/N at} & \multicolumn{1}{l}{Date of} & \multicolumn{1}{c}{Exp. time} \\ \multicolumn{1}{c}{}& \multicolumn{1}{c}{}& \multicolumn{1}{c}{(hh mm ss)}& \multicolumn{1}{c}{($\degr$ $\arcmin$ $\arcsec$)}& \multicolumn{1}{c}{(mag)}& \multicolumn{1}{c}{ } & \multicolumn{1}{c}{ }& \multicolumn{1}{c}{ } & \multicolumn{1}{c}{(km s$^{-1}$)} & \multicolumn{1}{l}{6000 \AA } & \multicolumn{1}{l}{observation} & \multicolumn{1}{c}{(sec)} \\ \hline NGC 1647 & 2 & 04 46 05.10 & $+$18 48 02.67 & 07.47 &$+$1.50 &$+$3.59 &$+$0.84 &$-$07.8$\pm$0.2 & 300 & 25-12-2013 & 2$\times$1200 \\ & 105 & 04 46 35.90 & $+$19 29 39.32 & 08.45 &$+$1.60 &$+$4.00 &$+$0.98 &$-$07.8$\pm$0.2 & 190 & 25-12-2013 & 2$\times$1800 \\ NGC 1664 & 17 & 04 51 02.11 & $+$43 38 45.70 & 11.41 &$+$1.08 &$+$2.67 &$+$0.67 &$+$09.2$\pm$0.1 & 150 & 25-12-2013 & 3$\times$1800 \\ & 75 & 04 51 19.72 & $+$43 42 16.54 & 11.06 &$+$0.95 &$+$2.44 &$+$0.67 &$+$06.4$\pm$0.3 & 100 & 24-12-2013 & 3$\times$1800 \\ NGC 2099 & 34 & 05 52 15.10 & $+$32 31 40.92 & 10.90 &$+$1.14 &$+$2.81 &$+$0.69 &$+$10.3$\pm$0.3 & 120 & 21-11-2013 & 2$\times$1800 \\ & 64 & 05 52 13.04 & $+$32 34 16.60 & 11.15 &$+$1.36 &$+$2.97 &$+$0.76 &$+$08.6$\pm$0.2 & 160 & 21-11-2013 & 2$\times$1800 \\ NGC 2281 & 55 & 06 48 15.09 & $+$41 04 22.23 & 08.86 &$+$0.96 &$+$2.27 &$+$0.57 &$+$19.6$\pm$0.2 & 150 & 27-12-2013 & 1$\times$1200 \\ & 63 & 06 48 21.72 & $+$41 18 08.36 & 07.24 &$+$1.37 &$+$3.06 &$+$0.83 &$+$19.3$\pm$0.2 & 280 & 27-12-2013 & 1$\times$1200 \\ NGC 2287 & 75 & 06 45 43.01 & $-$20 51 09.59 & 07.43 &$+$1.28 &$+$2.82 &$+$0.78 &$+$24.0$\pm$0.1 & 400 & 20-11-2013 & 1$\times$1800 \\ & 97$^{a}$ & 06 46 04.84 & $-$20 36 24.91 & 07.80 &$+$1.16 &$+$2.59 &$+$0.70 &$+$22.9$\pm$0.1 & 370 & 20-11-2013 & 1$\times$1800 \\ & 107$^{a}$ & 06 46 33.28 & $-$20 48 42.63 & 07.79 &$+$1.15 &$+$2.59 &$+$0.66 &$+$26.3$\pm$0.1 & 370 & 20-11-2013 & 1$\times$1800 \\ NGC 2345 & 34$^{a}$ & 07 08 21.84 & $-$13 10 23.25 & 09.94 &$+$1.50 &$+$4.21 &$+$0.97 &$+$63.6$\pm$0.4 & 120 & 25-12-2013 & 2$\times$1800 \\ & 43 & 07 08 26.32 & $-$13 11 14.44 & 10.70 &$+$1.81 &$+$4.54 &$+$1.09 &$+$58.0$\pm$0.4 & 180 & 27-12-2013 & 2$\times$1800 \\ & 60 & 07 08 30.37 & $-$13 13 52.54 & 10.48 &$+$1.82 &$+$4.51 &$+$1.11 &$+$57.4$\pm$0.3 & 160 & 25-12-2013 & 2$\times$1800 \\ NGC 2437 & 29$^{a}$ & 07 41 37.61 & $-$14 43 12.98 & 10.86 &$+$1.19 &$+$2.92 &$+$0.76 &$+$52.7$\pm$0.2 & 190 & 27-12-2013 & 2$\times$1800 \\ & 174$^{a}$ & 07 41 51.52 & $-$14 54 29.28 & 10.70 &$+$1.11 &$+$2.50 &$+$0.60 &$+$45.8$\pm$0.2 & 160 & 27-12-2013 & 2$\times$1800 \\ & 242$^{a}$ & 07 41 19.43 & $-$14 48 47.45 & 10.19 &$+$1.12 &$+$2.74 &$+$0.69 &$+$57.2$\pm$0.2 & 160 & 25-12-2013 & 2$\times$1800 \\ NGC 2548& 1218 & 08 13 35.43 & $-$05 53 02.18 & 09.64 &$+$0.94 &$+$2.18 &$+$0.52 &$+$08.7$\pm$0.2 & 110 & 24-12-2013 & 2$\times$1200 \\ & 1296$^{a}$ & 08 13 44.83 & $-$05 48 00.89 & 09.27 &$+$0.77 &$+$2.11 &$+$0.58 &$+$13.7$\pm$0.2 & 180 & 25-12-2013 & 2$\times$1500 \\ & 1560$^{a}$ & 08 14 17.03 & $-$05 54 00.62 & 08.19 &$+$1.10 &$+$2.77 &$+$0.73 &$-$00.5$\pm$0.2 & 200 & 25-12-2013 & 2$\times$1200 \\ & 1628 & 08 14 28.12 & $-$05 42 16.14 & 09.47 &$+$1.02 &$+$2.39 &$+$0.62 &$+$08.0$\pm$0.2 & 240 & 17-11-2011 & 2$\times$1800 \\ NGC 2632 & 212 & 08 39 50.71 & $+$19 32 26.93 & 06.60 &$+$0.96 &$+$2.19 &$+$0.79 &$+$35.5$\pm$0.1 & 400 & 20-11-2013 & 1$\times$1500 \\ & 253 & 08 40 06.42 & $+$20 00 28.04 & 06.39 &$+$0.98 &$+$2.16 &$+$0.54 &$+$34.0$\pm$0.1 & 300 & 20-11-2013 & 1$\times$1200 \\ & 283 & 08 40 22.09 & $+$19 40 11.78 & 06.39 &$+$1.02 &$+$2.20 &$+$0.60 &$+$34.7$\pm$0.1 & 390 & 20-11-2013 & 1$\times$1500 \\ NGC 6633 & 100 & 18 27 54.73 & $+$06 36 00.33 & 08.31 &$+$1.13 &$+$2.64 &$+$0.66 &$-$29.1$\pm$0.2 & 350 & 15-07-2013 & 2$\times$1800 \\ & 119 & 18 28 17.64 & $+$06 46 00.05 & 08.98 &$+$1.04 &$+$2.52 &$+$0.60 &$-$29.1$\pm$0.1 & 160 & 15-07-2013 & 1$\times$1800 \\ NGC 6940 & 67 & 20 34 04.12 & $+$28 16 48.62 & 10.96 &$+$1.13 &$+$2.44 &$+$0.62 &$+$07.8$\pm$0.2 & 130 & 13-10-2013 & 2$\times$1800 \\ & 69 & 20 34 05.75 & $+$28 11 18.38 & 11.63 &$+$1.13 &$+$2.69 &$+$0.62 &$+$07.9$\pm$0.2 & 160 & 13-10-2013 & 2$\times$1800 \\ & 139 & 20 34 47.61 & $+$28 14 47.25 & 11.35 &$+$1.09 &$+$2.56 &$+$0.66 &$+$07.5$\pm$0.2 & 180 & 13-10-2013 & 1$\times$1800 \\ NGC 7209 & 77 & 22 05 09.93 & $+$46 31 25.27 & 10.11 &$+$1.13 &$+$2.68 &$+$0.66 &$-$18.5$\pm$0.1 & 340 & 19-11-2013 & 2$\times$1800 \\ & 89 & 22 05 17.62 & $+$46 29 00.64 & 09.45 &$+$1.35 &$+$3.64 &$+$0.90 &$-$18.8$\pm$0.1 & 250 & 19-11-2013 & 2$\times$1800 \\ & 95$^{a}$ & 22 05 22.23 & $+$46 32 04.40 & 10.60 &$+$1.08 &$+$2.69 &$+$0.62 &$-$07.1$\pm$0.2 & 150 & 19-11-2013 & 2$\times$1800 \\ \hline \end{tabular} \flushleft Note: $^{a}$ Spectroscopic binary. \end{table*} | One aim of our studies of the chemical compositions of red giants in OCs is to provide insights into Galactic chemical evolution (GCE). Based on the high resolution spectra of red giants in twelve OCs and standard LTE differential abundance analysis relative to the Sun, we presented stellar parameters and the first estimates of chemical abundances of 26 species from Na to Eu sampling all the major processes of stellar nucleosynthesis -- $\alpha$ and $r$-process elements synthesized in Type II supernovae from massive stars, iron group elements synthesized in Type Ia supernovae, and the $s$-process elements synthesized in asymptotic-giant-branch (AGB) stars. We combined with our total sample of 28 clusters from this and previous papers with a sample of 51 OCs drawn from the literature whose chemical abundances and Galactocentric distances were remeasured to establish a common abundance and distance scales, respectively. This compilation provides a homogenised sample of 79 OCs spanning a narrow range in metallicity of $\sim$ $-$0.5 to 0.3 dex but a range of a few Myr to 9 Gyr in ages and 5.0 to 24.0 kpc in R$_{\rm gc}$. Our sample alone constitute about 35\% of OCs explored so far with high-resolution spectroscopy and we are the first to provide a homogeneous high-resolution abundance analysis of elements from Na to Eu for twenty-eight OCs in the radial range R$_{\rm gc}$=\,7.7 to 11.3 kpc; only a few OCs in the literature present abundance estimates for the heavy elements from Y to Eu. Following the kinematic criteria, we assigned OCs to the thin disc, thick disc or halo stellar populations and examined the age-abundance relations and the variation of [Fe/H] as a function of R$_{\rm gc}$ in the Galactic disc. We showed that the resolution of the lack of age$-$metallicity relation for OCs, as noted previously by other studies, lies in the incomplete coverage in metallicity compared to that of field stars which cover the range of $-$1.0 to $+$0.4 dex in metallicity and an age of a few Myr to 13 Gyr. Using our homogeneous sample of 79 OCs, we confirm the results of previous studies that OCs present a constant steep decline of metallicity out to R$_{\rm gc}$ of 12 kpc and a further flattening out to the entire radial extent of the Galactic disc. But our analysis is the first to demonstrate clearly that such bimodality accompanied by a sudden change in the slope of radial metallicity distribution of OCs at 12 kpc arise from the selection effects; at R$_{\rm\,gc}<\,$12 kpc, all the sampled clusters lie close to the Galactic midplane ($\lvert\,z\rvert<\,$0.5 kpc), younger than 1.5 Gyr with kinematics typical of thin disc and are constituting a constant steep decline of [Fe/H] with R$_{\rm\,gc}$ while the OCs populating the disc beyond 12 kpc are older with ages from 1.0 to 8.0 Gyr, metal-poor by [Fe/H] $<-$0.2 dex with thick disc kinematics and located away from the midplane (0.5$\,<\lvert\,z\rvert<\,$2.5 kpc) and constitute a shallow gradient over the entire radial extent of the Galactic disc followed by a change of slope at 12 kpc. We further compared the gradients traced by OCs with that of field stars, Cepheids and with chemodynamical model predictions. We demonstrated clearly that the OCs, field stars and Cepheids yield consistent radial gradients if the comparison samples are drawn from the similar vertical slices. The radial metallicity gradient traced by the field stars and Cepheids lying close to the midplane ($\lvert\,z\rvert<\,$0.5 kpc) are steeper while that measured for field stars located away from the midplane (0.5$\,<\lvert\,z\rvert<\,$2.0 kpc) are shallower over the entire radial extent of the disc and are comparable to similar gradients measured for the sample of OCs. The chemodynamical models of Minchev et al. (2013) produce a steep gradient for samples of stars younger than 2 Gyr and a flat trend for stars older than 2 Gyr throughout the radial range 5 to 16 kpc of the disc. In their simulations, a signature of flattening of the radial gradient arise naturally due to the radial mixing of stars but the young population is hardly affected as the radial mixing operates at a slow pace. As a result, the younger populations are expected to show a steeper gradient than that of the older ones. Taking into account the affects of radial mixing, these simulations fairly predict a steep gradient of $-$0.058 dex kpc$^{-1}$ (R$_{\rm\,gc}=$6$-$11 kpc) for the stellar populations younger than 2 Gyr and the shallow gradient of old populations which are in fair agreement with such gradients observed for the tracers such as OCs, field stars and Cepheids. Finally, we demonstrated through the computation of birthplaces of OCs that the sample of clusters (ages$<$ 1.5 Gyr and R$_{\rm\,gc}<\,$12 kpc) constituting a steep radial metallicity gradient of slope $-$0.052$\pm$0.011 dex kpc$^{-1}$ have circular orbits with birthplaces very close their present locations in the Galactic midplane. In contrast, the older clusters populating the outer disc (R$_{\rm\,gc}>\,$12 kpc) with a measured shallow slope of $-$0.015$\pm$0.007 dex kpc$^{-1}$ have relatively eccentric orbits but with birthplaces around R$_{\rm\,gc}$ of 11 kpc close to the midplane. Our analysis is the first to demonstrate clearly that the orbital eccentricity of all but the three outer disc clusters Be 29, Be 75 and Saurer 1 has taken them to present locations in the Galactic disc from their birthplaces inward of 12 kpc from a medium enriched with very similar metallicity. Moreover, these outer disc clusters also make large excursions away from the Galactic midplane, a pleasing result in support of their survival for longer periods of time. The older, distant clusters Be 29 and Saurer 1 at a height of 2 kpc and 1.7 kpc, respectively, are of extra-galactic origin while Be 75 may be an old genuine Galactic cluster on a perturbed orbit or the astrometric data employed in orbit calculations might be erroneous and require further scrutiny. In the spirit of speculation to encourage further spectroscopic analyses, the suggestion is made that the variation of metallicity of OCs with Galactocentric distance is a linear function with a constant negative slope over the entire radial extent of the disc but close to the Galactic midplane. The radial gradients become flatter as one moves away from the Galactic midplane and for the older OCs. To test such speculations, it would be useful to explore the chemical content of as many clusters as possible including the young and old OCs covering a wide range in R$_{\rm gc}$ from 5 to 18 kpc in the disc and a scale height of 0.0 to 3 kpc above the midplane and then to compare the abundance gradients of OCs with those derived from the field stars and Cepheids. \vskip1ex {\bf Acknowledgements:} We thank the anonymous referee for comments which have improved the Paper. We are grateful to the McDonald Observatory's Time Allocation Committee for granting us observing time for this project. DLL wishes to thank the Robert A. Welch Foundation of Houston, Texas for support through grant F-634. ABSR thanks Dr. Zhen-Yu Wu for generously providing his code for the integration of open cluster orbits. This research has made use of the WEBDA database, operated at the Institute for Astronomy of the University of Vienna and the NASA ADS, USA. This research has also made use of Aladin. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration (NASA) and the National Science Foundation (NSF). | 16 | 9 | 1609.02619 |
1609 | 1609.02333_arXiv.txt | Extensive photometric and spectroscopic observations are presented for \sn, a type IIP supernova (SN) exploding in the nearby galaxy \host. The observations are performed in optical and ultraviolet bands, covering from $-20$ to +400 days from the peak light. The stringent detection limit from prediscovery images suggests that this supernova was actually detected within about 1 day after explosion. Evolution of the very early-time light curve of \sn\ is similar to that predicted from a shock breakout and post-shock cooling decline before reaching the optical peak. Our photometric observations show that SN 2014cx has a plateau duration of $\sim 100$ days, an absolute $V$-band magnitude of $\sim -16.5$ mag at $t \approx 50$ days, and a nickel mass of $0.056 \pm 0.008$ M$_{\odot}$. The spectral evolution of SN 2014cx resembles that of normal SNe~IIP like SN 1999em and SN 2004et, except that it has a slightly higher expansion velocity ($\sim 4200$ \kms\ at 50 days). From the cooling curve of photospheric temperature, we derive that the progenitor has a pre-explosion radius of $\sim 640$ \rsun, consistent with those obtained from SNEC modeling ($\sim 620$ \rsun) and hydrodynamical modeling of the observables ($\sim 570$ \rsun). Moreover, the hydrodynamical simulations yield a total explosion energy of $\sim 0.4 \times 10^{51}$ erg, and an ejected mass of $\sim 8$ M$_{\odot}$. These results indicate that the immediate progenitor of \sn\ is likely a red supergiant star with a mass of $\sim 10$ M$_{\odot}$. | \label{sec:intro} Type IIP supernovae (SNe IIP) represent the most common subtype of stellar explosions, constituting about one third of all SNe \citep{2011MNRAS.412.1441L}. This subtype of SNe is thought to arise from core-collapse (CC) explosions of massive red supergiants (RSGs) with an initial mass of 8--25 \msun\ according to theoretical models of stellar evolution \citep{2003ApJ...591..288H}. On the other hand, direct analyses of supernova position on pre-explosion images give a much narrower range for the progenitor mass, e.g., 8.5--16.5 \msun\ \citep{2007ApJ...661.1013L, 2009ARA&A..47...63S}. Compared to other CC SNe, SNe IIP are characterized by prominent hydrogen features in their optical spectra \citep[e.g.,][]{1997ARA&A..35..309F} and an extended plateau phase in their light curves. During the plateau phase, their luminosity remains almost constant as a result of the energy balance between the hydrogen recombination and expansion cooling. The plateau feature distinguishes SNe~IIP from the Type IIL subclass, for which the light curve exhibits a linear decline (in mag day$^{-1}$) after the peak \citep{1979A&A....72..287B}. Recent statistical analyses show that the light-curve properties of SNe II may have a continuous distribution (\eg\ \citealt{2014ApJ...786...67A, 2015ApJ...799..208S, 2016MNRAS.459.3939V}), although there are also studies suggesting a distinct division between type IIP and IIL SNe \citep{2012ApJ...756L..30A, 2014MNRAS.442..844F, 2014MNRAS.445..554F}. Over the years, numerous SNe IIP have been well studied, such as SN 1999em \citep{2002PASP..114...35L}, SN 2004et \citep{2006MNRAS.372.1315S}, SN 2005cs \citep{2009MNRAS.394.2266P}, and SN 2013ej \citep{2015ApJ...807...59H}. These studies reveal a large spread in luminosities, plateau durations, expansion velocities, and nickel masses for SNe IIP (\eg\ \citealt{2003ApJ...582..905H}), which can be well understood with current explosion models (\eg\ \citealt{2009ApJ...703.2205K, 2010MNRAS.408..827D, 2011ApJ...741...41P, 2013MNRAS.434.3445P}). Nevertheless, early time observations are still sparse for SNe IIP, which are vital to constrain the explosion time and hence determine the properties of their progenitor stars \citep{2003MNRAS.346...97N, 2009MNRAS.395.1409S}. In particular, very early light curves of SNe IIP may be affected by a short, sharp blast of light as a result of shock breakout of the stellar surface, as predicted in core collapse explosion of massive stars \citep{1977ApJS...33..515F, 1978ApJ...223L.109K}. \sn\ represents such a CC SN that is captured within about one day after the explosion. \sn\ was independently discovered on UT 2014 Sept. 2 by \citet{2014CBET.3963....1N} and \citet{2014ATel.6436....1H} in the nearby SBd galaxy \host. Based on the astrometry from the USNO-A2.0 Catalogue \citep{1998usno.book.....M}, the J2000 coordinates of the SN are derived as $\alpha$ = 00$^h$59$^m$47.83$^s$ and $\delta = -07\arcdeg 34 \arcmin 19.3 \arcsec$, approximately 21.7$\arcsec\,$N and 33.7$\arcsec\,$W from the center of \host\ \citep{2014ATel.6436....1H}. \sn\ was reported as a young SN II based on both optical \citep{2014ATel.6440....1E} and near-infrared (NIR; \citealt{2014ATel.6442....1M}) spectra taken at about one day after the discovery. It was further classified as a Type IIP event according to the photometric observations by \cite{2015ATel.7084....1A}. We note that another SN~IIP, SN 2011dq, also exploded in \host. The distance to \host\ is estimated to be $18.0\pm3.6$ Mpc (distance modulus $\mu=31.27\pm0.43$ mag) by the Tully-Fisher method \citep{2014MNRAS.444..527S}; here we adopt this value for \sn. In this work, we present the results of our optical and UV observations of the type IIP supernova \sn\ that was discovered at a very young age. The observations and data reduction are addressed in Section \ref{sec:obs}, the photometric and spectroscopic evolution are described in Section \ref{sec:lc} and \ref{sec:spec}, respectively, and analysis of the progenitor properties of \sn\ via the photospheric temperature cooling curve and hydrodynamical modeling is given in Section \ref{sec:model}. The main results are summarized in Section \ref{sec:sum}. | \label{sec:sum} In this paper, we present extensive UV and optical photometry and optical spectroscopy of \sn\ in \host, spanning the period from $-30$ d to $+404$ d from the maximum light. The explosion time is constrained to be MJD = 56,901.89 with an accuracy of $\pm 0.5$ day. The characteristics of the light curves, such as the rise time, duration of the plateau phase, post-peak decline, and bolometric luminosity, suggest that \sn\ is a normal type IIP supernova. The KAIT unfiltered and LCOGT r-band light curves seem to experience two brightening components, with the first likely related to shock breakout of the supernova. The plateau duration is $\sim 100$ days, similar to that of our comparison SNe~IIP. The value of $M_V$ at mid-plateau phase ($\sim 50$ d) is $-16.48 \pm 0.43$ mag for \sn, lying between the luminous SNe~IIP ($\sim -17$ mag, SN 2004et) and subluminous SNe~IIP ($\sim -15$ mag, SN 2005cs). The mass of \nickel\ using the tail luminosity and steepness methods yield a value of 0.056 \msun, similar to that of SN 1999em and SN 2004et. The spectroscopic evolution of \sn\ shares a similarity with the typical Type IIP SNe 1999em and 2004et. The early-time spectra exhibit a nearly featureless continuum with only hydrogen Balmer lines and He~I visible. As the SN evolves, the continuum becomes redder and the metal lines emerge, becoming the dominant features during the photospheric phase. During the nebular phase, the spectra are dominated by strong emission lines. The continuum at 3800--5000 \AA\ is relatively blue, which might be caused by either late-time CSM interaction or a scattered-light echo. The value and evolution of the expansion velocity derived from Fe~II \ld5169 are similar to those of SN 2004et, but $\sim 1000$ km s$^{-1}$ higher than the expansion velocity of SN 1999em. By modeling the observables of \sn\ as derived from our observations, we estimate that this explosion produces a total energy of $0.4 \times 10^{51}$ ergs and an ejected mass of $\sim$ 8.0 \msun. The progenitor star is calculated to have a radius of $4 \times 10^{13}$ cm ($\sim 574$ \rsun), which agrees well with that estimated from the early photospheric temperature evolution ($643 \pm 60$ \rsun) and $g$-band rise time -- radius relation from SNEC ($619 \pm 10$ \rsun). The values above are consistent with a core-collapse scenario from a typical RSG having an initial mass of 9.5--10 \msun. | 16 | 9 | 1609.02333 |
1609 | 1609.02569_arXiv.txt | We report spatially resolved (FWHM$\sim3.8-4.6''$) mid-IR imaging observations of the planetary nebula (PN) NGC 7027 taken with the 2.5-m telescope aboard the Stratospheric Observatory for Infrared Astronomy (SOFIA). Images of NGC 7027 were acquired at 6.3, 6.6, 11.1, 19.7, 24.2, 33.6, and 37.1 $\mu\mathrm{m}$ using the Faint Object Infrared Camera for the SOFIA Telescope (FORCAST).The observations reveal emission from Polycyclic Aromatic Hydrocarbon (PAH) and warm dust ($T_D\sim90$ K) from the illuminated inner edge of the molecular envelope surrounding the ionized gas and central star. The DustEM code was used to fit the spectral energy distribution of fluxes obtained by FORCAST and the archival infrared spectrum of NGC 7027 acquired by the Short Wavelength Spectrometer (SWS) on the Infrared Space Observatory (ISO). Best-fit dust models provide a total dust mass of $5.8^{+2.3}_{-2.6}\times10^{-3}$ $\mathrm{M}_\odot$, where carbonaceous large ($a=1.5$ $\mu$m) and very small ($a \sim12\AA$) grains, and PAHs ($3.1\AA<a<12\AA$) compose 96.5, 2.2, and 1.3 $\%$ of the dust by mass, respectively. The 37 $\mu$m optical depth map shows minima in the dust column density at regions in the envelope that are coincident with a previously identified collimated outflow from the central star. The optical depth minima are also spatially coincident with enhancements in the 6.2 $\mu$m PAH feature, which is derived from the 6.3 and 6.6 $\mu$m maps. We interpret the spatial anti-correlation of the dust optical depth and PAH 6.2 $\mu$m feature strength and their alignment with the outflow from the central star as evidence of dust processing and rapid PAH formation via grain-grain collisions in the post-shock environment of the dense ($n_H\sim10^5\,\mathrm{cm}^{-3}$) photo-dissociation region (PDR) and molecular envelope. | Polycyclic aromatic hydrocarbon (PAH) molecules containing fewer than $\sim500$ carbon atoms are abundant and ubiquitous components of interstellar medium (ISM) characterized by their broad emission features at 3.3, 6.2, 7.7, 8.6, 11.3, and 12.7 $\mu$m (Tielens 2008 and ref. therein). Despite being smaller than $\sim10\AA$ in radius, observations of PAH features indicate that they account for $\sim10\%$ of carbon in the ISM (e.g. Allamandola et al. 1989, Draine \& Li 2007, Compi{\`e}gne et al. 2011). The details of PAH formation and destruction are, however, currently unclear. One of the most pervasive questions on interstellar PAHs is how they are replenished and/or survive destructive shocks driven by supernova explosions given that their theoretical destruction timescale is over an order of magnitude shorter than their production timescale from stars (e.g. Cherchneff, Barker, \& Tielens 1992; Micelotta et al. 2010a, b). Even though significant repopulation of PAH-sized grains is predicted to occur in $\sim100$ km $\mathrm{s}^{-1}$ shocks due to grain-grain collisions and the fragmentation of large graphitic grains (Jones et al. 1996), recent theoretical work shows that neither ``parent" nor ``daughter'' PAHs can survive the passage of shocks with velocities $\gtrsim100$ km $\mathrm{s}^{-1}$ (Micelotta et al. 2010a, b). This is in direct contrast with the detection of PAH features in shocked regions (e.g. Tappe al. 2006, Engelbracht et al. 2006, Armus et al. 2007). The circumstellar environment of carbon-rich post-Asymptotic Giant Branch (AGB) stars and young planetary nebulae (PNe) provide ideal laboratories to study the evolution of PAHs in the shocked regions impacted by collimated outflows from the central stars. Carbon-rich AGB stars, which are descendants of stars with an initial mass $<8$ $\mathrm{M}_\odot$, exhibit cool, dense outflows that are favorable environments for PAH formation (Latter 1991). These PAHs are believed to seed the growth of amorphous carbon and are therefore important components in the chemical pathways towards dust formation (Tielens 2008). Interestingly, PAH features are rarely observed towards carbon-rich AGB stars, which may be due to the cool effective temperatures of the AGB photosphere that are unable to excite the PAH vibrational and stretching modes. The PAH features are, however, present in spectra of carbon-rich post-AGBs and PNe where the circumstellar dust is heated by a radiation field blue from the hot degenerate core that is harder than the radiation field in the AGB-phase. NGC~7027 is a young and carbon-rich PNe, one of the brightest and best-studied of its breed (e.g. Gillett, Low, and Stein 1967; Becklin, Neugebauer, \& Wynn-Williams 1973; Moseley 1980). The compact $\sim10$ arcsec diameter central region is very bright at infrared and radio wavelengths, while more recent studies have identified an extended molecular shell, which is likely the envelope of the progenitor AGB star (Fig.~\ref{fig:XRAYFCIm}). Proper motion studies of features in the radio continuum in comparison with line of sight velocities show that the distance is $\sim1$ kpc and the dynamical age only $\sim1200$ yrs (Zijlstra et al. 2008). Infrared spectroscopy of the nebula reveals prominent PAH emission features as well as a high carbon abundance (Beintema et al. 1996, Bernard Salas 2001). Given the young age and relatively high mass ($\sim0.7$ $\mathrm{M}_\odot$, Zijlstra et al. 2008), it is not surprising that NGC~7027 is excited by a very hot white dwarf, which has a temperature in excess of 200,000 K and an estimated luminosity of 7700 $\mathrm{L}_\odot$ (Latter et al. 2000). Notably, the morphology of gas and dust in NGC~7027 exhibits significant deviations from spherical symmetry, which has been interpreted as evidence of interaction between the nebula and collimated high-velocity outflows from the central star (Graham et al. 1993, Kastner et al. 1994, Cox et al. 1997). Outflows from central stars are commonly observed in PNe and are believed to be one of the mechanisms responsible for shaping multi-polar morphologies (e.g. Sahai \& Trauger 1998; Sahai, Morris, \& Villar 2011). Kinematic observations of NGC~7027 indeed reveal asymmetric expansion velocities that suggest the passage of multiple collimated outflows through the ionized and molecular gas composing the nebula (Latter et al. 2000, Cox et al. 2002, L{\'o}pez et al. 2012). Three collimated outflows from the central star are identified by Cox et al. (2002) from kinematic information inferred from $H_2$ and Br-$\gamma$ radial velocities measurements. The outflows exhibit position angles of $-53^\circ$, $4^\circ$, and $-28^\circ$, which are referred to as outflows 1, 2, and 3, respectively (See Fig.~\ref{fig:XRAYFCIm}). Extended X-ray emission found along the axis of outflow 1 (Kastner et al. 2001) suggest that it is the most recent outflow ($<1500$ yr old). High red- and blue-shifted Br-$\gamma$ line velocities along outflow 1 of $\pm\,55$ km s$^{-1}$ with respect to the systemic velocity of the nebula ($v_\mathrm{LSR}\sim25$ km s$^{-1}$; Cox et al. 2002, Nakashima et al. 2010) reinforce this interpretation. Bains et al. (2003) provide evidence of the interaction between outflow 1 and the nebula from their interpretation of a bright knot of red-shifted ionized gas at the north-west of the nebula, consistent with the axis of outflow 1. They find that the knot is bright due to higher temperatures as opposed to a density enhancement and claim the knot is associated with the receding, far-side of the nebula that is observable due to a ``breach" in the near-side from a high velocity outflow. Recent kinematic models of NGC~7027 from CO, $H_2$, and Br-$\gamma$ velocity measurements substantiate this hypothesis and find evidence of a ``hole" in the structure of all three emission components only along outflow 1 (Nakashima et al. 2010). The nebular structure along outflows 2 and 3 are only found to exhibit holes in the CO and $H_2$ components, which may alternatively be due to UV radiation and photodissociation as opposed to outflow interaction. This alternative mechanism for hole formation via photodissociation does not apply to Br-$\gamma$, which implies the hole along outflow 1 exhibits the strongest evidence of interaction with a high-velocity collimated outflow. In this paper, we present mid-IR imaging of the young and carbon-rich PN NGC~7027 at wavelengths from 6 to 40 $\mu$m with $\sim4''$ resolution using the Faint Object Infrared Camera for the SOFIA Telescope (FORCAST). Compact PNe are ideal targets for SOFIA because of their high infrared surface brightness (e.g. Spuck et al. 2013, Werner et al. 2014). These are the first observations which extensively resolve the emission from the nebula at the wavelengths of its peak emission around 33.6 and 37.1 $\mu$m. The goal of our analysis is to study the morphology and energetics of the PAH and warm dust emission of NGC 7027 in the context of interactions with possible collimated outflows from the central star. In Sec. 2 we describe our observations with SOFIA and the data reduction. In Sec. 3, we report our results on comparing the mid-IR and 6.2 $\mu$m PAH feature morphology to archival HST imaging data, studying the color temperature and optical depth of the warm dust, and fitting dust models to our mid-IR spectral energy distribution. Lastly, in Sec. 4 we discuss the heating of dust in the nebula via trapped Lyman-$\alpha$ photons and report evidence of dust processing and rapid PAH formation along a recent outflow from the central star. \begin{figure}[t] \centerline{\includegraphics[scale=0.35]{XRayFCIm_Fig4}} \caption{False color image of NGC 7027 combining archival visible and near-IR observations with HST and x-ray observations with Chandra. Molecular hydrogen, Paschen-$\alpha$, and V-band emission are shown as red, green, and blue, respectively. X-ray contours are overlaid with levels corresponding to 20, 40, 60, and 80 \% of the peak x-ray flux. The bi-directional arrows labeled 1-3 correspond to the direction of the outflows identified by Cox et al. (2002). Outflow 1 (green) is believed to be the most recent and/or powerful outflow and shows a similar orientation to the bi-polar x-ray emission.} \label{fig:XRAYFCIm} \end{figure} | We have reported mid-IR images of the warm dust and PAH emission from the $\sim1000$ yr-old PN NGC 7027. The key contribution of the SOFIA/FORCAST observations of NGC 7027 were the spatially resolved 33.6 and 37.1 $\mu$m images that trace the peak emission of large grains and the 6.3 and 6.6 $\mu$m images that trace the 6.2 $\mu$m PAH feature in the nebula. The mid-IR morphology of the nebula closely resembles that of the ionized inner edge of the molecular envelope as traced by Paschen-$\alpha$ line emission (Fig.~\ref{fig:PNImall}A). At the NW and SE edges of the nebula there is extended and diffuse mid-IR emission that deviates from the azimuthal symmetry of the nebula and is coincident with similar extended features in molecular hydrogen and V-band images. Notably, this asymmetry is coincident with outflow 1 from the central star (Cox et al. 2002; Fig.~\ref{fig:XRAYFCIm}). Color temperature maps derived from the 19 and 37 $\mu$m images indicate that hot dust ($T_\mathrm{d}\gtrsim100$ K) is present in the cavity. Temperatures are also slightly higher along the outflow 1 ($T_\mathrm{d}\sim100$) relative to regions equidistant from the central star where $T_\mathrm{d}\sim70$ (Fig.~\ref{fig:PNCTOD}A). The 37 $\mu$m optical depth map exhibits a similar morphology to the CO (1-0) emission (Graham et al. 1993b) tracing the dense, molecular envelope and clearly reveals optical depth minima along outflow 1 (Fig.~\ref{fig:PNCTOD}B). Our DustEM emission model reproduced a close fit to the ISO/SWS spectrum and SOFIA/FORCAST mid-IR photometry of NGC 7027 (Fig.~\ref{fig:PNSED}). We adjusted the dust properties of three independent components to fit the emission. The dust components were amorphous carbon large grains (LGs, $a\sim1.5$ $\mu$m) and very small grains (VSGs, $a\gtrsim12\AA$), and PAHs ($3.1\AA<a<12\AA$). The best-fit revealed a total dust mass of $5.8^{+2.3}_{-2.6}\times10^{-3}$ $\mathrm{M}_\odot$, where LGs, VSGs, and PAHs compose 96.5, 2, and 1.5 $\%$ of the dust by mass, respectively. Analysis of the total IR luminosity ($L_\mathrm{IR}\sim5.6\times10^3$ $\mathrm{L}_\odot$) and the hard radiation field of the $\sim200,000$ K and $7.7\times10^3$ $\mathrm{L}_\odot$ central star shows that dust in the nebula cannot be heated directly by the central star. Given the free-free emission measurements from radio observations that are unobscured from the effects of local extinction, we determined that the total luminosity in trapped Lyman-$\alpha$ photons is consistent with the observed IR luminosity. These results imply that dust in the nebula is heated by the trapped Lyman-$\alpha$ photons re-radiated in the de-excitation of ionized hydrogen. Lastly, we claim that the spatial anti-correlation between the optical depth and the 6.2 $\mu$m PAH emission feature strength (Fig.~\ref{fig:PNLCIm}A) is evidence for recent dust processing and rapid PAH formation along outflow 1. Despite the size of the LGs in the nebula, LGs can be destroyed by outflow 1 dust to non-thermal sputtering from ``betatron acceleration" along magnetic field lines embedded in the dense, radiatively cooled post-shock regions. The size of the LGs also makes them susceptible to frequent grain-grain collisions that fracture the LGs and redistributes mass to VSGs and PAHs. Our observations demonstrate that PAHs can indeed be formed rapidly in shocks and that grain-grain collisions are likely an efficient mechanism for PAH formation. There are several follow-up questions that warrant further study on dust processing and PAH formation in shocks: how efficient is PAH formation in shocks via grain-grain collisions? Do PAHs survive after forming? Can the enhancements in PAH/VSG abundances be linked to the morphology, outflow history, and/or age of PNe? Interestingly, a recent mid-IR imaging study of the $\sim2500$ yr old bipolar PN M2-9 (Werner et al. 2014) revealed that equal masses of small ($a<0.1$ $\mu$m) and large ($a>1$ $\mu$m) are present in its lobes, where they suggest collisional processing has influenced the grain size distribution. Additional spatially resolved mid-IR observations of other PNe will be important to establish the relation between central star outflows and mass redistribution in dust. Obtaining mid-IR spectra will be equally as important. Although our mid-IR observations with SOFIA are able to resolve the prominent 6.2 $\mu$m PAH emission feature, there are degeneracies in the SED modeling of these regions due to the lack of spatially resolved spectral coverage. As can be seen in Fig.~\ref{fig:PNSED}, the relative PAH and VSG abundances are difficult to fit without spectral coverage between $6\lesssim \lambda \lesssim 12$ $\mu$m. The question of the survival of newly formed PAHs is an important question to address especially given recent theoretical studies that indicate PAHs should not be able to survive $\gtrsim100$ km $\mathrm{s}^{-1}$ shocks (Micelotta et al. 2010a, b). These newly formed PAHs, however, will be difficult to trace as they move further away from the central star due to the decrease in the incident flux of optical/UV photons that excite the PAH features. Future high angular resolution, ground and space-based IR observatories such as the thirty meter-class telescopes and the upcoming 6.5-m James Webb Space Telescope (JWST; expected launch date Oct 2018) will therefore be ideal platforms for exploring the formation and evolution of PAHs in various astrophysical contexts. \emph | 16 | 9 | 1609.02569 |
1609 | 1609.02796_arXiv.txt | {The census of the Solar neighbourhood is still incomplete, as demonstrated by recent discoveries of many objects within 5--10\,pc from the Sun. The area around the mid-plane and bulge of the Milky Way presents the most difficulties in searches for such nearby objects, and is therefore deficient in the known population. This is largely due to high stellar densities encountered. Spectroscopic, photometric and kinematic characterization of these objects allows better understand the local mass function, the binary fraction, and provides new interesting targets for more detailed studies. We report the spectroscopic follow-up and characterisation of 12 bright high PM objects, identified from the VISTA Variables in V\'ia L\'actea survey (VVV). We used the 1.9-m telescope of the South African Astronomical Observatory (SAAO) for low-resolution optical spectroscopy and spectral classification, and the MPG/ESP 2.2m telescope Fiber-fed Extended Range Optical Spectrograph (FEROS) high-resolution optical spectroscopy to obtain the radial and space velocities for three of them. Six of our objects have co-moving companions. We derived optical spectral types and photometric distances, and classified all of them as K and M dwarfs within 27 -- 264\,pc of the Sun. Finally, we found that one of the sources, VVV\,J141421.23-602326.1 (a co-moving companion of VVV\,J141420.55-602337.1), appears to be a rare massive white dwarf that maybe close to the ZZ\,Ceti instability strip. Many of the objects in our list are interesting targets for exoplanet searches.} {proper motions -- stars: low-mass -- (stars:) white dwarfs -- (stars:) binaries: visual -- (Galaxy:) solar neighbourhood -- techniques: spectroscopic. } | M-dwarfs account for over 70\% of stellar systems in the solar vicinity (Henry \etal 1997). Most of them are single ($\sim$60-70\%; Fischer \& Marcy 1992; Bergfors \etal 2010), making them more likely to host (potentially habitable) planets (e.g. Kraus \etal 2012). For comparison, the single star fraction is $\sim$54\% for solar-type stars (Duquennoy \& Marcy 1991; Raghavan \etal 2010) and it is $\sim$0\% for massive stars (Preibisch \etal 1999). This makes M-type dwarf stars the most numerous potential planet hosts of all the stellar classes (Lada 2006). Furthermore, all exoplanet detection methods (radial velocity, transits, direct imaging with Adaptive Optics and astrometry) are more sensitive to planets with host stars of lower masses. A number of exoplanet search programs are aggressively targeting M-dwarfs with radial velocities ({\it e.g.} M2K; Apps \etal 2010) and with transits (RoPACS and MEarth; Irwin \etal 2014). The first exoplanet to be imaged was orbiting a brown dwarf (BD) at $\sim 70$\,pc (Chauvin \etal 2005), and there is a on-going debate for an astrometrically discovered third planetary mass body in a nearby BD pair at $\sim 2.3$\,pc (Boffin \etal 2014; Sahlmann \& Lazorenko 2015). One of the most powerful methods to identify nearby stars is through identifying proper motion (PM), and by exploiting a PM search at infrared wavelengths, offers an additional advantage: cool objects are intrinsically brighter at those wavelengths than in the optical because their spectral energy distributions peak at $\lambda$$>$1\,$\mu$m (L\'epine \& Gaidos 2011) published an all-sky catalog of M-dwarfs with apparent near-InfraRed (near-IR) magnitude $J$$<$10. They selected 8889 stars from the on-going SUPERBLINK ({\it e.g.} L\'epine \& Shara 2005) survey of stars with $\mu$$>$40\,mas\,yr$^{-1}$, supplemented at the bright end with the TYCHO-2 catalogue. Recently, Lepine \etal (2013) presented a spectroscopic catalog of the 1564 brightest ($J$$<$9) M-dwarf candidates in the northern sky. The majority of surveys avoid Galactic plane and bulge, or are substantially incomplete near to these regions. However, these regions offer considerable latent potential for new discoveries of nearby low-mass stars and brown dwarfs. This is especially true for nearby or bright examples that have been overlooked in previous searchs due to confusion caused by high stellar densities and background contaminant objects (Folkes \etal 2012; Luhman 2013; Scholz 2014). A fortuitous aspect of discoveries at low Galactic latitudes is that these regions typically offer many suitable reference stars for high-Strehl ratio AO follow-up observaitions. This is the third paper of a project, after Beam\'in \etal (2013) and Ivanov \etal (2013), to generate a uniform catalog of high-PM objects within the VVV footprint, and characterise them with spectroscopic follow-up observations. It is organised as follows: the next section describes the sample selection, and the new observations. Section 3 describes the spectral type estimation, the distance measurements, and reports on the co-moving companions. Finally, in section 4 we present summary and conclusions. | We obtained spectroscopic follow-up observations of twelve new high PM objects found by the VVV survey during the initial testing of our searching method, and we also looked for possible new wide binary companions. We derived their optical spectral types and photometric distances. All of the analysed objects are K and M dwarfs located at 27--264\,pc from the Sun and are bright enough for further follow-up and search of planets using state of the art and upcoming NIR instruments. Also, all objects are in the densest regions of the Milky Way, surrounded by a pletora of bright stars, very suitable for AO imaging. That makes our targets ideal for searches of close neighbours. From the other side, the surrounded stars are ideal comparison stars for precise relative photometry, variability and transit studies. VVV~J141421.23-602326.1, a co-moving companion of VVV~J141420.55-602337.1, is a candidate for being a rare massive ZZ~Ceti type pulsator. Further spectroscopic and photometric follow-up is needed to better constrain nature and age of this object. \Acknow{We gratefully acknowledge use of data from the ESO Public Survey programme ID 179.B-2002 taken with the VISTA telescope, and data products from the Cambridge Astronomical Survey Unit. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by NASA and NSF.This research has benefitted from the M, L, T, and Y dwarf compendium housed at DwarfArchives.org. Support for MG, RK, JCB, DM, and JB is provided by the Ministry of Economy, Development, and Tourisms Millennium Science Initiative through grant IC120009, awarded to The Millennium Institute of Astrophysics, MAS. MG acknowledges support from Joined Committee ESO and Government of Chile 2014. RK, DM and JB are supported by FONDECYT grants No. 1130140, 1130196 and 1120601, respectively. Both PV and AYK acknowledge support from the National Research Foundation of South Africa. JB, RK, MG and VV are supported by CONICYT REDES140042. JCB acknowledge support from CONICYT FONDO GEMINI - Programa de Astronom{\'i}a del DRI, Folio 32130012.} | 16 | 9 | 1609.02796 |
1609 | 1609.04948_arXiv.txt | {} {This paper continues our study of the foreground population to the Orion molecular clouds. The goal is to characterize the foreground population north of NGC 1981 and to investigate the star formation history in the large Orion star-forming region. We focus on a region covering about 25 square degrees, centered on the $\epsilon$ Orionis supergiant (HD 37128, B0\,Ia) and covering the Orion Belt asterism.} {We used a combination of optical (SDSS) and near-infrared (2MASS) data, informed by X-ray (\textit{XMM-Newton}) and mid-infrared (WISE) data, to construct a suite of color-color and color-magnitude diagrams for all available sources. We then applied a new statistical multiband technique to isolate a previously unknown stellar population in this region.} {We identify a rich and well-defined stellar population in the surveyed region that has about 2\,000 objects that are mostly M stars. We infer the age for this new population to be at least 5\, Myr and likely $\sim10$\,Myr and estimate a total of about 2\,500 members, assuming a normal IMF. This new population, which we call the Orion Belt population, is essentially extinction-free, disk-free, and its spatial distribution is roughly centered near $\epsilon$ Ori, although substructure is clearly present.} {The Orion Belt population is likely the low-mass counterpart to the Ori OB Ib subgroup. Although our results do not rule out Blaauw's sequential star formation scenario for Orion, we argue that the recently proposed blue streams scenario provides a better framework on which one can explain the Orion star formation region as a whole. We speculate that the Orion Belt population could represent the evolved counterpart of an Orion nebula-like cluster.} | The Orion star formation complex is the closest massive star-forming region to the Sun and has generated about $10^4$ low- and high-mass stars for at least the last $\sim\,12$ Myr (e.g., \cite{Blaauw}; \cite{Brown1994}; \cite{Bally2008}; \cite{Muench2008}; \cite{Briceno2008}). The entire region, also known as the Orion OB\,I association, covers an area of approximately 10\degr\,$\times$\,20\degr\ on the sky and harbors a half dozen subgroups containing well-known OB stars and giant molecular clouds (see Fig.~\ref{fig:Orion}). The proximity of the region ($\sim$\,400\;pc; \cite{Hirota2007}; \cite{Menten2007}; \cite{Sandstrom2007}; \cite{Bally2008}) makes it one of the most significant star formation laboratories in astronomy. Indeed, much has been learned in Orion about star formation, for example, clues to the evolution and destruction of clouds, the physics and dynamics of the interstellar medium (ISM), and the role that OB associations and high-mass stars play in the cycling of gas between various phases of the ISM. It is very remarkable, although understandable given its size, that for a region of such fundamental importance most attention has been devoted to the embedded and dusty stellar populations (age $\leq$3 Myr) emerging from the molecular clouds complexes Orion A and Orion B \cite[e.g.,][]{ELada_1991, Allen, Megeath2012, Gutermuth09, daRio10, Spezzi15}, whilst only a few studies have tackled the Orion star-forming region as a whole. The Orion OB\,I association \citep{Blaauw} is composed of several stellar subgroups of different ages, gas, and dust amount. Blaauw divided Orion's association into four groups. Figure \ref{fig:Orion} presents a widefield image of the Orion Constellation superimposed with ellipses denoting the approximate boundaries of these four OB\,I subgroups. These groups appear to show a spatial-temporal relation that is suggestive of a sequence of star formation events, from dust free Ia subgroup to still dust embedded Id. This led \cite{Blaauw} to propose a sequential star formation scenario, where a previous generation of stars is responsible for the formation of a new one via positive feedback; this idea was later quantified by \cite{Elmegreen_Lada} and has remained very popular in the literature. Although there are differences in the estimated ages or exact sizes of the various groups, most of the published works in the region agree that the Orion OB\,Ia group toward the north is the oldest with an age of $\sim$\,8-10\,Myr (\citealt{Bally2008}, with a distance $\sim$\,350\,pc) or even 12\,Myr as originally proposed by \cite{Blaauw}. This group is also dust free. The OB\,Ib subgroup, containing the stars around the Orion Belt asterism, is located at a distance of $\sim$\,400\,pc \citep{Bally2008}. This subgroup has an estimated age of $\sim$\,3-6\,Myr \citep{Bally2008} or $\sim$\,1.7\,Myr \citep{Brown1994}, although the lower estimate is inconsistent with the age of the three supergiants ($\zeta$ Ori, $\epsilon$ Ori, and $\delta$ Ori) that form the naked-eye Belt. According to their spectral types, these three stars must be at least 5 Myrs old. The 3-6\,Myr old OBI\,c subgroup consists of stars around the Sword (about 4$^{\circ}$ below the Belt asterism). The older stars in the OB\,Ic are superimposed on the much younger and still embedded subgroup OB\,Id, which is associated with the Orion nebula and including the Trapezium stars, M43, NGC 1977, and the OMC1, 2, and 3 regions in the integral shaped filament along with the northern part of the Orion A molecular cloud (age $<$ 2\,Myr, d$\sim$\,420pc, \cite{Bally2008}). Although initially each subpopulation was assumed to be a distinct episode in a large star formation event, it was early realized that the subgroups are partially superimposed along our line of sight and several authors have described the boundaries between subgroups, their characteristics, and some discrepancies with the sequential star formation scenario \cite[e.g.,][]{Warren1978,deGeus1990,Brown1994,Gomez1998}. Unfortunately, the three-dimensional arrangement of star-forming regions, in particular massive ones, is far from simple and is essentially unknown for any massive star-forming region given the current distance accuracies. \begin{figure}[h!] \centering \includegraphics[width=\hsize]{Orion-crop.jpg} \caption{Widefield image of the Orion OB\,I stellar associations as described by A. \cite{Blaauw} and revised in \cite{Bally2008}. North is up, and east is left. Background image: R. Bernal Andreo, \url{www.deepskycolors.com}.} \label{fig:Orion} \end{figure} Recently, \cite{revisited1} and \cite{revisited2} presented evidence for a young and massive foreground population ($\sim5$ Myr) that is detached from the Orion A cloud but seen in projection toward it. They argue that this foreground population was formed about $4-5$ Myr ago in a different, but perhaps related, event in the larger Orion star formation complex and not in the existing Orion A molecular cloud like, for example, the Orion nebula cluster. This foreground population includes in part Blaauw's OB Ic population, but does not include the younger $\sigma$-Ori cluster, as suggested in \cite{Bally2008}. An intriguing result of their study was that the Orion A foreground population seemed to extend to the north, toward OB Ib, beyond the limits of their survey. This was later confirmed in \cite{Meingast2016} in their ESO-VISTA near-infrared (NIR) imaging of the entire Orion A cloud, who found a southern boundary to the foreground population, but not an obvious boundary toward the north. This raises the question of how well Ori OB\,Ic and Ib are separated spatially, if at all, and how they fit in a sequential star formation scenario. In a related study, \cite{streams} revisited the Hipparcos catalog and studied the spatial distribution in 3D of OB stars that are closer than 500\,pc from the Sun. Their analysis reveals that massive OB stars form large-scale structures that are well defined and elongated, which they refer to as ``blue streams''. The spatial coherence of these blue streams, and the monotonic age sequence over hundreds of parsecs, suggest that they are made of young stars. The two main blue streams are the Sco-CMa stream, including the Sco-Cen association, and the surprising Orion stream, originating in the Orion clouds and extending to regions as close to Earth as $\sim200$ pc, but likely even closer. In this scenario, the foreground population presented in \cite{revisited1} and \cite{revisited2} could be part of the Orion blue stream. Given this new scenario for the distribution of young stars in the local neighborhood, and in particular the realization that the Orion OB\,I association may be part of the Orion stream, there is a clear need to gather more information about the region on larger scales than collected previously and for regions further away from the molecular clouds. The main goal of this paper is to extend the work of \cite{revisited1} and \cite{revisited2} to the north and, (i) further investigate the extent of the young foreground population presented, (ii) investigate the relation between OB Ic and Ib populations, and (iii) contribute to the reconstruction of the star formation history of the Orion complex. We focus our study on Blaauw's subgroup OB\,Ib by studying almost 30 square degrees of sky centered on Orion's Belt. Within the limitations of our data, we compare the recently proposed Orion blue stream scenario with that of Blaauw's classical sequential star formation. The overdensity of blue massive stars in the Orion Belt region was first pointed out in Galileo's Sidereus Nuncius in 1609 as an example of how the telescope could resolve stars that are not visible by the human eye. The stellar overdensity was also recognized in 1931 by Swedish astronomer Per Collinder in his catalog of open clusters \citep{Collinder}. He distinguished the Orion's Belt asterism, comprised of the three famously aligned bright stars: Alnitak ($\zeta$ Ori, HD\,37742J, O9.7\,Ib+B0\,III), Alnilam ($\epsilon$ Ori, HD\,37128, B0\,Ia), and Mintaka ($\delta$ Ori, HD\,36486, B0\,III +O9\,V) as Collinder 70 (Col 70). Still, and even though it is immediately recognizable to the naked eye, the Orion Belt stellar population is paradoxically poorly known. \cite{Caballero_Solano2008} observed two circular areas of 45 arcmin radius each, centered on Alnilam and Mintaka and found 136 low-mass stars displaying features of extreme youth, and a total of 289 young stars in the surveyed area. They concluded that the two regions could be analogs to the $\sigma$ Ori cluster, but more massive, extended, and slightly older. Since the seminal work of Blaauw, it has been suggested that the age, distance, and radial velocity of the stellar components of subgroup OB\,Ib may not be consistent with a simple sequential star formation scenario (e.g., \cite{Hardie1964}; \cite{Warren1978}; \cite{Guetter1981}; \citealt{Gieseking1983}). In particular, the eastern part of the subgroup, which includes Alnitak, the Horsehead Nebula, the Flame Nebula (associated with NGC\,2024 in the Orion B cloud), and the H\,II region IC\,434, would be the farthest and youngest subgroup. The fourth brightest star in Orion's Belt is $\sigma$ Ori (48\,Ori, HD\,37468, O\,9.5\,V), the brightest source in the well-studied $\sigma$ Orionis Cluster \citep{Walter1997}, which has been assigned to OB\,Ib based on its spatial proximity. Still, two solid cases can be made against $\sigma$ Ori belonging to OB\,Ib. These two cases, which are discussed later in this paper, are, first, the age of the $\sigma$-Ori cluster (3 Myr; \cite{Caballero2008}) is younger than most of the stars in the Belt region and, second, the radial velocity of stars toward the cluster shows that the young $\sigma$-Ori cluster consists of two spatially superimposed components that are kinematically separated by 7 km/s in radial velocity \citep{Jeffries2006}. In the review of \cite{Bally2008} the $\sigma$-Ori cluster appears as member of OB\,Ic, as in Figure~\ref{fig:Orion}. This paper is organized as follows. The next section briefly describes the data used in this study. Section \ref{sec:results} and \ref{sec:properties} present our results, centering on the discovery of a large population of young stars around $\epsilon$ Ori. In Section \ref{sec:discussion} we discuss our results and we summarize them in Section \ref{sec:summary}. | \label{sec:discussion} We have found a population of about 2\,500 M stars (with 789 candidates with high probability) that are roughly coeval and extinction free and are distributed across $\sim3$ square degrees toward the Orion Belt asterism with an age of about 10 Myr. Photometry alone poorly constrains the distance to this population or its line-of-sight extent. This newly identified population can be as far as $\sim380$ pc (but in front of the Orion B cloud) or as close as 250 pc. In the closer case the OBP could be the low-mass counterpart to the well-known Orion supergiants at distances around 250 pc. The new population, the OBP, is likely the low-mass counterpart of Blaauw's Ori OB Ib subgroup. Relevant to this discussion, \cite{Jeffries2006} performed radial velocity observations of low-mass stars toward a relatively large field toward $\sigma$ Ori and found two spatially superimposed components that are kinematically separated by 7 km/s in radial velocity and with different mean ages. These authors suggest an age of about 10 Myr for the older component (their ``group 1''), which has a mean radial velocity of 23.8 km/s. \cite{Jeffries2006} suggested that the older ``group 1'' was made by stars from the OB Ia subgroup, but the results in this paper suggest that this second component is most likely comprised of the stars in the Orion Belt population, or the OB Ib subgroup. Figure 2 of their paper further supports this statement as one can see how the field closer to the OBP (the NW field) is mostly dominated by stars belonging to the 23.8 km/s group. We cross-checked our list of targets against the sources in \cite{Jeffries2006} and although their study is centered on $\sigma$ Ori, we found that most of the matches with the OBP belong to the older group 1; albeit this finding also matches group 2, which probably suggests that our selection method is not accurate enough to clearly separate the two different populations. Overall, our results reinforce the idea that overlapping populations at different evolutionary states and distances coexist along lines of sight toward the Orion clouds, as suggested in \cite{revisited1} and \cite{revisited2}. \subsection{Is the Orion sequential star formation scenario in trouble?} Blaauw's original idea of sequential star formation calls for a star formation event being directly responsible for the genesis of the next event. In Orion, it was proposed that the spatial-temporal sequence of events proceeded as follows \citep[e.g.,][]{Bally2008}: \vspace{0.5cm} \noindent {\small Ia ({$\sim 12$ \rm{M}yr)} $\rightarrow$ Ib ($\sim 10$ \rm{M}yr) $\rightarrow$ Ic ($\sim 5-7$ \rm{M}yr) $\rightarrow$ Id ($\sim 1-3$ \rm{M}yr).} \vspace{0.1cm} \noindent In recent decades, evidence has been accumulating suggesting that this attractive scenario suffers from several shortcomings. \cite{Brown1994} found that subgroup Ib is younger than Ic and, to address the problem of an obvious break in the spatial-temporal sequence, these authors argued that the sequential star formation scenario is still plausible if the Ic population had moved from its putative birthplace closer to the Ia population; this move has yet to be quantified. Nevertheless, if the OBP is indeed the low-mass counterpart of Ib, then the results in this paper are in tension with \cite{Brown1994} as we find that the age of the OBP is similar to the canonical age of Ib (around 10 Myr), apparently solving the break in the spatial-temporal sequence. Another problem for the sequential star formation scenario is the superposition of populations with different ages, as they do not easily fit a star formation sequence that covers about 100 pc from west to east. Evidence for such overlapping stellar populations has been accumulating in the literature \cite[e.g.,][]{Gomez1998,WarrenHesser1977,Jeffries2006,revisited1,revisited2}. For example, what event triggered the formation of the 1-3 Myr old $\sigma$ Ori cluster, seen along the same line of sight as the $\sim10$ Myr old OBP? Given that the Id subgroup is still forming stars and that Ia is too removed/old to be the trigger, one faces two options in a sequential star formation scenario: the trigger was either a) Ib or b) Ic, a subgroup closer to $\sigma$ Ori in age but not in projection. If a) then one needs to explain the roughly 7-8 Myr delay in the formation of $\sigma$ Ori. If b) one needs to explain the apparent break in spatial sequence (Ic is about 20 pc away from $\sigma$ Ori in projection, so probably in reality more). Option a) seems unlikely as Ib would have to trigger the formation of Ic to the southeast 5-7 Myr ago and the $\sigma$ Ori cluster 2-3 Myrs ago toward its background, as seen from Earth. Regarding option b), a possible solution to the break of the spatial-temporal sequence is to evoke that $\sigma$ Ori was formed elsewhere. This was recently suggested, in a different context by \citep{Ochsendorf_2015}. In the Ochsendorf-Tielens scenario $\sigma$ Ori was formed to the south (in galactic coordinates, see their Figure 1) of the Ic population and moved north toward the GS206-17+13 shell. It is hard to imagine how the feedback from Ic to the south would trigger the formation of the $\sigma$ Ori cluster and cause it to move north. In summary, neither option seems satisfactory. \subsection{The Orion blue stream scenario} Recently, \cite{streams} suggested a new scenario for the interpretation of the distribution of OB stars in the local neighborhood. In a reanalysis of the Hipparcos catalog, these authors found that the distribution of OB stars followed large-scale structures that are well-defined and elongated, which they refer to as blue streams. The roughly constant width of the streams, together with a monotonic age sequence over hundreds of parsecs, suggests that they are the outcome of a large star formation event. They describe the existence of three streams in the local 500 pc neighborhood, one of these is the Orion stream, originating at the position of the Orion clouds and extending to regions as close to Earth as $\sim200$ pc, but likely even closer. This scenario imposes a well-defined age sequence as it assumes that young stars stream away from their place of birth, currently the Orion A and B molecular clouds. The further a population is from its birth place, the older it should be. Given the current position of the Sun in the Galaxy, the Orion stream appears projected along its length for an observer on Earth, which implies that stellar populations at different ages and distances should appear superposed. Because of the particular projection effect the Orion stream, this new scenario does not require a spatial sequence, unlike Blaauw's sequential scenario for Orion. The new blue streams scenario appears to accommodate well the available observational data of the Orion star-forming region as a whole. As discussed above, there is plenty of evidence in the literature for superposition of populations with different ages along the direction to the Orion clouds. For example, the OBP fits well this new view of Orion as a stream projected along its length; a roughly 10 Myr old population is seen in projection toward a significantly younger $\sigma$ Ori, and an even younger NGC2024 cluster, still embedded in the Orion B cloud. The OBP, we argue, is closer to Earth than $\sigma$ Ori and the cloud interacting with it via the H${\rm II}$ region. We also argue that there is some evidence that the OBP (Ib subgroup) is closer and older than NGC1980/NGC1981 (Ic subgroup), which is closer and older than the $\sigma$ Ori cluster. This age and distance relation is in good agreement with the blue streams scenario presented in \cite{streams}. Finally, the Ia subgroup and the 25 Ori cluster \citep[e.g.,][]{Briceno2008,Downes2014,Downes2015}, which are not addressed in this work, would also be part of the Orion stream. If older than the OBP, they should correspond to the nearest components of the Orion stream. But this remains to be confirmed, as the OBP could be older, hence nearer, and could be the low-mass counterpart to the well-known Orion supergiants at about 250 pc from Earth. The streams scenario provides another advantage: it does not require that populations move substantially from each other, as proposed in \cite{Brown1994}, to solve the apparent break in the spatial-temporal sequence in Blaauw's scenario. In the streams scenario, the OB subgroups should have a space motion toward the same general direction, so a prediction of the streams scenario is that the proper motions between subgroups should be relatively small. In the streams scenario, Blaauw's OB subgroups could represent different components of the same stream with different ages and distances; these are all formed at about 400 pc by clouds long gone with the exception of subgroup Id, the ONC, embedded in Orion A cloud, and NGC2024 embedded in the Orion B cloud. \subsection{Is the OBP the future of the ONC?} Can the OBP be the evolved counterpart of an ONC-like cluster that was formed about 10 Myrs ago or is it an altogether different type of object? A striking property of the OBP population is that it is distributed over a large area of the sky and its low stellar density is very different from other well-known stellar clusters in Orion, such as the ONC, $\sigma$, $\lambda$, or $\iota$ Ori clusters. For example, the average stellar surface density in the Orion nebula cluster ($\sim200$ stars/pc$^{2}$; \cite{Hillenbrand_Hartmann_1998}) is about an order of magnitude higher than that of the OBP. The volume density of stars in the core of the ONC ($2-3\times10^4$ stars/pc$^{3}$, \cite{Hillenbrand_Hartmann_1998}) is about three orders of magnitude higher than that of the OBP. Remarkably, both the ONC and the OBP have a similar number of stars. Assuming that the OBP is not a gravitationally bound population, and that it is expanding freely since it got rid of its parental molecular cloud early in its formation, it would have taken about 5 Myr for the OBP\ to expand from 2 pc to 7 pc radius at about 1 km/s expansion velocity, or 10 Myr for 0.5 km/s velocity. These rough estimates are not implausible according to models of an ONC-type cluster expanding after gas removal \citep[e.g.,][]{Kroupa_2001}, and so the possibility that the OBP might represent an evolved ONC cannot be discarded with current data. High-resolution spectroscopic observations or accurate proper-motions measurements are needed for a more quantitative answer to this question; such measurements and observations do not exist at the moment. In order to find the spatial extension of the foreground stellar population to the Orion A cloud found in \cite{revisited1} and \cite{revisited2}, and to investigate the relation between Blaauw's OB\,Ic and Ib subgroups, we analyzed a circular area with a radius of 3$^\circ$ centered on $\epsilon$ Orionis (HD 37128, B0Ib), covering the Orion Belt region. The main results of this investigation are as follows: \begin{itemize} \item We found two large stellar overdensities in the Orion Belt region: one centered on the well-known $\sigma$-Ori cluster, and a new, richer but more extended overdensity close to $\epsilon$ Ori. We compared the stellar density in the surveyed region with a control field and estimated an upper limit for the size of the new overdensity of about 2\,345$\pm$215 sources. \item Optical and near-IR color-magnitude diagrams reveal a well-defined sequence above the Galactic field, which is suggestive of a large young stellar population that is approximately coeval and not affected by interstellar extinction. We used a new statistical multiband technique to select objects associated with the sequence detected in the color-magnitude diagram, and compiled a catalog of 783 probable members. Essentially, all of these objects have the colors of M stars. The selected sources are close, in projection, to $\epsilon$ Ori, but distributed in a roughly elliptical region ($1^\circ \times 3^\circ $) showing spatial substructure. \item This new population, that we call the Orion Belt population, is likely the low-mass counter part to the Ori OB Ib subgroup. We found a negligible amount of bona fide young stellar objects in the Orion Belt population (less than 2 \% for all available youth tracers (\textit{XMM-Newton}, KISO, and WISE surveys). This allows us to infer the minimum age of the cluster to be $\sim5$\,Myr. We estimate an age of about $\sim10$\,Myr for the OBP. \item We do not find evidence for an interaction between the selected members and the clouds, which together with the overall absence of extinction suggests that the new population lies in the foreground of Orion~B. We estimate the distance to this newly identified population to be between $\sim$250 and $\sim$380\,pc. \item Although our results do not rule out Blaauw's sequential star formation scenario for Orion, we argue that the current available evidence is shifting against it. We find, instead, that the blue stream scenario proposed in \cite{streams} provides a better framework on which one can explain the Orion star formation region as a whole. \item We speculate that the Orion Belt population could represent the evolved counterpart of a Orion nebula-like cluster. At least high-resolution spectroscopic data would be needed to make a more solid statement about the origin of this newly identified population. \end{itemize} Finally, although we argue that the OBP fits the blue stream scenario best, we caution that independent work is needed to confirm the existence of the blue streams. Nevertheless, giving the tantalizing proximity and youth of the new stellar population presented in this work, there is a need for a dedicated spectral and dynamic characterization of the OBP. This population could become a benchmark region for future searches of brown dwarfs and planetary mass objects and the low-mass end of the IMF, as well circumstellar disk evolution and planet formation. The final ESA Gaia catalog, to be released around 2023, will include much if not all of the OBP candidates presented in this work, and will be able to shed much light on the origin of the OBP, the existence and role of the Orion blue dtream, and star formation in Orion. \begin{appendix} | 16 | 9 | 1609.04948 |
1609 | 1609.06726_arXiv.txt | {} {The aim of this work is twofold: first, to assess whether the population of elliptical galaxies in cluster at $z\sim1.3$ differs from the population in the field and whether their intrinsic structure depends on the environment where they belong; second, to constrain their properties 9 Gyr back in time through the study of their scaling relations.} {We compared a sample of 56 cluster elliptical galaxies selected from three clusters at $1.2<z<1.4$ with elliptical galaxies selected at comparable redshift in the GOODS-South field ($\sim30$), in the COSMOS area ($\sim180$), and in the CANDELS fields ($\sim220$). To single out the environmental effects, we selected cluster and field elliptical galaxies according to their morphology. We compared physical and structural parameters of galaxies in the two environments and we derived the relationships between effective radius, surface brightness, stellar mass, and stellar mass density $\Sigma_{R_e}$ within the effective radius and central mass density $\Sigma_{1kpc}$ , within 1 kpc radius.} {We find that the structure and the properties of cluster elliptical galaxies do not differ from those in the field: they are characterized by the same structural parameters at fixed mass and they follow the same scaling relations. On the other hand, the population of field elliptical galaxies at $z\sim1.3$ shows a significant lack of massive ($\mathcal{M}_*> 2 \times 10^{11}$ M$_\odot$) and large ($R_e > 4-5$ kpc) elliptical galaxies with respect to the cluster. Nonetheless, at $\mathcal{M}_*< 2 \times 10^{11}$ M$_\odot$, the two populations are similar. The size-mass relation of cluster and field ellipticals at $z\sim1.3$ clearly defines two different regimes, above and below a transition mass $m_t\simeq2-3\times10^{10}$ M$_\odot$: at lower masses the relation is nearly flat (R$_e\propto\mathcal{M}_*^{-0.1\pm0.2}$), the mean radius is nearly constant at $\sim1$ kpc and, consequenly, $\Sigma_{R_e}\simeq\Sigma_{1kpc}$ while, at larger masses, the relation is R$_e\propto\mathcal{M}_*^{0.64\pm0.09}$. The transition mass marks the mass at which galaxies reach the maximum stellar mass density. Also the $\Sigma_{1kpc}$-mass relation follows two different regimes, above and below the transition mass ($\Sigma_{1kpc}\propto{\mathcal{M}_*}^{0.64\ >m_t}_{1.07\ <m_t}$) defining a transition mass density $\Sigma_{1kpc}\simeq2-3\times10^3$ M$_\odot$ pc$^{-2}$. The effective stellar mass density $\Sigma_{R_e}$ does not correlate with mass; dense/compact galaxies can be assembled over a wide mass regime, independently of the environment. The central stellar mass density, $\Sigma_{1kpc}$, besides being correlated with the mass, is correlated to the age of the stellar population: the higher the central stellar mass density, the higher the mass, the older the age of the stellar population. } {While we found some evidence of environmental effects on the elliptical galaxies as a population, we did not find differences between the intrinsic properties of cluster and field elliptical galaxies at comparable redshift. The structure and the shaping of elliptical galaxies at $z\sim1.3$ do not depend on the environment. However, a dense environment seems to be more efficient in assembling high-mass large ellipticals, much rarer in the field at this redshift. The correlation found between the central stellar mass density and the age of the galaxies beside the mass suggests a close connection of the central regions to the earliest phases of formation. } | The existence of correlations among some properties of the population of galaxies and the environment in which they reside is well established. The composition of the population of galaxies, that is, its morphological mix, is different according to the environment where the population belongs. A clear example is the well-known morphology-density relationship, according to which early-type galaxies, originally classified as elliptical and lenticular galaxies, preferentially populate high-density environments and vice versa \citep{oemler74,dressler80,postman84}. This environmental effect has been confirmed both in the local \citep[e.g.,][]{tran01, goto03,holden07,bamford09} and intermediate redshift Universe \citep[e.g.,][]{fasano00,treu03,smith05, vanderwel07,pannella09,tasca09,tanaka12}. In spite of the many observations supporting the above evidence, the mechanisms responsible for this morphological segregation are still debated. Galaxies in different environments can undergo different physical processes. For instance, contrary to field galaxies, cluster galaxies are affected by the dense and hot intracluster medium. The ram pressure can overcome the gravitational forces keeping the gas anchored to the potential well, at least of the less massive galaxies, removing their gas and quenching their star formation. Actually, the quenching efficiency seems to be higher in denser environments \citep[e.g.,][]{haines13, vulcani15}. This mechanism affects galaxies in a different way according to their mass and shape \citep[see e.g.,][for recent reviews]{boselli06, boselli14}. Many observations suggest that the formation epoch of galaxies depends mainly on their mass but, in the local universe, there is some evidence that cluster early-type galaxies form earlier than field galaxies of the same mass \citep[e.g.,][]{kuntschner02,gebhardt03,thomas05} and that the environment can play an important role in the late phases of their evolution \citep{thomas10}. On the other hand, some works seem to conclude that while galaxy mass regulates the timing of galaxy formation, the environment regulates the timescale of their star formation histories \citep[e.g.,][]{tanaka10,rettura11} and controls the fraction of star-forming galaxies \citep[e.g.,][]{muzzin12}. Actually, the star formation timescale appears shorter for galaxies in dense environments than for those in low-density fields \citep[e.g.,][]{thomas05,rettura11,tanaka13}. At intermediate redshift ($z\sim0.8$), however, the effect of the environment becomes less evident and it seems to vanish, with cluster and field early-type galaxies characterized by similar stellar population properties \citep[e.g.,][]{lonoce14} following the same color evolution and scaling relations \citep[see][for a review]{renzini06}. Actually, it is not clear whether and when the environment affects the properties and hence the evolution of galaxies at a given morphology. From the theoretical point of view, it is expected that field and cluster bulge-dominated galaxies display different structures and hence follow also different scaling relations. Bulge-dominated galaxies should result from a sequence of major and minor mergers through which most of their stellar mass is assembled \citep[e.g.,][]{delucia06, khochfar11, shankar13}. Minor mergers are considered very efficient in increasing the size of galaxies \citep[e.g.,][]{naab09,vandokkum10a}. Since mergers are expected to be more frequent in denser environments, they should produce larger galaxies than similarly massive counterparts in the field \citep[e.g.,][]{shankar13}. Some recent simulations actually predict a clear environmental dependence of the structure of bulge-dominated galaxies with their median size larger by a factor 1.5-3, moving from low to high-mass halos \citep{shankar14b}. From an observational standpoint, many recent studies focused on the environmental dependence of the mass-size relation of early-type galaxies. In the local universe, some works point toward the absence of an environmental dependence for this relation \cite[e.g.,][]{guo09, weinmann09,huertas13b, shankar14b} while other studies suggest that cluster early-type galaxies are slightly smaller than their field counterparts \citep[e.g.,][]{valentinuzzi10, poggianti13a}. Few studies at intermediate redshift point toward the absence of environmental effect on the size distribution of early-type galaxies, either morphologically- spectroscopically-, or color-selected \citep[e.g.,][]{maltby10,rettura10,kelkar15}. At higher redshift there are rather controversial results. For instance, while \cite{raichoor12} find that morphologically selected early-type galaxies in cluster at $z\sim1.2$ are more compact than in the field, the opposite is found by \cite{cooper12} at similar redshift and by \cite{papovich12} at slightly higher redshift, both works based on different selection criteria and data quality. Different selection criteria, different redshift ranges, different quality of the data, and hence different accuracy in the derivation of the structural and physical parameters (size, stellar mass, age) of galaxies may be the reasons for, at least, some of the above discrepancies. A critical issue in this kind of analysis is, indeed, the morphological selection of galaxies instead of using selection criteria related to the stellar population properties such as colors and/or star formation. Since stellar population and structural evolution do not appear synchronous, criteria based on stellar population properties select galaxies with different morphological mixes at different redshift and in different environments. A clear example is given by the significant different morphological mix observed in the red sequence population of cluster galaxies, largely populated by red disc-dominated (passive) galaxies at $z\sim1,$ and by ellipticals and lenticular in the local universe \citep[e.g.,][]{depropris15, mei09,moran07}. Analogously, the selection of passive galaxies based on color-color diagnostic plots or on their low specific star formation rate (sSFR) produces samples with a different mix of morphological types. For instance, the fraction of elliptical galaxies in the passive galaxy population is found to significantly change with mass at a given redshift and in redshift at fixed mass \citep[e.g.,][]{moresco13, tamburri14, huertas13, huertas15} and consequently the mean properties of the sample vary \citep[e.g.,][]{bernardi10,mei12}. The different composition of the samples thus selected prevents the singling out of possible environmental effects on the properties of a given morphological type. In this paper we aim to study in a coherent and homogeneous way the dependence of the population of elliptical galaxies, and of their properties, on the environment. Here, we focus our attention on cluster and field elliptical galaxies at $z\sim1.3$, while we refer to a forthcoming paper for the environmental effects on their evolution. We study a sample of 56 cluster elliptical galaxies selected in the three clusters: XMMJ2235-2557 at $z=1.39$ \citep{rosati09}, RDCS J0848+4453 at $z=1.27$ \citep{stanford97}, and XLSS-J0223-0436 at $z=1.22$ \citep{andreon05,bremer06}. We compare their properties with those of a sample of 31 field elliptical galaxies selected in the GOODS-South field according to the same criteria. When possible, we make use also of a larger sample of about 180 elliptical galaxies selected in the same way from the COSMOS catalog at slightly lower redshift, and of a sample of about 220 ellipticals selected from CANDELS. To single out the effect of the environment, we have tried to minimize all the sources of uncertainty discussed above: we selected galaxies in a narrow redshift range, $1.2<z<1.4$, to avoid significant evolutionary effects; we selected cluster and field elliptical galaxies on the basis of their morphology to compare samples with the same composition in the two environments; we derived morphology and structural parameters from Hubble Space Telescope (HST) images at the same wavelength (with the exception of CANDELS). Finally, the same wavelength coverage for all the galaxies allowed us to derive their physical parameters (stellar mass, age) with the same degree of uncertainty. In Sec. 2 we describe the data and the samples. In Sec. 3 we derive the structural (effective radius, surface brightness) and the physical (stellar mass, absolute magnitude, and age) parameters for our galaxies. In Sec. 4 we compare the population of cluster elliptical galaxies and their properties with those in the field. In Sec. 5 we derive the Kormendy relation of cluster and field ellipticals at $z\sim1.3$ while, in Sec. 6, we derive the size-mass relation. Section 7 is focused on the stellar mass density of galaxies. In Sec. 8, we summarize our results and present our conclusions. Appendix \ref{ap:scaling} summarizes the best fitting relations reported in the text obtained with the least squares method and reports also those obtained using the orthogonal regression. Throughout this paper we use a standard cosmology with $H_0=70$ Km s$^{-1}$ Mpc$^{-1}$, $\Omega_m=0.3$, and $\Omega_\Lambda=0.7$. All the magnitudes are in the Vega system, unless otherwise specified. | In this paper we focused our investigation on the dependence of the population of elliptical galaxies at $z\sim1.3$ and of their properties on the environment. We constructed two main samples of elliptical galaxies at the same redshift, the first one composed of 56 galaxies selected in three clusters at $1.2<z<1.4$, the second one composed of 31 galaxies selected in the GOODS-South field at the same redshift. A third and larger ($\sim180$ galaxies) sample of field galaxies has been extracted from the COSMOS catalogs at slightly lower redshift ($1.0<z<1.2$) and used when it has been possible in some of the comparisons. The selection of galaxies has been made on the basis of a pure morphological criterion based on the visual inspection of their luminosity profile in the ACS-F850LP image and of the residuals resulting from the profile fitting with a regular Sersic profile. The narrow redshift range adopted minimized the evolutionary effects, while the morphological selection of galaxies produced samples with the same composition, allowing us to single out the effect of the environment at a given morphology. We compared the physical and structural properties of the population of elliptical galaxies in the two environments and we derived and compared the relationships among effective radius, surface brightness, stellar mass, and stellar mass density. Our main results can be summarized as follows: \begin{itemize} \item The structure and the properties of cluster elliptical galaxies do not differ from those in the field. Cluster and field elliptical galaxies have the same median effective radius, the same mass normalized radius, and the same stellar mass density at fixed mass. \item Cluster and field elliptical galaxies at $z\sim1.3$ follow the same Kormendy (size-surface brightness) relation with a slope $\beta\simeq3.0$. They also follow the same size-mass relation and the same size-mass density relations. The comparison of our data with those from the literature at the same redshift and with comparable selection criteria shows excellent agreement. \item The population of cluster elliptical galaxies differs from the one in the field for high-mass and large elliptical galaxies. Indeed, there is a significant lack of massive ($\mathcal{M}_*>2\times10^{11}$) M$_\odot$ and large (R$_e>4-5$ kpc) elliptical galaxies in the field with respect to the cluster. Nonetheless, at $\mathcal{M}_*< 2 \times 10^{11}$ M$_\odot$, the two populations are similar. There seems to be also a lack of elliptical galaxies in the field which are as old as the oldest ellipticals in cluster. \end{itemize} Hence, we do not find a dependence of the structure of elliptical galaxies on the environment where they belong, contrary to what has been predicted by some recent simulations. The above results show that the structure and the shaping of elliptical galaxies at $z\sim1.3$ do not depend on the environment where they belong. However, they suggest that a dense environment is more efficient in assembling high-mass elliptical galaxies. From the study of the scaling relations, we obtained the following results: \begin{itemize} \item The size-mass relation is only coarsely best fitted by a single power law of the form R$_e\propto\mathcal{M}_*^{0.5}$. We find a transition mass $m_t\sim2-3\times 10^{10}$ M$_\odot$, in agreement with local studies, that defines two different regimes of the size-mass relation. Above this mass, the relation is steeper (R$_e\propto\mathcal{M}_*^{0.64}$), while below this mass the relation gets flat (R$_e\propto\mathcal{M}_*^{-0.13}$), and the effective radius is nearly constant at $\sim$1 kpc. The transition mass marks the mass at which a galaxy can reach the maximum stellar mass density within the effective radius. \item The stellar mass density within the effective radius is tightly correlated with the effective radius. This size-effective mass density relation is best fitted by $\Sigma_{R_e}\propto R_e^{-1.2}$ in agreement with the Kormendy relation. Galaxies with high or low effective stellar mass densities can be realized independently of the environment in which they reside and almost independently of their mass. The data show that a galaxy cannot occupy any locus in the [$\Sigma_{Re},\mathcal{M}_*$] plane but that a Zone of Exclusion, as the one defined by the early-type galaxies in the local universe, exists at $z\sim1.3$. \item The central stellar mass density within 1 kpc radius is tightly correlated with mass and well fitted by $\Sigma_{1kpc}\propto\mathcal{M}_*^{0.64}$, as previously found, at masses larger than $m_t$. At lower masses, the relation steepens with $\Sigma_{1kpc}\propto\mathcal{M}_*^{1.07}$. We find that the central mass density is also correlated with the age of the stellar population such that the higher the central stellar mass density, the older the age, the higher the mass. \end{itemize} High mass galaxies are characterized by correspondingly high central stellar mass densities and old stellar population. The scaling of the central mass density with mass suggests that the process of assembly and of shaping of elliptical galaxies does not depend on the mass, at least in the mass range considered. These correlations taken all together, are indicative of the close connection of the central regions of the galaxies to the earliest phases of formation. The central regions of elliptical galaxies most probably store the information on their assembly, retain memory of the initial conditions, and are strictly connected to the bulk of their stellar mass growth. The outer regions (effective stellar mass density and effective radius) instead, store the information on the (subsequent) events that the galaxy may have experienced and that have affected the outer structure, but not its mass. | 16 | 9 | 1609.06726 |
1609 | 1609.01106_arXiv.txt | We study supersonic Evershed downflows in a sunspot penumbra by means of high spatial resolution spectropolarimetric data acquired in the \ion{Fe}{1}~617.3~nm line with the CRISP instrument at the Swedish 1-m Solar Telescope. Physical observables, such as Dopplergrams calculated from line bisectors and Stokes~$V$ zero-crossing wavelengths, and Stokes~$V$ maps in the far red wing, are used to find regions where supersonic Evershed downflows may exist. We retrieve the LOS velocity and the magnetic field vector in these regions using two-component inversions of the observed Stokes profiles with the help of the SIR code. We follow these regions during their lifetime to study their temporal behavior. Finally, we carry out a statistical analysis of the detected supersonic downflows to characterize their physical properties. Supersonic downflows are contained in compact patches moving outward, which are located in the mid and outer penumbra. They are observed as bright, roundish structures at the outer end of penumbral filaments that resemble penumbral grains. The patches may undergo fragmentations and mergings during their lifetime, even some of them are recurrents. Supersonic downflows are associated with strong and rather vertical magnetic fields with a reversed polarity compared to that of the sunspot. Our results suggest that downflows returning back to the solar surface with supersonic velocities are abruptly stopped in dense deep layers and produce a shock. Consequently, this shock enhances the temperature and is detected as a bright grain in the continuum filtergrams, which could explain the existence of outward moving grains in the mid and outer penumbra. | \label{sec:introduction} The penumbra harbors a plethora of mass motions due to the presence of magnetoconvection in sunspots \citepads{1965ApJ...141..548D}. Among the most intriguing and less understood are supersonic\footnote{The sound speed in the photosphere is 7.2~km~s$^{-1}$.} downflows. The existence of strong Evershed downflows in the penumbra is suspected since long. They were first reported by \citet{bumba60} as line flags, where the line is not shifted as a whole but instead has a very strong asymmetry (flag). Such extreme cases of line asymmetries were interpreted in terms of a superposition of two unresolved structures: an almost unshifted strong component, which corresponds to the "mean" Evershed effect \citepads{1909MNRAS..69..454E}, and a weaker strongly displaced component (satellite) where the line-of-sight (hereafter, LOS) velocity reaches up to 7-8~km~s$^{-1}$ (\citeads{1960IAUS...12..403S}, \citeads{1961AnAp...24....1S};~\citeads{1964ApNr....8..205M};~\citeads{1971SoPh...18..220S}; ~\citeads{1995A&A...298L..17W}). In the early 2000s, inversion codes allowed a much more precise determination of flows in sunspot penumbrae through the use of full Stokes line profiles. \citetads{2001ApJ...549L.139D} and ~\citetads{2004A&A...427..319B}, for example, inferred the existence of supersonic downflows from the inversion of spectropolarimetric measurements in the infrared showing Stokes~$V$ profiles with strong asymmetries and multiple lobes. However, the downflows could not be imaged directly due to insufficient spatial resolution of the observations (about 1\arcsec). The advent of Hinode satellite \citepads{2007SoPh..243....3K} brought spectropolarimetric data with improved spatial resolution (0\farcs32). Using a set of magnetograms in the red far wing of \ion{Fe}{1}~630.2~nm, \citetads{2007PASJ...59S.593I} found the sinks of the Evershed flow and circular polarization measurements enter the scene. According to these authors, the bisector technique yields LOS velocities of 4-7~km~s$^{-1}$ in such sinks. Thereafter, \citetads{2010ASSP...19..193B} was the first to show a reversed two-lobed Stokes~$V$ profile associated to a LOS velocity of about 9 km~s$^{-1}$ (in concordance with those obtained by \citetads{2009A&A...508.1453F} from bisectors in pixels with rare circular polarization signals). Independently to the finding of \citetads{2010ASSP...19..193B}, the spatial resolution acquired by the Hinode satellite is not high enough to spatially resolve structures harboring supersonic downflows and Stokes~$V$ profiles are usually irregular, showing 3 or 4 lobes. In this regard, \citetads{2013A&A...557A..24V} has been the first to infer the LOS velocity and the magnetic field vector associated to supersonic downflows from spatially-coupled inversions \citepads{2012A&A...548A...5V}. However, they inferred extreme LOS velocity and magnetic field strength values in some cases (around 20~km~s$^{-1}$ and 7~kG, respectively). Furthermore, the Stokes profiles represented in their Figure~3 show very weak linear polarization signals and a very complex (pathological) Stokes~$V$ profile that has been inverted using only a component. \begin{figure*}[!t] \centering \includegraphics[trim = {3cm 0.5cm 3cm 1cm}, clip, width = 0.85\textwidth]{fig1.eps} \caption{Top row: continuum intensity image of the main sunspot of AR 11302 observed on 28 September 2011, 10:07:36~UT (left panel) and Dopplergram as derived from the line bisectors at the 70\% intensity level (right panel). Bottom row: the corresponding Stokes~$V$ map at +420~m\AA\, and the LOS velocity map given by the Stokes~$V$ zero-crossing points. Squares enclose three regions that will be studied in Section~\ref{sub:temporal}. The contours outline the penumbral region. The black arrow points to the disk center. Each major tickmark represents 10\arcsec.} \label{fig:intensity_observables} \end{figure*} Despite all previous efforts, we still do not have direct proofs of the existence of supersonic downflows in the penumbra. We do not know if they leave spectral signatures on all four Stokes profiles, or only in Stokes~$I$ and~$V$. Moreover, an important (and missing) element is the temporal evolution: to know how they appear, evolve and disappear is essential to understand the nature of supersonic downflows. Our motivation is entirely focused on the aspects mentioned above. Specifically, we study in-depth for the first time the morphology and temporal evolution of supersonic downflows in a sunspot penumbra, as well as characterize their physical properties. This goal has been achieved thanks to a time sequence of high spatial resolution (0\farcs13) and high cadence spectropolarimetric data of the \ion{Fe}{1}~617.3~nm spectral line acquired with the CRISP instrument at the Swedish 1-m Solar Telescope under excellent seeing conditions. This paper is organized as follows. The observations and the data reduction process are briefly described in Section~\ref{sec:observations_datareduction}. This is followed by Section~\ref{sec:dataanalysis} where we define the proxies used to detect areas susceptible of supersonic Evershed downflows. In Section~\ref{sec:supersonic}, we characterize the shape and the spatial distribution of the Stokes profiles emerging from patches harboring supersonic downflows and, finally, we describe how inversions were performed. In Section~\ref{sec:results}, we outline the temporal behavior of supersonic Evershed downflows by analyzing the temporal evolution of three spatially-resolved examples. After that, the physical properties of the detected supersonic downflows are enumerated. These results are discussed and compared with previous studies (Section~\ref{sec:discussion}). Finally, Section~\ref{sec:conclusion} offers a summary of our results and a conclusion. | \label{sec:conclusion} In this paper we have characterized the properties of supersonic Evershed downflows and we have described their temporal evolution for the first time. This has been possible thanks to the high spatial resolution and the excellent seeing quality of our spectropolarimetric data of a spot observed at only 6.8\degree\, from disk center. To detect supersonic Evershed downflows, we have considered information from the continuum intensity filtergrams, the LOS velocities given by the Stokes~$V$ zero-crossing wavelength, and the far-wing magnetograms, together with the LOS velocity and the magnetic field vector inferred from two-component inversions of the observed Stokes profiles. Supersonic Evershed downflows occur in the mid and outer penumbra. They are contained in compact patches that move outward and usually do not show strong far-wing Stokes~$V$ signals. They are often observed as bright and roundish features located at the outer end of a single filament or a more complex filamentary structure in intensity that resemble outward moving penumbral grains. However, sometimes it is not easy to identify the related filament in the intensity images. The lifetimes of the detected supersonic Evershed downflows vary from 1 to more than 5 minutes, showing a variable behavior. Most of the supersonic downflows show a LOS velocity between 7.5~and 9.5~km~s$^{-1}$, with a median value of 8~km~s$^{-1}$. We have detected peak LOS velocities of about 15~km~s$^{-1}$ in rare occasions. Their magnetic fields tend to be more vertical (by~30\degree) than those in the pixels surrounding them. The polarity of the magnetic field at the position of the supersonic Evershed downflow is reversed compared to that of the sunspot. The magnetic field strength of the supersonic downflows has a median value of 1.5~kG, which is a typical value within the penumbra. Furthermore, supersonic downflows tend to be cospatial with bright pixels, the median value of the continuum intensity being 0.91. Differences on the physical parameters are more significant when comparing supersonic pixels with those of the penumbral surroundings. From an analysis of the temporal evolution of the patches, we find that after appearing in the penumbra their LOS velocity increases and then fades. Subsequently, two behaviors are possible: the patch disappears and there are no more supersonic Evershed downflows, or the patch remains visible and a recurrent supersonic downflow is observed some frames later. During the evolution of the patch, the regions with stronger flows also harbor the more vertical magnetic field, but this does not mean that there is a one-to-one relation between the LOS velocity and the magnetic field inclination. Furthermore, we find that patches crossing the outer penumbral border experience LOS velocity and magnetic field enhancements, with no significant changes in the inclination. Once they leave the spot, such enhancements suddenly disappear and the patches fade in the quiet Sun. In addition, supersonic patches undergo mergings and fragmentations. When a merging occurs, the LOS velocity as well as the magnetic field inclination and the continuum intensity increase. Furthermore, there is a relation between magnetic field strength and inclination during interactions. This scenario suggests that the Evershed flow associated with different filaments can return back to the surface at the same position. In the case of fragmentations, the intensity structure brightens and is accelerated outward. At that moment, the supersonic Evershed downflow heads the feature and the patch breaks. The relation of supersonic patches with bright intensity features moving outward suggests that the nature of inner and outer penumbral grains is different. The Evershed flow returns back to the solar surface at the end of flow channels. As the downflow reaches supersonic velocities in the dense deep layers, it stops abruptly and produces a shock. Consequently, there is a temperature enhancement that increases the continuum intensity, which is observed as a bright penumbral grain moving outward. \\ \noindent \emph{Acknowledgements.} Financial support by the Knut and Alice Wallenberg Foundation and by the Spanish Ministerio de Econom\'{\i}a y Competitividad through grants ESP2013-47349-C6-1-R and ESP2014-56169-C6-1-R, including a percentage from European FEDER funds, are gratefully acknowledged. This paper is based on data acquired at the Swedish 1-m Solar Telescope, operated by the Institute for Solar Physics of Stockholm University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof\'{\i}sica de Canarias. This research has made use of NASA's Astrophysical Data System. | 16 | 9 | 1609.01106 |
1609 | 1609.08027_arXiv.txt | { The detection of ultrahigh-energy (UHE) neutrino sources would contribute significantly to solving the decades-old mystery of the origin of the highest-energy cosmic rays. We investigate the ability of a future UHE neutrino detector to identify the brightest neutrino point sources, by exploring the parameter space of the total number of observed events and the angular resolution of the detector. The favored parameter region can be translated to requirements for the effective area, sky coverage and angular resolution of future detectors, for a given source number density and evolution history. Moreover, by studying the typical distance to sources that are expected to emit more than one event for a given diffuse neutrino flux, we find that a significant fraction of the identifiable UHE neutrino sources may be located in the nearby Universe if the source number density is above $\sim10^{-6}\,\rm Mpc^{-3}$. If sources are powerful and rare enough, as predicted in blazar scenarios, they can first be detected at distant locations. Our result also suggests that if UHE cosmic-ray accelerators are neither beamed nor transients, it will be possible to associate the detected UHE neutrino sources with nearby UHE cosmic-ray and gamma-ray sources, and that they may also be observed using other messengers, including ones with limited horizons such as TeV gamma rays, UHE gamma rays and cosmic rays. We find that for a $\gtrsim5\sigma$ detection of UHE neutrino sources with a uniform density, $n_s\sim{10}^{-7}-{10}^{-5}~{\rm Mpc}^{-3}$, at least $\sim100-1000$ events and sub-degree angular resolution are needed, and the results depend on the source evolution model. } \begin{document} | The first detections of TeV - PeV ($10^{12}-10^{15}\,$eV) neutrinos by the IceCube collaboration \citep{Aartsen:2013bka, Aartsen:2013jdh, Aartsen:2014gkd, 2015PhRvD..91b2001A,Aartsen:2015ita,Aartsen:2015rwa} opened up the era of high-energy neutrino astronomy (see \cite{Waxman:2013zda,Meszaros:2014tta,Murase:2014tsa,Ahlers15} for reviews). The energies above a few PeV are still uncharted territory, but the existence of EeV ($10^{18}\,$eV) neutrinos is guaranteed because they will be produced by the interaction of ultrahigh-energy cosmic rays (UHECR, charged particles with energies greater than $10^{18.5}$ eV) directly in the source environment, or during their propagation in the intergalactic medium. While the flux level of neutrinos produced at the source depends on the data and modeling, for the neutrinos generated during the propagation of UHECRs in the intergalactic medium (the so-called {\it cosmogenic} neutrinos) predictions are less uncertain although the chemical composition has a significant effect on the expected neutrino flux \cite{KAO10}. UHECRs were first detected decades ago, but their sources remain unknown. One reason is that the trajectories of these charged particles are bent by extragalactic and Galactic magnetic fields, and are thus difficult to trace back. Neutrinos, in contrast, propagate over cosmological distances under the influence of only gravity. Neutrinos produced by interactions of ions with matter or radiation are expected to have energies $\sim 3-5$\% of the original hadron energy \citep{2009herb.book.....D}, which means that EeV neutrinos will be unambiguous probes of the accelerators of UHECRs \footnote{Larger statistics at TeV-PeV energies could lead to a sooner discovery of point sources in this energy range, which however, are not guaranteed to be UHECR sources.}. Nonetheless, detecting an ultrahigh-energy (UHE) neutrino does not guarantee the identification of an UHECR ``accelerator''. Since there is no horizon for high-energy neutrino propagation (unlike for gamma-rays or cosmic-rays), it is likely that a range of astrophysical objects will lie within the solid angle associated with a detected neutrino, making it difficult to determine the actual source. Source catalogs obtained by electromagnetic observations could be helpful, but high-energy catalogs are usually incomplete due to flux limitations. Besides, a source association may be difficult if sources of high-energy neutrinos are different from that observed in other wavebands, or if the events come from multiple types of sources. A more robust way of finding the UHE sources would be to identify the bright sources on top of a diffuse background. As in traditional astronomy, implications of a point-source search have been studied in the literature of neutrino astronomy \citep{Lipari:2008zf,Silvestri:2009xb,Murase:2012df,2014PhRvD..90d3005A, 2016arXiv160701601M}. Importantly, in the UHE range, the atmospheric neutrino background can be safely neglected. Thus, the success of such a search depends on the angular resolution and sensitivity of the detectors, but also on the source population luminosity density (or energy budget) and emissivity history. Many existing and projected experiments have been proposed to detect EeV neutrinos, including the Antarctic Impulsive Transient Antenna (ANITA\citep{2010arXiv1003.2961T}) , the ANTARES telescope \citep{2011NIMPA.656...11A}, the Askaryan Radio Array (ARA \citep{2012APh....35..457A}), the Antarctic Ross Ice-Shelf ANtenna Neutrino Array (ARIANNA \cite{Barwick:2006tg}), the Cubic Kilometre Neutrino Telescope (KM3NeT \citep{2016JPhG...43h4001A}), the ExaVolt Antenna (EVA \cite{2011APh....35..242G}), the Giant Radio Array for Neutrino Detection (GRAND \cite{Martineau-Huynh:2015hae}), IceCube \cite{IceCube10}, IceCube-Gen2 \citep{2014arXiv1412.5106I}, the JEM-EUSO Mission \citep{2013arXiv1307.7071T}, Cherenkov Telescope Array (CTA \cite{Gora:2016mmy}), CHerenkov from Astrophysical Neutrinos Telescope (CHANT \cite{Neronov:2016zou}), and Neutrino Telescope Array (NTA \cite{2014arXiv1408.6244S}). These experiments employ very different techniques, and have a wide range of effective areas, angular resolutions, and detection efficiencies \citep{2016arXiv160708232C, 2016arXiv160708781S}. In this work, we investigate the requirements for a future EeV neutrino detector to identify an UHE neutrino source. In Section~\ref{sec:sigma} we explore the prospects for pinpointing sources using detectors with a range of angular resolutions and effective areas. In Section~\ref{sec:UHECR} we examine the possibility of association between UHE neutrinos and UHECR sources. We conclude in Section~\ref{sec:discussion} with a discussion of the impact of our results in light of existing experiments. Finally, we point out that our work can also apply to neutrinos in the energy range of 100 TeV - EeV where the effect of atmospheric neutrinos is negligible. | \label{sec:discussion} We have investigated the requirements for a future EeV neutrino detector to identify a neutrino point source. We find that for non-evolving sources with $n_s\sim{10}^{-7}-{10}^{-5}~{\rm Mpc}^{-3}$, $\gtrsim100-1000$ events and sub-degree angular resolution are needed for a $\gtrsim5\sigma$ detection of UHE neutrino sources. This detection would also give relevant clues to the origins of UHECRs. The results are sensitive to the redshift evolution model, and the similar numbers are obtained for the SFR evolution with $n_s\sim{10}^{-9}-{10}^{-7}~{\rm Mpc}^{-3}$. By examining the typical distance to sources that can be identified by a UHE neutrino detector, we show that for source population with a number density above $\sim 10^{-6}\,\rm Mpc^{-3}$, a significant fraction of the brightest sources may be in the nearby Universe. Therefore if the sources are neither beamed nor transient, it would be possible to associate the detected sources with nearby objects observed using other messengers, including messengers with limited horizons. On the other hand, if sources are rare and powerful as predicted in blazar scenarios, they can be first found at distant locations. Note that UHECRs above the GZK energy should be suppressed for distant sources, but $\sim{10}^{19}$~eV cosmic rays may reach the Earth and their powerful sources may be relevant for anisotropy searches in the UHE range. So far we have only considered steady sources which therefore do not evolve over time. UHECRs from a transient event are expected to arrive at Earth with a spread of arrival times, which we designate by $\delta t$, because UHECR paths are deflected by magnetic fields between the source and the Earth. Therefore, for transient sources, the source rate per volume $\rho_{\rm s}$ can be converted to an apparent number density $n_{\rm s}$ via the UHECR arrival time spread: $\rho_{\rm s}\sim n_{\rm s}/\delta t$ \citep{Murase_Takami09,Takami:2011nn}. The time spread is $\delta t\sim 10^4\,{\rm yrs} \,(D/100 {\rm \,Mpc})\,(\delta\theta/1^\circ)^2$, where $\delta\theta$ is the angular deflection experienced by the particle during the propagation (e.g., \cite{KL08b}). The apparent number density may still govern the potential of pinpointing the sources of cosmogenic neutrinos, while the luminosity of EeV neutrinos from the sources can be much higher because the duration of UHE neutrino emission is short. Example transient sources of EeV neutrinos are GRB afterglows \cite{Waxman:1999ai,Dermer:2000yd,Murase:2007yt,Razzaque:2013dsa}, young magnetars and pulsars \cite{Murase09,FKMO14_letter}. Searches for transient UHE neutrino sources are relevant especially if the composition is dominated by protons and light nuclei. Unlike astrophysical UHE neutrinos, cosmogenic neutrinos are expected to suffer some, typically degree-level, angular displacement with respect to their source directions, due to the deflections suffered by the primary UHECRs that produce them. Therefore it may be difficult to pinpoint the actual source location even with a perfect detector. On the other hand, our work suggests that once the detected number of cosmogenic neutrinos reaches a few hundred, crucial information can be gained regarding the characteristics of the sources of UHECRs. Importantly, any small scale auto-correlations in the arrival directions of cosmogenic neutrinos, or the absence thereof, would constrain the number density and source evolution of UHECR sources, in a measurement that is complementary to the equivalent study of UHECR arrival directions, and PeV neutrinos, as presented in e.g., Refs.~\cite{1997ApJ...483....1W, 2000PhRvL..85.1154D, AugerBound, 2016arXiv160701601M}. In this work we only use the spatial information of neutrino events to look for sources. Therefore the significance of the searches does not depend on the energy spectrum of the events. However, a different energy spectrum could impact the implications through two ways: 1) by changing the expected total number of events (see equation~\ref{eq:Nev}), and 2) by changing the contribution of distant sources in a redshifted energy band. The influence of different spectral templates is demonstrated in Ref.~\cite{2016arXiv160701601M}. The current work also applies to astrophysical neutrinos at lower energies including the energy range covered by IceCube. Above $\gtrsim50-100$~TeV, coincidental pairs of muon neutrinos from the atmospheric background are negligible in most part of the sky \citep{2016arXiv160701601M}. For lower-energy tracks or shower events with a poorer angular resolution, a marginalization over energies or an energy-dependent likelihood~\citep{Fang:2016hyv} is necessary to avoid the confusion from the atmospheric events. Our 90\% confidential level ($\sim1.6\sigma$ by the conventional conversion with a Gaussian) results shown in Figures~\ref{fig:sigma} and \ref{fig:sigma7} are consistent with the number density constraints presented by Ref.~\citep{2016arXiv160701601M} (note that the number of $\gtrsim60$~TeV neutrinos observed in the six-year observation is $\sim100$ in the half sky.). Our results in Figures~\ref{fig:sigma},\ref{fig:sigma7}, and \ref{fig:fracMultiplet} also predict the point-source search potential of future statistics from IceCube and ANTARES, as well as next-generation detectors like IceCube-Gen2 \citep{2014arXiv1412.5106I} and KM3NET \citep{2016JPhG...43h4001A}. | 16 | 9 | 1609.08027 |
1609 | 1609.03932_arXiv.txt | We apply a novel spectral graph technique, that of \emph{locally-biased semi-supervised eigenvectors}, to study the diversity of galaxies. This technique permits us to characterize empirically the natural variations in observed spectra data, and we illustrate how this approach can be used in an exploratory manner to highlight both large-scale global as well as small-scale local structure in Sloan Digital Sky Survey (SDSS) data. In particular, we use this method in a way that simultaneously takes into account the measurements of spectral lines as well as the continuum shape. Unlike Principal Component Analysis, this method does not assume that the Euclidean distance between galaxy spectra is a good global measure of similarity between all spectra, but instead it only assumes that local difference information between similar spectra is reliable. Moreover, unlike other nonlinear dimensionality methods, this method can be used to characterize very finely both small-scale local as well as large-scale global properties of realistic noisy data. The power of the method is demonstrated on the SDSS Main Galaxy Sample by illustrating that the derived embeddings of spectra carry an unprecedented amount of information. By using a straightforward \emph{global} or \emph{unsupervised} variant of our method, we observe that the main features correlate strongly with star formation rate and that they clearly separate active galactic nuclei. In addition, computed parameters of the method can be used to describe line strengths and their interdependencies. By using a \emph{locally-biased} or \emph{semi-supervised} variant of our method, we are able to focus on typical variations around specific objects of astronomical interest. We present several examples illustrating that this approach can enable new discoveries in the data as well as a detailed understanding of very fine local structure that would otherwise be overwhelmed by large-scale noise and global trends in the~data. | \label{sec:intro} The physical properties of the Universe and the internal mechanisms of galaxies are ultimately intertwined in astronomical observations. Characterizing the diversity of galaxies is vital not only for understanding their evolution but also to unravel the nature of dark energy in the context of our cosmological models. While today's large-scale spectroscopic surveys provide a plethora of data, novel data analysis methods are needed to help extract astronomical insight from these data. Current data analysis approaches in this area generally fall into one of two categories. In the first category, the observed spectra are fitted by semi-analytic models, e.g., \cite{bc03}, to infer model-based parameters. These parameters in turn provide a model-dependent physical coordinate system with absolute scales such as age or metallicity. Challenges for these methods typically include systematic biases due to imperfect models as well as correlated parameters. In the second category, one adopts a more empirical approach, where galaxies are analyzed in relation to other galaxies based on the original measurements, i.e., based on the observed spectra. A major challenge for these more empirical methods is the conceptual problem of how best to compare empirical spectra, e.g., which features of a spectrum are most important for identifying similarities between two spectra. The approach we describe in this paper falls into this second category, and it aims to address the fundamental issue of measuring similarity between galaxy spectra, both with respect to the large-scale global properties of the empirical data, as well as with respect to finer-scale local structure that might be overwhelmed by global data analysis tools that focus on global properties of the~data. A canonical example of a global data analysis tool that adopts this more empirical approach is Principal Component Analysis (PCA), which is widely-used to find the globally-dominant linear trends in the data. PCA was first applied to galaxies by \cite{connolly1995spectral}, who found that a significant fraction of the variance in the spectra can be captured by only three components. In other words, the analyzed spectra could be well approximated by a linear combination of three \emph{eigenspectra}. The coefficients serve as summaries of the high-dimensional spectra, and in this coordinate system galaxies could be meaningfully compared to one another. PCA has been used in many research areas, including photometric redshift estimation \cite{connolly1999orthogonal,budavari2000creating}, sky subtraction \cite{wild05}, as well as classification of galaxies and quasars \cite{francis92,connolly99, yip04a, yip04b}. The Sloan Digital Sky Survey (SDSS) \cite{york00} has adopted the method in its data reduction pipeline, and it automatically derives the first five eigencoefficients (called \texttt{eCoeff\_0} -- \texttt{eCoeff\_4}). These are considered the state of the art in describing the continuum shape of the spectra. Figure~\ref{fig:intro-embeddings-pca} shows the mixing angles $\theta$ and $\phi$ of the three leading eigencoefficients for the Main Galaxy Sample (MGS) in SDSS Data Release 7 \cite{abazajian2009seventh}. These coordinates are defined as in~\cite{yip2004distributions} by \begin{equation} \label{eq:theta-phi-defn} \phi = \tan^{-1}\left( \frac{\texttt{eCoeff\_1}}{\texttt{eCoeff\_0}} \right), \quad \theta = \cos^{-1}(\texttt{eCoeff\_2}). \end{equation} In the embedding illustrated in Figure~\ref{fig:intro-embeddings-pca}, every point is a galaxy, and nearby points (i.e., galaxies or points that are near each other in the two-dimensional representation) have similar observations, by construction of the empirical coordinate system. In particular, ``red and dead'' galaxies appear at the top of the plot, while star-forming blue ones are on the lower right. \begin{figure} \begin{center} \includegraphics[width=.40\textwidth]{theta_phi_alpha.jpg} \end{center} \caption{Embedding of the Main Galaxy Sample of the SDSS Data Release 7 on the mixing angles $\theta$ and $\phi$ of the first three eigen-coefficients. } \label{fig:intro-embeddings-pca} \end{figure} There are, however, shortcomings of the simplified view of the data provided by PCA. First of all, a significant fraction of the galaxies are actually removed from or scattered out of the plot as a part of the usual data analysis pipeline used to generate the plot; we only really see the core of the distribution in a figure such as this one. In addition, the lack of structure in this visualization is surprising, especially considering the large amount of high-quality data and the wide range of galaxy types in the data. Finally, the interpretation of the axes is difficult, as they are linear combinations of actual data elements and not actual data elements. Extensions and variants of PCA have been proposed to overcome these challenges, including non-negative matrix factorization \cite{nmf}, the use of robust statistics \cite{budavari2009reliable}, and CX/CUR matrix decompositions~\cite{CUR_PNAS,yip2014objective}. While these methods have alleviated some of the issues associated with PCA, the fundamental limitation of its assumed linear model remained. Perhaps the biggest conceptual change in the area was introduced by \cite{vanderplas2009reducing}, who applied the Locally Linear Embedding (LLE) method of \cite{roweis2000nonlinear}. This more sophisticated empirical approach attempts to identify and exploit local structure in the data, and thus it broke away from the straightforward global linear model underlying PCA. While there are other related nonlinear approaches~\cite{tenenbaum2000global,belkin2003laplacian}, LLE in particular attempts to provide an angle-preserving mapping that assigns coordinates to galaxies such that each galaxy is approximately a linear combination of its nearest neighbors, with the same weights as in the observed space. The power and practical usefulness of LLE (as well as other related nonlinear methods~\cite{tenenbaum2000global,belkin2003laplacian,coifman2006diffusion}), however, is known to be severely diminished in many practical situations. The reasons for this are many: due to the difficulty of these methods in dealing with non-uniform point densities; since the global objective function used to enforce angle preservation or other neighborhood information can damage small-scale or local structure in the data; since these methods are quite sensitive to realistic noise in the data; and since these methods are very sensitive to the ``details'' of constructing the nearest neighbor graph, e.g., to the functional form of the nearby distance and the choice of parameters used to define nearness. (This is in spite of a large body of theory stating that in idealized situations these details do not matter.) In addition to exploiting the strong algorithmic and statistical theory underlying our main method~\cite{HM14_JRNL,mahoney2012local}, dealing appropriately with these and other related practical graph construction issues will be central to our approach, and thus we postpone further discussion of it until Sections~\ref{sec:methods} and~\ref{sec:knobs}. A third empirical approach that is worth mentioning is based completely on line measurements. Recall that the high resolution in wavelength often allows the identification and measurement of different spectral lines, and that it is common to plot spectra in terms of carefully-chosen line ratios. That is, while not usually described as an embedding method, the typical use of line measurements often involves embedding or mapping the data to a low-dimensional space. (The state-of-the-art in this area actually uses PCA as a pre-processing step to subtract the continuum, in order to measure better the line strength relative to this baseline~\cite{tremonti}.) In particular, the BPT diagrams~\cite{baldwin1981classification} plot different line ratios on a logarithmic scale, enabling, e.g., the classification of galaxies~\cite{brinchmann04,kewley2006}. For example, Figure~\ref{fig:intro-embeddings-bpt} shows several BPT diagrams of the SDSS MGS. In Figure~\ref{fig:intro-embeddings-bpt1}, the characteristic V-shape of the embedding on the ratios $\textrm{N}_\textrm{II}/\textrm{H}_\alpha$ vs.~$\textrm{O}_\textrm{III}/\textrm{H}_\beta$ is clearly visible, despite significant scatter that is partly due to noisy measurements of the individual lines. Little to no structure is evident in Figure~\ref{fig:intro-embeddings-bpt2} and~\ref{fig:intro-embeddings-bpt3}, whose $x$-axes plot different line strengths. The insight conveyed by the BPT plots can be considered complementary to that of the PCA results, which is primary based on the continuum shape. Again, though, the lack of fine-scale structure in these visualizations is somewhat surprising, given the quality and diversity of the data. \begin{figure} \begin{center} \subfigure[]{ \includegraphics[width=.30\textwidth]{{bpt_NiiHa_OiiiHb_alpha}.jpg} \label{fig:intro-embeddings-bpt1} } % \subfigure[]{ \includegraphics[width=.30\textwidth]{{bpt_SiiHa_OiiiHb_alpha}.jpg} \label{fig:intro-embeddings-bpt2} } % \subfigure[]{ \includegraphics[width=.30\textwidth]{{bpt_OiiHa_OiiiHb_alpha}.jpg} \label{fig:intro-embeddings-bpt3} } \end{center} \caption{Embeddings based on several pairs of BPT line-ratio diagrams of the SDSS DR7 MGS. } \label{fig:intro-embeddings-bpt} \end{figure} In this paper, we present a novel approach to studying galaxies that combines elements of several aforementioned techniques but that moves away from the limiting assumptions in their underlying mathematical models. Our method, which is an extension of \emph{semi-supervised eigenvectors}~\cite{HM14_JRNL} from \emph{locally-biased machine learning}~\cite{mahoney2012local}, uses ideas from spectral graph theory; and we study the properties of the SDSS Main Galaxy Sample data, using both a traditional global variant as well as a more recently-developed local variant of these methods. The global variant reduces to a version of Laplacian eigenmaps~\cite{belkin2003laplacian} and related methods such as diffusion maps~\cite{coifman2006diffusion}, while the local variant exploits recent work on local spectral graph partitioning to engineer locality into the basic global method, while at the same time preserving the strong algorithmic and statistical properties of the global variant~\cite{HM14_JRNL,mahoney2012local}. Among other things, the method efficiently handles the continuum shape and the spectral lines simultaneously, and the method can be used to explore the data in a qualitatively more refined manner in order to obtain better insights for galaxy studies. Since the method is unfamiliar in this application domain, in Section~\ref{sec:methods}, we briefly review the approach; and in Section~\ref{sec:knobs}, we illustrate how several of the key ``knobs'' of the method behave on these data, as a function of data modeling design decisions. Then, in Section~\ref{sec:global}, we apply the global variant of the method to study the large-structure of the data; and in Section~\ref{sec:local}, we use the local variant of the method to study much finer-scale structure in the data. Finally, in Section~\ref{sec:conclusion}, we present a brief conclusion. | \label{sec:conclusion} We have presented a novel technique based on locally-biased semi-supervised eigenvectors to gain insight into the similarity of galaxies, and we have demonstrated the recovered nonlinear structure on the Main Galaxy Sample of the Sloan Digital Sky Survey. By constructing low-dimensional embeddings which respect local connectivity, we are able to visualize a range of astronomical phenomenon, e.g., the process of stellar evolution from hot, blue galaxies to cool, red ones. Unlike previously-used methods such as PCA or other recently-popular nonlinear dimensionality reduction methods, our method can focus on and highlight in a refined way local properties and/or global properties of the data. Depending on the choice of knobs of the method, the embedding maps we create contain all galaxies, and they can be used to identify either large-scale global structure in the data or small-scale local structure in the data. The main parameters that we empirically derive clearly correspond to changes in the continuum shape and the strengths of spectral lines, and these help disambiguate traditional BPT plots and isolated AGNs. We observe that there are no disjoint groups of galaxies to indicate natural classes, but instead there are smooth continuous transitions between classes; and we observe that in many cases outliers (sometimes artifacts) are quite different than any other spectra in the data set. Moreover, the locally-biased versions of these embeddings (which use semi-supervised seed labels to construct semi-supervised eigenvectors) demonstrate that the method can be used to enhance the embeddings such that we can better focus on galaxies of interest, e.g., rare galaxy types that can be overwhelmed by global~methods. To summarize, our method of locally-biased semi-supervised eigenvectors may be viewed as a new type of \emph{computational microscope} for astronomical data. It can not only reproduce known properties of the data and identify outliers and artifacts, but it can also be used to enhance subtle trends around selected galaxies to facilitate new discoveries. Obvious future work should focus on applying this new methodology to improve galaxy classification, to derive continuous spectral models, e.g., for photometric redshift estimators, and to understand better galaxy evolution with the help of stellar population synthesis models. | 16 | 9 | 1609.03932 |
1609 | 1609.06510_arXiv.txt | {% The standard method to estimate the mass of a cosmic ray is the measurement of the atmospheric depth of the shower maximum ($X_\text{max}$). This depth is strongly correlated with the mass of the primary because it depends on the interaction cross section of the primary with the constituents of the atmosphere. Measuring the electric field, emitted by the secondary particles of an extensive air shower (EAS), with the Auger Engineering Radio Array (AERA) in the 30-80 MHz band allows the determination of the depth of shower maximum on the basis of the good understanding of the radio emission mechanisms. The duty cycle of radio detectors is close to 100\%, making possible the statistical determination of the cosmic-ray mass composition through the study of a large number of cosmic rays above 10$^{17}$ eV. In this contribution, $X_\text{max}$ reconstruction methods based on the study of the radio signal with AERA are detailed. } | \label{sec-4} The distribution of the deviation of the reconstructed $X_\text{max}$ from the FD measurements for the set of hybrid showers ($\Delta X_\text{max} $) are fitted by a gaussian function and the mean difference and standard deviation at 1 $\sigma$ confidence level are summarized in Table \ref{tab-1} for each method. The offsets could be related to the underlying simulation codes. The reconstruction of $X_\text{max}$ from the radio signal will be a decisive asset for the estimation of the mass composition of cosmic rays due to the high duty cycle of the radio stations. The next steps are the determination of the composition from radio without any bias and the combination of the strength of the radio methods. To be efficient, the reconstruction needs, to be precise and with a fast calculation time. The validity of the methods can only be proven through the comparison with the FD measurements. The influence of the air density and refractivity at the time of the measurements on the deviations are under study. The $X_\text{max}$ resolution from radio methods is currently close to 40 g/cm$^2$. \begin{table}[h] \centering \caption{Summary of the requirements and results of the radio methods} \label{tab-1} % \begin{tabular}{|c|c|c|c|c|} \hline method & A & B & C & D \\\hline requirements & RD direction & RD direction & SD direction & SD direction \\ & & & SD energy & SD energy \\\hline calculation time (one event) & 8 hours & - - - - - - - - - - & ~1 week & ~1 hour\\\ simulated antennas& 168 & 0 & 160 & 60 \\\ number of showers& 40 p + 10 Fe & 0 & 20 p + 10 Fe & 30 p + 30 Fe \\\hline mean($X_\text{max,RD} - X_\text{max,FD}$) (g/cm$^2$)& $43.6 $ & $-12 $ & $4.4 $ & $13.47$ \\\hline $\sigma$ ($X_\text{max,RD} - X_\text{max,FD}$) (g/cm$^2$) &43.4 & 59.3 & 47 & 41.03 \\\hline \end{tabular} \end{table} \vspace{-0.cm} | 16 | 9 | 1609.06510 |
|
1609 | 1609.02635.txt | We use data from the Sydney-AAO Multi-Object Integral Field Spectrograph (SAMI) Galaxy Survey and the Galaxy And Mass Assembly (GAMA) survey to investigate the spatially-resolved signatures of the environmental quenching of star formation in galaxies. Using dust-corrected measurements of the distribution of H$\alpha$ emission we measure the radial profiles of star formation in a sample of \SFNUM star-forming galaxies covering three orders of magnitude in stellar mass ($\rm{M}_{*}$; $10^{8.1}$-$10^{10.95} \, \mathrm{M}_{\odot}$) and in $5^{th}$ nearest neighbour local environment density ($\Sigma_{5}$; $10^{-1.3}$- $10^{2.1} \, \mathrm{Mpc}^{-2}$). We show that star formation rate gradients in galaxies are steeper in dense ($\log_{10}(\Sigma_{5}/\mathrm{Mpc^{2}})>0.5$) environments by $0.58 \pm 0.29 \, \mathrm{dex} \, \mathrm{r_{e}}^{-1}$ in galaxies with stellar masses in the range $10^{10} < \mathrm{M_{*}}/\mathrm{M_{\odot}} < 10^{11}$ and that this steepening is accompanied by a reduction in the integrated star formation rate. However, for any given stellar mass or environment density the star-formation morphology of galaxies shows large scatter. We also measure the degree to which the star formation is centrally concentrated using the unitless scale-radius ratio ($r_{50,H\alpha}/r_{50,cont}$), which compares the extent of ongoing star formation to previous star formation. With this metric we find that the fraction of galaxies with centrally concentrated star formation increases with environment density, from $\sim 5 \pm 4 \%$ in low-density environments ($\log_{10}(\Sigma_{5}/\mathrm{Mpc^{2}})<0.0$) to $30\pm 15 \%$ in the highest density environments ($\log_{10}(\Sigma_{5}/\mathrm{Mpc^{2}})>1.0$). These lines of evidence strongly suggest that with increasing local environment density the star formation in galaxies is suppressed, and that this starts in their outskirts such that quenching occurs in an outside-in fashion in dense environments and is not instantaneous. | The process of star formation is critical to the evolution of galaxies. The rate of star formation past and present has a significant effect on the optical colours and morphology of a given galaxy \citep{DG92}. It has become apparent that the environment within which a galaxy is situated plays an important role in that galaxy's development \citep[e.g.][]{HubbleAndHumason31,Oemler74,Dressler80}. The presence of a relationship between galaxy environment and star-forming properties suggests that the transformation from star-forming to quiescent, a process called quenching, could be affected by the environment. A number of mechanisms have been proposed that could cause quenching to occur. These mechanisms generally involve the removal of the gas supply that fuels star formation. Ram pressure stripping has been identified as a potential method for reducing the amount of available gas in a galaxy disc \citep{Gunn72} and in the surrounding halo \citep{McCarthy08}. The role of ram pressure stripping in quenching star formation in cluster galaxies has been well established. The best evidence for ram pressure stripping acting to quench star formation within clusters comes from unresolved measurements of neutral hydrogen in cluster galaxies \citep[e.g.][]{Giovanelli&Haynes85,Solanes2001,Cortese2011}. The signatures of ram pressure stripping include the confinement of star formation to the central regions of the galaxy \citep{Koopmann04b,Cortese2012}, and the presence of a tail visible in H$\alpha$, neutral hydrogen, or both \citep{Balsara1994}. \cite{Boselli2006} utilised multiwavelength imaging to constrain the stellar population gradients within a galaxy to confirm that ram pressure stripping acts on late type galaxies in clusters. This process is also capable of acting in compact groups \citep{Rasmussen08} as well as in less concentrated environments \citep{Bekki2009,Merluzzi13}. \cite{Nichols11} successfully explain the gas fractions of dwarfs in the Milky Way $+$ M31 system by modelling gas removal by ram pressure from these objects analytically. However, \cite{Rasmussen08} find that ram pressure stripping alone cannot explain the gas deficiencies in more massive group galaxies, with masses of approximately $4\times 10^{10} \, \rm{M}_{\odot}$ or greater. It has been pointed out that other transport processes could be responsible for the removal of gas from galaxy discs in dense environments. In particular, turbulent viscous and inviscid stripping have been suggested to be significant mechanisms for gas removal in dense environments \citep{Nulsen1982,Roediger13}. When the velocity of the galaxy through the intergalactic medium is subsonic, viscous and turbulent mixing between the two media can become important in liberating gas from galaxies. The timescale for removing the gas from galaxies undergoing viscous stripping is shorter than for ram-pressure stripping \citep{Nulsen1982}. Modelling has shown that the morphological signatures of these mechanisms may be similar to those of ram-pressure stripping \citep[e.g.][]{BoselliAndGavazzi2006,Roediger14}. The timescale on which a disc galaxy would deplete its gas through ordinary star formation is on the order of a few Gyr \citep{Miller79, Tinsley80}. In order to explain the existence of gas-rich disc galaxies in the Universe today, it is necessary to suppose that the gas within their discs has been replenished by infall from the intergalactic medium \citep[e.g.][]{Larson72}. If this supply is cut off, then quenching will occur in a process called strangulation. This is most likely to transpire when the gas envelope surrounding spiral galaxies is swept away as the galaxy and its halo enter a cluster or group and fall through the denser intergalactic medium of these environments \citep{Larson80}. Strangulation is likely to quench star formation over a period of several Gyrs \citep{McCarthy08} once the gas in the disc ceases to be replenished. \cite{Peng15} estimate that strangulation is responsible for quenching approximately 50\% of the passive galaxy population today, though they do not investigate the dependence of this fraction on environment density. Strangulation is predicted to occur in galaxy groups by \cite{Kawata2008}. Some simulations \citep[e.g.][]{Bekki02} suggest that the formation of anaemic spirals, systems with uniformly suppressed star formation \citep{vandenbergh91,Elmegreen02}, can be explained by strangulation. This process is not expected to produce the same spatial distribution of star formation as other processes such as ram-pressure stripping while quenching is taking place. %This process is not expected to change the distribution of star formation in a galaxy as gas is not stripped or ejected from the disc in this scenario. Galaxies in high-density environments such as clusters or galaxy groups may also experience tidal interactions with their nearest neighbours, or the group or cluster gravitational potential. These tidal interactions often have the effect of driving gas towards the centre of the galaxy and triggering circumnuclear starbursts \citep[e.g.][]{Heckman90}. These starburst episodes can deplete the gas reservoir of the galaxy, causing it to become quenched. Simulations \citep[e.g.][]{Hernquist89,Moreno2015} have shown that gravitational instabilities driven by tidal interactions will drive a large fraction of the gas in a galaxy towards the centre and enhance star formation on short timescales. This will have the effect of producing galaxies with centrally-concentrated star formation shortly after an interaction. \cite{Faber07} argued that gas-rich major mergers between blue-sequence galaxies can induce a period of starburst, which rapidly depletes the interstellar gas within these systems. Following this merger, the remnant moves to the red sequence. However, \cite{Blanton06} notes that between $z=1$ and $z=0$ the number of blue-sequence galaxies is reduced by less than 10\%. This implies that major-merger-driven quenching cannot have been the dominant mechanism for decreasing star formation in the second half of the Universe's history. Recently, large-scale spectroscopic and photometric surveys have been able to make significant progress in understanding the relationship between galaxy environment and star formation. Modern multi-object spectroscopic surveys such as the 2 Degree Field Galaxy Redshift Survey \citep[2dFGRS;][]{Colless01}, the Sloan Digital Sky Survey \citep[SDSS;][]{York2000} and the Galaxy and Mass Assembly survey \citep[GAMA;][]{Driver11,Hopkins13} have allowed the determination of star formation rates in several hundred thousand galaxies. However, the gain in sample size afforded by these single-fibre spectroscopic surveys is offset by the fact that the star formation in each galaxy is reduced to an estimate of the integrated total, which can be affected by aperture bias. Consequently, this observational technique has been unable to identify galaxies that are in the process of quenching, and arguments involving the timing and frequency of quenching must be invoked to determine which mechanisms are producing the observed trends. A spectroscopic study of $521$ clusters in the SDSS by \cite{vonderLinden2010} indicated that the star formation rates of galaxies decline slowly during the infall into a cluster, with the most rapid quenching only occurring at the centres of clusters. The inferred quenching timescales were therefore long, roughly a few Gyrs, which is comparable to the cluster crossing time. This conclusion is in agreement with \cite{Lewis02} and \cite{Gomez03} who note the existence of galaxies with low specific star formation rates at large distances from the centres of clusters and conclude that rapid environmental quenching alone cannot explain the population of galaxies seen in the local Universe. This picture appears to be inconsistent with the results of other work. Using the optical colours of galaxies, \cite{Balogh04} argue that once galaxy luminosity is controlled for, the dominant environmental trend is the changing ratio of red to blue galaxies, with very little change in the average colour of the blue galaxies. Similarly a spectroscopic investigation by \cite{Wijesinghe12} using GAMA data showed no correlation between the star formation rates of star-forming galaxies and the local environment density. In this sample, the trend was visible only when the passive galaxy population was included in the analysis, suggesting that it is the changing fraction of passive galaxies that is responsible for the observed environmental trends, and that quenching must therefore be either a rapid process or no longer proceeding in the local Universe. The ambiguity between fast and slow-mode quenching may in part arise from the different definitions of passive and star forming used by various teams \citep[see e.g.][]{Taylor15}, as well as the different methods of quantifying environment density employed. Moreover, these large-scale surveys are not able to investigate the spatial distribution of star formation in galaxies that are in the process of being quenched. The spatial properties of star formation in dense environments have been studied using narrow-band imaging of the H$\alpha$ distribution within galaxies \citep[e.g.][]{Koopmann04a, Koopmann04b,Gavazzi2006,Bretherton2013}. In the Virgo cluster, \cite{Koopmann04b} note that approximately half of $84$ observed spiral galaxies have spatially-truncated star formation while less than $10\%$ are anaemic (have globally reduced star formation). \cite{Welikala2008} and \cite{Welikala2009}, using spatially-resolved photometry from SDSS, suggested that the suppression of star formation in dense environments occurs predominantly in the centres of galaxies. \cite{Welikala2009} observed that the integrated star formation rates in star-forming galaxies decline with density, implying that the observed environmental trends cannot be completely explained by the morphology-density relation. In lower density group and field environments, \cite{Brough13} used optical integral field spectroscopy to examine the radial distribution of star formation for $18$ $10^{10}\, \mathrm{M}_{\odot}$ galaxies and found no correlation between the star formation rate gradient and the local density. Several studies have also suggested that the stellar mass, bulge mass or other internal properties of a galaxy have a greater influence on whether it is quenched than does the environment \citep[e.g.][among others]{Peng2010,Bluck14,Pan15}. \cite{Peng2010} argued that quenching can be explained by two separable processes that depend on mass and environment. Mass quenching could be achieved by several mechanisms including AGN feedback, which either removes gas from the galaxy disc directly or prevents it accreting from the halo, or by feedback that is related to star-formation such as supernova winds. While it is likely that all quenching mechanisms operate to some extent, it remains uncertain how dominant each mode is at a given environment density and galaxy mass. In this paper we investigate the radial distribution of star formation in galaxies observed as part of the Sydney-AAO Multi-object Integral Field Spectrograph (SAMI) Galaxy Survey \citep{Croom12, Allen15, Bryant15, Sharp15}. The application of spatially-resolved spectroscopy to this problem represents an important step towards a better understanding of the quenching processes in galaxies. With this technique applied to a large sample of galaxies, the spatial distribution of star formation in galaxies can be resolved and the direct results of the various quenching mechanisms can be observed. The broad range of stellar masses and environments targeted by SAMI make it an excellent survey for studying the spatial signatures of environmental quenching processes. In Section~\ref{TargetSelection} we introduce the data used, our target selection and important ancillary data. Section~\ref{DataReduction} details the data reduction techniques employed by SAMI and the subsequent analysis of the flux-calibrated spectra. Results are presented in Section~\ref{Results} with a discussion and conclusion given in Sections~\ref{Discussion} and~\ref{Conclusion} respectively. Throughout this paper we assume a flat $\Lambda$CDM cosmology with $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{M}=0.27$ and $\Omega_{\Lambda}=0.73$ and adopt a \cite{Chabrier2003} stellar initial mass function. | We have used SAMI integral-field spectroscopy to study the spatial distributions of ongoing star formation in a sample of \SFNUM\ star-forming galaxies as a function of their local environment densities. We have shown that the derived integrated H$\alpha$ star formation rates from integral field spectroscopy will be underestimated by approximately $9\%$ if a single average dust correction is applied to the galaxy. The non-linear nature of dust extinction demands that the attenuation must be accounted for locally within different regions of a galaxy. Failure to perform a local spaxel-by-spaxel dust extinction correction results in a systematic reduction of the total measured SFR of $\sim 8 \%$ in the most star-forming galaxies. Analysis of the star formation rate surface density radial profiles of \SFNUM\ galaxies in our sample has shown a large variation in the radial distributions of star formation in our galaxies. For any given mass or environment density the normalisation of a galaxy's star formation rate radial profile can vary by more than a factor of $10$. For our star-forming sample we note that the normalisation of the star formation increases almost linearly with stellar mass in accordance with the known star-formation rate versus stellar mass relation. There is no significant relationship between the central star formation rate surface density and the local environment density, but we note a $2\sigma$ significance steeping of the H$\alpha$ radial profiles in high density environments in galaxies with masses in the range $9.92 < \log_{10}(\mathrm{M}_{*}/\mathrm{M}_{\odot}) < 10.94$. This is consistent with the relationship between the specific star formation rates and $\Sigma_{5}$ in Figure \ref{ssfr_m_env_frac}, and implies that the environmental suppression of star formation must occur first in the outskirts of galaxies When the central star-formation rate surface density in a galaxy is controlled for, we see a significant steepening of the profiles in higher density environments. For galaxies in the highest mass bin ($9.92 < \log_{10}(\mathrm{M}_{*}/\mathrm{M}_{\odot}) < 10.94$) we observed the median normalised star formation rate profiles to steepen from $-0.54 \pm 0.18 \, \mathrm{dex} \, r_{e}^{-1}$ to $-1.09 \pm 0.26 \, \mathrm{dex} \, r_{e}^{-1}$. We have also shown that $(30 \pm 15)\%$ of galaxies in high-density [$\log_{10}(\Sigma_{5}/\rm{Mpc}^{2})>0.5)$] environments exhibit a centrally concentrated star formation distribution. In low-density environments only $(5 \pm 4)\%$ show similar star formation morphologies. This is taken as evidence that the occurrence of outside-in quenching is more common in dense environments and is consistent with the findings of studies in clusters \citep[e.g.][]{Koopmann04b,Koopmann06}. This conclusion is supported by the observed radial gradients in the D$_{n}4000$ flux ratio. Galaxies which exhibit the most extremely centrally concentrated H$\alpha$ also have stronger D$_{n}4000$ towards their edges, in contrast to systems which show extended H$\alpha$ morphologies. As the D$_{n} 4000$ spectral feature is sensitive to star formation quenching on timescales of up to $\sim1$ Gyr, we estimate that this outside-in quenching must occur faster than this, though careful modelling and a more comprehensive study of the spectral features will be required to determine the timescales accurately. The observed correlations between the star formation morphology of galaxies and the local environment density are significant but weak. These trends appear to be dominated by the intrinsic variation in galaxy properties across all environments. For this reasons sample sizes of over $225$ galaxies per mass bin will be required to ensure the statistical significance of the results of future studies. In lower density environments ($\log_{10}(\Sigma_{5}/Mpc^{2}) < 0.5$), the smaller fraction of passive systems indicates that quenching must be occurring at a lower rate than in higher density regions. In these environments galaxies with the lowest specific star formation rates appear to have spatially extended H$\alpha$ morphologies which we take to be the qualitative signature of different mechanisms (such as starvation) acting to suppress the star formation in galaxies. Galaxies of higher stellar mass are more likely to show the signatures of outside-in quenching. Even when the effect of correlation between mass and environment density is controlled for, this relationship persists. We must therefore conclude one of several things: 1) that higher mass galaxies are more susceptible to environmental effects than low mass galaxies, 2) the environmental mechanisms that quench high mass and low mass galaxies are qualitatively different, or 3) that the outside-in quenching of star formation in massive galaxies proceeds at a lower rate in higher mass galaxies and occurs almost instantaneously in lower mass galaxies. We believe that the last two of these scenarios are most plausible. A future study will utilise a larger sample size to investigate in more detail the role of galaxy mass on environmental quenching, and incorporate a more detailed analysis of the galaxy group properties that drive the various quenching processes. | 16 | 9 | 1609.02635 |
1609 | 1609.08519_arXiv.txt | NGC 288 is a diffuse Galactic globular cluster, it is remarkable in that its low density results in internal accelerations being below the critical MOND $a_{0}$ acceleration throughout. This makes it an ideal testing ground for MONDian gravity, as the details of the largely unknown transition function between the Newtonian and modified regimes become unimportant. Further, exact analytical solutions exist for isothermal spherical equilibrium structures in MOND, allowing for arbitrary values of the anisotropy parameter, $\beta$. In this paper we use observations of the velocity dispersion profile of NGC 288, which is in fact isothermal, as dynamical constraints on MONDian models for this cluster, where the remaining free parameters are adjusted to fit the observed surface brightness profile. We find the optimal fit requires $\beta =0$, an isotropic solution with a total mass of $3.5 \pm 1.1 \times 10^{4} M_{\odot}$. | \label{intro} Starting with observations of $\omega$Cen in Scarpa et al. (2003), it became apparent that Galactic globular clusters have projected velocity dispersion radial profiles which do not fall monotonically with radius along Newtonian expectations for isolated systems. Rather, after an initial radial drop, projected velocity dispersion profiles, $\sigma(R)$, settle to constant asymptotic values. This was then confirmed for the cases of M15 and NGC 6171 by Scarpa et al. (2004a,b), and then extended to NGC 7099 in Scarpa et al. (2007). Since, this generic feature has been corroborated for a growing sample of Galactic GCs by various independent groups, e.g. Lane et al. (2009), Lane et al. (2011). One of the most interesting features of this asymptotic flattening in $\sigma(R)$, is that the radii at which they become flat, closely corresponds to those where the typical stellar acceleration falls below the critical MOND acceleration of $a_{0}=1.2 \times 10^{-10}m s^{-2}$ e.g. Scarpa et al. (2007), Hernandez \& Jimenez (2012). This last has been interpreted as evidence in favour of MOND scenarios (e.g. Hernandez et al. 2013), tidal disruption from the overall Galactic gravitational field (e.g. K\"{u}pper et al. 2010), or dynamical evolution processes internal to the Globular Clusters themselves (e.g. Kennedy 2014 under Newtonian gravity). { The study of globular cluster dynamics as probes of possible variations in the form of gravity and/or details of the effects of internal dynamical evolution and tidal interactions with the Galaxy has been a topic of substantial interest over the past few years. Dynamical modelling under MOND has been preformed by Sollima \& Nipoti (2010), Sanders (2012) and Wu \& Kroupa 2013 under MOND, finding results in support of a MONDian interpretation, while Lane et al. (2010) find Newtonian models yield accurate descriptions. The available data samples are also growing, e.g. Kimmig et al (2015) and Lardo et al. (2015) preform kiematical samplings of growing sets of globular clusters, with the recent study by Baldwin et al. (2016) giving for the first time, proper motion kinematic profiles for a number of Galactic globular clusters.} In support of the MONDian interpretation however, is the fact that the amplitude of the asymptotic $\sigma$ values closely scales with the fourth root of the total baryonic mass of the clusters (Hernandez et al. 2013), in accordance with MONDian predictions. Within a Newtonian interpretation, this last appears as an unexplained coincidence. Further, given observed proper motions, Hernandez et al. (2013) also showed that Newtonian tidal radii at perigalacticon for the clusters in question, are on average a factor of 4 larger than those where the flattening appears. By MONDian gravity we refer to any modified theory where at $a>a_{0}$ scales standard Newtonian gravity is recovered, while for $a<a_{0}$ MONDian dynamics ensue e.g. { TeVeS of Bekenstein (2004), some of the $F(R)$ theories, Capozziello \& De Laurentis (2011), The extended Newtonian gravity of Mendoza et al. (2011) or the covariant $F(\chi)$ of Mendoza et al. (2013).} For such theories, beyond a radius given by $R_{M}= (GM /a_{0})^{1/2}$, centrifugal equilibrium velocities become flat at $V=(GM a_{0})^{1/4}$. Similarly, for pressure supported systems, beyond around $R_{M}$, velocity dispersion profiles will stop falling along Newtonian expectations and flatten out at a level of $\sigma_{\infty}=V/\sqrt{3}$, with $M$ the total baryonic mass for an astrophysical system e.g. Milgrom (1984), Hernandez \& Jimenez (2012). In more general terms, GCs are ideal testing grounds for gravity theories, as in the absence of any detectable gas or dust, the total baryonic mass is composed exclusively of well studied stars with measured metallicities and colour magnitude diagrams (CMDs). These have also been subject to detailed stellar population synthesis modeling tailored to each individual GC, in terms of metallicities and ages (e.g. McLaughlin \& van der Marel (2005). So far, GC MOND dynamical models have concentrated in modeling observed $\sigma(R)$ and projected surface brightness profiles simultaneously, under the assumption of isotropic orbits, i.e., an anisotropy parameter $\beta=0$, and particular forms of the MOND $\mu$ function which mediates the transition between the Newtonian and MOND regimes, e.g. Haghi et al. (2009), or including the effects of orbital anisotropy for particular clusters, e.g. Sollima \& Nipoti (2010) for NGC 2419. No MONDian modelling including anisotropy for the interesting cluster NGC 288 has been performed to date. Here, we use general anisotropic MOND exact analytic solutions for self-gravitating spherical stellar clusters to model simultaneously the projected velocity dispersion and surface brightness profiles of NGC 288. This particular cluster is interesting, as its low volume density result in internal accelerations below $a_{0}$ throughout. This last makes the details of any transition function between the Newtonian and MONDian regimes largely unimportant, a unique case which can be studied for consistency (or otherwise) of a MONDian scenario through analytic dynamical models including $\beta \neq 0$. As expected under MONDian schemes, the observed velocity dispersion profile is flat throughout, at a value of $2.3 \pm 0.15 km s^{-1}$. This is reminiscent of the two classes of sprial galaxies that exist, high and low surface brightness galaxies, with the low density former ones being ``dark matter dominated'' throughout. In section 2 we present the observations used to derive the flat velocity dispersion profile for NGC 288, which are then used in section 3, together with exact analytic MONDian dynamical models and the observed V-band surface density profile compilation for NGC 288 from Trager et al. (1995), to solve for a maximum likelihood dynamical model. Thus, we obtain both best fit values and confidence intervals for the mass to light ratio, the central density of the cluster, and the $\beta$ parameter. It is interesting that the preferred model has a total mass of $3.5 \pm 1.1 \times 10^{4} M_{\odot}$, { which implies a mass to light ratio of $1.09 \pm 0.37$, consistent with independent estimates of this ratio for the present day stellar population of this cluster of $1.42^{+0.37}_{-0.29}$ by Kruijssen \& Mieske (2009)}. Even allowing for arbitrary values of $\beta$, the preferred solution strongly suggests $\beta=0$. Finally, section 4 presents our conclusions. | 16 | 9 | 1609.08519 |
|
1609 | 1609.03117_arXiv.txt | {Power-law frequency distributions of the peak flux of solar flare X-ray emission have been studied extensively and attributed to a system of self-organized criticality (SOC). In this paper, we first show that, so long as the shape of the normalized light curve is not correlated with the peak flux, the flux histogram of solar flares also follows a power-law distribution with the same spectral index as the power-law frequency distribution of the peak flux, which may partially explain why power-law distributions are ubiquitous in the Universe. We then show that the spectral indexes of the histograms of soft X-ray fluxes observed by GOES satellites in two different energy channels are different: the higher energy channel has a harder distribution than the lower energy channel, which challenges the universal power-law distribution predicted by SOC models and implies a very soft distribution of thermal energy content of plasmas probed by the GOES. The temperature ($T$) distribution, on the other hand, approaches a power-law distribution with an index of 2 for high values of $T$. Application of SOC models to statistical properties of solar flares needs to be revisited. | Power-law frequency distributions exist ubiquitously in nature, e.g., the magnitude of earthquakes \citep{Gutenberg:Richter:1954}, the frequency that a word is used in literature \citep{Zipf:1949}. More frequency distributions following a power law can be found in \citet{Clauset:etal:2009} and \citet{Aschwanden:2011}. For a power-law frequency distribution, the number of events $dN$ scales with the magnitude of the event $x(>0)$ as a power-law function: \begin{equation} dN=Ax^{-\delta}dx\,, \label{equ:power-law-diff} \end{equation} where the coefficient $A>0$ and the power-law index $\delta$ are constant. Usually the distribution deviates from a power-law function towards the low end of the magnitude $x$. This deviation can be attributed either to the breakdown of the power-law scaling or to some observational bias \citep{Li:etal:2013}. There have been quite a number of statistical works on the X-ray emission from solar flares. The X-ray peak flux has a power-law frequency distribution with a power-law index varying from 1.6 to 2.1 for different studies \citep[e.g.,][]{Hudson:etal:1969, Drake:1971, Shimizu:1995, Lee:etal:1995, Feldman:etal:1997, Shimojo:Shibata:1999, Veronig:etal:2002, Yashiro:etal:2006, Aschwanden:Freeland:2012}. Without a background subtraction, \citet{Veronig:etal:2002} found that the peak soft X-ray (SXR) flux of flares obey a power-law distribution over three orders of magnitude from the flare GOES class C2.0 to X20. \citet{Feldman:etal:1997} divided flares observed by the GOES into different groups according to the background level and used the background-subtracted peak flux of flares for statistics. They found that the power law distribution can be extended down to A1.0 class flares. Based on the aforementioned statistical studies, \citet{Aschwanden:Freeland:2012} summarized in their Table 2 the total number of flares observed by the GOES, the flux range where the frequency distribution is consistent with a power-law, and the corresponding power-law indices. They found that these observations can be explained with a fractal-diffusive avalanche model \citep{Aschwanden:2012, Du:2015}. Although the peak flux of large flares can be easily obtained due to their high values, the value of the peak flux for small flares is always contaminated by the background emission, instrumental noise, and potential flare identification bias. The latter may be overcome by using the histogram of the flare flux, which incorporates properties of the light curve with the peak flux distribution, to study the statistics. \citet{Zhang:Liu:2015} recently showed that the characteristics of the soft X-ray light curve do not vary with the value of the peak flux. The histogram of the flare flux is therefore intimately related to the peak flux distribution. A series of GOES satellites have taken a huge amount of SXR flux measurement of the Sun over 40 years. In this paper, the data obtained from 1981 to 2012 are used. Moreover to reduce the effect of the selection bias, instead of identifying individual flares, we will include all data points at an original GOES time cadence of 3~s before 2009 and 2~s after 2009 to study the statistics of the differential histograms of the GOES fluxes, which are different from but intimately connected to the frequency distribution of the GOES peak fluxes studied before. The GOES data reduction is presented in Section 2. The histograms are shown in Section 3. In Section 4, we explore the origin of the power law distribution of the differential histogram and its deviation from a power law towards the low value end of the flux. Conclusions and discussions are presented in Section 5. | In this paper we have shown theoretically that if the shape of the flare light curve is not correlated with the peak flux, the differential histogram of the flare flux shares the same power-law distribution as the frequency distribution of the peak flux. Observationally, we have investigated the statistics of all usable GOES 1 - 8~\AA~ and 0.5 - 4~\AA~ flux observed from 1981 to 2012 to minimize the effect of selection bias on frequency distributions. There are two major findings in our work. (1) The histograms of two GOES channels obey power laws with different indices. The index of the power law for the 0.5 - 4~\AA~ GOES flux is harder than the one for the 1 - 8~\AA~ flux. And these two indices do not change with the sampling cadence. (2) A ``bump''-like structure is clearly seen in all the histograms of the 1 - 8~\AA~ flux. It could be interpreted by the element superposition model proposed in this paper. The original element frequency distribution of the entire 32-year data is a power law with an index of 2.09. This index is harder than the one derived from the fitting with the maximum likelihood method. The best-fit parameters of superposed sources $N_\mathrm{bin}$ is correlated with the level of solar activity. The GOES 1 - 8~\AA~ peak flux of flares without background subtraction has been found to follow a power law with an index greater than 2.1, e.g., \citet{Veronig:etal:2002} obtained an index of 2.11, and \citet{Yashiro:etal:2006} derived an index of 2.16. However, when the background is subtracted, the 1 - 8 \AA~peak flux of flares produced a harder index a bit below 2 \citep{Lee:etal:1995, Feldman:etal:1997, Aschwanden:Charbonneau:2002}. This difference between with and without background subtraction can possibly be interpreted by the element superposition model. The case with background subtractions is equivalent to the case with less number of superposed elements, as background subtraction would remove all other elementary sources other than flares. Actually, the frequency distribution of the background-subtracted flare peak flux could be linked to the upper portion of the distribution of the elementary sources. According to the theory of self-organized criticality (SOC) \citep{Aschwanden:2012}, \citet{Aschwanden:Freeland:2012} predicted that the peak flux of flares in SXR has a power-law frequency distribution with an index of 2. From GOES 1 - 8~\AA~ flux measurements, the frequency distributions of background-subtracted flare peak flux during 1975 to 2011 produced a power-law index of $1.98\pm0.11$\citep{Aschwanden:Freeland:2012}. Our result of the power-law index of 2.09 (see Figure~\ref{fig:distri_superp_model}) is a little bit softer than the theoretical value. The obtained power-law index in the 0.5 - 4~\AA~waveband using the maximal-likelihood fitting method (see Figure~\ref{fig:distr_sampling_time}) is about 1.92. The single event in this waveband is also a superposition of a number of elementary sources. If we use the superposition model to derive the original element frequency distribution, we would probably get a lower power-law index than 1.92, which was derived using the maximum likelihood method. As the ``bump'' structure is not very pronounced in this waveband, we did not use the superposition model for a further fitting. According to the SOC theory, the hard X-ray (HXR) peak flux of flares have a power-law index of 1.67. The emission from 0.5 - 4~\AA~ probably contains some contribution from HXR. Therefore, the theoretical index might be between 1.67 to 2. Our power-law index of $< 1.92$ is consistent with the theoretical expectation. These results also indicate that the power-law distribution of X-ray fluxes from solar flares involves convolution of complex physical processes over a broad scale range and may not be simply attributed to some scaling indexes of simple mathematical models. We have to note that the power law index of 2.09 we have derived is the statistical result over 32 years. We have tried to minimize the bias of event selection, as existing in flare statistics \citep{Parnell:Jupp:2000, Aschwanden:Charbonneau:2002, Li:etal:2013}. However, if we go to the distribution in each year as shown in Figure~\ref{fig:1981_2012}, the power-law index for the 1 - 8~\AA~flux distribution ranges from 1.86 to 2.75, and for the 0.5 - 4~\AA~flux, it ranges from 1.61 to 2.29. Therefore, the power-index is quite time dependent. In particular, in year 2005 and 2008, the power indices are very different from the values in other years. In our simple element superposition model, the original power-law distribution has a lower cutoff $x_0$. It is not a necessity, other forms of lower-end deficiency can be used, such as saturation. As mentioned in the introduction, 300 samples can cover the flux in two orders of magnitude in the power-law distribution with an index of 2. For our 322 millions of data points, in principle they can cover eight orders of magnitude data. However, in Figure~\ref{fig:distr_sampling_time} the apparent power law only covers two orders of magnitude. After removing the superposition effect, the frequency distributions of elementary sources could be able to cover five orders of magnitude data. Due to the instrumental saturation of the very high GOES flux, we can not know the exact upper limit of the flux. By extrapolating the frequency distribution of elementary sources to the flux greater than $10^{-2}\,\mathrm{W\,m^{-2}}$ (equivalent to a X100 class flare), we find that this extremely high flux may occur 1000 times per 32 years. As our sampling frequency is 1/3 Hz most of time, it corresponds a time period of 3000 s, similar to the lifetime of a large flare. That is to say, we should be able to observe a X100 class flare within 32 years. Have we observed this kind of super flares? We do not know. It may hide in the saturated data. | 16 | 9 | 1609.03117 |
1609 | 1609.01324_arXiv.txt | In this talk, we first briefly review the isospin dependence of the total nucleon effective mass $M^{\ast}_{J}$ inferred from analyzing nucleon-nucleus scattering data within an isospin dependent non-relativistic optical potential model, and the isospin dependence of the nucleon E-mass $M^{\ast,\rm{E}}_{J}$ obtained from applying the Migdal-Luttinger theorem to a phenomenological single-nucleon momentum distribution in nuclei constrained by recent electron-nucleus scattering experiments. Combining information about the isospin dependence of both the nucleon total effective mass and E-mass, we then infer the isospin dependence of nucleon k-mass using the well-known relation $M^{\ast}_{J}=M^{\ast,\rm{E}}_{J}\cdot M^{\ast,\rm{k}}_{J}$. Implications of the results on the nucleon mean free path (MFP) in neutron-rich matter are discussed. | \label{sec.I} To ease the following discussions, we first recall the basic definitions and relations of the three distinct nucleon effective masses used typically in non-relativistic descriptions of nuclear matter and give a few examples of model predictions. The k-mass $M^{\ast,\rm{k}}_{J}$ and E-mass $M^{\ast,\rm{E}}_{J}$ of a nucleon $J=n/p$ characterizes, respectively, the space and time non-locality of nuclear interactions. They are normally obtained from the momentum and energy dependence of the single-nucleon potential $U_{J} (\rho,\delta, k,E)$ in nucleonic matter of density $\rho$ and isospin asymmetry $\delta\equiv (\rho_n-\rho_p)/\rho$ via \cite{Jeu76,Mah85,jamo} \begin{equation} \frac{M^{\ast,\rm{E}}_{J}}{M}=1-\frac{\partial U_{J}}{\partial E}~\rm{and}~ \frac{M^{\ast,\rm{k}}_{J}}{M}=\left[1+\frac{M_{J}}{|\v{k}|}\frac{\partial U_{J}}{\partial|\v{k}|}\right]^{-1} \end{equation} where $M$ is the average mass of nucleons in free-space. Once an energy-momentum dispersion relation is assumed using the on-shell condition $E=k^2/2M+U_{J} (\rho,\delta, k,E)$, an equivalent single-particle potential either local in space or time can be obtained. The so-called total effective mass $M^{\ast}_{J}$ \begin{eqnarray}\label{em1} \frac{M^{*}_{J}}{M}&=&1-\frac{\d U_{J}(\v{k}(E),E,\rho,\delta)}{dE}\Bigg|_{E(k_{\rm{F}}^{J})}\\\nonumber &=&\left[1+\frac{M}{\hbar^2k_{\rm{F}}^J}\frac{\d U_{J}(\v{k},E(\v{k}),\rho,\delta)}{\d|\v{k}|}\Bigg|_{k_{\rm{F}}^J}\right]^{-1} \end{eqnarray} is then used to characterize equivalently either the momentum or energy dependence of the single-nucleon potential. We emphasize that once nucleons are put on shell, the total effective mass is the only effective mass one can extract from either the first or the second part of the above equation. As we shall discuss later, one then has to use other approaches to evaluate the E-mass and k-mass. The total effective mass is a measure of the energy level density. The well-known relationship \begin{equation}\label{METK} M^{\ast}_{J}=M^{\ast,\rm{E}}_{J}\cdot M^{\ast,\rm{k}}_{J} \end{equation} among the three kinds of nucleon effective masses can be derived by noticing that \cite{Jeu76} \begin{equation} \frac{dE}{dk}\equiv \frac{k}{M_J^*}=\frac{k}{M}+\frac{\partial U}{\partial k}+\frac{\partial U}{\partial E}\cdot \frac{dE}{dk}. \end{equation} In the above, $k_{\rm{F}}^J=(1+\tau_3^J\delta)^{1/3}\cdot k_{\rm{F}}$ with $k_{\rm{F}}=(3\pi^2\rho/2)^{1/3}$ being the nucleon Fermi momentum in symmetric nuclear matter at density $\rho$, $\tau_3^{J}=+1$ or $-1$ for neutrons or protons. \begin{figure}[htb] \includegraphics[scale=0.9,width=8cm]{Fig1-Muther.eps} \caption{Effective k-masses (solid lines) and E-masses (dashed lines) of neutrons (red) and protons (blue) derived from the BHF self-energies using the CD-Bonn interactions for nucleonic matter with an isospin asymmetry of 0.5 at saturation density. Taken from Ref. \cite{mu04} of Hassaneen and M\"uther.} \label{Muther} \end{figure} Many microscopic many-body theories using various interactions have been used in calculating all three kinds of nucleonic effective masses, see, e.g., ref. \cite{LiChen15} for a recent review. Shown in Fig. \ref{Muther} and Fig. \ref{Baldo} are examples of Brueckner-Hartree-Fock (BHF) predictions of the E-mass, k-mass and total effective mass of neutrons and protons in asymmetric nuclear matter using some of the most widely used nuclear interactions. The recent focus of many studies has been on the splitting of the neutron-proton effective masses and its dependence on the isospin asymmetry and density of the neutron-rich medium encountered in heavy-ion collisions and in some astrophysical situations \cite{mrs,page06}, such as in neutron stars and neutrino spheres of supernova explosions. A thorough understanding of the nucleon effective masses is critical for us to better understand many interesting issues in both nuclear physics and astrophysics. Generally, most of the models predict that in neutron-rich medium, neutrons have a k-mass and total effective mass higher than those for protons, and protons have a higher E-mass than neutrons at their respective Fermi surfaces. However, depending on the models and interactions used, the predictions can change dramatically. For example, some of the widely used Skyrme interactions predict that protons have a higher total effective mass than neutrons in neutron-rich matter. Thus, it is very important to extract reliable information about the nucleon effective masses from experiments \cite{Jun15,Lynch16}. While conclusions from recent analyses of heavy-ion experiments using transport models are still quite ambiguous even about the sign of the neutron-proton total effective mass splitting \cite{Jun15,Lynch16}, it is very encouraging that analyses of nucleon-nucleus and electron-nucleus scatterings can constrain clearly at least the sign of the neutron-proton total and E-mass splitting, respectively, at saturation densities \cite{XHLi,CaiLi16a}. In this talk, we shall briefly review theses results and then infer from them the neutron-proton k-mass splitting at saturation density. For more details, please see the original publications in Refs. \cite{XHLi,CaiLi16a,LiBA13}. \begin{figure}[htb] \includegraphics[scale=0.8,width=\columnwidth]{Fig2-Baldo.ps} \caption{Proton (p, full line) and neutron (n, dotted line) total effective masses as a function of density for different values of the isospin asymmetry parameter $\beta$ from a BHF calculation by Baldo et al. in Ref. \cite{Baldo16}. } \label{Baldo} \end{figure} | \label{sec.VI} In summary, due to the space-time non-locality of nuclear interactions single-nucleon potentials are momentum and/or energy dependent. Three distinct nucleon effective masses are normally used to character the momentum/energy dependence of nucleon potentials. How do they depend on the density and isospin asymmetry of the medium? How are they different for neutrons and protons? These have been among the longstanding questions in nuclear physics. Answers to these questions have many interesting ramifications in both nuclear physics and astrophysics. In this talk, we briefly reviewed some of our recent efforts to answer these questions. In particular, we showed that the total effective mass $M^{\ast}_{J}$ for neutrons is higher than that for protons in neutron-rich matter at saturation density based on a comprehensive analysis of existing nucleon-nucleus scattering data. While the E-mass $M^{\ast,\rm{E}}_{J}$ for neutrons is less than that for protons in neutron-rich matter from applying the Migdal-Luttinger theorem to a phenomenological single-nucleon momentum distribution in nuclei constrained by recent electron-nucleus scattering experiments. Combining information about the isospin dependence of both the nucleon total effective mass and E-mass, we inferred the isospin dependence of nucleon k-mass. The latter is important for determining the nucleon MFP in neutron-rich matter. We also noticed some open questions regrading the nucleon effective masses to be further explored both theoretically and experimentally. | 16 | 9 | 1609.01324 |
1609 | 1609.03051_arXiv.txt | We investigated the physical properties of molecular clouds and star formation processes around infrared bubbles which are essentially expanding \HII\ regions. We performed observations of 13 galactic infrared bubble fields containing 18 bubbles. Five molecular lines, \cof, \cos, \cot, \hcn, and \hco, were observed, and several publicly available surveys, GLIMPSE, MIPSGAL, ATLASGAL, BGPS, VGPS, MAGPIS, and NVSS, were used for comparison. We find that these bubbles are generally connected with molecular clouds, most of which are giant. Several bubble regions display velocity gradients and broad shifted profiles, which could be due to the expansion of bubbles. The masses of molecular clouds within bubbles range from 100 to 19,000 $M_\odot$, and their dynamic ages are about 0.3-3.7 Myr, which takes into account the internal turbulence pressure of surrounding molecular clouds. Clumps are found in the vicinity of all 18 bubbles, and molecular clouds near four of these bubbles with larger angular sizes show shell-like morphologies, indicating that either collect-and-collapse or radiation-driven implosion processes may have occurred. Due to the contamination of adjacent molecular clouds, only six bubble regions are appropriate to search for outflows, and we find that four of them have outflow activities. Three bubbles display ultra-compact \HII\ regions at their borders, and one of them is probably responsible for its outflow. In total, only six bubbles show star formation activities in the vicinity, and we suggest that star formation processes might have been triggered. | % \label{sect:intro} The \emph{spitzer} Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE; \citet{2003PASP..115..953B,2009PASP..121..213C}) identifies almost 600 bubbles~\citep{2006ApJ...649..759C,2007ApJ...670..428C}, extended to 5106 by the Milky Way Project (MWP)~\citep{2012MNRAS.424.2442S}. Bubbles were defined by the 8.0 micron emission~\citep{1984A&A...137L...5L}, which contains 7.7 micron and 8.6 micron polycyclic aromatic hydrocarbon (PAH) features, including the continuum~\citep{2010ApJ...713..592E}. \citet{2006ApJ...649..759C} claimed that 25\% of 322 bubbles outside 10\deg\ of the Galactic center coincide with known radio \HII\ regions, and~\citet{2010A&A...523A...6D} extended this proportion to 86\% based on a sample of 102 bubbles selected from the catalog provided by \citet{2006ApJ...649..759C}. Most bubbles are Photodissociation Regions (PDRs) ~\citep{1997ARA&A..35..179H,1999ApJ...527..795K} produced by high-mass stars ionizing atomic or molecular hydrogen. \citet{2010A&A...523A...6D} proposed a simple model for \HII\ region evolution involving two main phases, rapid ionization of the neutral medium followed by a long expansion. During the second phase, shock and ionization fronts form, and neutral material and cold dust collected between them. There are some studies supporting that the expanding of \HII\ regions are three-dimensional. \citet{2010ApJ...709..791B} find three-dimensional structures throughout a sample of 43 bubbles, using the observations of CO (J=3-2) and \hcosi\ (J=4-3) observed by the James Clerk Maxwell Telescope (JCMT). With more sensitive \emph{Herschel} data, \citet{2012A&A...542A..10A} detected emission from "near-side" and "far-side" of bubbles along the line of sight, which suggesting three-dimensional structures for bubbles. Molecular lines, usually characterized with particular critical densities, excitation temperatures, and kinematics information, are superb tools for studying molecular clouds including those around infrared bubbles. Observations of molecular lines contain physical, chemical, and dynamic information which is essential to the study of molecular clouds. CO, a good tracer of molecular clouds due to its low excitation temperature and low critical density~\citep{2001ApJ...547..792D,2015ARA&A..53..583H}, has been widely used to study infrared bubbles. For example, the Galactic Ring Survey (GRS) $^{13}$CO data~\citep{2006ApJS..163..145J} has been present in many bubble papers~\citep{2014A&A...569A..36X,2014A&A...565A...6S,2010A&A...513A..44P,2014MNRAS.438..426H}. Other molecular lines, such as \hcn\ and \hco\ which are probes of dense clumps and cores~\citep{2010ApJS..188..313W,2005ApJ...622..346C} can also been found throughout bubble studies~\citep{2012ApJ...755...71K,2010ApJ...709..791B}. Moreover, \hcos\ and \cof\ are good tracers of bipolar outflows, which are ubiquitous phenomena in star forming regions~\citep{2001ApJ...552L.167Z,2004MNRAS.351.1054R}, and \cofs\ has also been used to identify outflows around infrared bubbles~\citep{2010ApJ...709..791B}. A number of papers argue that star formation processes can be triggered by expanding \HII\ regions, and mainly two mechanisms for the triggering are proposed: collect-and-collapse~\citep{1977ApJ...214..725E}, and radiation-driven implosion (RDI)~\citep{1989ApJ...346..735B}. \citet{1977ApJ...214..725E} proposed that star formation can be triggered by the propagation of ionization and shock fronts through a molecular cloud complex. \citet{2005ApJ...623..917H} analyzed the dynamical expansion of \HII\ regions and the outer PDRs around a high-mass star by solving the UV and FUV radiation transfer and the thermal and chemical processes, using time dependent hydrodynamics. They find that a molecular gas shell with a mass in the order of $10^4$ $M_\odot$ can be shaped in $\sim$1 Myr, and this triggering process is of great importance for star formations of next generation. Seventeen candidate regions for the collect-and-collapse process were identified by \citet{2005A&A...433..565D}, and a large number of young stellar objects (YSOs) were found in the vicinity of bubbles \citep{2008ApJ...681.1341W,2009ApJ...694..546W,2010ApJ...716.1478W}. \citet{2012ApJ...755...71K} found a strong positional correlation between high-mass YSOs (HMYSOs) and \HII\ regions with MWP bubbles at separations of $\textless$2 bubble radii. A statistical study of YSOs around 322 \emph{Spitzer} mid-infrared bubbles has been done by \citet{2012MNRAS.421..408T}, and they found a significant overdensity of Red MSX-Source Survey (RMS) YSOs around the bubbles. These results support that expanding \HII\ regions can provide an effective mechanism to form next generation stars. However, this scenario is not universal. \citet{2012A&A...542A..10A} found the cold gas lies in a ring instead of a sphere, indicating a flattened molecular clouds which could be not greatly compressed by expanding shock fronts. In this case, the formation of new stars could be hindered. \citet{2015MNRAS.450.1199D} investigated the most commonly used signposts and found none of them significantly improved the chances of correctly identifying a given star as triggered. They argued that we should be cautious in interpreting observations of star formation near feedback driven structures in terms of triggering. If bubbles can truly trigger star formation, then the characteristics of star formation, i.e., collapse, outflows, or masers, should be found around them. Although these phenomena are not specific for triggered star formation, the more star formation process we find around bubbles, the safer we can say they are triggered or accelerated by expanding \HII\ regions. Identifying YSOs is the most commonly used method to study triggered star formation. However, it is difficult to determine the distances, ages, and masses of YSOs. If we could find outflows or specific masers, such as OH, H$_2$O, and CLASS \textsc{ii} CH$_3$OH \citep{1995MNRAS.272...96C}, which are direct signposts of star formation, we can at least confirm that star formation processes are indeed present around bubbles, although some of these processes are probably spontaneous. \citet{2009ApJ...702.1615C} made a class \textsc{i} and \textsc{ii} \meth\ masers survey of approximately 20 HMYSO outflow candidates, referred to as Extend Green Objects (EGOs) which are candidates of massive outflows, and three of them are associated with bubbles. \citet{2010ApJ...709..791B} proposed 12 eye-based outflow candidates, and however, stronger evidence is required to confirm them because moments maps are inadequate to confirm outflows due to the complex environments around bubbles. We studied 13 infrared bubble regions selected from the bubble catalog of~\citet{2006ApJ...649..759C}, including 18 bubbles. For each bubble, five molecular lines were observed, revealing the dynamic and physical features of molecular clouds around bubbles. Among the five lines, C$^{18}$O (J=1-0), HCN (J=1-0), and HCO$^+$(J=1-0) are not widely present in previous observations. Several publicly available surveys are also involved, such as the GLIMPSE~\citep{2003PASP..115..953B,2009PASP..121..213C}, the Multiband Imaging Photometer for \emph{Spitzer} (MIPS) Galactic Plane Survey (MIPSGAL)~\citep{2009PASP..121...76C,2015AJ....149...64G}, the APEX Telescope Large Area Survey of the Galaxy (ATLASGAL)~\citep{2009AA...504..415S}, the Bolocam Galactic Plane Survey (BGPS)~\citep{2011ApJS..192....4A,2010ApJS..188..123R}, the VLA Galactic Plane Survey (VGPS)~\citep{2006AJ....132.1158S}, the Multi-Array Galactic Plane Imaging Survey (MAGPIS)~\citep{2006AJ....131.2525H}, and the NRAO VLA Sky Survey (NVSS)~\citep{1998AJ....115.1693C}. Details of these surveys are described in Table \ref{Tab:surveys}. Combing these surveys, we performed a multi-wavelength analysis of the 18 bubbles, focusing on their physical properties and star formation processes around them. \begin{table}[H] \footnotesize \begin{center} \caption{ Observation parameters of surveys. } \label{Tab:surveys} \begin{tabular}{lccccccccc} \tableline\tableline Survey &Wavelengths &Resolutions &Facilities & References \\ \tableline GLIMPSE& 3.6, 4.5, 5.8, 8.0 \mum & $\sim2''$ &\emph{Spitzer} & \citet{2003PASP..115..953B,2009PASP..121..213C} \\ MIPSGAL& 24 \mum &6$''$&\emph{Spitzer} &\citet{2009PASP..121...76C,2015AJ....149...64G} \\ ATLASGAL & 0.87 mm& $19.2''$ &APEX & \citet{2009AA...504..415S} \\ BGPS &1.1 mm &33$''$ &Bolocam &\citet{2011ApJS..192....4A,2010ApJS..188..123R}\\ VGPS &20 cm&60$''$&VLA& \citet{2006AJ....132.1158S} \\ MAGPIS&20 cm&6$''$&VLA& \citet{2006AJ....131.2525H} \\ NVSS &20 cm&45$''$&VLA&\citet{1998AJ....115.1693C}\\ \tableline \end{tabular} \end{center} \end{table} | We have presented an investigation of 13 regions containing 18 infrared bubbles, using three CO isotopic lines and two high density molecular lines, \hcos\ and \hcn. Some profiles of bubble regions show broad redshifted \cofs\ profiles, arc structures, and \cofs\ velocity gradients, indicating they are probably interacting with the molecular clouds around them. Most of the bubbles are associated with dense molecular gas traced by \hcos\ and \hcns, and using \cots, we identified 24 molecular clumps near 18 infrared bubbles, most of which are gravitationally bound. These facts are generally in agreement with either the collect-and-collapse or the RDI model. A search of outflow was carried out in six bubble regions. Four bubble regions, N14, N37, N55, and N133, show outflow activities, while only N55 shows a bipolar structure. The energy of outflows indicates the driven source of this outflow is probably high-mass, and however, no EGO is found towards these outflow candidates. These outflows are convincing evidence that the star formation processes are occurring around bubbles. Besides outflows, ultra-compact \HII\ regions are found on the border of N55, N49, and N123. Among the 18 selected bubbles, six bubbles possess star formation activities nearby, and the detection ratio of outflows and UC \HII\ regions near bubbles is about 0.3. We speculate that star formation processes might have been triggered around these bubbles. However, higher spatial resolution observations are needed to confirm these star formation activities. | 16 | 9 | 1609.03051 |
1609 | 1609.03860_arXiv.txt | The brightest southern quasar above redshift $z=1$, HE 0515$-$4414, with its strong intervening metal absorption-line system at $\zab=1.1508$, provides a unique opportunity to precisely measure or limit relative variations in the fine-structure constant ($\varal$). A variation of just $\sim$3\,parts per million (ppm) would produce detectable velocity shifts between its many strong metal transitions. Using new and archival observations from the Ultraviolet and Visual Echelle Spectrograph (UVES) we obtain an extremely high signal-to-noise ratio spectrum (peaking at $\SN\approx250$\,pix$^{-1}$). This provides the most precise measurement of $\varal$ from a single absorption system to date, $\varal=-1.42\pm0.55_{\rm stat}\pm0.65_{\rm sys}$\,ppm, comparable with the precision from previous, large samples of $\sim$150 absorbers. The largest systematic error in all (but one) previous similar measurements, including the large samples, was long-range distortions in the wavelength calibration. These would add a $\sim$2\,ppm systematic error to our measurement and up to $\sim$10\,ppm to other measurements using Mg and Fe transitions. However, we corrected the UVES spectra using well-calibrated spectra of the same quasar from the High Accuracy Radial velocity Planet Searcher (HARPS), leaving a residual 0.59\,ppm systematic uncertainty, the largest contribution to our total systematic error. A similar approach, using short observations on future, well-calibrated spectrographs to correct existing, high $\SN$ spectra, would efficiently enable a large sample of reliable $\varal$ measurements. The high $\SN$ UVES spectrum also provides insights into analysis difficulties, detector artifacts and systematic errors likely to arise from 25--40-m telescopes. | The Standard Model of particle physics is incomplete because it cannot explain the values of fundamental constants, or predict their dependance on parameters such as time and space. Therefore, without a theory that is able to properly explain these numbers, their constancy can only be probed by measuring them in different places, times and conditions. Furthermore, many theories which attempt to unify gravity with the other three forces of nature invoke fundamental constants that are varying \citep[see][]{2011LRR....14....2U}. In this work we focus on constraining the variability of the fine-structure constant, $\alpha \equiv e^2/ \hbar c$, which represents the coupling strength of the electromagnetic force. In the last 15 years there have been many attempts to measure this constant in absorption systems along the lines of sight to distant quasars. The approach called the ``Many Multiplet" (MM) method, pioneered by \citet{1999PhRvL..82..888D} and \citet{1999PhRvL..82..884W}, compares the relative velocity spacing between different metal ion transitions and relates it to possible variation in $\alpha$. For example, considering just a single transition, variation in $\alpha$ is related to the velocity shift $\Delta v_i$ of a transition \begin{equation} \varal\equiv\frac{\alpha_{\rm obs}-\alpha_{\rm lab}}{\alpha_{\rm lab}}\approx\frac{-\Delta v_i}{2c}\frac{\omega_i}{q_i}\,, \end{equation} where $c$ is the speed of light, $q_i$ is sensitivity of the transition to $\alpha$ variation, calculated from many body relativistic corrections to the energy levels of ions and $\omega_i$ is its wavenumber measured in the laboratory. There have been two MM method studies of large absorber samples: the Keck High Resolution Echelle Spectrometer (HIRES/Keck) sample of 143 absorption systems \citep{1999PhRvL..82..884W, 2001PhRvL..87i1301W, 2001MNRAS.327.1208M, 2003MNRAS.345..609M, 2004LNP...648..131M} and the Ultraviolet and Visual Echelle Spectrograph on the Very Large Telescope (UVES/VLT) sample of 154 absorption systems \citep{2011PhRvL.107s1101W, 2012MNRAS.422.3370K}. These studies reported weighted mean values of $\varal=-5.7\pm1.1$\,parts per million (ppm) and $2.29\pm0.95$\,ppm, respectively. \citet{2011PhRvL.107s1101W} and \citet{2012MNRAS.422.3370K} also combined both samples and found a $4.1\sigma$ statistical preference for a dipole-like variation in $\alpha$ across the sky. Contrary to these studies, several other attempts to measure $\alpha$ from individual absorbers \citep[][]{2004A&A...415L...7Q, 2006A&A...451...45C, 2005A&A...434..827L, 2006A&A...449..879L, 2007A&A...466.1077L} or smaller samples of absorption systems \citep[][]{2004PhRvL..92l1302S, 2004A&A...417..853C} reported $\varal$ consistent with zero. However, these analyses contained shortcomings that produced biased values and/or considerably underestimated error bars (\citeauthor{2007PhRvL..99w9001M}, \citeyear{2007PhRvL..99w9001M}, \citeyear{2008MNRAS.384.1053M}, cf. \citeauthor{2007PhRvL..99w9002S}, \citeyear{2007PhRvL..99w9002S}). For example, a reanalysis including remodelling of the \citet[][]{2004A&A...417..853C} spectra by \citet{2015MNRAS.454.3082W}, yielded a weighted mean of $\varal = 2.2\pm2.3$\,ppm which, given the distribution of the 18 quasars across the sky, was not inconsistent with the evidence for dipole-like $\alpha$ variation. The crucial observational question is whether a systematic effect, or some combination of multiple systematic effects, could mimic the results from the two large studies. Although early studies \citep{2001MNRAS.327.1223M, 2003Ap&SS.283..577M} did not report systematic effects that could strongly affect measurements of $\varal$ obtained from the large statistical samples, more recent analyses \citep[e.g.][]{2013MNRAS.435..861R} identified long-range distortions of the wavelength scale established from traditional ThAr calibration lamp spectra. These distortions were discovered via ``supercalibration" techniques in which solar or solar-like spectra from asteroids or ``solar twin" stars were used to establish an alternative wavelength scale. Earlier, similar supercalibration checks by \citet{2008A&A...481..559M}, \citet{2010ApJ...708..158G} and \citet{2010ApJ...723...89W} did not reveal such long-range distortions. However, such systematic effects were confirmed by \citet{2015MNRAS.447..446W} who analysed 2 decades of archival solar twin spectra from UVES and HIRES. They found that long-range distortions were common and provided a compelling explanation for the non-zero values of $\varal$ found in the large samples of quasar spectra from Keck and VLT, particularly the latter. From ``The UVES Large Program for testing Fundamental Physics", which was specifically designed to measure $\varal$, several recent measurements have been made. \citet{2013A&A...555A..68M} provided the tightest constraint on $\varal$ from an individual absorber, with a statistical precision of 2.4\,ppm and a systematic error estimated to be $\pm$1\,ppm. However, long-range distortions were not corrected in this study. Recently, \citet{2014MNRAS.445..128E} made 9 measurements of $\varal$ in 3 absorbers using 3 different telescopes, and used supercalibrations to correct for the long-range distortions. Their combined result indicated no variation in $\alpha$ at the $\approx$4\,ppm precision level, including estimates of any remaining systematic effects. This represents the only distortion-corrected measurement of $\varal$ so far. Therefore, there is a clear need to both make new, reliable measurements of $\varal$ and to further explore potential systematic effects. This is particularly important in view of the higher-quality spectra that will soon be available from better-calibrated spectrographs and larger telescopes. In this paper we report the most robust and precise constraint on $\alpha$-variation in the well-studied $\zab=1.1508$ absorption system towards QSO HE 0515$-$4414. This is the brightest quasar at redshift above $z=1$ in the southern sky, so it offers the best opportunity to minimize statistical errors and study systematic effects beyond those discernible in lower-fidelity spectra. So far, three constraints on $\varal$ have been attempted in this absorber \citep[][hereafter Q04, L06, C06]{2004A&A...415L...7Q, 2006A&A...449..879L, 2006A&A...451...45C}. The L06 result was revised in \citet[][hereafter M08b]{2008EPJST.163..173M}. We discuss the results from these studies in the context of our new results in \Sref{comp_pre_mea}, suffice it to say here that there are several motivations to measure $\varal$ in this absorption system again. Firstly, significantly more observational spectra are now available, including 3 times the number of UVES exposures previously used. \Tref{obs1} summarises all the available spectra obtained with the UVES spectrograph used in this study. When all the UVES spectra are combined, it results in the highest $\SN$ spectrum taken with a high-resolution optical spectrograph of a quasar, to our knowledge. This allows a thorough investigation of systematic effects which provides an insight into the problems likely to be faced when similar high-quality spectra are obtained from upcoming new spectrographs and telescopes. \newcommand\oldtabcolsep{\tabcolsep} \setlength{\tabcolsep}{0.5em} \begin{table*} \caption{Summary of the observations from UVES/VLT used in this work. `Project ID' represents the internal ESO project number in which exposures were taken. `Exposure time' represents the total observing time in each project, divided into dichroic (when both arms of the spectrograph were in use) and single-arm observations. The blue and red \SN\ correspond to neighbouring continuum to the \ion{Fe}{ii}\,1608 and \ion{Mg}{ii} doublet transitions, respectively. The seeing is the average during the project, with its range in parentheses. Years 1999a and 1999b correspond to 1$\times$2 and 1$\times$1 CCD binning, respectively. In year 1999a, the exposures in the 346-nm setting were taken in the blue arm only, while the 437-nm and 860-nm exposures were taken together in dichroic mode.}\label{obs1} \centering \begin{tabular}{@{}l c c c c c c c c c c c c c c c } \hline Project ID/Year & \multicolumn{5}{c}{Number of exposures} & \multicolumn{2}{c}{Exposure time [s]} & \multicolumn{2}{c}{\SN} & \multicolumn{5}{c}{Slit width ["]} & Seeing ["] \\ & 346 & 437 & 520 & 580 & 860 & Dichroic & Single-arm & Blue & Red & 346 & 437 & 520 & 580 & 860 & \\ \hline 60.A-9022(A)/1999a & $3$ & $2$ & $0$ & $0$ & $2$ & $9500$ & $14200$ & 80 & & $0.8$ & $0.8$ & & & $0.7$ & $0.86$ $(0.6$--$1.4)$ \\ 60.A-9022(A)/1999b & $0$ & $0$ & $1$ & $3$ & $0$ & & $16600$ & & 103 & & & $0.7$ & $0.7$ & & $0.61$ $(0.5$--$0.8)$ \\ 66.A-0212(A)/2000 & $7$ & $6$ & $0$ & $4$ & $9$ & $53100$ & & 70 & 121 & $0.8$ & $0.8$ & & $0.8$ & $0.8$ & $0.61$ $(0.4$--$0.8)$ \\ 072.A-0100(A)/2003 & $0$ & $0$ & $0$ & $2$ & $0$ & & $6120$ & & 68 & & & & $0.7$ & & $0.81$ $(0.8$--$0.9)$ \\ 079.A-0404(A)/2007 & $3$ & $0$ & $0$ & $3$ & $0$ & $9900$ & & 22 & 56 & $0.7$ & & & $0.7$ & & $1.77$ $(1.7$--$1.8)$ \\ 082.A-0078(A)/0809 & $16$ & $0$ & $0$ & $16$ & $0$ & $46400$ & & 79 & 158 & $0.6$ & & & $0.5$ & & $1.00$ $(0.6$--$1.7)$ \\ \hline \end{tabular} \end{table*} \setlength{\tabcolsep}{\oldtabcolsep} Secondly, the previous work on this absorber has used only 6 \ion{Fe}{ii} transitions to constrain $\varal$, of which only one transition, $\lambda1608$, has a velocity shift in the opposite direction from the other transitions if $\alpha$ varies. If there exists a systematic error (e.g.~a calibration error) that causes the $\lambda1608$ line to shift, it would be seen as a non-zero $\varal$. Artificial shifts between this and other \ion{Fe}{ii} transitions are possible, particularly because it falls in a different arm of the UVES spectrograph in this absorption system. Artificial velocity shifts may also be due to the aforementioned long-range distortions. Therefore, in this work we include all transitions identified by \citet[][]{2014MNRAS.438..388M} that are useful for measuring $\varal$, which fall in the wavelength region of our spectra and which are not blended with absorbers from different redshifts. This somewhat increases the information available to reduce the statistical error on $\varal$, with the main motivation being the possibility of greater resistance to systematic effects. Thirdly, it is apparent that the absorption profile models used in previous work significantly `under-fitted' the spectra, with 36 or fewer velocity components. According to \citet[][figure 8]{2008MNRAS.384.1053M}, under-fitting the spectra may lead to a substantial systematic error in the measured $\varal$ value. For this reason, we try to avoid this problem in our model by fitting as much of the statistically significant velocity structure as possible. This implies using 106 velocity components -- see \Sref{fitt} -- which is much larger than used in any previous MM analysis. We explore the possibility that we may still be under-fitting or somewhat `over-fitting' the spectra in \Sref{systematic_err}. Finally, we use complementary observations from the High Accuracy Radial velocity Planet Searcher (HARPS) on the ESO 3.6 m La Silla telescope of the same object to recalibrate the wavelength scales of the UVES spectra. The supercalibration studies of \citet{2011A&A...525A..74M} and \citet{2015MNRAS.447..446W} demonstrated that the HARPS wavelength scale has much smaller (if any) long-range distortions than UVES and HIRES. By directly comparing the HARPS and UVES spectra we effectively transfer the relatively accurate HARPS wavelength scale onto the UVES spectra -- see \Sref{DC_method}. Additionally, we use the HARPS and complimentary Bench-mounted High Resolution Optical Spectrograph (bHROS) on the Gemini South Telescope, of very high resolution, $R\sim140000$, but with lower signal-to-noise ratio ($\SN$), to search for extra velocity structure (\Sref{unresolved}). Complementary observations from the HARPS and bHROS are described in Sections 2.2 and 2.3, respectively. \section[]{Observations, data reduction and calibration} HE 0515$-$4414 is a bright quasar with $V\approx 14.9$\,mag and redshift $\zem=1.71$ which was first identified in \citet[][]{1998A&A...334...96R}. We have used a large set of publicly available observations taken with three instruments: the UVES/VLT, HARPS/ESO$-$3.6\,m and bHROS/Gemini. We also publish the reduced spectra in \citet{srdan_kotus_2016_51715}. \subsection{UVES/VLT} \label{UVES_ORC} The largest data set comprises 90 separate exposures from UVES/VLT \citep{2000SPIE.4008..534D}, which were observed between 1999 and 2009 in five different projects. The main $\varal$ constraint is from this dataset because it collectively has much higher $\SN$ ratio (per \kms) than the other spectra. In all of the UVES exposures, the slit was at the paralactic angle, projected perpendicular to the horizon. Three exposures in 2003 were taken with an iodine cell in the light path \citep[][]{2010ApJ...723...89W} and have significantly lower $\SN$ ratio and the expected forest of $\textrm{I}_{2}$ absorption features; they have been excluded from this analysis. Ten exposures do not have ``attached" ThAr exposures, which means that the echelle or cross-disperser gratings were likely reset between the quasar and ThAr exposures. This can cause velocity shifts and possibly distortions in the wavelength scale of a quasar spectrum and so we exclude these exposures as well. After these selections, the total number of useful UVES exposures reduces to 77 which have attached calibrations for measuring $\varal$ (\Tref{obs1}). These 77 exposures can also have velocity shifts and distortions in their wavelength scales, which we will account for in \Sref{DC_method}. Most of the exposures were taken by observatory staff (i.e.~``service mode"), except the three exposures taken in 2007, which were taken by visiting astronomers (i.e.~``visitor mode"). The exposure time of individual exposures is in the range between 2700 and 5400 s with a mean of $\sim$1\,h. The continuum around the bluest transition that we use in our study, \ion{Fe}{ii}\,1608 which falls near 3460\,\AA\ in the 346-nm setting, has a \SN\ in the range between 9 and 26\,per 1.3-\kms\ pixel in individual exposures. The \SN\ in the continuum near the \ion{Mg}{ii} doublet at $\approx$6030\,\AA\ in the red 580-nm setting is much higher, ranging between 28 and 62\,per pixel in individual exposures. The object is very bright so it was not necessary to observe it only during dark nights. Therefore, all exposures have moon illumination between $0$ and $100$\,per cent, with a mean of $43$\,per cent. \begin{figure*} \centering \includegraphics[width=0.99\textwidth,natwidth=610,natheight=642]{spectrum.pdf} \caption{Example transitions from the combined UVES spectrum of HE 0515$-$4414. The upper panel shows the \ion{Fe}{ii}\,1608 transition, which represents the blue part of the spectrum, with ${\rm S/N}\sim 140$\,per 1.3-\kms\ pixel. The middle and lower panels show the \ion{Fe}{ii}\,2374 and 2382 transition pair and the \ion{Mg}{ii} doublet, respectively, with ${\rm S/N}\sim 240$\,per pixel; these represent the red part of the spectrum. Blue, green and red colours in each panel represent the three fitting regions in our analysis, referred to as the left, central and right region. The combined spectrum here includes all UVES exposures except those with 2$\times$1 binning and is available in \citet{srdan_kotus_2016_51715}.}\label{spectrum} \end{figure*} For data reduction we used the ESO UVES Common Pipeline Library (CPL 4.7.8). Initially it bias-corrects and flat-fields the quasar exposures. The quasar flux is then extracted with an optimal extraction method over the several pixels in the cross dispersion direction where the source flux is distributed. The wavelength calibration, which is a very important step for our measurements, was performed using an attached ThAr lamp exposure, taken immediately after each quasar exposure, and the air wavelengths and calibration procedure described by \citet{2007MNRAS.378..221M}. Instead of using the spectra which are automatically redispersed onto a linear wavelength scale by the CPL code, we used only un-redispersed flux and corresponding error arrays for individual echelle orders in the rest of reduction procedure. After the wavelength calibration, the air wavelength scale of individual echelle orders, in all quasar exposures, was corrected to vacuum using the (inverse) Edlen (1966) formula. It was then converted to the Solar System heliocentric reference frame using the date and time of the mid-point of the exposure integration, using a custom code, {\sc uves\_popler} \citep{michael_murphy_2016_44765}. This code was also used to redisperse the flux from individual exposures onto a common log-linear wavelength scale with dispersion 1.3\,\kms\,$\text{pixel}^{-1}$ for 1$\times$1 binning and $2.5$\,\kms\,$\text{pixel}^{-1}$ for 2$\times$1-binned exposures. The rebinned flux arrays from all exposures were scaled to match that of overlapping orders and then combined with inverse-variance weighting and outlier rejection. We also used {\sc uves\_popler} to automatically fit a continuum to the spectrum using low-order polynomial fits to overlaping $2000$\,\kms\ wide sections. This continuum was generally acceptable, though some local adjustments were made using customised low-order polynomial fits in the vicinity of our transitions of interest. {\sc uves\_popler} automatically rejects pixels at the edges of echelle orders which have uncertainties in flux above some threshold, but we have rejected all pixels at the edges of echelle orders if they overlapped with our transitions of interest. We have also manually rejected some pixels from individual exposures around cosmic rays and other obvious artifacts. The final reduced spectrum, with all exposures combined except those with 2$\times$1 binning, is publicly available in \citet{srdan_kotus_2016_51715}. It covers all wavelengths from 3051 to 10430\,\AA\ except for a small inter-chip gap at 8537--8664\,\AA. Its resolving power is higher at redder wavelengths -- $\sim$63500 at $\la$4500\,\AA\ compared with $\sim$75000 at $\ga$5000\,\AA\ -- because a smaller slit width was typically used for the UVES red arm, though the many exposures contributing to the final spectrum had a range of resolving powers at all wavelengths, ranging between 53500--70000 at $\la$4500\,\AA\ and 62000--93500 at $\ga$5000\,\AA. We discuss the resolving power in more detail in \Sref{res_pow}. Example sections of the spectrum are shown in \Fref{spectrum} which cover 5 of the transitions used in our MM analysis. The continuum around the bluest transition, \ion{Fe}{ii}\,1608 shown in the top panel, falling near 3460\,\AA\ in the 346-nm setting, has a \SN\ of $\approx$140\,per 1.3-\kms\ pixel. The strongest transitions of \ion{Fe}{ii} and \ion{Mg}{i}/{\sc ii}, which dominate our analysis, fall in the wavelength range 5000--6400\,\AA\ (middle and bottom panels), are covered by many exposures in the 580-nm setting and all have a \SN\ of $\approx$240\,per pixel. The \SN\ peaks at $\approx$250\,per pixel around the \ion{Mg}{ii} doublet falling around 6015\,\AA. To our knowledge, these represent the highest \SN\ values for a quasar absorption system in an echelle spectrum at $z>1$. Our spectra were observed over ten years, in five separate projects, with different charge-coupled device (CCD) on-chip binning and a variety of slit widths. This variety determines a range of different nominal resolving powers. This, combined with the very high $\SN$ of the spectra, means that modelling the absorption profiles accurately enough to measure $\varal$ is only possible if the combined spectrum is separated into five `sub-spectra' -- i.e.\ combined sub-sets of exposures taken in different `epochs' with different resolving power and on-chip binning. These sub-spectra are summarised in \Tref{obs1}. Observations from project 60.A-9022(A) are separated into two sub-spectra, 1999a and 1999b, because exposures in 1999a have 2$\times$1 binning, while those in 1999b are not binned. Observations from projects 072.A-0100(A) and 079.A-0404(A) are combined together into sub-spectrum 0307 because they used the same slit width. Other observations were combined according to the project; we refer to these as sub-spectra taken in epochs 2000 and 0809. We do not notice any obvious deviations in the absorption profile shapes between the different epochs. However, this possibility could be investigated further with the published sub-spectra. \subsection{HARPS/ESO-3.6\,m} \label{HARPS_ORC} HARPS \citep{2003Msngr.114...20M} is a fiber-fed spectrograph with resolving power of $R\approx112000$. It is contained in an enclosure in which very stable conditions are maintained, such as very low and stable pressure and constant temperature. It is calibrated with ThAr lines and uses optical fibres, instead of a slit, to introduce light into the spectrograph. However, some of the systematic effects seen in the UVES and HIRES spectrographs also seem to be present in the HARPS wavelength scale. These systematic effects were first identified in the frequency comb study of \citet{2010MNRAS.405L..16W} as short-range distortions that are repeated from echelle order to echelle order. \citet{2013A&A...560A..61M} measured the amplitude of these `intra-order' distortions in each echelle order to be $\pm$40\,\ms. \citet{2013A&A...560A..61M} did not identify any significant long-range distortions, as were found subsequently in UVES and HIRES. While, \citet{2015MNRAS.447..446W} identified small $\sim$45\,\ms\,per 1000\,\AA\ long-range distortions in their analysis of HARPS solar twin spectra, they are most likely caused by systematic errors in the solar FTS spectrum used as the reference spectrum in that analysis. The HARPS spectrum used in this work consists of 18 exposures, with exposure times between 1\,h and 1.75\,h taken during six nights in 2003 and 2009 in ESO Projects 60.A-9036(A) and 072.A-0244(A), with a total exposure time of 93000\,s. A ThAr calibration exposure, taken before each night, was used to derive a wavelength solution for all quasar exposures on that night. The HARPS wavelength scale is stable to within just 15\,cm\,s$^{-1}$ over several hour time-scales \citep{2010MNRAS.405L..16W} so this approach does not limit the calibration uncertainty. The observations used one fibre on the quasar while the other was on a nearby sky position to allow sky subtraction. The quasar, sky and ThAr flux, together with an estimate of the blaze correction from flat-field exposures, were all automatically extracted by the standard HARPS data reduction software. This software also derived the wavelength calibration solution from the extracted ThAr flux. These products were then combined using {\sc uves\_popler} in the same way as the UVES exposures. However, the sky-subtraction, construction of a flux error spectrum, and blaze correction are not performed by the HARPS reduction software, so these steps were performed within {\sc uves\_popler}. The sky flux was redispersed onto the same common wavelength grid and subtracted from the quasar flux in the same echelle order before the (sky-subtracted, blaze-corrected) flux from all orders and exposures was combined. The error spectrum for each exposure was estimated, assuming Gaussian statistics, from the combination of the quasar CCD electron counts and an estimate of the read noise in each extracted pixel. The final reduced HARPS spectrum that we use in our analysis, which we make publicly available in \citet{srdan_kotus_2016_51715}, covers the wavelength range between 3791 and 6905\,\AA\ with a small gap between 5260 and 5338\,\AA\ due to the physical gap between the HARPS blue and red CCD chips. The \SN\ around 5050\,\AA\ is $\approx$29\,per 0.85-\kms\ pixel but at the position of the \ion{Mg}{ii} doublet it peaks at $\approx$33\,per pixel. \subsection{bHROS/Gemini} \label{bHROS_ORC} bHROS \citep{2008psa..conf..297M} was bench-mounted in the gravity-invariant pier of the Gemini South Telescope and fed by a 0\farcs9\,$\times$\,0\farcs9 optical fibre that, via an image slicer, projected a 0\farcs14-wide pseudo-slit into the spectrograph to achieve a resolving power of $R\approx140000$. A total of 37\,$\times$\,3600-s exposures of HE0515$-$4414 were obtained during bHROS ``science verification" in November and December 2006 and January 2007 during variable conditions. Each quasar exposure was followed immediately by a ThAr lamp exposure. Unfortunately, one of the two CCD chips failed before the observations and the efficiency of bHROS appeared well below specifications during them. The former meant that observations in 2 separate wavelength settings were required to cover most strong transitions in the $\zab=1.1508$ absorber. The spatial profile of light from the image slider was spread over a large number of pixels ($\approx$50), so the contribution of CCD read noise was significant and further reduced the \SN\ of the spectra considerably. ThAr lamp observations also revealed instabilities that caused $\sim$0.5\,\kms\ shifts in the spectrum over several hours. These factors dramatically reduced the \SN\ of the spectra and rendered them useless for directly measuring $\varal$. However, the very high resolution may still assist in revealing additional velocity structure, and we explore this possibility in \Sref{sys_v_str}. No dedicated data reduction pipeline is available for bHROS spectra, so we used our own custom codes, based loosely on the {\sc reduce} suite of routines \citep{2002A&A...385.1095P}. The two-dimensional shifts between quasar exposures was determined from their corresponding ThAr exposures and the heliocentric velocity at the mid-point of the quasar exposure integration. These shifts were used to combined the low-\SN, individual (dark current, bias and flat-field corrected) quasar exposures into a single, higher-\SN\ one upon which the flux extraction procedure was performed. The flux in each echelle order containing a transition of interest was extracted with a custom optimal extraction routine which accounted for bad pixels and cosmic rays. Error arrays were estimated in the same process. No sky-subtraction was possible because only a single fibre was available. However, the observations were conducted with typically $\sim$25--50\,per cent moon illumination, so the sky background flux is expected to be negligible. The final, reduced bHROS spectrum is publicly available in \citet{srdan_kotus_2016_51715} for the transitions of \tran{Mg}{i}{2852}, \tran{Mg}{ii}{2796}, 2803 and \tran{Fe}{ii}{2383}, 2586 \& 2600 that could be covered in two separate wavelength settings. The spectra around these transitions were extracted to a linear wavelength grid with 0.012\,\AA\ dispersion and the final \SN\ for, e.g., the \tran{Fe}{ii}{2382} at 5123\,\AA\ is $\approx$24 per 0.70-\kms\ pixel. | In this work we have measured the relative variation in the fine-structure constant in the $\zab=1.1508$ absorber towards the quasar HE 0515$-$4414. Because of its brightness ($V\approx14.9$\,mag) this quasar has been observed frequently with UVES, so a large number of archival spectra are available ($\approx$30\,h). Here we add 13\,h of new UVES exposures to obtain a total $\SN\approx250$\,per 1.3-\kms\ pixel (at its peak), the highest for any echelle spectrum of a quasar at $\zem>1$. This, and the large number of narrow features in the \ion{Mg}{i/ii} and \ion{Fe}{ii} absorption profiles, provide a very small statistical uncertainty on $\varal$ of 0.55\,parts per million (ppm). Most importantly, we have corrected a large systematic error in the UVES spectra, from long-range distortions of the wavelength calibration, by using a HARPS spectrum of the same quasar. Left uncorrected, these would have caused a spurious shift in $\varal$ of $\approx$2.1\,ppm. However, by directly comparing the UVES and HARPS spectra the correction leaves a residual systematic uncertainty of just 0.59\,ppm. This assumes that the distortions are linear with wavelength and that they have the same slope in the red and blue arms of the UVES spectrograph. Previous studies of these distortions in UVES spectra generally support these assumptions \citep{2015MNRAS.447..446W}. Other systematic errors, mainly from short-range (i.e.\ intra-order) distortions and uncertainties in the absorber's velocity structure, contribute a further 0.26\,ppm uncertainty. A series of consistency checks suggest that our total systematic error budget of 0.65\,ppm is reliable and that astrophysical systematic errors, such as isotopic abundance variations, are unimportant. Our final result for this absorber is $\varal=-1.42\pm0.55_{\rm stat}\pm0.65_{\rm sys}$\,ppm. This is consistent with no variation in the fine-structure constant and is the most precise measurement from a single absorption system to date. The precision is comparable to the ensemble precision from the large Keck and VLT samples of absorption systems studied by \citet{2004LNP...648..131M} and \citet{2012MNRAS.422.3370K}. It is unlikely that measurements of $\varal$ in other, individual absorption systems, will match the precision obtained here until new 25--40-m telescopes become available. Indeed, this work provides a preview of effects that must be addressed in the very high $\SN$ spectra from those future telescopes. For example, accurate knowledge of the resolving power was required to model the absorption profile in this work, and CCD and/or data-reduction artifacts were evident in our $\SN\sim250$\,pix$^{-1}$ spectrum. Finally, given that all previous $\varal$ measurements \citep[except those of][]{2014MNRAS.445..128E} will have been significantly affected by long-range distortions, it is crucial to obtain new measurements which are corrected for (or resistant to) this important systematic effect. In this context, the upcoming ESPRESSO spectrograph on the VLT is very important because it should provide spectra free of this effect (and the intra-order distortions). This work provides two insights for using ESPRESSO most efficiently for precise and reliable $\varal$ measurements. Firstly, the recalibration approach demonstrated in this work could be applied: existing, high-$\SN$ UVES spectra could be recalibrated with relatively short ESPRESSO observations ($\sim$2--3\,h) of the same quasars. This would produce a sample of $\sim$20 reliable, high-precision measurements in $\sim$20\,per cent of the time required to build the same $\SN$ with all new ESPRESSO spectra. Secondly, comparison of the UVES and higher-resolution HARPS and bHROS spectra of HE 0515$-$4414 implies that the absorber comprises many closely-packed, narrow velocity components; the increased resolution provides little additional information about the velocity structure or a significant increase in precision on $\varal$. Therefore, if the velocity structures of most other absorbers also comprise many closely-packed, narrow components, it is unlikely that a resolving power $R>100000$ will benefit varying-$\alpha$ measurements, especially if $\SN$ is compromised to obtain the higher resolution. | 16 | 9 | 1609.03860 |
1609 | 1609.08175_arXiv.txt | {The Magellanic System (MS), consisting of the Large Magellanic Cloud (LMC), the Small Magellanic Cloud (SMC) and the Magellanic Bridge (MBR), contains diverse sample of star clusters. Their spatial distribution, ages and chemical abundances may provide important information about the history of formation of the whole System. We use deep photometric maps derived from the images collected during the fourth phase of The Optical Gravitational Lensing Experiment (OGLE-IV) to construct the most complete catalog of star clusters in the Large Magellanic Cloud using the homogeneous photometric data. In this paper we present the collection of star clusters found in the area of about 225 square degrees in the outer regions of the LMC. Our sample contains 679 visually identified star cluster candidates, 226 of which were not listed in any of the previously published catalogs. The new clusters are mainly young small open clusters or clusters similar to associations. }{Catalogs Star Clusters: general Surveys} | The Magellanic System is an ideal astrophysical laboratory for studying the structure and evolution of galaxies. It is located close enough to the Galaxy so that millions of stars can be easily resolved. This can be used in a range of scientific projects exploring stellar populations, their evolution, ages, metallicity (Jacyszyn-Dobrzeniecka \etal 2016, Skowron \etal 2014). One of the methods for such studies is the investigation of star clusters. The Large Magellanic Cloud contains a large sample of these systems. The spatial distribution od clusters, their age, chemical composition, structural parameters and dynamical evolution may provide valuable information about the LMC formation history. The Magellanic System has a very rich and diverse structure. The positions of centroids of both Clouds depend on which stellar population is used for their estimation (Cioni, Habing, Israel 2000, Deb and Singh 2014). Moreover, the Magellanic Clouds % look asymmetrical, with denser parts located toward the Bridge (Klein \etal 2014, Scowcroft \etal 2015). Analysis of the spatial distribution of the star clusters in the MS enables the comparison to the general stellar populations. This may give some hints of what caused the asymmetry of the LMC and the SMC. The correlation between age, size, metallicity and spatial distribution can bring new information about the MS history (Palma \etal 2016, Piatti \etal 2014). To make all these studies possible, however, a complete collection of star clusters is needed, derived from homogeneous observational data, possibly from a single photometric survey. So far, the largest catalog of extended objects (excluding background galaxies) in the Magellanic System was published by Bica \etal (2008) and contains a compilation of all the previously published catalogs. The most important contribution to this sample was due to the catalog based on the OGLE-II data published by Pietrzy{\'n}ski \etal (1999). That catalog, however, covered only the central part of the LMC: 5.8 square degrees -- only 3\% of the area observed during the current OGLE-IV phase (Udalski, Szyma{\'n}ski and Szyma{\'n}ski 2015). The parameters (such as size, age, metalicity) of some clusters listed in the Bica catalog (Bica \etal 2008) were estimated by several groups, including Glatt, Grebel and Koch (2010), Palma \etal\ (2016), Choudhury, Subramaniam, Piatti (2015) and Piatti \etal\ (2002, 2003a,b, 2009, 2014, 2015). The OGLE survey data were also used for this purpose, first by Pietrzy{\'n}ski and Udalski (2000) using OGLE-II data (Udalski, Kubiak and Szyma{\'n}ski, 1997) and recently, by Nayak \etal\ (2016), using OGLE-III data (Udalski \etal 2008). The OGLE-IV survey covers practically the entire Magellanic System inclu\-ding both galaxies and the Magellanic Bridge (about 650 square degrees). The large collection of high resolution individual images enabled constructing deep images of the entire region with a depth comparable to images collected by a 4-meter class telescope. Uniform OGLE-IV photometric data provides a unique, homogeneous data set for the search of star clusters and associations. It will allow us to create a complete catalog of star clusters in the MS and their basic parameters. This paper presents the first part of the catalog based on the OGLE-IV data. The central part of the LMC has already been observed or analyzed by many other projects so we decided to start our exploration with the outer parts of the galaxy. We have found 679 star clusters in the outer regions of the LMC. Among those, 226 objects were not listed in any of the previous catalogs, 438 objects were listed in the Bica catalog and 15 objects were listed in the Dark Energy Survey (DES) publication (Pieres \etal 2016). As some extended objects cannot be unambiguously classified, we have performed a cross-match of our sample to both star clusters and associations from the Bica catalog. Almost all known objects which were in the area of the analyzed OGLE fields were detected by our algorithm, proving the effectiveness of algorithm and the completeness of the sample. There are five Bica objects which we do not see in our images and one DES object which is located in the gap between the OGLE subfields. | We have presented a catalog of star clusters in the outer regions of the Large Magellanic Cloud based on the OGLE-IV deep photometric maps. We found a total of 679 star clusters, including 226 new objects which were not listed in any of the previous catalogs, 438 clusters listed in Bica \etal 2008 and 15 objects listed in the Dark Energy Survey publication (Pieres \etal 2016). For all of them the equatorial coordinates and cross-identification with previous catalogs are provided. The detection method presented in this paper is very effective. With our algorithm we found almost all previously known clusters in this characteristic sparse region of the LMC and increased the total number of these objects by 50\%. This paper is the first of a series of publications. In the coming papers we will present a complete catalog of star clusters in the whole Magellanic System with their parameters determined, for the first time, from a single, homogeneous photometric data set collected by the OGLE project. | 16 | 9 | 1609.08175 |
1609 | 1609.04671_arXiv.txt | We investigate a new method to search for keV-scale sterile neutrinos that could account for Dark Matter. Neutrinos trapped in our galaxy could be captured on stable $^{163}$Dy if their mass is greater than 2.83~keV. Two experimental realizations are studied, an integral counting of $^{163}$Ho atoms in dysprosium-rich ores and a real-time measurement of the emerging electron spectrum in a dysprosium-based detector. The capture rates are compared to the solar neutrino and radioactive backgrounds. An integral counting experiment using several kilograms of $^{163}$Dy could reach a sensitivity for the sterile-to-active mixing angle $\sin^2\theta_{e4}$ of $10^{-5}$ significantly exceeding current laboratory limits. Mixing angles as low as $\sin^2\theta_{e4} \sim 10^{-7}$ / $\rm m_{^{163}\rm Dy}\rm{(ton)}$ could possibly be explored with a real-time experiment. | 16 | 9 | 1609.04671 |
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1609 | 1609.01781_arXiv.txt | We present a numerical and analytical study of incompressible homogeneous conducting fluids using a helical Fourier representation. We analytically study both small- and large-scale dynamo properties, as well as the inverse cascade of magnetic helicity, in the most general minimal subset of interacting velocity and magnetic fields on a closed Fourier triad. We mainly focus on the dependency of magnetic field growth as a function of the distribution of kinetic and magnetic helicities among the three interacting wavenumbers. By combining direct numerical simulations of the full magnetohydrodynamics equations with the helical Fourier decomposition we numerically confirm that in the kinematic dynamo regime the system develops a large-scale magnetic helicity with opposite sign compared to the small-scale kinetic helicity, a sort of triad-by-triad $\alpha$-effect in Fourier space. Concerning the small-scale perturbations, we predict theoretically and confirm numerically that the largest instability is achieved for the magnetic component with the same helicity of the flow, in agreement with the Stretch-Twist-Fold mechanism. Vice versa, in presence of a Lorentz feedback on the velocity, we find that the inverse cascade of magnetic helicity is mostly local if magnetic and kinetic helicities have opposite sign, while it is more nonlocal and more intense if they have the same sign, as predicted by the analytical approach. \\ Our analytical and numerical results further demonstrate the potential of the helical Fourier decomposition to elucidate the entangled dynamics of magnetic and kinetic helicities both in fully developed turbulence and in laminar flows. \\ \noindent {\it Key words:} dynamo $-$ magnetohydrodynamics (MHD) $-$ turbulence | \label{sec:intro} Turbulent flows are ubiquitous on Earth and in the sky, e.g.~\citep{Frisch95,Pope00}. In many astrophysical and geophysical cases \citep{Belenkaya09}, the fluid is also conducting and one needs to control the entangled dynamics of velocity and magnetic fields; this is the case of magnetohydrodynamic (MHD) turbulence \citep{Moffatt78,Biskamp03,Verma04}. The spectrum of possible configurations and applications is vast, depending on the presence of particular external forcing mechanisms, mean flows or fields and boundaries \citep{Priest00,Priest14,Plihon14,Stieglitz01,Gailitis00,Gailitis01,Goodman02,Nornberg06,Frick10}. Here we wish to address basic properties of all MHD configurations connected with the interactions leading to the transfer of total energy, magnetic helicity and kinetic helicity across scales. For MHD, the presence of three inviscid invariants, total energy, magnetic and cross-helicity, makes the problem of predicting spectral properties difficult \citep{Iroshnikov64,Kraichnan65a,Matthaeus89,Goldreich95,Boldyrev05a,Boldyrev06}. Even the direction of the different transfers is not completely under control, only empirical results exist \citep{Biskamp03}. Moreover, because of obvious applied and fundamental issues, predicting or controlling the growth rate of a magnetic field is a key question, connected to the famous dynamo problem \citep{Moffatt69,Krause80,Brandenburg03,Brandenburg05,Tobias13}. Since the magnetic energy is not conserved, the magnetic field may stretched, folded or advected even in absence of external input and dissipation. A huge amount of literature has been devoted to the identification of the key dynamical and statistical ingredients needed to promote or deplete such a growth and to control the growth rate. Magnetic and kinetic helicities are among the key quantities that play a role in such a phenomenon: \begin{align} \label{eq:Hm} H_m(t) &=\int_V d \bx \ \ba(\bx,t) \cdot \bb(\bx,t) \ , \\ \label{eq:Hk} H_k(t) &=\int_V d \bx \ \bu(\bx,t) \cdot \bomega(\bx,t) \ , \end{align} where $\bu$, $\bb$, $\ba$ and $\bomega$ are the velocity, the magnetic field, the magnetic vector potential and the vorticity, respectively. The third ideal invariant of the MHD equations is given by cross-helicity: \begin{align} \label{eq:Hc} H_c(t) &=\int_V d \bx \ \bu(\bx,t) \cdot \bb(\bx,t) \ , \end{align} that is connected to the degree of Alfv{\'e}nization of the system, i.e. to the presence of waves traveling in the direction of the mean global or local magnetic field \citep{Dobrowolny80,Biskamp93}. In this paper we will focus on the importance of magnetic and kinetic helicities (helicities in short in what follows) for the growth rate of a large-scale magnetic field, always considering the case of almost vanishing cross helicity (see also the concluding remarks about possible generalization of our work to include also the latter). To be as simple as possible we will concentrate only on periodic and homogeneous conditions. The analytical and numerical work is based on the helical Fourier decomposition developed for Navier-Stokes equations by the pioneering work of \citet{Waleffe92} and \citet{Constantin88}. This decomposition is exact and has led to an important breakthrough in the understanding of the entangled energy-helicity dynamics in three-dimensional Navier-Stokes turbulence \citep{Waleffe92,Biferale13}. It has only recently been extended to MHD \citep{Lessinnes09,Linkmann16}, and it promises to be a key tool also for problems where the physics is controlled by the interactions among magnetic or kinetic helical waves \citep{Cho11,Galtier03,Galtier05,Galtier14}. Moreover, the helical Fourier basis is also the natural decomposition to be used in numerical simulations, either to analyze the data or to perform explicit numerical experiments by projecting the equations on a given subset of Fourier modes, in order to highlight the physics of some particular interacting waves. This procedure has already been carried out for the Navier-Stokes equations with some surprising results \citep{Biferale12,Biferale13a,Sahoo15,Alexakis16a} connected to the discovery of a sub-class of (kinetic) helical modes transferring energy backwards in a fully three-dimensional turbulent flow, i.e. the identification of those Fourier interactions responsible for the energy backscatter. Simulations of homogeneous magnetohydrodynamic turbulence in a periodic box without a background magnetic field have been used as prototypical systems to study the circumstances under which large-scale magnetic field growth occurs, such as the $\alpha$-effect \citep{Steenbeck66,Brandenburg01} and the inverse cascade of magnetic helicity \citep{Frisch75,Pouquet76,Balsara99,Brandenburg01,Alexakis06,Mueller12,Malapaka13}. Concerns have been raised in the literature about the effectiveness of the $\alpha$-effect in generating large-scale magnetic fields with strong amplitude because of the detrimental feedback that fast-growing small-scale magnetic fields have on the growth rate of the large-scale magnetic field \citep{Vainshtein92,Cattaneo96}. This is the problem of catastrophic $\alpha$-quenching, and much theoretical efforts have been made in order to find dynamo models which are able to circumvent this problem. In this paper we focus on the statistical and dynamical factors that might promote or deplete the growth of a large-scale magnetic field. We do this by using a systematic dissection of the three-dimensional MHD equations in helical Fourier modes as pioneered by \citet{Linkmann16}. We first analyze the temporal evolution of three velocity and three magnetic helical Fourier modes at wavenumbers, $k,p,q$, the basic brick of any quadratic non-linear transfer. Considering all helical combinations, there are 64 possible different subsets of closed dynamical systems that represent the minimal backbone of interacting modes conserving all inviscid invariants. A graphical representation of such a basic system is given in Figure~\ref{fig:two-triad-system}. A stability analysis of a subfamily of these dynamical systems was carried out by \citet{Linkmann16}, considering the most general equilibria. This led to the identification of linear instabilities that could be associated with forward and inverse transfer of total energy and magnetic helicity. Here we restrict our attention to equilibria and instabilities that can be connected to cases of astrophysical interest, i.e. kinematic dynamo regimes and inverse cascade effects. We further extend the analysis carried out by \citet{Linkmann16} by specifying the magnetic field growth rates (if any) and providing clear predictions on the expected helical signatures of the dominant instabilities. This enables us to link specific dynamical properties connected to the kinematic dynamo action and/or the inverse cascade of magnetic helicity with the geometrical structure of the triad (local versus non-local Fourier interactions) and with its kinematic contents (helicities). Interestingly, the entangled dynamics of velocity field with the magnetic fields is already extremely rich at the level of these most basic interactions. \\ The second part of the paper is devoted to develop for the first time a thoughtful numerical validation and benchmark of the previous theoretical analysis by Direct Numerical Simulation (DNS) of the full MHD equations with and without a small-scale forcing on the magnetic field. Forcing the magnetic field allows us to switch from a situation where the magnetic field is initially in the kinematic dynamo regime to a case where Lorentz force is always acting at all scales, thanks to the strong injection of magnetic fluctuations. We always analyze velocity and magnetic fluctuations in terms of their helical Fourier components such as to be able to directly match the theoretical predictions based on the simplified single-triad dynamics. We changed the helical properties of the magnetic forcing (when applied) to break the mirror symmetry with different injection mechanisms. We study two different configurations, with large- or small-scale injection of kinetic energy and helicity, corresponding to turbulent or laminar regimes. Furthermore, we also present some ad hoc simulations by restricting the dynamics of the velocity field to evolve only on modes with one given sign of helicity, which induces a strong breaking of the mirror symmetry already at the level of the equations of motion.\\ The combined analytical and numerical analysis lead to the | \label{sec:conclusions} We studied the dynamics of helical triad interactions in homogeneous MHD turbulence both analytically and numerically. We have shown that the helical Fourier decomposition of the full MHD equations is a key tool to better disentangle different inertial transfer processes in the fully coupled dynamics. First, we extended the set of helical triad interactions \citep{Lessinnes09,Linkmann16} to the most general MTI systems, and we clarified in which cases the stability analysis of the subset of triadic interactions carried out in \citep{Linkmann16} is sufficient to capture all possible linear instabilities. We further analysed two cases of astrophysical interest concerning the emergence of large scale magnetic fields, i.e. dynamo action and the inverse cascade of magnetic helicity, by extending the results of \citet{Linkmann16} to provide qualitatively testable predictions on the helical contents of the resulting magnetic growth at large scales and small scales. Subsequently, we carried out two series of suitably designed numerical experiments in order to test the theoretical results. In Sections~\ref{sec:lsdynamo} and \ref{sec:ssdynamo} we clarified which of the linear instabilities identified by \citet{Linkmann16} is the leading one and which global helical signature should be expected for the magnetic field. In Section~\ref{sec:selfinteraction} we focused on linear instabilities which can be associated with the inverse cascade of magnetic helicity: a helical magnetic equilibrium at a given scale can only be unstable with respect to like-signed helical magnetic perturbations at larger scales \citep{Linkmann16}. We found that the level of kinetic helicity affects the growth rates associated with the inverse magnetic helicity cascade: the inverse cascade of magnetic helicity is more efficient in a helical flow where magnetic and kinetic helicity are of the same sign. We point out that the perturbation problem considered for strongly magnetised flows was restricted to large-scale magnetic perturbations only, and it did not include the effect of the Lorentz force on the flow. In principle, these two dynamical effects are intimately related. A more refined analysis that distinguishes between nonlinear dynamo and inverse cascade effects, for instance, therefore requires a simultaneous analysis of the Lorentz force, which in turn requires a distinction between homo- and heterochiral MTI systems. Similarly, the possible influence of the characteristic scale of the flow on the local or non-local nature of the inverse cascade would also require a study of the general MTI system. Due to the structure of the chosen equilibria our analysis did not consider the effect of non-negligible cross-helicity on the evolution of the magnetic field. Let us remark that the effects induced by an equilibrium solution with a non-trivial cross-helicity can also be handled analytically as shown by \citet{Linkmann16}. \\ \noindent All theoretical results have been derived from a linear stability analysis of the basic triadic structure of the MHD equations where only three modes interact. However, in any physical MHD configuration all modes interact, therefore it is not immediately obvious if the theoretical results obtained from the single MTI systems correctly predict the behavior of the full dynamics \citep{Moffatt14}. Therefore we carried out two series of numerical experiments, where series D discussed in Section~\ref{sec:linear} corresponds to dynamo simulations and series IC in Section~\ref{sec:nonlinear} to simulations with an inverse magnetic helicity cascade. The numerical results confirm the theoretical predictions presented in the previous paragraph. Our dynamo results agree qualitatively with the dynamo simulations by \citet{Brandenburg01} except for some quantitative difference in the dimensionless growth rates that is due to different forcing strategies and scales. \noindent It is important to note that the triadic dynamo instabilities analyzed here do not result from any further modeling assumptions such as scale separation or the first-order smoothing approximation \citep{Moffatt78,Krause80,Brandenburg05}. Instead, they are present in the basic dynamics of the MHD equations restricted to a small number of degrees of freedom. Furthermore, the triadic $\alpha$-type dynamo instabilities found here may not have the same limitations as the classical $\alpha$-effect of mean-field electrodynamics that originate from the more efficient growth of the small-scale magnetic field compared to the large-scale magnetic field \citep{Vainshtein92}. In the latter case, one considers the mean-field induction equation of the $\alpha$-dynamo $ \partial_t \vec{B}_0 = \alpha \nabla \times \vec{B}_0 \ , $ where the magnetic field $\vec{B}$ has been decomposed into a large-scale mean and a small-scale fluctuating part $\vec{B} = \vec{B}_0 + \bb$, and the coefficient $\alpha$ is given by $ \alpha = \frac{1}{3}\Big(-\langle \bu \cdot \vec{\omega} \rangle + \langle \bb \cdot \vec{j} \rangle \Big) \ , $ with the angled brackets denoting an appropriate average \citep{Brandenburg05} and $\vec{j} = \nabla \times \bb$ the current density. If the growing small-scale magnetic field $\bb$ and the small-scale flow $\vec{u}$ have like-signed helicities, then the coefficient $\alpha$ decreases, thus quenching the growth rate of the large-scale magnetic field $\vec{B}_0$. Because the evolution of the small-scale magnetic field $\bb$ is faster than that of large-scale magnetic field $\vec{B}_0$, the coefficient $\alpha$ could be quenched before leading to large-scale magnetic field growth. Therefore concerns have been raised in the literature about whether the classical $\alpha$-effect is efficient enough to generate large-scale magnetic fields. This is especially problematic at high magnetic Reynolds numbers, where it eventually leads to catastrophic $\alpha$-quenching \citep{Cattaneo96}. Our analysis is concerned with instabilities of the ideal MHD equations, and the corresponding growth rates do not depend on the magnetic Reynolds number $Rm$. Hence the dynamo instabilities we found are in principle present even at large $Rm$. A process similar to $\alpha$-quenching and/or saturation can also be studied within our approach, but it may not be catastrophic because the the growth rates are independent of $Rm$. The small-scale instabilities we found in Section~\ref{sec:ssdynamo} preferentially lead to a growing small-scale magnetic field with the same sign of helicity as the flow. Eventually, the small-scale magnetic field will back-react on the flow, and the linear instability leading to the $\alpha$-like large-scale dynamo may be removed. At this point we cannot be more precise, as a rigorous assessment of dynamo quenching is outside the scope of the stability analysis carried out here. This requires mixed equilibria, while only purely mechanical or electromagnetic equilibria were analyzed here. However, we point out that dynamo quenching and the effect of the Lorentz force can also be assessed by a similar kind of stability analysis. This analysis, which requires a different set of equilibria and is technically more complex, is currently in progress and will be reported elsewhere. \\ \noindent The $\alpha$-effect has further limitations due to its intrinsic scale separation. \citet{Boldyrev05} showed by considering the Kazantsev model \citep{Kazantsev68} that fast-growing eigenmodes exist at all scales, which are not included in the $\alpha$-dynamo due to the required scale separation. Although the Kazantsev model assumes the velocity field to have Gaussian statistics and is as such is not applicable to turbulence, the important point is that a correct description of dynamo action should involve all scales. We point out that scale separation is not necessary for the derivation of the triad-by-triad dynamo instabilities. Nevertheless, information about local and nonlocal dynamics can be obtained by varying the shape of the wavevector triad, and we find that strongly nonlocal triads lead mostly to $\alpha$-type dynamo action. A recent numerical investigation into % the efficiency of the kinematic dynamo depending on the energy injection scale showed that intermediate-scale forcing results in the most efficient dynamo \citep{Sadek16}. This may be qualitatively interpreted by noting that the triadic dynamo growth rates depend on the geometry of the triad, i.e., the locality and nonlocality of the triadic interactions. The growth rate first increases with decreasing equilibrium (or forcing) scale, but this trend is reversed once the scale separation becomes very large, which suggests an intermediate range of forcing scales where a triadic dynamo may indeed be most efficient. \\ \noindent Beyond the confirmation of the theoretical results further interesting observations can be made from the numerical work. First, the instability of a mechanical helical equilibrium associated with large-scale kinematic dynamo action appears to persist in the nonlinear regime, as shown in Figures~\ref{fig:dynamo} and \ref{fig:upbpbm} in Section~\ref{sec:linear}. This suggests that a very strong magnetic field is necessary in order to quench the dynamo. Second, as shown in Figures~\ref{fig:r2ic-spectra}(c) and \ref{fig:r3ic-spectra}(d) in Section~\ref{sec:nonlinear}, the transfer of magnetic to kinetic energy due to the feedback of the Lorentz force on the flow at the small scales is sensitive to the sign of magnetic helicity: The velocity field modes with the same sign of helicity as the magnetic field increase in intensity. On the theoretical side, we found an important difference between the large-scale magnetic field growth due to dynamo instabilities compared to instabilities of the inverse cascade type, which occur due to nonlinear self-interaction in strongly magnetized flows. Dynamo action produced large-scale magnetic fields of opposite helicity compared to the small-scale flow while the inverse cascade of magnetic helicity was most efficient in the helical magnetic field sector with the same sign of helicity as the flow. This may lead to transitional behavior with increasing small-scale magnetic field strength, eventually changing the helical signature of the large-scale magnetic field. Furthermore, it suggests the possible existence of a dynamo quenching mechanism at the basic triad level. \\ | 16 | 9 | 1609.01781 |
1609 | 1609.09490_arXiv.txt | We present a promising new technique, the $g$-distribution method, for measuring the inclination angle ($i$), the innermost stable circular orbit (ISCO), and the spin of a supermassive black hole. The $g$-distribution method uses measurements of the energy shifts in the relativistic iron line emitted by the accretion disk of a supermassive black hole due to microlensing by stars in a foreground galaxy relative to the $g$-distribution shifts predicted from microlensing caustic calculations. We apply the method to the gravitationally lensed quasars \rxj\ ($z_{s}$ = 0.658, $z_l$ = 0.295), \qj\ ($z_{s}$ = 1.294, $z_l$ = 0.317), and \sdss\ ($z_{s}$ = 1.734, $z_l$ = 0.68). For RX~J1131$-$1231 our initial results indicate that $r_{\rm ISCO}$~$\simlt$~8.5~gravitational radii ($r_{\rm g}$) and $i$ $\simgt$ 76$^{\circ}$. We detect two shifted Fe lines, in several observations, as predicted in our numerical simulations of caustic crossings. The current ${\Delta}E$-distribution of \rxj\ is sparsely sampled but further X-ray monitoring of RX~J1131$-$1231 and other lensed quasars will provide improved constraints on the inclination angles, ISCO radii and spins of the black holes of distant quasars. \noindent \\ \\ | One technique for measuring the innermost stable circular orbit (ISCO) and spin parameter $a$ ($a = Jc/GM_{BH}^{2}$, where $J$ is the angular momentum) of AGN relies on modeling the relativistically blurred Fe K$\alpha$ fluorescence lines originating from the inner parts of the disk (e.g., Fabian et al. 1989; Laor 1991; Reynolds \& Nowak 2003). This relativistic iron line method has be applied to about 20 relatively bright nearby Seyferts where the line is detectable with a high signal-to-noise ratio (Reynolds 2014 and Vasudevan et al. 2016). The sample sizes are starting to become large enough where the distribution of the spin parameter can be calculated and compared to simulated ones such as those presented in Volonteri et al. (2013). Even then, the Fe K${\alpha}$ line in most Seyferts is typically very weak, and constraining the spin and accretion disk parameters of Seyferts requires considerable observing time on {\sl XMM-Newton} and {\sl Chandra}. Moreover, the accuracy of the relativistic Fe line method for constraining the spin of a black hole from the broadened red wing of the Fe line profile is also questioned (e.g., Miller et al 2009; Sim et al. 2012). We note, however, that independent measurements of the size of the corona from microlensing and reverberation mapping indicate that the X-ray source is compact, consistent with the lampost model assumed in the relativistic Fe iron line method. Additional support of the relativistic Fe line method is provided by 3--50 keV observations with {\sl NuSTAR}, such as the recent observations of Mrk~335 that indicate a spin parameter of $>$~0.9 at 3 $\sigma$ confidence. The high-energy {\sl NuSTAR} spectra can help to constrain the reflection component and better distinguish between models (Parker et al. 2014). Most of the measured spin parameters in Seyfert galaxies are found to be $\simgt$ 0.9 (e.g., Reynolds 2014 and references therein). This may be the result of a selection bias in flux-limited samples (Vasudevan et al. 2016). Specifically, high spin black holes are more luminous and hence brighter for a given accretion rate, and hence will simply be more highly represented in flux limited surveys. Recent observations and simulations (Fabian 2014; Keck et al. 2015; Vasudevan et al. 2016) also suggest that rapidly spinning black holes will tend to have stronger reflected relative to direct X-ray emission, making it easier to measure the spin parameter in these objects. The relativistic iron line method has been applied to the gravitationally lensed quasars \rxj\ (Reis et al. 2014) and Q2237$+$0305 (Reynolds et al. 2014) and well as to a stacked spectrum of 27 lensed 1.0 $<$ $z$ $<$ 4.5 quasars observed with {\sl Chandra} (Walton et al. 2015). Specifically, the relativistic disk reflection features were fit with standard relativistic Fe~K$\alpha$ models to infer inclination angles and spin parameters of $i = {{15^{\circ}}^{+9^{\circ}}_{-15^{\circ}}}$ and $a = 0.87^{+0.08}_{-0.15}$ for \rxj\ (Reis et al. 2014) and $i \simlt 11.5^{\circ}$ and $a = 0.74^{+0.06}_{-0.03}$ for Q2237$+$0305 (Reynolds et al. 2014). These studies have not, however, correctly accounted for the effects of gravitational microlensing. Gravitational microlensing is a well studied phenomenon in lensed quasars (e.g., see the review by Wambsganss 2006 and references therein) where stars near the lensed images produce time variable magnification of source components whose amplitude depends on the location and size of the emission region. In particular, in our analysis of the X-ray spectra of lensed quasars (Chartas et al. 2016), we have frequently observed structural changes in the Fe K$\alpha$ emission indicating that the line emission is being differentially microlensed. Thus, applying the relativistic Fe~K$\alpha$ line method to stacked spectra of lensed quasars, without accurately accounting for microlensing, is likely to lead to unreliable and unrealistic results. In Section~2 we present the X-ray observations and analyses of the {\sl Chandra} observations of \rxj, \qj, and \sdss. In Section~3 we discuss our recently developed technique based on microlensing to provide a robust constraint on the inclination angle, the location of the ISCO, and the spin parameter, and present an analytic estimate of the fractional energy shifts $g$~=~$E_{\rm obs}$/$E_{\rm rest}$ and numerical simulations of microlensing events. In Section~4 we present our results from modeling the observed distribution of $g$ and the distribution of the measured energy separations of shifted Fe~K$\alpha$ lines in cases where two shifted lines are detected in an individual spectrum. Finally, in Section~5 we rule out several alternative scenarios to explain the shifted iron lines and present a summary of our conclusions. Throughout this paper we adopt a flat $\Lambda$ cosmology with $H_{0}$ = 67~km~s$^{-1}$~Mpc$^{-1}$ $\Omega_{\rm \Lambda}$ = 0.69, and $\Omega_{\rm M}$ = 0.31 (Planck Collaboration et al. 2013). | Our systematic spectral analysis of all the available {\sl Chandra} observations of lensed quasars RXJ1131, QJ0158, and SDSS1004 has revealed the presence of a significant fraction of lines blueshifted and redshifted with respect to the energy of the expected Fe~K$\alpha$ fluorescent line. We interpret these energy shifts as being the result of ongoing microlensing in all the images. This seems logical given the prior detections of microlensing of the optical/UV (e.g., Blackburne et al. 2006; Morgan 2008; Fohlmeister et al. 2007; Fian et al. 2016, Motta et al. 2012) and X-ray continuum (e.g., Chartas et al. 2009, 2012, 2016; Dai et al. 2010; Chen et al. 2012) in all three sources. We consider several alternative scenarios and examine whether they can explain the observed shifted iron lines. (a) {\sl Non-microlensed emission from hot-spots and patches from an inhomogeneous disk.}\\ Redshifted Fe emission lines have been reported in observations of a few bright Seyfert galaxies (i.e., Iwasawa et al. 2004; Turner et al. 2004, 2006; Miller et al. 2006; Tombesi et al. 2007). Correlated modulation of redshifted Fe line emission and the continuum were reported in NGC 3783 (Tombesi et al. 2007). Specifically, the spectrum of NGC 3783 shows, in addition to a core Fe~K$\alpha$ line at 6.4~keV, a weaker redshifted wing and redshifted Fe emission line component. The redshifted line and wing appear to show an intensity modulation on a 27~ks timescale similar to that of the 0.3-10~keV continuum. Tombesi et al. (2007) argue that the lack of Fe line energy modulation disfavors the orbiting flare/spot interpretation for NGC 3783. We note that the relative intensity of the core Fe line to the redshifted Fe line component in NGC 3783 is about a factor of 9. The core Fe~K$\alpha$ line has an equivalent width of about 120~eV and the redshifted line of $\sim$13 eV. Redshifted lines are also reported to be present in the spectrum of Mrk 766 (Turner et al. 2004; 2006), where a weak component of the Fe line with an equivalent width of 15$^{+6}_{-5}$~eV shows a periodic variation of photon energy. The proposed scenario by Turner et al. (2006) is that the energy variation is caused by a hot spot on the disk within $\sim$ 100~$r_{\rm g}$ orbiting with a period of about $\sim$~165~ks. The average spectrum of Mrk~766 shows a broad iron line centre near 6.7~keV with an equivalent width of about 90~eV (possibly from reflection off an ionized disk) and a narrower component at 6.4keV (possibly from reflection from distant material). The shifted Fe lines detected in our gravitationally lensed quasar sample have very different properties from those reported in NGC 3783 and Mrk 766 while our our microlensing interpretation is consistent with the observed properties of the energy shifted lines. Specifically, the equivalent widths of the energy shifted lines in the lensed quasar sample range from EW = 500~eV to 3,000~eV compared to the EW=13~eV and EW=15~eV in the shifted lines detected in NGC 3783 and Mrk 766, respectively. Non-microlensed emission from hot-spots and patches from an inhomogeneous disk lie below the detection threshold of the individual spectra of the lensed quasars. The non-microlensed emission from the Fe~K$\alpha$ line at 6.4~keV is not detected in individual spectra of \rxj, but is detected in the stacked spectra of \rxj\ with E = 6.36$^{+0.07}_{-0.08}$~keV and EW = 154$^{+70}_{-80}$~eV (Chartas et al. 2012, 2016). In the cases of Seyfert galaxies NGC~3783 and Mrk~766 the redshifted weak Fe lines are always accompanied by a significantly stronger core component near 6.4~keV whereas this is not the case for the lensed quasars. Another point supporting the microlensing interpretation is the detection of a significant number of double lines in \rxj, as predicted in our numerical simulations of magnification caustics crossing an accretion disk. The intensities of these doubles are not consistent with non-microlensed Fe emission from hot-spots and patches from an inhomogeneous disk. (b) {\sl Possible intrinsic absorption mimicking two apparent lines in \rxj.} \\ Our analysis of the 4$\times$38 spectra of \rxj\ does not find any significant intrinsic absorption of the continuum spectra (Chartas et al 2012) and there are no absorption lines detected in the spectra that contain doubles. (c) {\sl Possible ionization of accretion disk.} \\ In Figure 9 of Chartas et al. (2012) we showed a stacked spectrum of image C of RXJ1131 covering a period of about seven years filtering for epochs where microlensing was not significant in this image. A significant Fe~K$\alpha$ line is detected in the stacked spectrum of image C at an energy of 6.36$^{+0.07}_{-0.08}$~keV with an equivalent width of 154$^{+70}_{-80}$~eV. This detected energy of the Fe~K$\alpha$ line in image C is consistent with the presence of a non-ionized disk in RXJ1131. The expected energy of the iron line for an ionized disk would be 6.67keV (He-like Fe) or 6.97keV (H-like iron). We conclude that the variability in the flux and energy of the iron line detected in the {\sl Chandra} observations of RXJ1131, QJ0158, and SDSS1004 is due to ongoing microlensing of all the images. Large magnification events are typically inferred from the departure of the time-delay corrected flux-ratios of images from a constant and images that show significant uncorrelated variability are significantly affected by microlensing. We note, however, that a relatively large number of the shifted Fe lines were found in images that do not show significant variability of their flux ratios or significant uncorrelated variability. This implies that the line profiles found using stacked spectra taken over multiple epochs, even when excluding spectra from images that show variability of their flux ratios, are also distorted by microlensing and that applying the relativistic Fe~K$\alpha$ method as used to analyze local Seyferts will not lead to reliable results. Any apparent broadening of the iron line in a stacked spectrum of a lensed quasar is a combination of both microlensing and the relativistic blurring seen in unlensed Seyferts. The average stacked spectrum of a microlensed quasar differs significantly, especially near the Fe K$\alpha$ line, from that of an unlensed one. The main reason for this difference is that the reflection and direct(coronal) components of a stacked spectrum of a microlensed quasar are not a simple uniform magnification of the reflection and direct components of a unlensed quasars. Caustic magnification patterns (e.g. see Figures $12-16$ of Kochanek 2004) move along a certain direction with respect to the source (that differs between images) and the caustic magnification is not uniform. Stacked spectra of lensed quasars will thus not result in uniformly magnified reflection and direct spectral components but will produce spectra that have been selectively magnified by caustics moving along a limited range of caustic crossing angles. The caustic magnification factor K will also vary between caustic crossings making the stacked spectrum deviate even more than a uniformly magnified quasar spectrum. We summarize our main conclusions from the X-ray observations of RXJ1131, QJ0158, and SDSS1004 as follows: (1) Redshifted and blueshifted Fe lines with rest-frame EWs ranging between 500~eV and 3,000~eV are detected in the individual epoch spectra of lensed quasars RXJ1131, QJ0158, and SDSS1004. We interpret these energy shifts as the result of microlensing of Fe line emission within $\sim$ 20 $r_{\rm g}$ of the black hole. (2) The $g$-distribution of the observed energy shifts in RXJ1131 is compared to analytic and numerical models and both models provide similar constraints on accretion disk parameter. Specifically, the maximum value of $g_{\rm max}$= 1.29$^{+0.04}_{-0.04}$ constrains the inclination angle to be $i \simgt 76^{\circ}$ and the minimum value $g_{\rm min}=0.59$ constrains $r_{\rm ISCO} \simlt 8.5 r_{g}$. One of the strengths of the $g$-distribution method is that the energies of the shifted Fe lines are more robustly detected than the extreme weak red wings of the relativistic Fe~K$\alpha$ line that are found mostly in nearby Seyfert galaxies. For example, the energies of the shifted lines are not sensitive to the modeling of the underlying continuum, while, the shape of the relativistic Fe~K$\alpha$ red wing is very sensitive to the continuum model. The $g$-distribution method can be applied to infer the inclination angles, ISCO radii and spins of distant quasars where the relativistic Fe~K$\alpha$ red wing is typically to faint to constrain these parameters in distant quasars. One of the weaknesses of the $g$-distribution method is that there are additional model parameters related to describing the gravitational lens. (3) Several spectra show two shifted Fe lines, which we refer to as doubles. The peak energies of the doubles are moderately correlated. Our numerical simulations reproduce the double lines during caustic crossings and we find that the distribution of the separations of the peak energies is strongly dependent on the spin parameter. The maximum of the observed ${\Delta}E$-distribution of \rxj\ peaks at $\sim$ 3.5 $\pm$ 0.2~keV suggesting a high spin parameter. Although a statistical analysis of the results still needs to be performed, inspection of the mean energy, peak energy, and shape of the distribution indicates a rather high value of the spin $a \simgt 0.8$. The available ${\Delta}E$-distribution of \rxj\ is sparsely populated and additional observations are required to better constrain the peak energy of this distribution and infer the spin parameter more accurately. (4) We find several correlations in the microlensed spectra of RXJ1131 with results summarized in Table 8. Specifically, we find a correlation between the rest-frame equivalent width of the iron lines and the generalized Doppler shift parameter $g$ of the iron line ($\tau$~=~0.28, $P > 99.9\%$ ). The flux of the shifted Fe line is found to be correlated with the flux of the continuum for Fe K$\alpha$ lines detected at $>$ 90\% confidence in all images of RXJ1131 ($\tau$~=~0.5, $P > 99.9\%$ ). The energies of the doubles in image A are also found to be correlated ($\tau$~=~0.6, $P > 98\%$ ). (5) Our numerical simulations of microlensing caustic crossings reproduce the observed distribution of energy of single and double shifts in microlensed spectra, and predict the correlations observed in the spectra. Scheduled future monitoring observations with {\sl Chandra} of the lensed quasars RXJ1131, QJ0158, SDSS1004, and Q 2237+0305 with sufficiently long exposure times to improve the significance of the detections will provide more representative and complete distributions of the generalized Doppler shift $g$ values of the Fe~K$\alpha$ line in these objects. The $g$-distribution method and modeling of doubles is expected to provide robust constraints on inclination angle, the ISCO radii and spins of the black holes of these distant lensed quasars. With the Large Synoptic Survey Telescope coming online in the near future we expect $\sim$ 4,000 new lensed systems to be discovered, opening up the possibility of measuring black hole and accretion disk parameters over a wide range of redhsifts and quasar Eddington ratios $L_{Bol}/L_{\rm Edd}$. | 16 | 9 | 1609.09490 |
1609 | 1609.00052_arXiv.txt | The geometry and intrinsic ellipticity distribution of ultra diffuse galaxies (UDG) is determined from the line-of-sight distribution of axial ratios q of a large sample of UDGs, detected by Koda et al. (2015) in the Coma cluster. With high significance the data rules out an oblate, disk-like geometry, characterised by major axi a=b$>$c. The data is however in good agreement with prolate shapes, corresponding to a=b$<$c. This indicates that UDGs are not thickened, rotating, axisymmetric disks, puffed up by violent processes. Instead they are anisotropic elongated cigar- or bar-like structures, similar to the prolate dwarf spheroidal galaxy population of the Local Group. The intrinsic distribution of axial ratios of the Coma UDGs is flat in the range of $0.4 \leq $a/c$ \leq 0.9$ with a mean value of $\langle a/c \rangle = 0.65 \pm 0.14$. This might provide important constraints for theoretical models of their origin. Formation scenarios that could explain the extended prolate nature of UDGs are discussed. | A new class of rather peculiar but frequent galaxies has been discovered, called Ultra-Diffuse Galaxies (UDG; van Dokkum et al. 2015a,b; Koda et al. 2015; van der Burg, Muzzin \& Hoekstra 2016; Roman \& Trujillo 2016). UDGs are quiescent stellar systems on the red galaxy sequence with exceptionally low stellar surface densities of order a few M$_{\odot}$ pc$^{-2}$, two orders of magnitudes below the typical surface densities of Milky-Way type objects. They have effective radii of 2-6 kpc, similar to giant galaxies. Their luminosities are however of order $10^8$ L$_{\odot}$, resembling dwarf galaxies. UDGs had been seen earlier (Impey, Bothun \& Malin 1988, Dalcanton et al. 1997). Only recently has it however become clear that they can be quite ubiquitous, contributing significantly to the total galaxy population at least in some galaxy clusters. This was shown by Koda et al. (2015; see also Yagi et al. 2016), who detected of order 800 UDGs in the Coma cluster with a radial distribution within the cluster that is similar to the giant galaxies outside a core radius of 300 kpc. The origin of these ghost galaxies represents an interesting mystery. Gas, entering a dark halo will always tend to settle into a rotationally supported disk which is the lowest energy state for given specific angular momentum. The stars that form from this gas disk should then also be distributed in a disk. Amorisco \& Loeb (2016) argue that the large radii of UDGs might result from the infall of very high-angular momentum gas. The stellar surface mass densities of UDGs are however smaller than the critical gas surface density, of order 10 M$_{\odot}$ pc$^{-2}$, required for molecular gas clouds to form and condense into stars (McKee \& Krumholz 2010). This indicates that the surface densities of the gas disks during the star formation phase were much higher than the observed stellar surface densities. It in turn means that the current diffuse state results at least partly from a phase of substantial gaseous mass loss, e.g. through ram pressure stripping or stellar feedback driven galactic winds (Agertz \& Kravtsov 2015; Yozin \& Bekki 2015), accompanied maybe by an expansion of the stellar disk. Di Cintio et al. (2016) show that multiple episodes of gas infall and blow-out could result in such an expansion of the stellar disk and, interestingly, the formation of a cored dark matter halo (Burkert 1995). This is in agreement with earlier numerical simulations of Ogiya \& Mori (2011, 2014) and Ogiya et al. (2014) who find that the non-linear response of dark halos to periodic events of gas in- and outflows could generate cored dark matter density distributions with scaling relations that are in agreement with observations (e.g. Burkert 2015; Kormendy \& Freeman 2016). Violent destruction of an early, fast rotating disk would also be consistent with the measurements of high stellar velocity dispersions in one of the largest Coma UDGs, Dragonfly 44 (van Dokkum et al. 2016; see also Mart\'inez-Delgado et al. 2016). Its ratio of velocity dispersion to rotational velocity is 0.2 and its mass-to-light ratio is 50 -100, indicating that the stellar system is not self-gravitating. The stars are just tracer particles within a dominant dark halo potential, despite the fact that at the time of their formation the gas disk must have been self-gravitating in order to form stars. Because if its kinematics, van Dokkum et al. (2016) classified this galaxy as a dispersion dominated, elliptical-like galaxy, rather than a disk galaxy. Interestingly, the individual galaxies studied by van Dokkum et al. (2016) and Mart\'inez-Delgado et al. (2016) live in low density field environments. The harsh conditions of a galaxy cluster might therefore not always be required for their formation. On the other hand, Mihos et al. (2015) discovered one UDG in the Virgo cluster that shows evidence for tidal disruption. Most UDGs however do not show such signatures and might have been shaped by internal processes, probably a combination of high-angular momentum gas infall (Amorisco \& Loeb 2016), combined with one or more violent episodes of gas loss (Di Cintio et al. 2016). If the stellar component of UDGs formed in a dense, self-gravitating disk-like gas component, the question arises whether information of this initial state is still hidden in the current structure of UDGs. Galactic disks usually are characterised by low Sersic indices n, of order unity. Indeed the surface density distribution of UDGs is well described by n $\approx 1$. Van Dokkum et al. (2015a) and Roman \& Trujillo (2016) however argue that the distribution of apparent axial ratios q is not flat, as expected for thin disks. UDGs could however be thick disks, puffed-up by perturbations as discussed earlier. Typical axial ratios are quite high with q $\approx 0.5$. They are therefore more likely spheroids than thick disks. Thick, axisymmetric disks resemble oblate spheroids. The observed ellipticities could however also be explained by prolate shapes, as e.g. expected if the UDGs are bars or anisotropic ellipticals, affected by tidal fields or if the stars just trace the potential of a prolate dark matter core. Determining the intrinsic geometry and the true intrinsic ellipticity distribution of UDGs could therefore add valuable information and constraints for a better understanding of their origin. This is the purpose of this paper. Koda et al. (2015) provide a histogram of the measured axial ratios of a large sample of UDGs in the Coma cluster. In Section 2 we deproject an updated histogram of the axial ratio distribution of Coma UDGs, provided by Jin Koda (private communication, see also Koda et al. 2015) and derive their true intrinsic axial ratios. We demonstrate that, with high significance, the Coma UDG population cannot have an oblate, disk-like geometry. Adopting prolate shapes, we find an intrinsic ellipticity distribution that in projection is in good agreement with the observations. A discussion of this result and conclusions follow in Section 3. | The deprojection of observed axial ratios of Coma UDGs leads to the conclusion that these galaxies are on average prolate, rather than oblate. Focussing on Figure 4 and adopting an oblate geometry (upper left panel), N is strongly negative for the largest axial ratio bin. The problem is much smaller for prolate geometries. One could of course argue that this large observed deficit of round galaxies, adopting oblate intrinsic shapes, is a result of some observational bias, such that oblate spheroids, seen face on, are more difficult to detect. In this case, at least 66 round UDGs, out of a total number of 768 UDGs, would be missing which appears unlikely. One could also argue that UDGs have a mixture of shapes, some being oblate, others prolate. Note however that even for prolate shapes projection effects lead to a somewhat larger number of round galaxies, compared with the observations. This $\sim 1 \sigma$ discrepancy exists independent of binning and might indeed indicate some observational bias against round ellipticals. Assuming that some of these galaxies are actually oblate would however make this problem worse. Minimizing the difference between observed round galaxies and theoretically expected round galaxies therefore requires that all UDGs are prolate. Interestingly, Coma UDGs and Virgo cluster dwarf ellipticals have a very similar distribution of shapes. Chen et al. (2010, see also Lisker et al. 2009) investigated the structural properties of 100 ACS Virgo Cluster Survey galaxies. The mean axial ratios of their faint galaxy sample is $\langle$ q $\rangle = 0.73 \pm 0.18$. This value is precisely the same as the projected mean axial ratio of the Coma UDGs with $\langle$ q $\rangle = 0.72 \pm 0.16$. In addition, Chen et al. (2010) find no Virgo dwarfs, flatter than 0.35. And, like the Coma UDGs, the number of dwarfs increases with increasing axial ratios, with a depression in the highest axial ratio bin, that is for round objects. Chen et al. (2010) note that this signature is not consistent with random projection of flat, disk-like geometries. Similar processes might therefore have shaped the Coma UDGs and the Virgo dwarf galaxy population. That UDGs are prolate, rather than oblate should provide important information about the processes that lead to their characteristic diffuse state. First of all, it indicates that they are dispersion dominated bars or cylinders, rather than rotation supported, thick disks, in agreement with the small $v_{rot}/\sigma$ values found by van Dokkum et al. (2016). If the stellar component is in virial equilibrium within a dominant dark halo potential this could be achieved either by assuming that the dark halo distribution itself is prolate and/or that the stellar system's velocity dispersion is anisotropic with larger dispersions parallel to the long axis, probably as a result of the processes that generated their elongated state. In this case, the population of projected round UDGs should have systematically larger observed velocity dispersions than flattened UDGs of the same stellar mass, as the line-of-sight of round UDGs would be parallel to the long axis. The situation would be the opposite for oblate galaxies where velocity dispersions are usually larger in the equatorial plane than perpendicular to it. It is interesting that prolate shapes would link the UDGs to their diffuse, low-mass relatives: Local Group dwarf spheroidal galaxies (dSphs). Hayashi \& Chiba (2015), for example, investigated the kinematical data of 12 dSphs and argue that these objects probably life in elongated, bar-like dark matter halos with their shapes reflecting the elongation of their confining dark halos. They also note that, if true, this elongation is inconsistent with $\Lambda$CDM models. This is interesting as it might indicate that internal processes that lead to the diffuse state of these galaxies also reshaped their inner dark halo components. The flat distribution of prolate axial ratios of UDGs in the range of 0.4 to 0.9, independent of their effective radii, provides important constraints for any theoretical model of their formation. Up to now, we do not know whether existing models can reproduce this result. For example, Ceverino et al. (2015) and Tomassetti et al. (2016) find in their high-resolution numerical simulations that prolate shapes of stellar systems and dark matter halos are generic for lower-mass galaxies with stellar masses of order M$_* \leq 10^9$ M$_{\odot}$. These authors focussed on early galaxies in the high cosmic redshift regime of z $\approx 2-4$. Whether this applies also to present-day galaxies and to the peculiar and extreme physical conditions that produce UDGs is however not clear. In addition, it has not been shown that the distribution is flat and in the observed range of axial ratios. If the stars formed in a fast rotating disk that was lateron puffed up by substantial mass loss one might expect that the system keeps a preferentially oblate shape. On the other hand, violently relaxing particle systems can experience radial orbit instability that generates a bar with an anisotropic velocity dispersion (Merritt 1987). Another process could be an encounter with a massive perturber or tidal effects caused by the Coma cluster potential. An interested suite of numerical simulations of tidally stirred disky satellite galaxies has been presented by Martinez-Delgado et al. (2013, see also Kazantsidis et al. 2017). They focussed on the population of dwarf spheroids orbiting Milky-Way-sized hosts and demonstrated that especially for cored dark matter halos (Burkert 1995) tidal effects could transform oblate, disky dwarfs into prolate spheroids. The same mechanism might actually be active for UDGs in cluster potentials although it is not clear yet whether it would lead to the observed ellipticity distribution with $\langle a/c \rangle = 0.65 \pm 0.14$. This scenario is also promising, as it might explain the peculiar orientiation of the UDG's long axis in Coma which is not random. Yagi et al. (2016) find that the UDG's major axis is preferentially radially aligned towards the cluster center. Note however, that UDGs have now also been found in low-density environments. Martinez-Delgado et al. (2016) report the discovery of an ultra-diffuse, quenched galaxy, DGSAT 1, in the outskirts of the Pisces-Perseus supercluster. They argue that DGSAT 1 might be a "backsplash" galaxy that passed through the center of the cluster Zw 0107+3212 (Gill et al. 2015) with high velocities and is now in the outskirts at distances of 1-2 virial radii. If field UDGs did not experience tidal effects they might actually be oblate, rather than prolate. In order to test this conjecture one would however need a sample of order a few 100 objects in order to determine their geometrical shape using the method, presented in this paper. Maybe, the stars in UDGs never formed in a disk configuration. Could star formation have occured in individual dense clouds, orbiting within the inner region of a cored, prolate dark halo? Like normal galaxies, UDGs have an extended globular cluster population (Beasley et al. 2016; Peng \& Lim 2016; Beasley \& Trujillo 2016; van Dokkum et al. 2016), similar to the halo globular clusters of the Milky Way. Did the stars and globulars form ex-situ in substructures that were lateron accreted by a dark halo that did not manage to form in-situ stars? In this case, UDGs would be failed galaxies with an accreted stellar halo component, similar to the Milky Way halo, but missing a disk component. Even more extreme would be the assumption that the more than 800 Coma UDGs are actually galaxies that are currently in the process of being tidally disrupted (Hozumi \& Burkert 2015; Ploeckinger et al. 2015). A prominent example of such a process is the S-shaped dwarf galaxy, observed in the Hydra I cluster (Koch et al. 2012). In order to distinguish between these scenarios more detailed observations, including an investigation of the kinematics and rotation of UDGs and the metallicity and age distribution of their stellar populations is required. | 16 | 9 | 1609.00052 |
1609 | 1609.02327_arXiv.txt | It is observationally and theoretically well established that a considerable amount of dust is efficiently formed in regions around asymptotic giant branch (AGB) stars. This process is classically considered as the primary source of dust grains in galaxies, and the typical formation timescale in the Milky Way is $\sim 3\times 10^9$\,yr. In contrast, supernova (SN) explosions in the interstellar medium (ISM) trigger shock waves that are able to quickly process dust grains and are considered the dominant mechanism of dust destruction in the ISM. Recent theoretical and observational work on interstellar dust destruction in shock waves led to an estimated dust lifetime much shorter than the assumed dust formation timescale from AGB stars. Although SNe are believed to be efficient interstellar dust destroyers, there is increasing observational evidence today for the formation of non-negligible quantities of dust grains in the ejecta of SNe. Given the relatively short timescale between the explosion of two SNe, this would lead to an effectively shorter timescale for dust formation. In this work we present a new code called GRASH\_Rev that treats dust processing in a supernova explosion. This code couples all the dust processing included in the GRASH\_EX code (Bocchio et al. 2014) with the dynamics and structure of the SN as modelled by Bianchi \& Schneider (2007, BS07) but extending it to include the full dynamics of dust grains within the ejecta and in the surrounding ISM. \vspace{-.2cm} | Given the rather large uncertainties on measurements of dust masses, our models appear to reproduce the dust masses in the four SNe well. The average effective dust yield is estimated to be (1.55 $\pm$ 1.48) 10$^{-2}$ M$_{\odot}$. When compared to dust destruction efficiencies in SN-driven interstellar shocks that were recently estimated by theoretical models (Bocchio et al. 2014; Slavin et al. 2015) and observations (Laki\'cevi\'c et al. 2015), this implies that SNe may be net dust destroyers, pointing to grain growth in the ISM as the dominant dust enrichment process both in local galaxies and at high redshifts. \vspace{-.2cm} \small % | 16 | 9 | 1609.02327 |
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1609 | 1609.02920_arXiv.txt | \\ We present the first estimate of the Ly$\alpha$ luminosity function using blind spectroscopy from the {\emph{Multi Unit Spectroscopic Explorer}}, MUSE, in the {\emph{Hubble Deep Field South}}. Using automatic source-detection software, we assemble a homogeneously-detected sample of 59 Ly$\alpha$ emitters covering a flux range of $-18.0 < {\rm{log_{10}}}\; (F) < -16.3\; ({\rm{erg\; s}}^{-1} {\rm{cm}}^{-2}$), corresponding to luminosities of $41.4 < {\rm{log_{10}}}\; (L) < 42.8\; ({\rm{erg\; s}}^{-1})$. As recent studies have shown, Ly$\alpha$ fluxes can be underestimated by a factor of two or more via traditional methods, and so we undertake a careful assessment of each object's Ly$\alpha$ flux using a curve-of-growth analysis to account for extended emission. We describe our self-consistent method for determining the completeness of the sample, and present an estimate of the global Ly$\alpha$ luminosity function between redshifts $2.91 < z < 6.64$ using the $1/V_{max}$ estimator. We find the luminosity function is higher than many number densities reported in the literature by a factor of $2-3$, although our result is consistent at the $1\sigma$ level with most of these studies. Our observed luminosity function is also in good agreement with predictions from semi-analytic models, and shows no evidence for strong evolution between the high- and low-redshift halves of the data. We demonstrate that one's approach to Ly$\alpha$ flux estimation does alter the observed luminosity function, and caution that accurate flux assessments will be crucial in measurements of the faint end slope. This is a pilot study for the Ly$\alpha$ luminosity function in the MUSE deep-fields, to be built on with data from the {\mbox{{\emph{Hubble Ultra Deep Field}}}} which will increase the size of our sample by almost a factor of $10$. % | \label{intro} The Ly$\alpha$ emission line is one of the most powerful probes of the early Universe, giving us insight into the very early stages of galaxy formation. Galaxies detected via their Ly$\alpha$-emission {\mbox{(LAEs; \citealt{Cowie&Hu1998})}} offer us a means to study high-redshift star-forming galaxies, even with continuum magnitudes too faint to be observed using current technology. These low-mass objects form the building blocks of {\mbox{$L$$^*$ galaxies}} in the local Universe ({\mbox{\citealt{Dayal2012}}}, \citealt{Garel2016}), meanwhile theoretical models suggest they may also play a significant role in driving cosmic reionisation e.g. \cite{Gronke2015a}, \cite{Dijkstra2016}, \cite{Santos2016}. Although Ly$\alpha$ physics is complex (e.g. \citealt{Verhamme2006}, \citealt{Gronke2015b}), we can begin to understand the physical processes underway at these epochs by measuring the luminosity function of LAEs -- a fundamental statistic of the population. The luminosity function tells us about the relative abundance of different luminosity objects in the overall distribution (e.g. see \citealt{Johnston2011}), and ultimately for deep enough samples, measurement of the faint-end slope will tell us if LAEs are numerous enough to be the primary sources of reionisation (see \citealt{Dressler2015} for a discussion). The largest samples of LAEs to date come from narrow-band selection, whereby a narrow filter (typically $< 100$ \AA), is used in combination with broad-band photometry to detect emission-line galaxies in distinct redshift ``slices''. This approach is very efficient, having led to samples of thousands of H$\alpha$, H$\beta$, [O{\sc{iii}}] and [O{\sc{ii}}] emitters out to $z\approx 2.0$ (\citealt{Sobral2009}, \citealt{Sobral2013}, \citealt{Drake2013}, \citealt{Drake2015}) as well as LAEs at $z > 3.0$ (\citealt{Rhoads2000}, \citealt{Ouchi2003}, \citealt{Hu2004}, \citealt{Ouchi2008}, \citealt{Yamada2012}, \citealt{Matthee2015}, \citealt{Konno2015}, \citealt{Santos2016}). These relatively shallow surveys have provided increasingly robust estimates of the Ly$\alpha$ luminosity function down to luminosities of log$_{10}$ $L\approx 42.0\; {\rm{erg\, s}}^{-1}$, in the redshift interval $\approx 2.0 < z < 7.0$. Typically these studies estimate values of the characteristic number density and luminosity of the sample, although the faint-end slope remains unconstrained. Spectroscopic studies provide an alternative approach, allowing the identification of LAEs without any need for ancillary data, but typically surveying far smaller volumes. In addition to targetted spectroscopy, one can place long-slit spectrographs blindly on sky, but the results often suffer from severe slit-losses and a complicated selection function. (See also survey results from low-resolution slitless spectroscopy \citealt{Kurk2004}, \citealt{Deharveng2008} and IFU studies \citealt{vanBreukelen2005}, \citealt{Blanc2011}). In recent years, spectroscopic surveys have begun to push Ly$\alpha$ samples to lower flux limits than ever before, complementing wide, shallow, studies with very deep integrations. The two deepest such surveys to date come from \cite{Rauch2008} and \cite{Cassata2011} reaching 1 dex deeper than their narrow-band counterparts. \cite{Rauch2008} used a 92 hour long-slit exposure with the ESO VLT FORS2 instrument, detecting single-line emitters of just a few $\times 10^{-18} {\rm{erg\,s}}^{-1} {\rm{cm}}^{-2}$ corresponding to Ly$\alpha$ luminosities of $\approx 8 \times 10^{40} {\rm{erg\,s^{-1}}}$ for LAEs in the range $2.67 < z < 3.75$. The authors note however that their luminosities could be underestimated by factors of $2 –- 5$ due to slit losses, and the identification of many of their single-line emitters is somewhat uncertain. Another notable study came from the VIMOS-VLT Deep Survey (VVDS; \citealt{Cassata2011}) finding $217$ LAEs with secure spectroscopic redshifts between $2.00 < z < 6.62$, and fluxes reaching as low as $F = 1.5 \times 10^{−18} {\rm{erg\, s}}^{-1} {\rm{cm}}^{-2}$. The detections came from a combination of targetted and serendipitous spectroscopy however, and again resulted in a complex selection function and slit losses. Nevertheless, the number of emitters in their sample allowed the authors to split the data into three redshift bins, to look for any sign of evolution in the observed luminosity function. They ultimately found no evidence in support of evolution, consistent with the previous results of \cite{vanBreukelen2005}, \cite{Shimasaku2006}, and \cite{Ouchi2008}. Finally, at the highest redshifts, the first robust constraints on the faint end of the Ly$\alpha$ luminosity function came from \cite{Dressler2015}. They found a very steep value of the faint-end slope at $z = 5.7$, using targets selected via ``blind long-slit spectroscopy'', further reinforcing the significance of intrinsically faint LAEs in the early Universe (see also \citealt{Dressler2011} and \citealt{Henry2012}). The low--luminosity LAE population is now at the forefront of research, meaning that the accurate recovery of total LAE fluxes is of high priority for upcoming work. Indeed, some studies have already suggested that all LAEs exhibit extended, low-surface-brightness Ly$\alpha$ emission coming from the surrounding circum-galactic medium. The detection of this emission is difficult, and requires very sensitive measurements indeed. \cite{Momose2014} built on the work of \cite{Matsuda2012} by stacking LAE detections in $5$ redshift slices between $\approx 2.2 < z < 6.6$ resulting in the detection of extended Ly$\alpha$ emission around normal star-forming galaxies across this entire epoch. They found typical exponential scale lengths of $\approx5 -– 10$ kpc, but the emission was not detectable around any individual galaxy (see also \citealt{Yuma2013} for indivisual detections of metal-line blobs at lower z). The Multi Unit Spectroscopic Explorer (MUSE; \citealt{Bacon2010}) on the Very Large Telescope (VLT) allows us to carry out blind spectroscopic selection of LAEs between redshifts $\approx 3.0 < z < 6.5$ with a homogeneous selection function. The efficiency of this approach to detect line emission allows us to use MUSE as a detection machine for the kind of star-forming galaxies we wish to trace, also enabling an accurate assessment of total Ly$\alpha$ fluxes. \cite{Bacon2015}, hereafter B15, presented a blind-spectroscopic analysis of the {\mbox{{\emph{Hubble Deep Field South}} (HDFS)}}, and the resultant catalogue showcased the detection power of MUSE. Indeed, B15 presented several galaxies detected via their line emission alone, that were otherwise undetectable in the deep broad-band HST imaging (I$_{814} > 29$ mag AB). Additionally, MUSE is able to overcome the effects of slit-loss that have so far hampered Ly$\alpha$ flux estimates from long-slit specroscopy, allowing us to perform a careful evaluation of the total Ly$\alpha$ flux from each galaxy. For instance, \cite{Wisotzki2016} used a curve-of-growth analysis on $26$ isolated halos in the B15 catalogue, and presented the first ever detections of extended Ly$\alpha$ emission around individual, high-redshift, star-forming galaxies. The objects presented were in the flux range $ 4.5 \times 10^{-18} {\rm{erg\, s}}^{-1} {\rm{cm}}^{-2}$ up to $3 \times 10^{-17} {\rm{erg\, s}}^{-1} {\rm{cm}}^{-2}$ across the redshift interval $2.96 < z < 5.71$, and halos were detected around $21$ of these objects. The omission of this low surface brightness contribution to the total Ly$\alpha$ flux has potentially led to a systematic underestimation of Ly$\alpha$ fluxes in the literature, and lends support to the importance of a re-assessment of the Ly$\alpha$ luminosity function. In this paper we present a pilot study for the LAE luminosity function using blind spectroscopy in the 1 square arcminute HDFS field. We use automatic detection software to present a homogeneously selected sample of 59 LAEs and estimate Ly$\alpha$ fluxes via a curve-of-growth analysis to account for extended Ly$\alpha$ emission. We have developed and implemented a self-consistent method to determine the completeness of our sample, allowing us to compute a global Ly$\alpha$ luminosity function using the $1/V_{max}$ estimator. The outline of this paper is as follows. In Section \ref{sect:data} we present our observations from MUSE and outline our method of catalogue construction and sample selection. In Section \ref{sect:flux} we describe our approach to estimating the Ly$\alpha$ flux, and in Section \ref{sect:compl} we present and discuss our completeness estimates for the sample. In Section \ref{sect:results} we present our estimation of the LAE luminosity function between $2.91 < z < 6.64$, and discuss our results in the context of observational literature as well as in comparison to the semi-analytic model of \cite{Garel2015}. In Section \ref{sect:disc} we examine the effect of using different flux estimates for LAEs and look for evolution over the redshift range of our observed luminosity function. Finally, we summarise our results in {\mbox{Section \ref{sect:concl}}}. \begin{figure*} \begin{center} \includegraphics[width=0.48\textwidth]{Fig1a.pdf} \includegraphics[width=0.48\textwidth]{Fig1b.pdf} \caption{A comparision between numbers of LAEs presented in Bacon et al. (2015) and detections recovered using the detection software {\sc{muselet}}. In the left-hand panel we show the redshift distribution of our detections overlaid on the redshift distribution of the B15 LAEs. This demonstrates an even recovery rate across the entire redshift range i.e. no redshift bias in our method of detection. In the right-hand panel we use the published flux estimates of B15 to show the distribution of fluxes recovered by {\sc{muselet}} vs the distribution for B15 LAEs. We successfully recover the majority of bright LAEs before incompleteness becomes more apparent below {\mbox{log$_{10}$ $F$ Ly$\alpha$ (B15) $= -17.32$}}. Bright LAEs which are not recovered by {\sc{muselet}} lie in the small parts of the cube with fewer than 50 percent of the final exposure time. The average sample completeness is overlaid (dashed and dotted lines) and its derivation is described in \mbox{Section \ref{sect:compl}}.} \label{fig:no_counts} \end{center} \end{figure*} The total co-moving volume between {\mbox{$2.91 < z < 6.64$}} equates to {\mbox{$10351.6$ Mpc$^3$}}. As parts of the cube are excluded from the search however (see Section \ref{subsect:det}) the total co-moving survey volume is reduced to {\mbox{$10144.57$ Mpc$^3$}}. Throughout this paper we assume a $\Lambda$CDM cosmology, \mbox{$H_0 = 70.0$} kms$^{-1}$ Mpc$^{-1}$, $\Omega_{m} =0.3$, $\Omega_\Lambda = 0.7$. | \label{sect:concl} In summary, we have presented a homogeneous, automatically-detected sample of $59$ LAEs in the HDFS using blind spectroscopy from MUSE. We validate the Ly$\alpha$ line through a careful matching to the deeper (heterogeneously constructed) catalogue of B15. We have shown that the method of Ly$\alpha$ flux estimation can significantly alter measured Ly$\alpha$ fluxes and investigated the effect this has on the luminosity function. We have designed a procedure self-consistent with our detection software to determine our selection function through recovery of fake point-source line emitters from deep MUSE datacubes, and compute a global Ly$\alpha$ luminosity function using a curve-of-growth analysis of the Ly$\alpha$ flux, and implementing the $1/V_{max}$ estimator. We compare our results to literature studies, and semi-analytic model predictions from \cite{Garel2015}, before finally examining the dataset for signs of evolution in the observed luminosity function.\\ Our main conclusions can be broadly summarised as follows: \begin{itemize} \renewcommand{\labelitemi}{$\bullet$} \item We automatically detect 59 LAEs in the HDFS across a flux range of $\approx -18.0 < {\rm{log\; F}} < -16.3\; ({\rm{erg s}}^{-1} {\rm{cm}}^{-2})$ using homogeneous and robust selection criteria, validating each LAE by matching to the deep catalogue of B15. \item Our global luminosity function between $2.91 < z < 6.64$ sits higher by a factor of $2-3$ than the literature in our most well-constrained bins, although $1\sigma$ errorbars overlap with the data of several literature studies at the same luminosity. \item The small drop in number density between our penultimate and faintest luminosity bin is likely to be entirely due to the limitations of our method; namely the effect of incomplete bins on the $1/V_{max}$ estimator, and our idealised completeness assessment where LAEs are treated as point sources. We will investigate this in Drake et al., (in prep) using the MUSE HUDF mosaic sample. \item Our luminosity function is in good agreement with the semi-analytical model of \cite{Garel2015} with the exception of our bin at log$_{10} \,L = 42.58$. The bin is once again a factor $\approx 3$ higher than the predictions, with a $1\sigma$ Poissonian error bar that just touches the $1\sigma$ error on the model predictions. \item Method of Ly$\alpha$ flux estimation plays a role in the determination of the Ly$\alpha$ luminosity function and becomes most significant when measuring the faint end slope. Care should be taken here as studies start to probe further into the low-luminosity LAE population. \item When splitting our data at the central redshift and comparing the two halves of the data we see no evidence for strong evolution in the observed Ly$\alpha$ luminosity function across the redshift range $2.91 < z < 6.64$. This is entirely consistent with results in the literature. \end{itemize} Our pilot study demonstrates the efficiency of MUSE as a detection machine for emission-line galaxies, and strongly motivates our analysis of the HUDF $9 \times 9$ square arcminute mosaic. The conservative nature of our selection process means that the objects presented here represent a robustly-selected sub-sample of the galaxies MUSE will ultimately detect and identify, and this is very encouraging for the potential of additional blank-field datasets from MUSE. | 16 | 9 | 1609.02920 |
1609 | 1609.01732_arXiv.txt | We present the improved visibility based Tapered Gridded Estimator (TGE) for the power spectrum of the diffuse sky signal. The visibilities are gridded to reduce the computation, and tapered through a convolution to suppress the contribution from the outer regions of the telescope's field of view. The TGE also internally estimates the noise bias, and subtracts this out to give an unbiased estimate of the power spectrum. An earlier version of the 2D TGE for the angular power spectrum $C_{\ell}$ is improved and then extended to obtain the 3D TGE for the power spectrum $P(\k)$ of the 21-cm brightness temperature fluctuations. Analytic formulas are also presented for predicting the variance of the binned power spectrum. The estimator and its variance predictions are validated using simulations of $150 \, {\rm MHz}$ GMRT observations. We find that the estimator accurately recovers the input model for the 1D Spherical Power Spectrum $P(k)$ and the 2D Cylindrical Power Spectrum $P(k_\perp,k_\parallel)$, and the predicted variance is also in reasonably good agreement with the simulations. | Observations of the redshifted neutral hydrogen (HI) 21-cm radiation hold the potential of probing a wide range of cosmological and astrophysical phenomena over a large redshift range $0 < z \lsim 200$ \citep{BA5,furla06,morales10,prichard12,mellema13}. There now are several ongoing experiments such as the Donald C. Backer Precision Array to Probe the Epoch of Reionization (PAPER{\footnote{http://astro.berkeley.edu/dbacker/eor}}, \citealt{parsons10}), the Low Frequency Array (LOFAR{\footnote{http://www.lofar.org/}}, \citealt{haarlem,yata13}) and the Murchison Wide-field Array (MWA{\footnote{http://www.mwatelescope.org}} \citealt{bowman13,tingay13}) which aim to measure the power spectrum of the 21-cm radiation from the Epoch of Reionization (EoR, $6 \lsim z \lsim 13$). Future telescopes like the Square Kilometer Array (SKA1 LOW{\footnote{http://www.skatelescope.org/}}, \citealt{koopmans15}) and the Hydrogen Epoch of Reionization Array (HERA{\footnote{http://reionization.org/}}, \citealt{neben16}) are planned to achieve even higher sensitivity for measuring the EoR 21-cm power spectrum. Several other upcoming experiments like the Ooty Wide Field Array (OWFA; \citealt{prasad,ali14}), the Canadian Hydrogen Intensity Mapping Experiment (CHIME{\footnote{http://chime.phas.ubc.ca/}; \citealt{bandura}), the Baryon Acoustic Oscillation Broadband, Broad Beam Array (BAOBAB{\footnote{http://bao.berkeley.edu/}; \citealt{pober13a}) and the Square Kilometre Array (SKA1 MID; \citealt{bull15}) target the post-Reionization 21-cm signal ($0 < z \lsim 6$). Despite the sensitive new instruments, the main challenge still arises from the fact that the cosmological 21-cm signal is buried in astrophysical foregrounds which are $4-5$ orders of magnitude brighter \citep{shaver99,dmat1,santos05, ali,paciga11,ghosh1,ghosh2}. A large variety of techniques have been proposed to overcome this problem and estimate the 21-cm power spectrum. The different approaches may be broadly divided into two classes (1.) Foreground Removal, and (2.) Foreground Avoidance. The idea in Foreground Removal is to model the foregrounds and subtract these out either directly from the data (eg. \citealt{ali}) or from the power spectrum estimator after correlating the data (eg. \citealt{ghosh1,ghosh2}). Foreground Removal is a topic of intense current research \citep{jelic08,bowman09,paciga11,chapman12,parsons12,liu12,trott1,pober13,paciga13, parsons14,trott16}. Various studies (eg. \citealt{adatta10}) show that the foreground contribution to the Cylindrical Power Spectrum $P(k_{\perp},k_{\parallel})$ is expected to be restricted within a wedge in the two dimensional (2D) $(k_{\perp},k_{\parallel})$ plane. The idea in Foreground Avoidance is to avoid the Fourier modes within the foreground wedge and only use the uncontaminated modes outside the wedge to estimate the 21-cm power spectrum \citep{vedantham12,thyag13,pober14,liu14a,liu14b,dillon14,dillon15,zali15}. In a recent paper \citet{jacob16} have compared several power spectrum estimation techniques in the context of MWA. Point sources dominate the low frequency sky at the angular scales $\le 4^{\circ}$ \citep{ali} which are relevant for EoR 21-cm power spectrum with the telescopes like the GMRT, LOFAR and the upcoming SKA. It is difficult to model and subtract the point sources which are located at the periphery of the telescope's field of view (FoV). The antenna response deviates from circular symmetry, and is highly frequency and time dependent at the outer parts of the telescope's FoV. The calibration also differs from the phase center due to ionospheric fluctuations. The residual point sources located far away from the phase centre cause the signal to oscillates along the frequency direction \citep{ghosh1,ghosh2}. This poses a severe problem for foreground removal techniques which assume a smooth behavior of the signal along the frequency direction. Equivalently, these distant point sources reduce the EoR window by increasing the area under the foreground wedge in $(k_{\perp},k_{\parallel})$ space \citep{thyag15}. In a recent paper, \citet{pober16} showed that correctly modelling and subtracting the distant point sources are important for detecting the redshifted 21-cm signal. Point source subtraction is also important for measuring the angular power spectrum of the diffuse Galactic synchrotron radiation \citep{bernardi09,ghosh150,iacobelli13}. Apart from being an important foreground component for the EoR 21-cm signal, this is also interesting in its own right. It is possible to suppress the contribution from the outer parts of the telescope's FoV by tapering the sky response through a suitably chosen window function. \citet{ghosh2} have analyzed $610 {\rm MHz}$ GMRT data to show that it is possible to implement the tapering by convolving the observed visibilities with the Fourier transform of the window function. It is found that this reduces the amplitude of the oscillation along the frequency direction. Our earlier work \citet{samir14} (hereafter Paper I) has introduced the Tapered Gridded Estimator (TGE) which places the findings of \citet{ghosh2} on a sound theoretical footing. Considering observations at a single frequency, the TGE estimates the angular power spectrum $C_{\ell}$ of the 2D sky signal directly from the measured visibilities while simultaneously tapering the sky response. As a test-bed for the TGE, Paper I considers a situation where the point sources have been identified and subtracted out so that the residual visibilities are dominated by the Galactic synchrotron radiation. This has been used to investigate how well the TGE is able to recover the angular power spectrum of the input model used to simulate the Galactic synchrotron emission at $150 \, {\rm MHz}$. While most of the analysis was for the GMRT, simulations for LOFAR were also considered. These investigations show that the TGE is able to recover the input model $C_{\ell}^M$ to a high level of precision provided the baselines have a uniform $uv$ coverage. For the GMRT, which has a patchy $uv$ coverage, the $C_{\ell}$ is somewhat overestimated using TGE though the excess is largely within the $1\sigma$ errors. This deviation is found to be reduced in a situation with a more uniform and denser baseline distribution , like LOFAR. Paper I also analyzes the effects of gain errors and the $w$-term. In a recent paper \citet{samir16} (hereafter Paper II) we have further developed the simulations of Paper~I to include the point sources. We have used conventional radio astronomical techniques to model and subtract the point sources from the central region of the primary beam. As detailed in Paper~II, it is difficult to do the same for the sources which are far away from the phase center, and these persist as residuals in the visibility data. We find that these residual point sources dominate the $C_{\ell}$ estimated at large baselines. We also show that it is possible to suppress the contribution from these residual sources located at the periphery of the FoV by using TGE with a suitably chosen window function. Removing the noise bias is an important issue for any power spectrum estimator. As demonstrated in Paper II, the TGE internally estimates the actual noise bias from the data and subtracts this out to give an unbiased estimate of the power spectrum. In the present work we report the progress on two counts. First, our earlier implementation of the TGE assumed a uniform and dense baseline $uv$ coverage to calculate the normalization coefficient which relates visibility correlations to the estimated angular power spectrum $C_{\ell}$. We, however, found (Paper I) that this leads to an overestimate of $C_{\ell}$ for instruments like the GMRT which have a sparse and patchy $uv$ coverage. In Section 2 of this paper we present an improved TGE which overcomes this problem by using simulations to estimate the normalization coefficient. Second, the entire analysis of Papers I and II has been restricted to observations at a single frequency wherein the relevant issue is to quantify the 2D angular fluctuations of the sky signal. This, however, is inadequate for the three dimensional (3D) redshifted HI 21-cm signal where it is necessary to also simultaneously quantify the fluctuations along the frequency direction. In Section 3 of this paper we have generalized the TGE to quantify the 3D 21-cm signal and estimate the spatial power spectrum of the 21-cm brightness temperature fluctuations $P(\k)$. We discuss two different binning schemes which respectively yield the spherically-averaged (1D) power spectrum $P(k)$ and the cylindrically-averaged (2D) power spectrum $P(\kpm,\kp)$, and present theoretical expressions for predicting the expected variance. We have validated the estimator and its variance predictions using simulations which are described in Section 4 and for which the results are presented in Section 5. Sections 6 presents the summary and conclusions. In this paper, we have used cosmological parameters from the (Planck + WMAP) best-fit $\Lambda$CDM cosmology (\citealt{ade15}). | Quantifying the statistical properties of the diffuse sky signal directly from the visibilities measured in low frequency radio-interferometric observation is an important issue. In this paper we present a statistical estimator, namely the Tapered Gridded Estimator (TGE), which has been developed for this purpose. The measured visibilities are here gridded in the $uv$ plane to reduce the complexity of the computation. The contribution from the discrete sources in the periphery of the telescope's FoV, particularly the sidelobes, pose a problem for power spectrum estimation. The TGE suppresses the contribution from the outer regions by tapering the sky response through a suitably chosen window function. The TGE also internally estimates the noise bias from the input data, and subtracts this out to give an unbiased estimate of the power spectrum. In addition to the mathematical formalism for the estimator and its variance, we also present simulations of $150 \, {\rm MHz}$ GMRT observations which are used to validate the estimator. We have first considered a situation where we have observation at a single frequency for which the 2D TGE provides an estimate of the angular power spectrum $C_{\ell}$. The work here presents an improvement over an earlier version of the 2D TGE presented in Paper I. This is important in the context of the diffuse Galactic synchrotron emission which is one of the major foregrounds for the cosmological 21-cm signal. Apart from this, the diffuse Galactic synchrotron emission is a probe of the cosmic ray electrons and the magnetic fields in the ISM of our own Galaxy, and this is an important study in its own right. It is necessary to also include the frequency variation of the sky signal in order to quantify the cosmological 21-cm signal. Here the 3D TGE provides an estimate of $P(\k)$ the power spectrum of the 21-cm brightness temperature fluctuations. We have considered two different binning schemes which provide the 1D Spherical Power Spectrum $P(k)$ and the 2D Cylindrical Power Spectrum $P(k_\perp,k_\parallel)$ respectively. In all cases, we find that the TGE is able to accurately recover the input model used for the simulations. The analytic predictions for the variance are also found to be in reasonably good agreement with the simulations in most situations. Foregrounds are possibly the biggest challenge for detecting the cosmological 21-cm power spectrum. Various studies (eg. \citealt{adatta10}) show that the foreground contribution to the Cylindrical Power Spectrum $P(k_{\perp},k_{\parallel})$ is expected to be restricted within a wedge in the $(k_{\perp},k_{\parallel})$ plane. The extent of this ``foreground wedge'' is determined by the angular extent of the telescope's FoV. In principle, it is possible to limit the extent of the foreground wedge by tapering the telescope's FoV. In the context of estimating the angular power spectrum $C_{\ell}$, our earlier work (Paper II) has demonstrated that the 2D TGE is able to suppress the contribution from the outer parts and the sidelobes of the telescope's beam pattern. We have not explicitly considered the foregrounds in our analysis of the 3D TGE presented in this paper. We however expect the 3D TGE to suppress the contribution from the outer parts and the sidelobes of the telescopes beam pattern while estimating the power spectrum $P(k_{\perp},k_{\parallel})$, thereby reducing the area in the $(k_{\perp},k_{\parallel})$ plane under the foreground wedge. The 3D TGE holds the promise of allowing us to reduce the extent of the foreground wedge by tapering the sky response. It is, however, necessary to note that this comes at a cost which we now discuss. First, we lose information at the largest angular scales due to the reduced FoV. This restricts the smallest $k$ value at which it is possible to estimate the power spectrum. Second, the reduced FoV results in a larger cosmic variance for the smaller angular modes which are within the tapered FoV. The actual value of the tapering parameter $f$ that would be used to estimate $P(k_{\perp},k_{\parallel})$ will possibly be determined by optimising between the cosmic variance and the foreground contribution. A possible strategy would be to use different values of $f$ for different bins in the $(k_{\perp},k_{\parallel})$ plane. It is also necessary to note that the effectiveness of the tapering proposed here depends on the actual baseline distribution, and a reasonably dense $uv$ coverage is required for a proper implementation of the TGE. We propose to include foregrounds in the simulations and address these issues in future work. We also plan to apply this estimator to $150 \, {\rm MHz}$ GMRT data in future. | 16 | 9 | 1609.01732 |
1609 | 1609.08339_arXiv.txt | {Galaxy clusters undergo mergers that can generate extended radio sources called radio relics. Radio relics are the consequence of merger-induced shocks that propagate in the intra cluster medium (ICM).} {In this paper we analyse the radio, optical and X-ray data from a candidate galaxy cluster that has been selected from the radio emission coming from a candidate radio relic detected in NRAO VLA Sky Survey (NVSS). Our aim is to clarify the nature of this source and prove that under certain conditions radio emission from radio relics can be used to trace relatively low-mass galaxy clusters.} {We observed the candidate galaxy cluster with the Giant Meterwave Radio Telescope (GMRT) at three different frequencies. These datasets have been analysed together with archival data from ROSAT in the X-ray and with archival data from the Gamma-Ray Burst Optical/Near-Infrared Detector (GROND) telescope in four different optical bands.} {We confirm the presence of a 1~Mpc long radio relic located in the outskirts of a previously unknown galaxy cluster. We confirm the presence of the galaxy cluster through dedicated optical observations and using archival X-ray data. Due to its proximity and similar redshift to a known Abell cluster, we named it \target{}. The galaxy cluster is amongst the least massive clusters known to host a radio relic.} {We showed that radio relics can be effectively used to trace a subset of relatively low-mass galaxy clusters that might have gone undetected in X-ray or Sunyaev-Zel'dovich (SZ) surveys. This technique might be used in future deep, low-frequency surveys such as those carried on by the Low Frequency Array (LOFAR), the Upgraded GMRT (uGMRT) and, ultimately, the Square Kilometre Array (SKA).} | \label{sec:introduction} Cosmic structure forms hierarchically, in a bottom-up scenario, with smaller structures merging to form bigger ones. On the largest scales, clusters of galaxies merge releasing energies on the order of $10^{64}$ erg on time scales of 1--2 Gyr \citep[e.g.][]{Hoeft2007}. During these events, large-scale shock waves with moderate Mach numbers ($\lesssim 4$) are created in the intra cluster medium (ICM). Shocks are collision-less features whose interactions in the hot plasma are mediated by electromagnetic fields. They act on the ICM accelerating a fraction of the thermal distribution of particles transforming them into non-thermal populations of cosmic rays (CRs) as initially proposed by \cite{Ensslin1998}. The mechanism initially thought to be able to accelerate CRs was diffusive shock acceleration \citep[DSA, e.g.][]{Drury1983}, however the efficiency of acceleration via the DSA mechanism in weak shocks is thought to be very low \citep{Malkov1995}. A possibility to overcome this problem is that shocks re-accelerate populations of already mildly relativistic electrons, instead of directly accelerating them from the thermal pool \citep[e.g.][]{Markevitch2005,Kang2011,Kang2012,Kang2016}. Some observational pieces of evidence for this scenario are available \citep{Bonafede2014a,Shimwell2015} but are not conclusive. Recently, particle-in-cell (PIC) simulations provided new insight into proton and electron acceleration in weak shocks in galaxy clusters. They indicate that shock drift acceleration (SDA) might be efficient in injecting particles into the DSA process \citep{Caprioli2014,Guo2014}. Regardless of the acceleration method, the empirical evidence of the presence of large-scale shocks in the ICM is given by radio relics\footnote{We use the commonly accepted name ``radio relic'' while other authors prefer the name radio gischt.} \citep[for a review see e.g.][]{Feretti2012,Brunetti2014}. These are extended radio sources in the periphery of galaxy clusters. Their spectrum is rather steep ($\alpha \lesssim -1$)\footnote{$S_\nu \propto \nu^\alpha$; where $S_\nu$ is the flux density at frequency $\nu$.} and their emission is strongly polarised, with an ordered magnetic field usually aligned with the relic extension \citep[e.g.][]{vanWeeren2010a,deGasperin2015a}. \begin{figure*}[!ht] \centering \subfloat[GMRT 148 MHz]{\includegraphics[width=.49\textwidth]{radio150.pdf}\label{fig:150}} \subfloat[GMRT 323 MHz]{\includegraphics[width=.49\textwidth]{radio325.pdf}\label{fig:325}}\\ \subfloat[GMRT 608 MHz]{\includegraphics[width=.49\textwidth]{radio610.pdf}\label{fig:610}} \subfloat[Spectral index (323-608 MHz)]{\includegraphics[width=.49\textwidth]{spidx.pdf}\label{fig:spidx}} \caption{Radio image of the radio relic found in \target{} at 148, 323 and 608 MHz obtained with the GMRT. Contours are at $(1,2,3,4,5,7,9,12,15,20,25) \times 3\sigma$ with $\sigma=1400$, 67 and 44~\mujybeam{}, respectively. The resolutions of the three images are: \beam{32}{15} (148 MHz), \beam{18}{8} (323 MHz), and \beam{9}{5} (608 MHz). In the 608 MHz image the low-resolution (\beam{19}{15}) contours at $3\sigma$ ($\sigma=81$~\mujybeam) are over-plotted in red. The last panel shows the spectral index map obtained as described in the text. Over-plotted in green are the regions used to calculate the spectral index in Fig.~\ref{fig:spidx-plot}.}\label{fig:radio} \end{figure*} Thanks to the new SZ measurements obtained with the Planck satellite and the South Pole Telescope, many new galaxy clusters were recently identified \citep{PlanckCollaboration2015,Bleem2015}. However, SZ measurements are biased towards detecting massive clusters. The other traditional method to locate galaxy clusters is through their Bremsstrahlung radiation from the hot ICM \citep[e.g.][]{Boehringer1995,Ebeling1998a}, which is also biased towards massive and cool-core systems \citep{Eckert2011}. Given the correlation between radio-relic power and hosting-cluster mass \citep{deGasperin2014c}, most of the radio relics discovered so far are located in quite massive galaxy clusters ($M_{500} > {\rm few} \times 10^{14}$~\Msun). Radio surveys can be used to detect previously unknown massive galaxy clusters by tracing powerful radio relics \citep{VanWeeren2012e}. \cite{Macario2014} also discovered a galaxy cluster thanks to the presence of extended radio emission, however the radio source could not be firmly classified. Less massive clusters has been also discovered by tracing diffuse radio emission in surveys \citep{Brown2011}. The use of radio relics to trace low-mass galaxy clusters is a promising technique in the light of new high-sensitivity, large field-of-view telescopes such as LOFAR, uGMRT and, ultimately, SKA. In this paper we report the discovery of a $\gtrsim 1$~Mpc long radio relic in the periphery of a previously unknown galaxy cluster. The presence of the galaxy cluster has been confirmed by optical and X-ray observations. Throughout the paper we adopt a fiducial $\Lambda$CDM cosmology with $H_0 = 70\rm\ km\ s^{-1}\ Mpc^{-1}$, $\Omega_m = 0.3$ and $\Omega_\Lambda = 0.7$. At the redshift of the target ($z\simeq0.2$) 1\arcsec = 3.3 kpc. Unless otherwise specified errors are at $1\sigma$. | \label{sec:conclusions} We reported the discovery of a radio relic located in the outskirts of a previously unknown galaxy cluster: \target{} ($z=0.20\pm0.02$). The relic nature of the source has been confirmed from morphological and spectral properties. The most interesting characteristic of this source is its high luminosity and unusually flatter spectral index which, assuming simple DSA prescription, corresponds to a rather high Mach number in the associated shock wave. The associated galaxy cluster is also peculiar for its low mass ($M_{500}=3.3\pm0.3 \times 10^{14}$~\Msun). This is the third least massive system known to date to host a radio relic. However, better X-ray data are required to firmly assess the global cluster characteristics. Furthermore, together with the extremely bright ($S_{\rm 1.4\ GHz} = 0.32 \pm 0.02$~Jy) radio relic in the ``toothbrush'' cluster \citep{VanWeeren2012e}, \target{} is the only other galaxy cluster to have been detected thanks to the presence of a radio relic. This detection proves that radio surveys can be used to locate a subset of relatively low mass merging galaxy clusters ($M_{500} \simeq {\rm few} \times 10^{14}$~\Msun) and that the occurrence of radio relics in low mass clusters may be more common than previously thought. \begin{figure} \centering \includegraphics[width=.49\textwidth]{spidx-plot.pdf} \caption{Black solid-line: radio spectral index between 323 and 608 MHz. Red dashed-line: radio flux density at 323 MHz across the source extension. The data-points correspond to the regions displayed in Fig.~\ref{fig:spidx}. Error-bars on the x-axes correspond to the region size.}\label{fig:spidx-plot} \end{figure} | 16 | 9 | 1609.08339 |
1609 | 1609.08613_arXiv.txt | We probe the higher-order clustering of the galaxies in the final data release (DR12) of the Sloan Digital Sky Survey Baryon Oscillation Spectroscopic Survey (BOSS) using the method of germ-grain Minkowski Functionals (MFs). Our sample consists of $410,615$ BOSS galaxies from the northern Galactic cap in the redshift range $0.450$--$0.595$. We show the MFs to be sensitive to contributions up to the six-point correlation function for this data set. We ensure with a custom angular mask that the results are more independent of boundary effects than in previous analyses of this type. We extract the higher-order part of the MFs and quantify the difference to the case without higher-order correlations. The resulting $\chi^{2}$ value of over $10,000$ for a modest number of degrees of freedom, $O(200)$, indicates a $100$-sigma deviation and demonstrates that we have a highly significant signal of the non-Gaussian contributions to the galaxy distribution. This statistical power can be useful in testing models with differing higher-order correlations. Comparing the galaxy data to the QPM and MultiDark-Patchy mocks, we find that the latter better describes the observed structure. From an order-by-order decomposition we expect that, for example, already a reduction of the amplitude of the MD-Patchy mock power spectrum by 5\% would remove the remaining tension. | The analysis of cosmic large-scale structure in the distribution of galaxies contributed significantly to establish the current cosmological concordance model. With the increasing degree of perfection of the standard two-point correlation analysis, the related measurements reached unprecedented precision% . However, the available and upcoming galaxy redshift survey data \citep{2013AJ....145...10D,2016AJ....151...44D,2013arXiv1308.0847L,2011arXiv1110.3193L} cannot merely be characterized by its two-point correlation properties. Already the filamentary structure that we see in those surveys indicates that the higher-order correlations play an important role. There have been different approaches to use this information for improving our knowledge about the constituents and evolution of the Universe. The most direct approach of course is to simply measure the higher-order correlations directly. This has been done for both simulated and observed data starting for early catalogs with \citet{1975ApJ...196....1P} and \citet{1978ApJ...221...19F}. Especially measurements of the three-point function have been standard for recent surveys \citep{2011ApJ...737...97M,2013MNRAS.432.2654M,2015arXiv151202231S} and provide useful complimentary information \citep{2015MNRAS.448....9S}. For higher orders, however, there are computational and conceptual challenges. On the computational side, increasing the order of the correlation function measured unleashes the curse of combinatorics. Measuring all pairs in current surveys is feasible, for all triangles there are methods to do it \citep{2015arXiv151202231S}, but beyond that it becomes increasingly hard. And then, from the conceptual side, even if measuring the full $n$-point function was possible, it would require modeling and covariances for a large number of data bins. Already for a modest number of bins % {} for each parameter of the $n$-point function this is a challenge. Therefore, it is worthwhile to consider other approaches that contain information on certain aspects of the higher-order clustering without referring to the full correlation functions. Since the beginning of large-scale structure analysis, a large number of such methods have been used, including moments of counts in cells, void probability \citep{Stratonovich1963,1979MNRAS.186..145W}, structure functions on minimal spanning trees, wavelet methods, the genus, etc. Recently also approaches that attempt to bring the higher-order information back into the two-point correlation function through non-linear transforms were used like % {} a logarithmic transformation \citep{2009ApJ...698L..90N,2011AAS...21823304N,2011ApJ...735...32W}. The method (re-)considered in this paper is using Minkowski Functionals% . They also also fall into the category of methods that provide condensed access to higher-order information. This information is encoded in the morphology of the density field. Minkowski Functionals access it by quantifying extended regions of the field through their geometrical properties like volume and surface (see Sec.~\ref{sub:Minkowski-Functionals}). For the construction of these extended regions, two different methods have been widely used in the past\footnote{See \citet{2010arXiv1006.4178A} for a recent alternative.}: In the first application of Minkowski Functionals for analyzing large-scale structure in the galaxy distribution in \citet{1994A&A...288..697M} (see \citet{1995lssu.conf..156B} for a short motivation and \citet{1996dmu..conf..281S} for a brief overview), the regions considered were defined by the union of balls around every galaxy location. We shall use the same method, referred to as germ-grain model, here and explain the details in Sec.~\ref{sub:Germ--grain}. After a couple of applications in the analysis of galaxy and cluster catalogs \citep{1996app..conf...83K,1998A&A...333....1K,2001A&A...373....1K,1997MNRAS.284...73K}, the germ-grain model has been less used in recent years. The second and most popular form of the application of Minkowski Functionals works with extended regions that are defined by iso-contours of the field under investigation. It has been used for the higher-order contributions to the dark matter overdensity field \citep{1996app..conf..251P,1997ApJ...482L...1S,1998ApJ...495L...5S,1998ApJ...508..551S,1999ApJ...526..568S,2003PASJ...55..911H,2004astro.ph..8428N,2014MNRAS.437.2488B,2013ApJS..209...19C}, for the weak lensing shear field \citep{2012PhRvD..85j3513K,2013PhRvD..88l3002P}, for local morphology of cosmic structure \citep{2014A&A...562A..87E} and for the morphology of reionization bubbles \citep{2006MNRAS.370.1329G,2016arXiv160202351Y}. Another very active field of recent application is the study of iso-temperature contour maps of the cosmic microwave background (CMB) \citep{1998MNRAS.297..355S,2000ApJ...544L..83S,2012MNRAS.425.2187H,2013MNRAS.429.2104D,2013MNRAS.428..551M,2013MNRAS.434.2830M,2014A&A...571A..23P,2014A&A...571A..25P,2016A&A...594A..16P}, where Minkowski Functionals helped to establish stringent bounds on the non-Gaussianity of the CMB. So the iso-contour Minkowski functionals are a well established tool in various domains of cosmology. However, as we shall discuss in Sec.~\ref{sub:Germ--grain}, the germ-grain model accesses different aspects of the higher-order distribution and can be better suited for the analysis of galaxy redshift surveys. We have demonstrated some of these advantages already in the analysis of the Sloan Digital Sky Survey (SDSS) Data Release seven (DR7) Luminous Red Galaxies (LRGs) in \citet{2014MNRAS.443..241W}. In the present paper, we concentrate more explicitly on the higher-order aspects of the Minkowski Functionals. To this end, we exploit the analytically known connection to the integrals over higher-order correlation functions. Using direct measurements of the two-point correlation function, we can isolate the contribution of higher-order correlation functions to the Minkowski Functionals. This allows us to quantify the evidence for non-Gaussianity of the galaxy distribution (see \citet{2012MNRAS.423.3209P} and \citet{2013MNRAS.435..531C} for related recent studies of non-Gaussianity with isodensity contour functionals). The analysis is demonstrated with the final data from the Baryon Oscillation Spectroscopic Survey (BOSS) \citep{2013AJ....145...10D} that is now publicly available within the DR12 of SDSS-III. Being the largest spectroscopic galaxy redshift data set yet observed, it provides an ideal testbed for the analysis and reveals the strong presence of higher-order correlations already at redshifts of $\approx0.55$. We proceed as follows: Sec.~\ref{sec:Minkowski-Functionals-in-the} recalls the properties of Minkowski Functionals in general and the germ-grain model in particular. We write out the connection of the functionals to higher-order correlations explicitly in \ref{sub:DensDef} and lay out in \ref{sub:Density-dependence}, how knowing the concrete form allows us to isolate the part of the functionals that is only dependent on higher-orders. We explain the subtleties of a Gaussian reference model in the discrete case in \ref{sub:Gauss-Poisson} and recall the basic properties of our Minkowski code in \ref{sub:SamplingCode}. In Sec.~\ref{sec:Data} we turn to the description of the SDSS DR12 data used. Sec.~\ref{sub:CMASS} summarizes the data and mocks we are using. In \ref{sub:Ceometry}, we address issues related to the imperfections in the data, constructing a custom angular mask that defines a high quality region. In \ref{sub:Radial} we turn to the radial distribution of the data and choose the analysis volume and the radial grid. In \ref{sub:Weights} we finally discuss the effect of the regions that can not be accounted for by our custom mask. Sec.~\ref{sec:HigherOrders} then turns to the core of the analysis. In \ref{sub:MFdensities} we present the Minkowski Functionals of the mocks and the data and explain how we transform the data to be able to isolate the higher-order only part. This part is then examined in \ref{sub:HighSubtract}. In \ref{sub:SeriesCoefficients} we demonstrate that the contributions to this higher-order part range at least up to the $6$-point function. To wrap up we discuss in \ref{sub:Voidfraction} an interesting aspects about the galaxy distribution using the base Minkowski Functional to study the volume in extreme voids in the observed field. We conclude in Sec.~\ref{sec:Conclusion}. | \label{sec:Conclusion} In this paper, we perform a new analysis that combines Minkowski Functionals with the standard two-point correlation function to measure the amount of non-Gaussianity of the density field. We find in Sec.~\ref{sub:HighSubtract} that the significance of the higher-order contribution to the MFs is of the order of $100\sigma$ (see Tab.~\ref{tab:ChiSqRadComb} and Fig.~\ref{fig:ModMinkSubtract}). This demonstrates that this part of the MFs can be well enough measured to give meaningful constraints on the higher-order properties of the galaxy distribution. Given this constraining power it is reassuring that we find from Fig.~\ref{fig:ModMinkSubtract} and Tab.~\ref{tab:ChiSqMockDat} that the MD-Patchy mocks, that were constructed to match the two- and three-point properties of the measured galaxy distribution, match fairly well to the higher-order MFs. This indicates that as far as the higher-orders accessible by the MFs are concerned, the mock galaxy catalogs are quite good. There is still some room for improvement as the combined deviation is as large as $5\sigma$ for some functionals (see Tab.~\ref{tab:ChiSqRadComb}), but this might be addressed by a $5\%$ shift in the normalization of the mock power spectrum (see Fig.~\ref{fig:ModMinkDensDep}). On the technical side, our careful boundary treatment minimizes the effects of the survey geometry on our measurement of the modified MFs (see Fig.~\ref{fig:MaskCorrection}). The corrected values lie within the error bars of a boundaryless periodic box. For the QPM mocks, we find a much larger deviation (beyond $10\sigma$) from the data than for MD-Patchy (see Tab.~\ref{tab:ChiSqMockDat}). As in this case already the two-point correlations are overnormalized in the range of scales that we are interested in (see Fig.~\ref{fig:MeanCorrFun}), a reduction of the amplitude of the initial Gaussian field could also help in this case. We show that in the the higher-order part of the MFs calculated for the SDSS DR12 CMASS galaxies contains at least information up to the integrated $6$-point function (see Tab.~\ref{tab:HigherFitChi}). This underlines the power of the MFs to bring condensed higher-order information into a simple functional form. \subsubsection*{Acknowledgements} AW thanks Antonio Cuesta for providing support with the QPM correlation functions. We thank Graziano Rossi and Thomas Buchert for valuable comments on the manuscript. The work of AW was supported by the German research organization DFG, Grant No. WI 4501/1-1. DJE is supported by U.S. Department of Energy grant DE-SC0013718 and as a Simons Foundation Investigator. 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 Office of Science. 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, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, 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, Max Planck Institute for Extraterrestrial Physics, 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. | 16 | 9 | 1609.08613 |
1609 | 1609.06324_arXiv.txt | \baselineskip = 11pt \leftskip = 0.65in \rightskip = 0.65in \parindent=1pc | 16 | 9 | 1609.06324 |
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1609 | 1609.01168_arXiv.txt | In the Milky Way, the thick disk can be defined using individual stellar abundances, kinematics, or age; or geometrically, as stars high above the mid-plane. In nearby galaxies, where only a geometric definition can be used, thick disks appear to have large radial scale-lengths, and their red colors suggest that they are uniformly old. The Milky Way's geometrically thick disk is also radially extended, but it is far from chemically uniform: \al-enhanced stars are confined within the inner Galaxy. In simulated galaxies, where old stars are centrally concentrated, geometrically thick disks are radially extended, too. Younger stellar populations flare in the simulated disks' outer regions, bringing those stars high above the mid-plane. The resulting geometrically thick disks therefore show a radial age gradient, from old in their central regions to younger in their outskirts. Based on our age estimates for a large sample of giant stars in the APOGEE survey, we can now test this scenario for the Milky Way. We find that the geometrically-defined thick disk in the Milky Way has indeed a strong radial age gradient: the median age for red clump stars goes from $\sim 9$ Gyr in the inner disk to 5 Gyr in the outer disk. We propose that at least some nearby galaxies could also have thick disks that are not uniformly old, and that geometrically thick disks might be complex structures resulting from different formation mechanisms in their inner and outer parts. | Thick disks have now been known to exist for more than 30 years, both in nearby galaxies \citep{Burstein1979,Tsikoudi1979} and in the Milky Way \citep{Gilmore1983}. However, there are different ways to define a thick disk: \begin{itemize} \item geometrically (or morphologically), based on decomposition of vertical density profiles \citep{Gilmore1983, Juric2008, Yoachim2006,Comeron2011}, or at a fixed height above the disk mid-plane \citep{Yoachim2008,Rekjuba2009, Cheng2012a} \item kinematically \citep{Morrison1990, Majewski1992,Bensby2003, Reddy2003,Adibekyan2012,Haywood2013} \item chemically, as the \al-rich sequence in the [\al/Fe] vs [Fe/H] plane \citep{Fuhrmann1998,Navarro2011,Adibekyan2012,Bovy2012} \item as the old part of the disk \citep{Haywood2013,Bensby2014,Xiang2015} \end{itemize} While all of these definitions can be applied in the Milky Way, only a geometric definition can be used for external galaxies. These geometrically thick disks are extended (they form a red envelope all around the thin disks) and have scale-lengths comparable to those of thin disks \citep{Yoachim2006,Pohlen2007,Comeron2012}. Their red colors (and the absence of radial color gradients) have led to the tentative conclusion that they are made of uniformly old stellar populations \citep{Dalcanton2002,Rekjuba2009}. However, the degeneracy between age and metallicity measured from broad-band photometry complicates further exploration of the age and chemical structure of thick disks in nearby galaxies. In the Milky Way, the geometrically defined thick disk has a large scale-length \citep[$\sim$ 3.5--4 kpc,][]{Ojha2001, Juric2008, Jayaraman2013}, in agreement with measurements for nearby galaxies. However, in the Milky Way this geometrically thick disk does not correspond to a uniform physical component in terms of chemical properties. Indeed, the \al-rich thick disk is centrally concentrated with a short scale-length of about 2 kpc \citep{Bensby2011, Cheng2012b, Bovy2012,Bovy2016}, and very few \al-rich stars are found in the outer disk of the Milky Way \citep{Nidever2014, Hayden2015}. This means that the chemically-defined thick disk and the geometrically-defined thick disk have a totally different structure. While this discrepancy has been mentionned by several authors \citep[e.g.,][]{Bovy2012,Jayaraman2013}, the reasons for the discrepancy itself have received less attention. In \cite{Minchev2015} we used numerical simulations to propose an explanation. We showed that in simulated disks the oldest stellar populations are indeed concentrated within the inner disk, while younger stellar populations have larger scale-lengths and smaller scale-heights \citep[see also][]{Martig2014}. However, we also showed that a thin-thick disk decomposition is still possible even in the outer disk, and that such geometrically-defined thick disks are very extended. This is because most mono-age populations flare in their outer regions, with the flaring radius increasing for younger populations (such a flaring was recently found by \citealp{Bovy2016} for \al-poor stars in the Milky Way). As a consequence, while the very center of the galaxy is dominated by old stars, the more extended parts of the thick disk are made of progressively younger stellar populations, so that a geometrically defined thick disk would have a radial age gradient going from old stars in the center to young stars in the outskirts of the galaxy. Such a radial age gradient has also been seen independently in simulations by \cite{Rahimi2014} and \cite{Miranda2015}. However, we lack a direct observational test of this theoretical picture using actual stellar ages instead of abundance proxies. Two studies (\citealp{Martig2016}, M16, and \citealp{Ness2016}, N16) have recently (and for the first time) determined ages for stars over a large volume of the Galaxy within the Apache Point Observatory Galactic Evolution Experiment (APOGEE) survey \citep{Majewski2015}. In this paper we use these two sets of stellar ages to show that in the Milky Way the geometrically-defined thick disk indeed shows a radial age gradient, as predicted by the simulations. In Section 2, we present the APOGEE data and the techniques used to derive ages. We then present in Section 3 our results on the age structure of the disk of the Milky Way. In Section 4, we discuss the robustness of our results, compare the Milky Way to nearby galaxies and conclude the paper with a discussion of the implications of our results for thick disk formation scenarios. | \subsection{Robustness of our results} The current implementation of our age determination technique does not allow for a measure of the age uncertainties on a star-by-star basis, which prevents us from performing a proper study of how our results are affected by age uncertainties. However, as described in M16, we used a leave-one-out cross validation algorithm to estimate that the r.m.s age error for our training set is $\sim$40\%. We find that the radial age gradients are still present if we create mock data samples by convolving our ages with a 40\%-wide Gaussian error. We also test if our results on the age gradient depend on the method used to determine stellar ages. We repeat our analysis of RC stars using ages obtained by N16 via \textit{The Cannon} (see top panel of Figure \ref{fig:Cannon}). With the N16 ages, the age gradient in the geometrically-defined thick disk is still present --- it is even steeper, with older ages for thick-disk stars in the inner disk. There is, however, a good general agreement between the two age determination techniques, which is reassuring. Finally, we check that our results are not an artifact related to the use of RC stars. This could arise either from the age values themselves (less robust for RC stars because ages are affected by mass loss during the RGB phase), or from the fact that RC ages are a biased sampling of the underlying total stellar population. We show in the bottom panel of Figure \ref{fig:Cannon} the age gradients for RGB stars (defined as giants in APOGEE DR12 but not in the RC catalogue). We use ages determined by N16 and distances from \cite{Ness2015}. The age gradients are also found for RGB stars, although the gradients are shallower and the shape of the radial trends is slightly different: this reflects the different age distribution of RC vs RGB stars, but also the $\sim$3 times larger distance uncertainties for RGB stars compared to RC stars. Using both a different set of stellar ages and a different type of stars, we thus confirm that the geometrically-defined thick disk is younger in its outer regions. We emphasize again that the median age we find for RGB and RC stars is in no way representative of the age of the underlying total stellar population, and as such cannot be directly compared to simulations. The main obstacle is not so much the survey selection function (as discussed in \citealp{Hayden2015}, the survey selection function does not depend strongly on metallicity and the sample of giants observed is representative of the underlying population of giants), but rather the complex age distribution of RGB and RC stars. The age distribution of RC and RGB stars tends to be biased towards younger ages, but the strength of the effect depends on the stellar evolutionary phase and the local star formation history \citep{Girardi2001, Bovy2014, Hayden2015}. Correcting for this age bias would require some complex modelling which is beyond the scope of this paper. We note however that the age bias does not affect our main result, i.e. the existence of an age difference between the inner and outer disk. \begin{figure} \centering \includegraphics[width=0.48\textwidth]{fig3a.pdf} \includegraphics[width=0.48\textwidth]{fig3b.pdf} \caption{Radial age gradients using stellar ages computed by N16 using \textit{The Cannon}, for RC stars on the top panel and for RGB stars on the bottom panel. This confirms the presence of strong radial age gradients at all heights above the mid-plane. The errors on distances are larger for RGB stars, so that the radial gradients are shallower than for RC stars.} \label{fig:Cannon} \end{figure} \subsection{The Milky Way compared to nearby galaxies} Our results show that the geometrically-defined thick disk in the Milky Way has a strong radial age gradient. This reconciles measurements of a short scale-length for the \al-rich disk with measurements of a large scale-length for the geometrically thick disk. Large scale-lengths are also measured for (geometrically) thick disks in external galaxies, but we do not know yet if these large scale-lengths have the same origin as in the Milky Way (in which case the disks would have a radial age gradient), or if these external geometrically thick disks are uniformly old components. Given the variety of formation histories for disk galaxies \citep[e.g.,][]{Martig2012}, it is likely that both types of thick disks exist. The most direct way to test this would be to identify which nearby galaxies also have an age gradient in their geometrically-defined thick disk. However, measuring ages for thick disks outside of the Milky Way is extremely challenging. A few Hubble Space Telescope (HST) studies have measured the properties of resolved stars in nearby edge-on disk galaxies, and found older stars at large scale-heights, but do not probe the radial structure of the thick disk \citep{Seth2005, Tikhonov2005, Mould2005}. An exception is \cite{Rekjuba2009}, who study RGB stars in NGC 891 with the HST and do not find any radial color or metallicity gradient along the thick disk. Spectroscopic studies are limited by the very faint surface brightness of the outer regions of thick disks. The Lick indices study of \cite{Yoachim2008b} was not able to probe the radial age structure of thick disk. Similarly, while Integral Field Unit spectroscopy is the ideal tool to probe thick disks, the VLT/VIMOS observations of \cite{Comeron2015} were limited by S/N, grouping the entire thick disk region in a single bin, and were thus unable to unveil its structure. Broad-band photometry can more easily reach deeper levels of surface brightness, but age and metallicity are degenerate and age determinations are quite approximate. \cite{Dalcanton2002} measured the B-R and R-K colors of thick disks, finding that thick disks typically have red colors (B-R $\sim$1.3--1.5) and no strong radial color gradient. This absence of a color gradient \citep[as also found by][]{Rekjuba2009} could naively be interpreted as an argument against an age gradient. However, these colors are very insensitive to age for populations older than $\sim$5 Gyr. We test this using the PARSEC isochrones \citep{Chen2014} combined with a Chabrier IMF. As an example, we show in Figure \ref{fig:colors} the variation of B-R and R-K colors for a single stellar population (SSP) of increasing age and a metallicity of -0.5 (roughly typical of a thick disk population). An age gradient from 10 to 5 Gyr along the thick disk would only give a small change in color of $\sim$0.1 mag. This means that current broad-band observations cannot exclude younger ages for the outer parts of thick disks and that deeper spectroscopic observations would be needed to probe the age structure of thick disks. \begin{figure} \centering \includegraphics[width=0.48\textwidth]{fig4.pdf} \caption{Integrated colors as a function of age for a SSP with a metallicity of -0.5 (based on a Chabrier IMF and the PARSEC stellar evolution models). The color difference between populations of 5 and 10 Gyr is of only $\sim$0.1 mag. If nearby edge-on galaxies had the same radial age gradient as the Milky Way, this would not be clearly obvious from their broad-band colors.} \label{fig:colors} \end{figure} \subsection{Final words: implications for thick disk formation scenarios} We leave to a future paper the detailed comparison between the age structure of the thick disk in the Milky Way and in simulated galaxies, which will require a careful modeling of the selection function of our RC sample. However, to first order, the observed age gradient in the geometrically thick disk is qualitatively consistent with what we find in our simulations \citep{Martig2014,Minchev2015} and also independently in simulations by \cite{Rahimi2014} and \cite{Miranda2015}. This suggests that complex age structures in thick disks might be a common feature of disk galaxy evolution. We leave for future work the understanding of the relation between the age structure of a thick disk and its detailed formation history. A first conclusion we can already draw is that geometrically thick disks (certainly in the Milky Way and maybe also in some external galaxies) might arise from a succession of events of different nature, and do not need to form all at once at high redshift. The inner parts might have formed in a violent phase at high redshift (either via disk instabilities or mergers), while the outer parts formed later, from the flaring of younger and more extended populations. To test this for external galaxies will require deep spectroscopic observations which measure either age or \aFe profiles along their thick disk (metallicity is not a good indicator of formation history: the Milky Way's geometrically thick disk has a flat metallicity gradient but a complex formation history). This should now be possible with instruments like MUSE on the VLT and will allow for direct tests of the similarity between the Milky Way and its neighbors. | 16 | 9 | 1609.01168 |
1609 | 1609.06054_arXiv.txt | We present the {stellar mass - size relation} for 49 galaxies within the $z$~=~1.067 cluster SPT-CL J0546$-$5345, with FWHM $\sim$80-120~mas $K_{\mathrm s}$-band data from the Gemini multi-conjugate adaptive optics system (GeMS/GSAOI). This is the first such measurement in a cluster environment, performed at sub-kpc resolution at rest-frame wavelengths dominated by the light of the underlying old stellar populations. The observed {stellar mass - size relation} is offset from the local relation by 0.21~dex, corresponding to a size evolution proportional to $(1+z)^{-1.25}$, consistent with the literature. The slope of the {stellar mass - size relation} $\beta$ = 0.74 $\pm$ 0.06, consistent with the local relation. The absence of slope evolution indicates that the amount of size growth is constant with stellar mass. This suggests that galaxies in massive clusters such as SPT-CL J0546$-$5345 grow via processes that increase the size without significant morphological interference, such as minor mergers and/or adiabatic expansion. The slope of the cluster {stellar mass - size relation} is significantly shallower if measured in $HST$/ACS imaging at wavelengths blueward of the Balmer break, similar to rest-frame UV relations at $z$~=~1 in the literature. The {stellar mass - size relation} must be measured at redder wavelengths, which are more sensitive to the old stellar population that dominates the stellar mass of the galaxies. The slope is unchanged when GeMS $K_s$-band imaging is degraded to the resolution of $K$-band HST/NICMOS resolution but dramatically affected when degraded to $K_s$-band Magellan/FourStar resolution. Such measurements must be made with AO in order to accurately characterise the sizes of {compact, $z$~=~1 galaxies}. | In the local universe, most luminous galaxies belong to one of two dominant populations: early- or late-type galaxies. The former are typically passive, red, spheroids, further classified into fast and slow rotators, while the latter are generally blue, star-forming disks. These familiar Hubble-type classifications do not apply as readily to high-redshift galaxies, the most massive of which are compact and red \citep{Szomoru2011,Talia2014}. Few compact systems exist at the present day \citep[][]{Trujillo2009,Taylor2010,Trujillo2012,Trujillo2014}, so it is logical to suppose that the most massive high-redshift galaxies must undergo significant size evolution to become present-day passive elliptical galaxies. For example, $z\sim 1$ galaxies are $\sim$ 2 times more compact when compared with $z=0$ \citep[e.g.][]{Daddi2005,vanDokkum2008,Damjanov2009}, and those at $z\sim 4$ are $\sim$ 6 times smaller \citep{Straatman2015}. Consequently the zero-point of the {stellar mass - size relation} decreases with increasing redshift \citep[{e.g.}][]{Buitrago2008,Nagy2011,Bruce2012,Law2012,Ownsworth2014}. The three main proposed mechanisms for this size increase are major mergers, minor mergers and adiabatic expansion. The slope of the {stellar mass - size relation} is typically seen to be constant with redshift \citep[e.g.][]{Damjanov2011,Newman2012,vanderWel2014}, requiring that the size increase not depend on stellar mass. This disfavours major mergers, which would increase the size of the most massive galaxies at a higher rate than less massive galaxies \citep{Khochfar2006}. Major mergers are also not predicted to be sufficiently frequent to explain the observed size evolution \citep{Bundy2009,Lotz2011}. However, the situation is not clear, as some authors do observe the slope of the {stellar mass - size relation} to change with redshift, e.g. \citet{HuertasCompany2013,Ryan2012,Delaye2014}. {These studies are generally undertaken in the rest-frame UV, so size measurements may be affected by clumps of recent star formation.} In the currently favoured minor merger paradigm described by \citet{Oser2012} and \citet{Toft2014}, compact galaxies first form out of collapsing gas, then later accrete gas-poor satellites in dry minor mergers. This is known as `two-phase galaxy formation' and was demonstrated to occur in the nearby universe by \citet{Forbes2011} with a study of globular clusters. At higher redshifts, other observational work such as \citet{vanDokkum2010,Barro2013,vanDokkum2013,Tadaki2014,vanDokkum2015} and simulations by \citet{Noguchi1999,Dekel2014,Wellons2015} have demonstrated the early formation phase, while the growth phase has been studied by \citet{Newman2012}, who found that the mergers could explain the size evolution but must happen on a rapid timescale, and by \citet{Morishita2016}, who observed sufficient numbers of satellites around a $z$~=~1.9 compact galaxy to explain the predicted size growth, once additional star formation is taken into account. In the adiabatic expansion model proposed by \citet{Fan2008,Fan2010}, galaxies experience a rapid mass loss event caused by AGN or supernova winds. After some time delay, expansion occurs in an amount proportional to the fraction of mass lost \citep{RangoneFigueroa2011}. Recent analysis by \citet{Wellons2016} of galaxies in the Illustris simulation \citep{Genel2014,Vogelsberger2014a,Vogelsberger2014b,Nelson2015} suggests that adiabatic expansion is responsible for less size growth than mergers are. Adiabatic expansion associated with the formation of new stellar populations {(but not due to AGN winds) during 2~$> z >$~1.2} was essentially ruled out by \citet{Damjanov2009}{, though that study was conducted at resolutions similar to the angular sizes of galaxies at those redshifts}. The {stellar mass - size relation} can be used to distinguish between major mergers and minor mergers / adiabatic expansion. {The accuracy of the relation depends most strongly on resolution and rest-frame wavelength.\footnote{Limiting magnitude appears less critical than spatial resolution for morphological classification; \citep[][]{Povic2015}. For cluster galaxies, sufficient spatial coverage of the cluster is also required in order to remove any environmental effect \citep[e.g.][]{Strazzullo2010}.} On one hand, s}ufficient resolution is required to measure effective radii and S{\'e}rsic \citep{Sersic1968} indices of the most compact galaxies. {On the other hand, and equally i}mportantly, the rest frame redwards of the 4000 \AA\/ spectral break is necessary for this measurement, as it is stellar mass that is the main driver of galaxy properties such as colour, age and specific star formation rate. Current efforts to measure high-redshift galaxy morphologies at high resolution {\citep[e.g.][]{HuertasCompany2013,Ryan2012,Delaye2014}} are limited to optical HST imaging \citep[e.g. FWHM 0.09 arcseconds in F814W $\sim I$-band,][]{Scoville2007}, but such filters are rest-frame UV at these redshifts. Consequently, they suffer from contamination of starburst events rather than tracing the bulk of the stellar mass. Near infrared imaging has necessarily poorer resolution, e.g. HST/WFC3 $\sim H$-band FWHM is 0.18 arcseconds \citep{Guo2011}, which is 1.48 kpc at $z=1$ and thus may not provide sufficient resolution for accurate profile fitting {\citep[e.g.][]{Damjanov2009}}. Due to this tradeoff between resolution and wavelength, there is a wide variety of rest-frame wavelengths at which the measurements are made, with correspondingly wide variety in the results. \citet{Carrasco2010} addressed this tradeoff with ground-based adaptive optics imaging, but were limited to just eight field galaxies observed one at a time due to the small effective field of view of the Gemini NIRI camera when corrected by the Altair laser guide star system. {Consequently, the current literature cannot be used to distinguish between major mergers, minor mergers and adiabatic expansion.} In this paper we study the massive, evolved galaxy cluster SPT-CL J0546$-$5345. We present the first high-resolution ground-based measurement of the {stellar mass - size relation} measured in the rest-frame light of the old stellar population at $z$~=~1, and demonstrate the importance of measuring the relation redward of the 4000 \AA\/ break by comparing with rest-frame UV {stellar mass - size relation}s for the same sample. In Section 2 we present our GeMS $K_s$-band imaging and describe our observing and data processing strategies, including successful correction of the optical distortion and faint source residual images. We also present our ancillary data of GMOS spectroscopy and HST archival F606W and F814 band imaging. Section 3 contains the galaxy profile fitting and cluster membership. The {stellar mass - size relation} is presented and discussed in Section 4. In Section 5 we summarise our conclusions. Throughout this work we assume a concordance cosmology with $\Omega_M$=0.264, $\Omega_\Lambda$=0.736 and H$_{0}$=71 km s$^{-1}$ Mpc$^{-1}$, for consistency with \citet{Brodwin2010}. Magnitudes are presented on the AB system. | We present the first high angular-resolution measurement of the rest-frame $Y$-band {stellar mass - size relation} for the galaxy cluster SPT-CL J0546$-$5345 at $z$~=~1.067. We demonstrate our strategies for addressing the data processing challenges associated with these complex imaging data in order to achieve a final stacked image with a mean PSF FWHM~$\sim$~110~mas (corresponding to a radius of 450 pc at the cluster redshift), with a uniform sky subtraction free from strong residual images that echo faint field sources. The spatial variation of the PSF is modelled as a two-dimensional Moffat profile fit to known stars across the field and interpolated at the location of each measured galaxy. Forty-nine cluster member galaxies are detected, with median S{\'e}rsic index $n$ = 3.8 $\pm$ 0.5. The {stellar mass - size relation} at $z\sim$ 1 is offset from that at $z\sim$ 0 by 0.21 dex, corresponding to a size evolution of $\gamma \propto (1+z)^{-1.25}$, which is consistent with previous results for minor mergers \citep{Nipoti2009}. The {stellar mass - size relation} exhibits a slope of $\beta$ = 0.74 $\pm$ 0.06, consistent with the local slope of $\beta$ = 0.70 $\pm$ 0.05 reported for quiescent low redshift cluster galaxies \citep{Guo2009}. This suggests that the cluster SPT-CL J0546$-$5345, an extremely massive cluster at $z\sim$ 1, has ceased its early, rapid growth dominated by major mergers, leaving the galaxies to increase in size via minor mergers and/or adiabatic expansion due to AGN mass loss winds. The {stellar mass - size relation} for the cluster members is also measured in the rest-frame $B$ and $U$ bands from archival $HST$/ACS F814W and F606W imaging. We show that those measurements are contaminated by knots of star formation that affect the light profiles. Galaxy effective radii are preferentially overestimated for low-mass systems and preferentially underestimated for high-mass systems when measured in the rest-frame UV wavelengths that are blueward of or straddle the $\lambda$ = 4000\AA\ break, with a stronger bias at shorter wavelengths. The effect of this is that the slope of the {stellar mass - size relation} is severely affected by the rest-frame wavelength range at which it is measured, although the zero-point is not affected such that a 10$^{11}$ M$_\odot$ galaxy has an effective radius of 2.45 kpc, consistent with the literature. This is a vivid illustration of the necessity of performing size measurements in the rest-frame underlying old stellar population in order to avoid bias by clumpy regions of star formation. We measure the {stellar mass - size relation} after Gaussian smoothing to typical imaging resolutions obtained with diffraction-limited HST/NICMOS $K$-band and seeing-limited Magellan/FourStar $K_s$-band. The NICMOS resolution gives a result consistent with our GSAOI imaging, but the seeing-limited FourStar resolution gives a significantly steeper slope. Hence the {stellar mass - size relation} cannot be accurately measured at $z \gtrsim$ 1 from the ground without the use of AO. Having demonstrated the potential of wide field AO observations to characterise the underlying stellar populations in high redshift cluster galaxies, in future work we will present additional clusters across the redshift range 1 $<z<$ 2 to assess the mechanisms for galaxy growth across this key epoch. | 16 | 9 | 1609.06054 |
1609 | 1609.06262_arXiv.txt | The data taken by the advanced LIGO and Virgo gravitational-wave detectors contains short duration noise transients that limit the significance of astrophysical detections and reduce the duty cycle of the instruments. As the advanced detectors are reaching sensitivity levels that allow for multiple detections of astrophysical gravitational-wave sources it is crucial to achieve a fast and accurate characterization of non-astrophysical transient noise shortly after it occurs in the detectors. Previously we presented three methods for the classification of transient noise sources. They are Principal Component Analysis for Transients (PCAT), Principal Component LALInference Burst (PC-LIB) and Wavelet Detection Filter with Machine Learning (WDF-ML). In this study we carry out the first performance tests of these algorithms on gravitational-wave data from the Advanced LIGO detectors. We use the data taken between the 3rd of June 2015 and the 14th of June 2015 during the 7th engineering run (ER7), and outline the improvements made to increase the performance and lower the latency of the algorithms on real data. This work provides an important test for understanding the performance of these methods on real, non stationary data in preparation for the second advanced gravitational-wave detector observation run, planned for later this year. We show that all methods can classify transients in non stationary data with a high level of accuracy and show the benefits of using multiple classifiers. | \label{section:introduction} The advanced Laser Interferometer Gravitational-Wave Observatory (aLIGO) detectors are two 4km interferometers at Hanford, Washington (H1) and Livingston, Louisiana (L1) \cite{2010CQGra..27h4006H, 0264-9381-32-7-074001}. The Italian 3km interferometer Virgo is expected to join the advanced detector network early in 2017 \cite{2008CQGra..25k4045A}. The detector duty cycle and sensitivity to astrophysical signals will be determined by noise sources created by the instruments and the environment. In particular, as the detector noise is non-Gaussian short-duration transients will limit the sensitivity of searches for transient astrophysical sources such as compact binary coalescences \cite{Allen:2005cz}. The first aLIGO observation run (O1) began autumn 2015. On the 14th September 2015 the aLIGO and Virgo teams detected gravitational waves from the binary black hole system GW150914 \cite{2016PhRvL.116f1102A}. A second binary black hole detection was made on the 26th of December \cite{PhysRevLett.116.241103}. An extensive study of the noise transients, which occurred in the data containing the detections, was carried out for the validation of the signals \cite{2016arXiv160203844T}. As the advanced detector network approaches its design sensitivity, the number of detections is expected to increase. Adding more detectors to the network increases the number of possible noise sources and the time it will take to identify their origin. Transients which occur in any one detector will limit the joint analysis time for the network. Understanding the sources of noise transients in the detectors will become increasingly more important with a latency of a few hours. The detectors contain many environmental and instrumental sensors, which produce auxiliary channels of data that can be used to monitor the detector behaviour and track the causes of short-duration noise artifacts. Auxiliary channels that are not sensitive to gravitational waves can be used to identify noise transients, also known as ``glitches", in the detector output and veto those events \cite{2011CQGra..28w5005S, 2014PhRvD..89l2001A, 2014arXiv1410.7764T}. Classification and categorization of transients using individual channels of data may provide valuable clues for the identification of their sources, which can aid in efforts to eliminate them \cite{2015CQGra..32u5012P, 1742-6596-243-1-012006}. So far classification has mainly been achieved by visual inspection of spectrograms of the transients, but automatic classification is essential for future detections of astrophysical gravitational-wave signals. Three methods for fast classification of transients have been developed for the analysis of aLIGO and Virgo data. They are Principal Component Analysis for Transients (PCAT), Principal Component LALInference Burst (PC-LIB) and Wavelet Detection Filter with Machine Learning (WDF-ML). Previous work has shown that these methods can classify artificial data sets with an efficiency up to $95\%$ \cite{2015CQGra..32u5012P}. In this paper we evaluate the performance of these algorithms using glitches in real data from aLIGO. In Section 2 we provide details of the detector data. In Section 3 we give a brief overview of the three different algorithms and details of any improvement they underwent since the previous study. In Section 4 we present the results for the three algorithms on glitches from aLIGO L1 and H1 detector data. This is followed by a discussion in Section 5 of our plans for future improvements and classification during the second aLIGO run (O2) and first Virgo observation run. \begin{figure}[!t] \begin{centering} \includegraphics[width=\textwidth]{plots/range.pdf} \label{fig:range} \caption{The mean binary neutron star inspiral range for the two aLIGO detectors during ER7. The Hanford detector had a higher range but also a higher glitch rate. The average range was 50-60 Mpc.} \end{centering} \end{figure} | \label{section:discussion} Non-Gaussian noise in the aLIGO and Virgo detectors can potentially mimic a gravitational-wave signal, reduce the duty cycle of the instruments and decrease the sensitivity of the detectors. Classification of different noise transient signals may help identify their origins and lead to a reduction in their number. We have developed three methods for noise classification and have previously demonstrated their performance on simulated transients in simulated Gaussian aLIGO noise \cite{2015CQGra..32u5012P}. However, as real noise from the advanced detectors is non-stationary and non-Gaussian, a better understanding of how our methods will perform during the upcoming observation runs of the advanced detectors is required. Although the detectors are typically more stable during observing runs than during ER7, we expect the types of glitches investigated in this work to be representative of the glitch classes in the upcoming observing runs. In the ER7 data from the L1 detector PCAT missed 90 transients and classified $95\%$ of the remaining transients correctly. PC-LIB missed 33 transients and classified $98\%$ of the remaining transients correctly. WDF-ML classified all transients and $95\%$ of them were correct. In the H1 data PCAT missed 120 transients and classified $99\%$ of the remaining transients correctly. PC-LIB missed 6 transients and classified $95\%$ of the remaining transients correctly. WDF-ML classified all transients and $92\%$ of them were correct. We conclude that our methods have a high efficiency in real non-stationary and non-Gaussian detector noise. The efficiency of the WDF-ML algorithm is reduced for the Hanford glitches because the duration of the transients becomes much larger than the analysis window, which reduces the efficiency of the overall classification. This could be prevented by applying a high duration cutoff to the transients found by the ETG before classification. Most high duration and SNR transients are removed by data quality vetoes. Conversely, short duration transients will be more important as they have a higher impact on the gravitational-wave search backgrounds. Since they are rarely removed by vetoes their accurate classification is crucial to improve gravitational-wave searches as an accurate categorization will allow us to search for couplings within the detector \cite{2016arXiv160203843T, 2016arXiv160203844T}. In the future the whitening performed by WDF-ML will be improved by using a technique known as adaptive whitening \cite{adapwhite:04}. Because of the different strengths and weaknesses of the different methods having multiple classifiers is a winning strategy. WDF-ML can classify lower frequency transients than the other two methods. PC-LIB is better able to classify longer duration transients due to its longer analysis window. PCAT can classify new types of transients as soon as they appear in the data and thus provide transient waveforms for PC-LIB's signal models. Further improvements could also be made by using a training set of pre-classified waveforms or exploring the use of dictionary learning algorithms for glitch classification \cite{dictlearn}. The aLIGO gravity spy project aims to build these data sets through a citizen science program \cite{2016AAS...22810902Z, Simpson:2014:ZOW:2567948.2579215}. \ack We thank Salvatore Vitale, Reed Essick and the Burst and DetChar groups of the LIGO Scientific Collaboration for helpful discussions of this work. DT and MC are partially supported by the National Science Foundation through award PHY-1404139. ISH and JP gratefully acknowledge the support of the UK Science and Technology Facilities Council (STFC) grant numbers ST/L000946/1 and ST/L000946/1. JP, ISH and EC also gratefully acknowledge the support of the Scottish Universities Physics Alliance (SUPA). ATF and JAF gratefully acknowledge the support of the Spanish MINECO (grant AYA2013-40979-P) and the Generalitat Valenciana (PROMETEOII-2014-069). This paper has been assigned LIGO document number LIGO-P1600263. | 16 | 9 | 1609.06262 |
1609 | 1609.01674_arXiv.txt | We analyze the broad-band X-ray spectrum (0.3--50\,keV) of the luminous Seyfert 1 / quasar PG\,1211+143 - the archetypal source for high-velocity X-ray outflows - using near-simultaneous {\it XMM-Newton} and {\it NuSTAR} observations. We compare pure relativistic reflection models with a model including the strong imprint of photoionized emission and absorption from a high-velocity wind \citep{Pounds16a, Pounds16b}, finding a spectral fit that extrapolates well over the higher photon energies covered by {\it NuSTAR}. Inclusion of the high $S/N$ {\it XMM-Newton} spectrum provides much tighter constraints on the model parameters, with a much harder photon index / lower reflection fraction compared to that from the {\it NuSTAR} data alone. We show that pure relativistic reflection models are not able to account for the spectral complexity of PG\,1211+143 and that wind absorption models are strongly required to match the data in both the soft X-ray and Fe\,K spectral regions. In confirming the significance of previously reported ionized absorption features, the new analysis provides a further demonstration of the power of combining the high throughput and resolution of long-look {\it XMM-Newton} observations with the unprecedented spectral coverage of {\it NuSTAR}. | The standard model of an active galactic nucleus (AGN) is driven by accretion onto a supermassive black hole (SMBH; $M_{\rm BH} \sim 10^{6-9}$\,$M_{\odot}$). AGN are powerful sources of X-rays which most likely originate close to the SMBH itself. The X-ray spectra of unobscured AGN display an array of spectral features. In particular, they are often dominated by a hard power-law component, thought to be produced when ultraviolet photons, emitted from an optically thick, geometrically thin accretion disc \citep{ShakuraSunyaev73}, are inverse-Compton scattered by a `corona' of hot electrons \citep{HaardtMaraschi93}. Additional spectral features typically include a soft excess $< 2$\,keV \citep{ScottStewartMateos12}, a `Compton reflection' component $> 10$\,keV \citep{NandraPounds94} and emission lines, the strongest of which is often Fe\,K$\alpha$ fluorescence at $\sim$6.4\,keV \citep{GeorgeFabian91}. Through systematic spectral studies with {\it ASCA}, {\it XMM-Newton}, {\it Chandra} and {\it Suzaku}, a significant fraction of AGN are now routinely observed to also exhibit strong signatures of ionized absorption in their X-ray spectra, with at least half of all AGN hosting photoionized ``warm'' absorbers (e.g. \citealt{ReynoldsFabian95, Blustin05}). These absorbers produce numerous narrow and blueshifted absorption features (e.g. \citealt{Kaspi02, CrenshawKraemerGeorge03, McKernanYaqoobReynolds07}), implying material outflowing at a velocity from several hundred to several thousand km\,s$^{-1}$. Through the study of blueshifted absorption lines of K-shell Fe, observations of higher-luminosity AGN have also revealed the presence of much more highly ionized absorbers originating in high-velocity disc winds ($v_{\rm out} \sim 0.1c$; e.g. PG\,1211+143; \citealt{Pounds03}, PDS 456; \citealt{ReevesOBrienWard03}). Subsequent studies of archival {\it XMM-Newton} and {\it Suzaku} data have shown that similar high-velocity, highly-ionized outflows may be relatively common in nearby, luminous AGN (e.g. \citealt{Tombesi10, Tombesi11, Gofford13}). In general, the derived outflow rates are comparable to the AGN accretion rate (up to several solar masses per year) and carry kinetic power on the order of a few per cent of the bolometric luminosity. Such high-velocity outflows are believed to play a key role in linking growth of the SMBH and the host galaxy \citep{King03, King10} and offer an appealing explanation for the observed $M$--$\sigma$ relation for galaxies (Ferrarese \& Merritt 2000; Gebhardt et al. 2000). The archetypal high-velocity-outflow source is PG\,1211+143, a luminous narrow-line Seyfert galaxy / quasar at a redshift of $z = 0.0809$ \citep{Marziani96}, which is both optically bright with a strong ``Big Blue Bump" and X-ray bright with a typical X-ray luminosity of $\sim$10$^{44}$\,erg\,s$^{-1}$. The source is well-known for its spectral complexity and, through an {\it XMM-Newton} observation in 2001, provided the first detection of a sub-relativistic outflow ($v_{\rm out} \sim 0.14c$) in a non-BAL (broad absorption line) AGN \citep{Pounds03, PoundsPage06}. Subsequent {\it XMM-Newton} observations of PG\,1211+143 were combined to model the photoionized absorption and emission spectra, quantifying the mass flux and energetics of the outflow and confirming the potential importance for galaxy feedback \citep{PoundsReeves07, PoundsReeves09, Pounds14}. An extended {\it XMM-Newton} observation in 2014 has since shown the highly-ionized wind to have several outflow components, with primary velocities of $v_{\rm out} \sim 0.066c$ and $\sim$$0.129c$ \citep{Pounds16a}. Surprisingly, the analysis of a 2014 {\it NuSTAR} observation of PG\,1211+143 \citep{Zoghbi15} found no evidence for the high-velocity outflow, leading to the suggestion that it may be a variable, transient feature - a result which would have important implications for the prevalence of high-velocity outflows in high-luminosity AGN and their wider significance for feedback. Here, we report on a combined analysis of the same {\it NuSTAR} data with the contemporaneous extended $\sim$630\,ks {\it XMM-Newton} observation obtained in 2014. | \label{sec:discussion} An extended {\it XMM-Newton} observation of PG\,1211+143 has revealed spectral structure in the Fe\,K band, unseen in previous (shorter) observations. Utilizing the high throughput of the EPIC-pn, \citet{Pounds16a} showed that highly-ionized outflow resolved two velocity components ($v_{\rm out} = 0.129$ and $0.066c$), while an analysis of the high-resolution RGS data in \citet{Pounds16b} found lower-ionization co-moving counterparts of that dual-velocity flow. Here, we have extended our modelling over a wider energy range, incorporating contemporaneous {\it NuSTAR} data. The power of this approach -- i.e. combining {\it XMM-Newton} and {\it NuSTAR} data -- has previously been demonstrated in a precise measurement of the higher-energy continuum and thereby helping to disentangle absorption and reflection in a number of AGN - e.g. NGC 1365 \citep{Risaliti13, Walton14, Rivers15}, NGC 5548 \citep{Kaastra14, Ursini15, Cappi16} and PDS 456 \citep{Nardini15}. In Section~\ref{sec:hard_band_relxill}, we fitted the hard X-ray spectrum with relativistic reflection, following the recent analysis of \citet{Zoghbi15}. Using the {\it NuSTAR} data alone, we confirm their solution, finding a steep photon index ($\Gamma \sim 2.4$) and high reflection fraction ($R \sim$ 2--3). However, by also including the simultaneous {\it XMM-Newton} spectrum, just in the 3--10\,keV range (which overlaps with the {\it NuSTAR} bandpass), a harder-$\Gamma$ / lower-$R$ solution is strongly preferred. Replacing the overlapping {\it XMM-Newton} data with the high $S/N$ stacked 2014 EPIC-pn spectrum further tightens the constraints, with $\Gamma = 1.95^{+0.07}_{-0.05}$ and $R = 1.05^{+0.47}_{-0.49}$. Such continuum parameters are now fully consistent with those reported by \citet{Pounds16a, Pounds16b}, where a detailed analysis of the long-look {\it XMM-Newton} observation is performed, taking into account the spectral imprints of the outflowing wind. In Section~\ref{sec:broad-band_model_reflection}, we attempted to fit the broad-band (0.3--50\,keV) spectrum with relativistic reflection, finding that, while \textsc{relxill} alone cannot simultaneously fit the excess of emission in the soft band and the significant Fe\,K emission complex, a two-reflector model (comprising a distant reflector and a relativistically blurred reflector) provides a good fit to the simultaneous rev2670+{\it Nu}7 spectrum. However, Figure~\ref{fig:2014_pn_relxill} illustrates the inability of this model to match the spectral features of the stacked 2014 EPIC-pn spectrum. Clearly, a pure reflection model cannot match fine structure in the high $S/N$ spectrum without significant modification - in particular, strong signatures of ionized absorption at $\sim$0.5--0.9\,keV and $\sim$7.5\,keV remain in the residuals, along with an excess of ionized emission at $\sim$6.7--6.9\,keV - spectral features resolved by the excellent statistics of the {\it XMM-Newton} data and identified in \citet{Pounds16a, Pounds16b}. This emphasises the value of long-look observations, where the $\sim$630\,ks stacked EPIC-pn spectrum reveals significant spectral structure, otherwise lost in the statistical noise of shorter observations. In Section~\ref{sec:broad-band_model} we explore an alternative broad-band 0.3--50\,keV {\it XMM-Newton} + {\it NuSTAR} fit, now including the ionized absorption/emission components required by the {\it XMM-Newton} spectrum. We find that the model presented in \citet{Pounds16a, Pounds16b} extrapolates well to the {\it NuSTAR} band and, overall, provides an excellent fit to the broad-band X-ray spectrum. Having thus accounted for the various components of the high-velocity outflowing absorber, we note that a contribution of relativistic reflection is allowed by the data, manifesting itself in the form of excess emission just red-ward of the Fe\,K complex, although its strength is rather modest ($R \sim 0.6$). Additionally, when fitted to the 0.3--50\,keV band, the absorption-dominated model requires a best-fitting photon index of $\Gamma \sim 1.8$, which is significantly different from the value of $\Gamma \sim 2.4$ arrived at by fitting the {\it NuSTAR} data alone. This conclusion is, of course, supported by our MCMC analysis which shows that a harder-$\Gamma$/lower-$R$ solution is preferred when including the simultaneous {\it XMM-Newton} data, regardless of the assumed continuum model. Finally, we comment on one of the headline results of the \citet{Zoghbi15} analysis, namely that the {\it NuSTAR} spectrum exhibits no evidence of the high-velocity outflow originally reported in \citet{Pounds03}. This claim was based on a line-detection search across the appropriate bandpass with emphasis placed on the $\sim$7.1\,keV absorption line, first detected in the 2001 observation. The non-detection of that absorption line in the {\it NuSTAR} spectrum led to the suggestion that the ultra-fast outflow in PG\,1211+143 may be transient, or at least highly variable, implying that the contribution of ultra-fast outflows to feedback could be substantially lower than previously thought. \citet{Pounds16a} find that this absorption is indeed variable, depending on both column density and ionization parameter, and had a mean equivalent width $\sim$3 times smaller during the 2014 {\it XMM-Newton} observation than it was when first detected in 2001. As such, it is unsurprising that the feature was not significantly detected by a single line search in the {\it NuSTAR} spectrum. More generally, while a line search has provided a useful method in the past for detecting individual strong resonance lines, a preferred approach is to use self-consistent photoionization models (e.g. with \textsc{xstar}), as we do here, which incorporate a physically realistic array of common-velocity lines and absorption edges from various ions (including additional higher-order lines) and imprint appropriate related continuum curvature, for a given $N_{\rm H}$, $\xi$ and $v_{\rm out}$. The improvement in the fit statistic is then computed from the summed contributions of all of the absorption features, where several individual weak features may have a significant impact overall (although we do note that even this approach would not have led to a significant detection of the outflow using the contemporaneous {\it NuSTAR} data alone). In summary, while \citet{Zoghbi15} detected no discrete absorption features in the {\it NuSTAR} spectrum, the extended {\it XMM-Newton} observations, also in 2014, confirm that, not only is the outflow still present, but that it exhibits more complex spectral structure than previously realized. In studying the properties of powerful AGN winds, now widely accepted as one of the most important outcomes of the past decade of X-ray observations, we have noted the importance of self-consistent modelling with physically-motivated absorption models which properly account for multiple ions and continuum curvature, and emphasise that (i) finding a common velocity for a line series can provide a highly significant result, and (ii) {\it XMM-Newton} remains a powerful observatory for detailed, high-resolution spectroscopy, particularly when used in tandem with the broad-band coverage of {\it NuSTAR}. | 16 | 9 | 1609.01674 |
1609 | 1609.03392_arXiv.txt | Improved orbits are presented for the visual binaries WDS 02366+1227, WDS 02434-6643, WDS 03244-1539, WDS 08507+1800, WDS 09128-6055, WDS 11532-1540, WDS 17375+2419, and WDS 22408-0333. The latest orbits for these binaries were demonstrating a great inconsistency between the systemic mass obtained through Kepler's Third Law and that calculated as a sum of their components' mass through standard mass-luminosity and mass-spectrum relationships. Our improvement allowed us to obtain consistent systemic masses for WDS 02434-6643 and WDS 09128-6055 without a need for changing the Hipparcos parallax. For the remaining 6 pairs, we suggest the use of their dynamical parallax as a reliable distance estimate unless more precise parallax is reported. Astrophysical and dynamical properties of individual objects are discussed. | Determination of accurate orbits in visual binaries with known distances represents a reliable method to obtain their systemic mass using Kepler's Third Law. However, its direct application sometimes leads to unrealistic mass values largely inconsistent with the mass sum of individual components derived from empirical mass-luminosity and mass-spectrum calibrations. A careful inspection of their orbits and/or estimated distances is needed in order to clarify the reasons and nature of such inconsistent data. Essentially (but not only), these are poorly known (or even erroneous) distance determination, low quality (preliminary) orbital solution, and/or the previously unknown multiplicity of the components. In a recent paper of Tamazian et al. (2016) we have selected a small sample of 17 such binaries and have presented brief comments regarding how such consistency can be fixed either by improving the orbit or by refining the distance. Here we report a detailed analysis of new orbits from this sample and discuss their astrophysical and dynamical characteristics as well. | 16 | 9 | 1609.03392 |
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1609 | 1609.07052_arXiv.txt | \noindent Some scenarios for planetesimal formation go through a phase of collapse of gravitationally bound clouds of mm-cm-sized pebbles. Such clouds can form for example through the streaming instability in protoplanetary disks. We model the collapse process with a statistical model to obtain the internal structure of planetesimals with solid radii between 10 and 1,000 km. In the collapse, pebbles collide and, depending on relative speed, collisions have different outcomes. A mixture of particle sizes inside a planetesimal leads to better packing capabilities and higher densities. In this paper we apply results from new laboratory experiments of dust aggregate collisions (presented in a companion paper) to model collision outcomes. We find that the internal structure of a planetesimal is strongly dependent on both its mass and the applied fragmentation model. Low-mass planetesimals have no/few fragmenting pebble collisions in the collapse phase and end up as porous pebble-piles. The amount of fragmenting collisions increases with increasing cloud mass, resulting in wider particle size distributions and higher density. The collapse is nevertheless ``cold'' in the sense that collision speeds are damped by the high collision frequency. This ensures that a significant fraction of large pebbles survive the collapse in all but the most massive clouds. Our results are in broad agreement with the observed increase in density of Kuiper belt objects with increasing size as exemplified by the recent characterization of the highly porous comet 67P/Churyumov-Gerasimenko. | \label{sec:intro} Planet formation takes place around young stars as $\mu$m-sized dust and ice particles grow to ever larger bodies \citep{safronov69}. This leads to planets of characteristic sizes $10^4$-$10^5$ km. The growth starts with particles sticking together by contact forces \citep[see review by][]{blum08}. Compactified pebbles of mm-cm sizes have poor sticking properties, but growth to planetesimals can be aided by the mutual gravity in pebble clouds that are concentrated in turbulent gas \citep[see review by][]{johansen14}. This leads to the formation of planetesimals with a distribution of sizes ranging from 10 to several 100 km \citep{johansen15,simon15}. Many details of the gravitational collapse phase are not yet fully understood. \citet{nesvorny10} pioneered the modelling of the collapse phase in N-body simulations of a large number of pebbles coming together by their mutual gravity. They found that pebble clouds with high internal angular momentum collapse into binary planetesimals. This can explain the high fraction of binaries observed in the classical cold population of trans-Neptunian objects \citep{noll08}. The two components in binary Kuiper belt objects appear to have the same colour and composition \citep{benecchi09}, suggesting that they formed together, since the Kuiper belt, overall, has a broad colour distribution. In the case of binary formation through three-body encounters \citep[e.g.][]{goldreich02} the components, most likely, should not have the same composition. \medskip \citet[hereafter WJJ]{wahlberg14} investigated the evolution of the particle size distribution during the collapse phase, based on laboratory experiments of particle collisions \citep{guttler10}. The initial pebble clouds were assumed to arise from the streaming instability \citep[e.g.][]{youdin05,johansen09,johansen12,bai10}. A major result of this paper was that the collapse process is very dependent on planetesimal mass. More massive clouds collapse faster and collisions between pebbles result in pebble fragmentation. This affects the internal structure (e.g. density and porosity) of the resulting planetesimal: low-mass planetesimals should be porous pebble-piles while higher-mass planetesimals are a denser mixture of dust and pebbles. This relation with density increasing with increasing size agrees with observations of Kuiper belt objects \citep{brown13}. However, other parameters such as composition, radioactive heating and collisions will also affect the structure, but the effect of those processes all depend on the initial porosity and packing efficiency. Therefore the outcome of pebble cloud collapse models can be used as starting point for calculations of the long-term thermal evolution of planetesimals \citep[e.g.][]{lichtenberg16}. \citet{lorek16} expanded the cloud collapse model in \citetalias{wahlberg14} to investigate the evolution of the density of the pebbles throughout the collapse. Bouncing collisions cause porous dust (ice) aggregates to become more compact. The authors used their results to constrain the range of initial conditions (cloud mass, dust-to-ice ratio and initial filling factor) that can produce comets and other observed bodies in the outer Solar System. They find that planetesimals with observed comet bulk density of $\sim$0.5 g cm$^{-3}$ can form either if the cloud is low-mass and initially have compact pebbles or if the cloud is massive independent on initial pebble porosity. \medskip Observational data on the structure of planetesimals in the outer Solar System has increased dramatically in the past years. The space probes Rosetta (ESA) and New Horizons (NASA) have both reached their respective targets, the Jupiter family comet 67P/Churyumov-Gerasimenko (hereafter 67P) and the dwarf planet Pluto. Rosetta has multiple instruments that provide measurements for understanding the origin of planetesimals. OSIRIS is an optical, spectroscopic and infrared system for imaging the nucleus of 67P from Rosetta \citep{keller07}. The ``goosebump'' structures in the walls of the deep pits have been suggested to represent the primordial pebbles that make up the bulk of the comet \citep{sierks15}, although the meter scale of those pebbles are in some disagreement with the particle sizes that are believed possible to form by coagulation in the outer regions of protoplanetary disks \citep{birnstiel12,lambrechts14}. High-resolution images returned by the Philae lander indicate a typical scale closer to 1 cm at the surface \citep{mottola15}, more in agreement with expectations. Measurements of shape and gravity field have yielded a bulk density of only 0.53 g cm$^{-3}$, so clearly 67P is very porous (70-75\% depending on assumed dust-to-ice ratio). The CONSERT radar \citep{kofman07} had a main aim to measure the internal structure of the comet. Gravity measurements and radar tomography \citep{patzold16,kofman15} indicate that 67P is approximately homogeneous on length scales $<$3 m and very porous. These results are consistent with 67P being a pebble-pile consisting of loosely packed primordial pebbles from the solar protoplanetary disk. The constituent particles of 67P can also be inferred from the dust particles that fly off the surface. Particles with radii between 2 cm and 1 m have been observed with OSIRIS photometry \citep{rotundi15}. The GIADA instrument \citep{colangeli07} has detected compact (suggesting thermal processing) dust grains of sizes $\leq$100 $\mu$m escaping the comet, but also fluffy, low density ($\rho\sim 10^{-3}$ g cm$^{-3}$) dust aggregates with radii $\sim$0.1-1 mm \citep{rotundi15,fulle15}. The COSIMA instrument \citep{kissel07} collected dust aggregates onto plates to visually analyse their internal structure. The particles collected are porous aggregates that fragment easily upon collision \citep{schulz15}. Low collision speeds ($<$1-10 m s$^{-1}$) and the analysis of the collected aggregates \citep{hilchenbach16} suggest that they are not composed of an ice-dust-mixture and originate from the ice-free surface layers of the comet. \citet{skorov12} presented a comet model consisting of a top layer of ice-free dust aggregates residing on an interior mixture of ice and dust aggregates. Pebble-sized dust aggregates are needed to explain observed comet activity, as the tensile strength of a surface of $\mu$m-sized dust is too high for water sublimation \citep{blum14,blum15}. \citet{gundlach15} applied the model to 67P and found that it can explain the release of observed $\sim$cm-m-sized dust aggregates from the comet surface. \citet{massironi15} found that 67P is likely a contact binary, inferred from the onion-like structure with shell surfaces centered on the center-of-mass of each separate lobe. Thus 67P may have originally been a binary cometesimal, as is commonly the result of the gravitational collapse model of \citet{nesvorny10}, that later merged gently to a bimodal structure. Altogether, Rosetta and Philae observations of 67P are fully consistent with formation through slow gravitational contraction of a dense cloud of pebbles. Observations of the comet 103P/Hartley 2 by the EPOXI spacecraft supports the theory of pebble-pile comets. The comet has, like 67P, a bimodal shape and a low density \citep[$\rho\sim$ 0.22-0.88 g cm$^{-3}$ depending on porosity and composition,][]{ahearn11}. EPOXI also found large particles ($\sim$cm-m) in the comet's coma. Investigations by \citet{kretke15}, assuming formation through gravitational collapse, suggest that these particles could be primordial pebbles from which Hartley 2 was formed. Pluto with its diameter of $\sim$2,400 km is an icy planetesimal on the extreme other end of the size range of Kuiper belt objects. The fly-by by New Horizons showed, surprisingly, that the surface of Pluto is young \citep{stern15}, indicating heating by either short-lived or long-lived radionuclides and recent interior restructuring. Other possible sources of heating (e.g. tidal effects) are, today, insignificant \citep{moore15}. New Horizons is now continuing its journey, through the Kuiper belt, towards the object 2014 MU69, a mid-sized Kuiper belt object (diameter $<$45 km) of the cold population \citep{porter15}. This object has an intermediate size between 67P and Pluto. Its size may be low enough to have avoided extensive particle fragmentation during the collapse \citepalias{wahlberg14}, in contrast to Pluto, and thus maintain its primordial structure the same way as 67P. The results from \citet{lorek16} predict that 2014 MU69 has a dust-to-ice ratio of $\sim$3-7 and constituent pebbles with a volume filling factor close to the maximum value of $\sim$0.4. \medskip A major simplification in the work of \citetalias{wahlberg14} was that pebble fragmentation during the collapse was always assumed to be the source of a cloud of $\mu$m-sized monomer particles. In this paper we therefore expand the model for simulating the collapse of pebble clouds with a more realistic fragmentation model. The critical fragmentation speed and fragment size distribution are based on new experimental results presented in a companion paper \citep[hereafter Paper I]{bukhari16}. With this improvement we get more physically correct properties of the resulting planetesimals and can better compare the results with observations of e.g. the next target of New Horizons. The paper is organized as follows. In \sref{sec:model} we summarize the model and numerical method used in \citetalias{wahlberg14}. The implementation of the results of \citetalias{bukhari16} (the outcome of fragmenting collisions) is described in \sref{sec:fragModel}. The simulations of the collapse of clouds are presented in \sref{sec:sims}. In \sref{sec:disc} we discuss the relevance of our results for the formation of planetesimals by hierarchical accumulation and the validity of neglecting gas drag in our simulations. A discussion of the results and a comparison with previous simulations are presented in \sref{sec:conc}, which is based upon the results of \citetalias{bukhari16}. | \label{sec:conc} In this paper we have added the results of the silicate dust aggregate collision experiments from \citet[][Paper I]{bukhari16} to the model of planetesimal formation in \citet[][WJJ]{wahlberg14}. The planetesimals are assumed to form by the gravitational collapse of pebble overdensities in protoplanetary disks thanks to energy loss in inelastic collisions between the particles in the cloud. The new model has a more realistic treatment of fragmenting collisions and mass transfer, using the laboratory collision experiments in \citetalias{bukhari16}. To investigate the sensitivity to various aspects of the model, we run three sets of simulations using slightly modified fragmentation models (see \sref{sec:IC}). As in \citetalias{wahlberg14} the collapse times are short and decrease with increasing planetesimal mass. The collapse speed is, however, limited by free-fall and massive planetesimals ($R_\t{solid}\gtrsim 100$ km) all collapse, roughly, on the free-fall time ($\sim$25 years). The free-fall limit causes massive clouds to undergo a cold collapse where particles move with speeds slower than virial equilibrium speed. Particle speeds decrease by collisional dampening as the collapse progresses (\fref{fig:vvir}), causing pebbles to survive collisions even in massive pebble clouds (\fref{fig:pebble_frac}). In the new model, compared to \citetalias{wahlberg14}, the primordial pebbles have a harder time to survive in the most massive planetesimals. In the left plot of \fref{fig:pebble_frac} we see that, for the full model, a Pluto-sized planetesimal only has $\sim$0-20\% of its mass in pebbles. In the \textit{WJJ}-model the same planetesimal has $\sim$50\% of the mass in particles with radii $\geq$0.5 mm. The main explanation for this is that in the \textit{WJJ}-model fragmentation is modelled as erosion: a fragmenting collision results in a large remnant and cloud of small dust particles. In the new model, however, we have a continuous fragment size distribution \erefp{eq:cumMass} and dust production is rare. Energy dissipation is more efficient with small particles \citepalias[Appendix A in][]{wahlberg14} so with the \textit{WJJ}-model energy is dissipated faster, the collision speeds rapidly become subvirial and fewer fragmenting pebble-pebble collisions occur. The experiments in \citetalias{bukhari16} find a wide range in steepness, $\alpha$, of the power-law describing the fragment size distribution \erefp{eq:cumNum}. To explore the sensitivity to $\alpha$, we made two sets of simulations for the full fragmentation model: one with a shallow slope (less mass in small fragments, $\alpha=0.5$) and one with a steeper slope (more mass in small fragments, $\alpha=0.9$). The plots in \fref{fig:pebble_frac} show that the difference between the two sets is small compared to the differences to the other fragmentation models. \fref{fig:pebble_frac} also shows that modelling the mass transfer correctly is important. With efficient mass transfer (\textit{100\% MT}-model) pebbles not only survive the collapse to a larger degree (left plot) but also grow orders of magnitude in size (right plot). The difference between the \textit{100\% MT}-model and the full model is that a collision with mass transfer (MT in \fref{fig:outcomeFig}) results in perfect merger, whereas only 10-30\% of the projectile mass is transferred in the full model. In the collapse, many collisions occur in the velocity regime 1 m s$^{-1}\leq v_\t{n}<v_\t{0.5}$ (\fref{fig:vvir}, \erefnp{eq:v05}) where mass transfer is possible for similar-sized particles (\fref{fig:outcomeFig}). In the full model such collisions result in a slightly larger target but most of the projectile mass in small fragments, while in the \textit{100\% MT}-model the result is one merged target. Our results confirm previous suggestions made in \citetalias{wahlberg14}, that low-mass planetesimals (few tens of km or smaller) should consist mainly of primordial pebbles, be very porous and have low internal strength. More massive planetesimals consist of a mixture of pebbles and smaller particles. They would then have a better packing capability and be more dense, in agreement with the size-density correlation observed for Kuiper belt objects \citep{brown13}. Our results use the collision experiments with silicate dust aggregates from \citetalias{bukhari16}, while the outer regions of the Solar System contain a large fraction of ices. Our simulations in which we model the particles as ice change the outcome significantly, since ice particles survive much higher collision speeds. This result may nevertheless change with the inclusion of CO and CO$_2$ ice, which have recently been shown to have equally poor sticking properties as silicates \citep{musiolik16}. Therefore our results obtained with silica particles could be a good proxy for the collapse of actual pebble clouds in the outer regions of protoplanetary disks. | 16 | 9 | 1609.07052 |
1609 | 1609.02488_arXiv.txt | {Gaps, cavities and rings in circumstellar disks are signposts of disk evolution and planet-disk interactions. We follow the recent suggestion that Herbig Ae/Be disks with a flared disk harbour a cavity, and investigate the disk around HD~97048.} {We aim to resolve the 34$\pm$ 4 au central cavity predicted by Maaskant et al. (2013) and to investigate the structure of the disk.} {We image the disk around HD~97048 using ALMA at 0.85~mm and 2.94~mm, and ATCA (multiple frequencies) observations. Our observations also include the \element[ ][12]{CO} J=1-0, \element[ ][12]{CO} J=3-2 and \element[+][]{HCO} J=4-3 emission lines.} {A central cavity in the disk around HD~97048 is resolved with a 40-46 au radius. Additional radial structure present in the surface brightness profile can be accounted for either by an opacity gap at ~90 au or by an extra emitting ring at ~150 au. The continuum emission tracing the dust in the disk is detected out to 355 au. The \element[ ][12]{CO} J=3-2 disk is detected 2.4 times farther out. The \element[ ][12]{CO} emission can be traced down to $\approx$ 10 au scales. Non-Keplerian kinematics are detected inside the cavity via the \element[+][]{HCO} J=4-3 velocity map. The mm spectral index measured from ATCA observations suggests that grain growth has occurred in the HD~97048 disk. Finally, we resolve a highly inclined disk out to 150 au around the nearby 0.5~$M_{\sun}$ binary ISO-ChaI 126.} {The data presented here reveal a cavity in the disk of HD 97048, and prominent radial structure in the surface brightness. The cavity size varies for different continuum frequencies and gas tracers. The gas inside the cavity follows non-Keplerian kinematics seen in \element[+][]{HCO} emission. The variable cavity size along with the kinematical signature suggests the presence of a substellar companion or massive planet inside the cavity.} | Protoplanetary disks are the birth environments of planetary systems. How these planets form in their disks is an ongoing topic of debate, which is informed by an increasing number of disks that show various degrees of dispersal such as opacity cavities (transitional disks) and opacity gaps (pre-transitional disks) \citep[e.g.][]{2011ARA&A..49...67W}. Examples of such disks with directly imaged cavities at (sub) mm wavelengths include HD~100546 \citep{2014ApJ...791L...6W}, Sz~91 \citep{2015ApJ...805...21C, 2016MNRAS.458L..29C}, LkCa~15 \citep{2006A&A...460L..43P, 2011ApJ...742L...5A}, HD~142527 \citep{2013Natur.493..191C} and SAO~206462 \citep{2009ApJ...704..496B}. The common denominator between these disks is that their structure can be described by one large cavity or a broad ring of dust grains at reasonably large radii and with large ring widths (at least tens of au in radii and width) or with a pile-up of large dust in narrow rings. The gaps and/or cavities in these disks are not empty: they contain both smaller dust grains, as traced by scattered light imaging \citep[e.g. ][]{2012ApJ...745....5K, 2014ApJ...781...87A}, and gas, as traced by rotational \citep{2015ApJ...798...85P, 2015A&A...579A.106V} and ro-vibrational carbon monoxide (CO) lines \citep{2009A&A...500.1137V, 2011ApJ...733...84P}. Recently, long baseline Atacama Large Millimeter Array (ALMA) observations of HL Tau \citep{2015ApJ...808L...3A} and TW Hya \citep{2016ApJ...820L..40A, 2016ApJ...819L...7N, 2016arXiv160500289T} have demonstrated that these disks show a rich substructure of many concentric rings and gaps at scales as small as 1 au when observed at very high spatial resolution. Indeed, it is possible that most disks contain similar detailed structures which have not yet observed \citep{2016ApJ...818L..16Z}. Disks around Herbig Ae/Be (HAeBe) stars have historically been split into group I and group II. Group I sources have been interpreted as hosting gas-rich protoplanetary disks with a flared, bright dust surface; whereas the dust in group II disks is assumed to have settled towards the mid-plane and are therefore weak in mid- to far-infrared emission \citep{2001A&A...365..476M, 2004A&A...417..159D}. Recent modeling of resolved observations of group I sources suggest that their bright infrared emission can be attributed to the large vertical walls that exists as a consequence of large dust cavities \citep{2012ApJ...752..143H,2013A&A...555A..64M}. HD~97048 is a 3 Myr \citep{2006Sci...314..621L}, 2.5 M$_\sun$ \citep{2007A&A...470..625D} HAeBe with spectral type A0 at a distance of 158$^{+17}_{-14}$ parsec \citep{2007A&A...474..653V}. Its Spectral Energy Distribution (SED) is classified as group I, is bright in the mid- to far-IR, and rich in PAH features, but lacks any sign of amorphous and crystalline silicate features \citep{2014A&A...563A..78M}. The mass accretion rate onto the star is low with an upper limit of log(M$_{acc}$) $\leq$ -8.16 M$_{\odot}$ yr$^{-1}$ \citep{2015MNRAS.453..976F}. The disk around HD~97048 is exceptionally bright and is one of only two HAeBe disks in which near-IR 1-0 S(1) \citep{2011A&A...533A..39C} and Mid-IR \citep{2009ApJ...695.1302M} H$_2$ emission has been detected. The outer disk has been resolved in the PAH bands and shows a typical flaring geometry, with a flaring index of 1.26$\pm$0.05 \citep{2006Sci...314..621L}, and an inclined disk geometry with the eastern side farther from us. The disk has been very well studied using CO emission, and shows an 11 au cavity in ro-vibrational emission \citep{2009A&A...500.1137V}, but no detection in overtone emission \citep[][]{2015A&A...574A..75V}. The rotational CO ladder is richly populated as detected by {\textit Herschel} \citep[][]{2013A&A...559A..84M, 2014MNRAS.444.3911V, 2016arXiv160402055F}. In the (sub) mm bands, this disk has only been detected using single dish observations and has never resolved \citep{1998A&A...336..565H,2011PhDT.......219P, 2014AJ....148...47H}. The disk has, however, been resolved by polarimetric differential imaging of polarized scattered light \citep{2012A&A...538A..92Q}, showing a bright disk surface between $\approx$ 0.1\arcsec–1.0\arcsec ($\approx$16–160 AU), but no evidence for a disk cavity. \citet{2013A&A...555A..64M} also resolve the disk in the Q band (20 $\mu$m) spectrum and find that a large gap should be present in the disk between 2 and 34$^{+4}_{-4}$ au. In this manuscript we present resolved (sub) mm observations of the disk around HD~97048 obtained with ALMA and the Australia Telescope Compact Array (ATCA)\footnote{The Australia Telescope Compact Array is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.}. We describe the observations and data reduction in \S \ref{sec:obs+data}, and present the results in \S \ref{sec:results}. We discuss these results and put them into context with previous results in \S \ref{sec:discussion}, and present our conclusions in \S \ref{sec:conclusion}. \begin{figure*} \centering \includegraphics[width=\hsize]{./HD97048_continuum_all.eps}% \caption{Images of \object{HD 97048} for the ALMA band 7 (left panel) and band 3 (central panel) and the combined ATCA 33+35 GHz (right panel) observations, reconstructed using uniform (ALMA) and natural (ATCA) weighting. The intensity scale for all images is in units of Jy/beam. Over plotted are contours with [3, 15, 100 and 1100] times the RMS value of respectively 0.20, 0.18 and 0.010 mJy beam$^{-1}$. The beam is shown in orange in the bottom left of each panel.} \label{fig:continuum_all} \end{figure*} | \label{sec:conclusion} We summarize the main conclusions of this work as follows: \begin{itemize} % \item The dust disk of HD~97048 is resolved and extends radially out to 2.25\arcsec, after which the surface density drops sharply. \item There is a dust cavity visible in the 853 $\mu$m continuum emission out to 43 $\pm$3 au. \item Beyond the outer disk inner rim, we find additional radial structure in the continuum surface brightness profile. This radial structure can be modeled with either a disk gap centered at $\approx$ 90 au \textit{or} an extra emitting ring centered at $\approx$ 150 au. \item The disk cavity is not detected at 9~mm because free-free emission from the star fills in the cavity at our spatial resolution. \item The \element[ ][12]{CO} J=1-0, \element[ ][12]{CO} J=3-2 and \element[+][]{HCO} 4-3 emission lines all are more extended than the disk continuum emission, up to a radius of 5.2\arcsec for the \element[ ][12]{CO} J=3-2 emission. The discrepancy in size between the mm dust emission and the associated line emission, together with the sharp outer edge of the dust disk, can be explained by inward radial drift of the larger dust grains. \item Both the \element[+][]{HCO} J=4-3 and the \element[ ][12]{CO} gas extend inside the dust cavity, down to distances as close as 13.4 au from the central star. \item The \element[+][]{HCO} intensity-weighted velocity map shows {\gc a velocity structure deviating from Keplerian motion expected from a co-planar disk} inside of the cavity. {\gc Possible explanations for this velocity structure are a distortion by an inclined inner disk and an extra, non-Keplerian, velocity component.} \item A cavity size varying with dust grain size and the {\gc apparent} non-Keplerian motions inside of the cavity both hint at an extra body orbiting inside the cavity. \item A planet with a mass of $\approx$ 0.7 M$_{jup}$ at a radial distance between 2.5 and 11 au could open up the disk cavity observed at 853 $\mu$m. \item We resolve emission between 2.9 and 16.7 mm from the 0.5~M$_{\odot}+0.5$~M$_{\odot}$ binary ISO-ChaI 126 that we interpret as originating from a disk. The spectral slope of this emission between 70 $\mu$m and 15.8 mm is 1.9 $\pm$ 0.2. Using standard assumptions, we calculate a total disk mass of 0.80\% of the total system mass. \end{itemize} | 16 | 9 | 1609.02488 |
1609 | 1609.02731_arXiv.txt | Chemical abundances of stars of different metallicity are used in studying the chemical evolution of galaxies, in particular, the Milky Way (MW) and dwarf satellite galaxies. We will use the iron abundance relative to the solar one, [Fe/H] = log~$(N_{\rm Fe}/N_{\rm H})_{star} - (N_{\rm Fe}/N_{\rm H})_\odot$, as a metallicity indicator. The further into the low-metallicity region we want to advance, the more distant, on average, objects should be observed. That is why the observational data for the [Fe/H] $< -3$ region in the MW are obtained mostly from giant stars, and only giants are accessible to high-resolution spectroscopy in distant globular clusters and dwarf satellite galaxies. The spectral line formation conditions in the atmospheres of, in particular, metal-deficient giants are far from the equilibrium ones. Nevertheless, in most cases, not only the abundances of chemical elements but also the atmospheric parameters (effective temperature $\Teff$, surface gravity log~g, and microturbulence $\xi_t$) are determined using the assumption of local thermodynamic equilibrium (LTE). For example, for the sample of stars from the classical paper by Cayrel~et~al. (2004), abundances of many elements were revised by Andrievsky et~al. (2007, Na; 2008, Al; 2009, Ba; 2010, Mg and K) and Spite et~al.(2012, Ca) based on the non-LTE line formation, but using, at the same time, log~g and [Fe/H] determined by Cayrel et~al. (2004) within LTE. As shown by Mashonkina et~al. (2011) and Bergemann et~al. (2012), LTE underestimates the abundance derived from Fe~I lines, and the effect increases with decreasing metallicity. This means that the ratio [X(non-LTE)/Fe(LTE)] is overestimated by a larger amount at smaller [Fe/H], and we obtain a distorted view of the change in the relative abundance X/Fe with metallicity (i.e., with time). Another source of errors in determining X/Fe is the surface gravity obtained in LTE from the Fe~I/Fe~II ionization equilibrium method. Since the Fe~I lines are subject to departures from LTE, while the non-LTE effects for the Fe~II lines remain negligible down to a very low metallicity, [Fe/H] $\simeq -5$, LTE leads to underestimated values of log~g. The LTE assumption is commonly used for stars in dwarf galaxies to determine both atmospheric parameters (Frebel et~al. 2010; Simon et~al. 2010, 2015) and elemental abundances (Tafelmeyer et~al. 2010; Gilmore et~al. 2013; Jablonka et~al. 2015). The only non-LTE paper is Sk{\'u}lad{\'o}ttir et~al. (2015), in which the sulfur abundance was determined for stars in the dwarf spheroidal (dSph) galaxy in Sculptor. The goal of this paper is to study the systematic errors due to the use of the simplifying classical LTE assumption in determining the abundances from lines of Ca~I, Ti~I, Ti~II, Fe~I, and Fe~II in cool giants in a wide metallicity range. We calculated the non-LTE abundance corrections for a large set of lines in the ranges of stellar parameters 4000~K $\le \Teff \le$ 5000~K, $0.5 \le$ log~g $\le 2.5$, $-4 \le$ [Fe/H] $\le 0$ and treated a code for three-dimensional interpolation that allows the non-LTE correction for an individual spectral line and given atmospheric parameters to be obtained online. All the data are publicly available and can be used in the studies of red giants to determine their atmospheric parameters from Fe~I and Fe~II lines and the calcium, titanium, and iron abundances. The non-LTE abundance corrections for Fe~I and Fe~II lines were computed in the literature (Lind et~al. 2012) in a wide range of stellar parameters, but this does not belittle the practical benefits of our work. The point is that the non-LTE results for Fe~I-Fe~II and for other atoms depend strongly on a free parameter used in non-LTE calculations. This is the scaling factor \kH\ to the formulas of Steenbock and Holweger (1984) that were derived based on the classic theory of Drawin (1968) and are used to calculate the rate coefficients for excitation and ionization of atoms by inelastic collisions with neutral hydrogen atoms. This approach has been repeatedly criticized (see, e.g., Barklem et~al. 2011) for the groundlessness of applying the Drawin (1968) theory to the calculation of inelastic collisions with H~I atoms, however, we continue to use it in calculating the statistical equilibrium, because for most atoms there are neither laboratory measurements nor calculations of the cross sections for these processes. As shown by Mashonkina et~al. (2016), applying the Al~I + H~I collision cross sections obtained in quantum-mechanical calculations gives an advantage in analyzing the Al~I lines in stellar spectra compared to the formulas from Steenbock and Holweger (1984), but, at the same time, if there are no accurate data, it is better to take into account, even if approximately, the inelastic collisions with H~I than to ignore them. Lind et~al. (2012) used the classic Drawin rates with \kH\ = 1 in their calculations. A different estimate (\kH\ = 0.5) was obtained by Sitnova et~al. (2015) when analyzing the Fe~I and Fe~II lines in a sample of dwarf stars in the range $-2.6 \le$ [Fe/H] $\le 0.2$. Our analysis of the iron lines in giants in the Sculptor dSph (see Section 3.3.1) confirmed the estimate by Sitnova et~al. (2015). Therefore, the non-LTE calculations for Fe~I-Fe~II were performed with \kH\ = 0.5. The paper is structured as follows. The methods of calculations are described in Sect.~2. The departures from LTE for Ca~I, Ti~I-Ti~II, and Fe~I-Fe~II lines depending on atmospheric parameters are studied in Sect.~3. Section 4 provides the methodical recommendations and instructions for interpolation of the non-LTE corrections for given line and atmospheric parameters. | 16 | 9 | 1609.02731 |
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1609 | 1609.00734_arXiv.txt | Gravity waves are primarily generated in the lower atmosphere, and can reach thermospheric heights in the course of their propagation. This paper reviews the recent progress in understanding the role of gravity waves in vertical coupling during sudden stratospheric warmings. Modeling of gravity wave effects is briefly reviewed, and the recent developments in the field are presented. Then, the impact of these waves on the general circulation of the upper atmosphere is outlined. Finally, the role of gravity waves in vertical coupling between the lower and the upper atmosphere is discussed in the context of sudden stratospheric warmings. | The lower atmosphere, where meteorological processes take place, is the primary source of internal atmospheric waves: gravity waves (GWs), planetary (Rossby) waves, and solar tides. These waves can propagate upward and influence the dynamics and thermal state of the middle and upper atmosphere \citep[see e.g., the reviews of][]{FrittsAlexander03, Lastovicka06, YigitMedvedev15}. Waves transfer their energy and momentum to the mean flow via breaking and dissipative processes, such as radiative damping, eddy viscosity, nonlinear diffusion, molecular diffusion and thermal conduction, and ion drag \citep{Yigit_etal08}. Sudden stratospheric warmings (SSWs) are spectacular events that disturb the circulation in the winter hemisphere. They affect not only the stratosphere, but their influence extends to the mesosphere and thermosphere. In the upper atmosphere, plasma processes, such as Joule and auroral heating, ion friction are important processes that shape the morphology and dynamics. Thus, interactions between the lower and upper atmosphere should be considered within the framework of the atmosphere--ionosphere system. Such coupled upper atmosphere--ionosphere system is subject to the following internal and external influences: \begin{itemize} \item Meteorological effects that encompass internal wave impacts and transient processes of lower atmospheric origin, \item Internal processes due to nonlinearity, \item Space weather effects that are associated with the solar and magnetospheric phenomena. \end{itemize} Among the meteorological effects, we distinguish a direct influence of internal GWs on the upper regions of the atmosphere. Although transient events such as SSWs are technically categorized as stratospheric processes, and, thus take place above the region of weather-dominated phenomena, they are often referred to as meteorological effects in the context of the upper atmosphere research. The thermosphere--ionosphere system is highly nonlinear. In the real atmosphere, ion and neutral parameters vary simultaneously, and the resulting changes in the heating ought to contain higher order terms, which is indicative of the nonlinear nature of the system \citep{YigitRidley11a}. The atmosphere-ionosphere system is subject to the influence of space weather, which can enhance these nonlinear processes and impact the upper atmosphere \citep[][and references therein]{Prolss11}. In this paper, we report on the recent advances in understanding the meteorological effects in the upper atmosphere, focusing primarily on the links between SSWs, small-scale GWs, and thermosphere--ionosphere dynamics. | \label{sec:conc} This paper has briefly reviewed the current state of knowledge and most recent developments with understanding the role of GWs in vertical coupling during SSWs. The observed upper atmosphere changes during SSWs have been presented. An emphasis was placed on the processes above the mesopause, and on how they can be studied with general circulation models. The geosciences community increasingly recognizes that the effects of lower atmospheric gravity waves extend beyond the middle atmosphere into the atmosphere-ionosphere system and are of global nature. Similarly, sudden stratospheric warmings were used to be looked upon as stratospheric phenomena, but now compelling observational evidences of their signatures in the thermosphere-ionosphere are being routinely provided. With the rapid progress in the field of atmospheric coupling, further key science questions on the role of GWs in coupling atmospheric layers arise: \begin{itemize} \item What are the spectra of gravity waves in the lower and upper atmosphere? How do they change during the different phases of SSWs? \item How well do GW parameterizations describe wave spectra and reproduce their effects during SSWs? \item What is the relative role of GW- and electrodynamical coupling between atmospheric layers in the variability of the atmosphere-ionosphere system during SSWs? \item What are the effects of GWs on the circulation and thermal budget of the upper atmosphere during major sudden stratospheric warmings? \item Do GWs in the upper atmosphere affect the development of sudden stratospheric warmings, or they are a mere reflection of processes in the lower atmosphere? \item Do GWs have a role in latitudinal coupling in the thermosphere during SSW events? \end{itemize} This is certainly an incomplete list of scientific questions, answering which requires concerted observational, theoretical, and modeling efforts from scientists of both lower and upper atmosphere communities. | 16 | 9 | 1609.00734 |
1609 | 1609.07322_arXiv.txt | It has been suggested that single and double jets observed emanating from certain astrophysical objects may have a purely gravitational origin. We discuss new classes of plane-fronted and pulsed gravitational wave solutions to the equation for perturbations of Ricci-flat spacetimes around Minkowski metrics, as models for the genesis of such phenomena. These solutions are classified in terms of their {\it chirality} and generate a family of non-stationary spacetime metrics. Particular members of these families are used as backgrounds in analysing time-like solutions to the geodesic equation for test particles. They are found numerically to exhibit {\it both} single and double jet-like features with dimensionless aspect ratios suggesting that it may be profitable to include such backgrounds in simulations of astrophysical jet dynamics from rotating accretion discs involving electromagnetic fields. | Many astrophysical phenomena find an adequate explanation in the context of Newtonian gravitation and Einstein's description of gravitation is routinely used (together with Maxwell's theory of electromagnetism and the use of time-like spacetime geodesics to model the histories of massive point test particles) to analyse a vast range of phenomena where non-Newtonian effects are manifest. However, there remain a number of intriguing astrophysical phenomena suggesting that our current understanding is incomplete. These include the large scale dynamics of the observed Universe and a detailed dynamics of certain compact stellar objects interacting with their environment. In this note we address the question of the dynamical origin of the extensive ``cosmic jets'' that have been observed emanating from a number of compact rotating sources. Such jets often contain radiating plasmas and are apparently the result of matter accreting on such sources in the presence of magnetic fields. One of the earliest models to explain these processes suggested that the gravitational fields of rotating black holes surrounded by a magnetised ``accretion disc'' could provide a viable mechanism \cite{blandford_z}. More recently, the significance of magneto-hydrodynamic processes in transferring angular momentum and energy into collimated jet structures has been recognised \cite{blandford,lynden,pringle}. Many of these models implicitly assume the existence of a magnetosphere in a {\it stationary} gravitational field and employ ``force-free electrodynamics'' in their development. To our knowledge, a dynamical model that fully accounts for all the observed aspects of astrophysical jets does not exist. \\ However in recent years there has been mounting evidence, both theoretical and numerical, suggesting that the genesis of jets may have a purely gravitational origin. By the genesis of such phenomena we mean a mechanism that initiates the plasma collimation process whereby electrically charged matter arises from initial distributions of neutral matter in a background gravitational field. In \cite{mashhoonPRD,mashhoonPLA}, the authors carefully analyse the properties of a class of Ricci-flat cylindrically symmetric spacetimes possessing time-like and null geodesics that approach attractors confining massive particles to cylindrical spacetime structures. Additional studies \cite{mashhoon_cosmic,mashhoon_peculiar,mashhoon_tidal} of the asymptotic behaviour of test particles on time-like geodesics with large Newtonian speeds relative to a class of co-moving observers have given rise to the notion of {\it cosmic jets} associated with different types of gravitational collapse scenarios satisfying certain Einstein-Maxwell field systems. There has also been a recent approach based on certain approximations within a linearised gravitational framework involving ``gravito-magnetic fields'' generated by non-relativistic matter currents \cite{poirier}. All these investigations auger well for the construction of models for astrophysical jets that include non-Newtonian gravitational fields as well as electromagnetically induced plasma interactions. \\ Although astrophysical jets involve both gravitational and electromagnetic interactions with matter it is natural to explore the structure of electrically neutral test particle geodesics in non-stationary, anisotropic background metric spacetimes as a first approximation to what is undoubtedly a complex dynamical process. In this paper we explore geodesics that relax any assumption of a cylindrically Killing-symmetric background spacetime metric. Furthermore, to this end we construct particular exact solutions to the {\it linearised Einstein vacuum equations} which are then used to numerically calculate time-like geodesics in non-stationary backgrounds. The use of the linearised Einstein vacuum equations facilitates the construction of families of complex eigen-solutions with definite {\it chirality} that are used to construct real spacetime metrics exhibiting families of time-like geodesics possessing particular jet-like characteristics on space-like hyper-surfaces. Test particles on such time-like geodesics exhibit, in general, a well defined sense of ``handed-ness'' in space that we argue may offer a mechanism that initiates a flow of matter into directed jet-like structures. In particular, we construct families of plane-fronted gravitational wave metrics and new non-stationary metrics having propagating {\it pulse-like} characteristics with bounded components in three-dimensional spatial domains. These are analogous to particular exact solutions of the vacuum Maxwell equations which we have recently shown can be used to model single-cycle electromagnetic laser pulses \cite{goto_lasers_JPA,goto_lasers_NIMB}. \\ In section~\ref{sect:LinEin}, our notation is established in the context of Einstein's vacuum equations on a spacetime manifold $\MAN{\wh{g}}$ with a Lorentzian metric $\wh{g}$ and their linearisation about a flat Minkowski metric $\eta$ on a region $\mathcal{U}\subset\MAN{\wh{g}}$. \\ Section~\ref{sect:GravWaves} describes a family of complex plane-fronted gravitational wave exact solutions to the linearised field equations that are eigen-tensors of the complex axial-symmetry operator about any particular direction in space. We call the eigenvalues of this operator ``chirality'' and show how the associated ``helicity'' of the real-part solutions with respect to their direction of propagation is correlated with certain properties of time-like geodesics associated with the real linearised spacetime metric. The simplest {\it harmonic} plane-fronted gravitational wave modes have helicity $\pm 2$ and, in general, induce a (modulated) non-circular helicoidal motion of massive test-particles that are initially arranged (with non-zero speeds) in a ring that lies in a plane orthogonal to the direction of propagation of such waves. The set composed of the individual particle paths form a {\it uni-directional} bi-jet-like array in space emanating from such a ring in a background harmonic wave of definite helicity.\\ By contrast in section~\ref{sect:compact}, we give a construction of exact solutions to the linearised equations with pulse-like perturbations that are bounded in all spatial dimensions. The simplest gravitational pulse-like solutions have chirality zero and are shown to generate a real non-singular Ricci curvature scalar field on $\mathcal{U}$ with well defined loci in spacetime emanating from a core containing a maximum at a particular event. Future-pointing time-like geodesics that emanate from events on rings containing massive test particles centred on this core can give rise to single and {\it oppositely directed} jet-like arrays in space, transverse to the plane of the rings (bipolar outflows).\\ It is assumed that exact analytic solutions to the linearised Einstein vacuum equations and their associated time-like geodesics will exhibit features that persist to some degree beyond the linearisation regime and, in particular, offer an approach to a better understanding of the genesis of observable cosmic jets in models that include charged matter with plasma interactions. Based on a numerical exploration of particular time-like geodesics associated with background metrics constructed from complex eigen-solutions with definite chirality, we conclude that it may be profitable to include these non-Newtonian gravitational backgrounds in simulations of cosmic jet dynamics from rotating accretion discs involving electromagnetic fields. \\ \newpage | 16 | 9 | 1609.07322 |
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1609 | 1609.09109_arXiv.txt | William P. Bidelman---Editor of these {\it Publications\/} from 1956 to 1961---passed away on 2011 May 3, at the age of 92. He was one of the last of the masters of visual stellar spectral classification and the identification of peculiar stars. I review his contributions to these subjects, including the discoveries of barium stars, hydrogen-deficient stars, high-galactic-latitude supergiants, stars with anomalous carbon content, and exotic chemical abundances in peculiar A and B stars. Bidelman was legendary for his encyclopedic knowledge of the stellar literature. He had a profound and inspirational influence on many colleagues and students. Some of the bizarre stellar phenomena he discovered remain unexplained to the present day. | 16 | 9 | 1609.09109 |
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