ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2" The corresponding ratios of (comoving) magnetic and radiation energy densities at the start of simulations are UpCR,)/U.(R,)= 0.006. 0.06. 0.6. 3 and 12."," The corresponding ratios of (comoving) magnetic and radiation energy densities at the start of simulations are $U_{\rm B}(R_{\rm n})/U_\gamma(R_{\rm n}) = 0.006$ , $0.06$ , $0.6$ , $3$ and $12$ ."3 The results ofour calculations are summarized in Table 1. and deseribed in detailbelow., The results ofour calculations are summarized in Table \ref{tbl-1} and described in detailbelow.4 Strong magnetic fields imply significant. synchrotron cooling of the high-energyc pairs injected by nuclearcollisions. which can compete with Compton cooling.," Strong magnetic fields imply significant synchrotron cooling of the high-energy$e^\pm$ pairs injected by nuclearcollisions, which can compete with Compton cooling."5 Since the synchrotron photons do not ereate secondary. pairs. the magnetic field has a suppressing effect on the&* cascade.," Since the synchrotron photons do not create secondary pairs, the magnetic field has a suppressing effect on the$e^\pm$ cascade."6 The multiplicity of each subsequent pair generation in a saturated, The multiplicity of each subsequent pair generation in a saturated7where the ratio of stellar (lux to accretion disk flux at 5000 iis about 10.000.,"where the ratio of stellar flux to accretion disk flux at 5000 is about $10,000$."8 Consequently. we do not include any optical light from an accretion disk.," Consequently, we do not include any optical light from an accretion disk."9 We (urn to the question of (he X-ray heating of (he supergiant star and its effect on the binary model., We turn to the question of the X-ray heating of the supergiant star and its effect on the binary model.10 The X-ray. heating is computed using the technique outlined in Wilson(1990)., The X-ray heating is computed using the technique outlined in \cite{wil+1990}.11. The X-ray. source geometry is assumed (o be a (hin disk in the orbital plane wilh a radius vanishinglv small compared to the semimajor axis (this structure should not be confused with much larger accretion disk (hat potentially could be a source of optical flux)., The X-ray source geometry is assumed to be a thin disk in the orbital plane with a radius vanishingly small compared to the semimajor axis (this structure should not be confused with much larger accretion disk that potentially could be a source of optical flux).12 Points on the stellar surface “see” the X-ray. source at inclined aneles. and the proper foreshortening is accounted for.," Points on the stellar surface “see” the X-ray source at inclined angles, and the proper foreshortening is accounted for."13" Measurement of the broadband X-ray luminosity of Cvenus X-1 (La,,: hereafter in units of 10"" erg 1, adjusted to the revised distance of 1.86 kpc) requires special instrumentation and considerations."," Measurement of the broadband X-ray luminosity of Cygnus X-1 $L_{\rm xbol}$; hereafter in units of $10^{37}$ erg $^{-1}$, adjusted to the revised distance of 1.86 kpc) requires special instrumentation and considerations."14 There are soft ancl hard states of Cygnus X-1 (e.g.. Gou et 22011).," There are soft and hard states of Cygnus X-1 (e.g., Gou et 2011)."15 The hard state is especially challenging because the effective temperature of the accretion disk is relatively low (2< 0.5 keV). while the hard power-law component (with photon index ~ 1.7) must be integrated past the cutoff energy 150 keV (Gierlinski et al.," The hard state is especially challenging because the effective temperature of the accretion disk is relatively low $T <$ 0.5 keV), while the hard power-law component (with photon index $\sim 1.7$ ) must be integrated past the cutoff energy $\sim 150$ keV (Gierlinski et al."16 1997: (πο Bel et 22006)., 1997; Cadolle Bel et 2006).17 Since the ground. based. observations (ie. photometric data and. radial velocity measurements) are distributed over many. vears. both the range and the long-term average of the X-ray luminosity must be estimated.," Since the ground based observations (i.e., photometric data and radial velocity measurements) are distributed over many years, both the range and the long-term average of the X-ray luminosity must be estimated."18 The archive of the Rossi X-vav Timing Explorer (RXTE) contains several thousand observations of Cyenus X-1 collected. in numerous monitoring campaigns conducted over the life of the mission., The archive of the Rossi X-ray Timing Explorer (RXTE) contains several thousand observations of Cygnus X-1 collected in numerous monitoring campaigns conducted over the life of the mission.19 We processed and analyzed 2343 exposure intervals (1996 January io 2011 February: mean exposure 2.2 ks) with the PCA instrument. and we then computed normalized light curves in [our energy bands (2-18 keV) in the manner described by Remillard AeClintock (2006).," We processed and analyzed 2343 exposure intervals (1996 January to 2011 February; mean exposure 2.2 ks) with the PCA instrument, and we then computed normalized light curves in four energy bands (2-18 keV) in the manner described by Remillard McClintock (2006)."20 In the harduiess-intensitv. diagram. the soft and hard states of (νο X-1 can be separated by the value of hard color (HC: ie.. the ratio of the normalized PCA count rates al 8.6-18.0. versus 5.0-8.6 keV). using a simple discrimination line at {16=0.7.," In the hardness-intensity diagram, the soft and hard states of Cyg X-1 can be separated by the value of hard color $HC$; i.e., the ratio of the normalized PCA count rates at 8.6-18.0 versus 5.0-8.6 keV), using a simple discrimination line at $HC = 0.7$."21 On this basis. we determine that Cvgnus N-1 is found to be in the hard state of the (ime.," On this basis, we determine that Cygnus X-1 is found to be in the hard state of the time."22 Zhang et ((1997) studied the broadband spectra of Cvgnus N-1 with instruments of RATE and the Compton Gamma Rav Observatory (CORO)., Zhang et (1997) studied the broadband spectra of Cygnus X-1 with instruments of RXTE and the Compton Gamma Ray Observatory (CGRO).23 During the hard and soft states ol 1996. thev found broad-hancl X-ray luminosities in the range 1.6-2.2 for the hard state and 2.2-3.3 for the soft state.," During the hard and soft states of 1996, they found broad-band X-ray luminosities in the range 1.6-2.2 for the hard state and 2.2-3.3 for the soft state."24 Samples of the hard aud intermediate states during 2002-2004 with the International Gamma-Ray Astrophysics Laboratory INTEGRAL) vielded Ly.) in the range 1.2-2.0 (Cadolle Bel et 22006)., Samples of the hard and intermediate states during 2002-2004 with the International Gamma-Ray Astrophysics Laboratory (INTEGRAL) yielded $L_{\rm xbol}$ in the range 1.2-2.0 (Cadolle Bel et 2006).25 Additional measurements of the soft state with, Additional measurements of the soft state with26the N-rav data in red and the radio data in green.,the X-ray data in red and the radio data in green.27 We fud no enhancement of the X-ray cussion near the radio relic. which is differeut from what was secu for Abell 133 (Fujitactal.2002).," We find no enhancement of the X-ray emission near the radio relic, which is different from what was seen for Abell 133 \citep{fsk+02}."28. (ναι the low surface brightuess of the N-vay eimidssion. it is dificult to determine if there is a deficit of N-ray cussion at the position of the racio relic. commonly referred to as an N-rav cavity or radio bubble.," Given the low surface brightness of the X-ray emission, it is difficult to determine if there is a deficit of X-ray emission at the position of the radio relic, commonly referred to as an X-ray cavity or radio bubble."29 The absolute astrometric accuracy of the radio aud N-ray maps is roughly 0755 for each., The absolute astrometric accuracy of the radio and X-ray maps is roughly 5 for each.30 The radio poiut source near the cluster center is not associated with galaxy IT. lustead. it seenis coimcideut with a barely resolved pair of ealaxies in the DSS and Two Micron. All ον Survey (2Q\TASS) images. listed in ἂν 2MASN. 1980067.," The radio point source near the cluster center is not associated with galaxy H. Instead, it seems coincident with a barely resolved pair of galaxies in the DSS and Two Micron All Sky Survey (2MASS) images, listed in as 2MASX $-$ 1930067."31 No velocity information is available for 2\IASX 1930067. making association with the cluster uncertain.," No velocity information is available for 2MASX $-$ 1930067, making association with the cluster uncertain."32 To euliauce the structures in the cluster core. we fitted the observed cluster cussion with a smooth elliptical isophotal model. usingthe IRAF/STSDAS taskellipse.," To enhance the structures in the cluster core, we fitted the observed cluster emission with a smooth elliptical isophotal model, using the IRAF/STSDAS task."33 We allowed the ellipticity. position angle. aud iuteusities of cach isophote to vary while keeping the ceutroids fixed to the position of galaxy IL The model was then subtracted from the data after multiplication of the model counts by 0.5 to avoid oversubtraction.," We allowed the ellipticity, position angle, and intensities of each isophote to vary while keeping the centroids fixed to the position of galaxy H. The model was then subtracted from the data after multiplication of the model counts by 0.5 to avoid oversubtraction."34 Iu Figure L. we show the subtracted imaee of the cluster in the region around the radio relie with the radio contours overlaid.," In Figure \ref{fig:comp}, we show the subtracted image of the cluster in the region around the radio relic with the radio contours overlaid."35 No excess enission Is associated with the radio relie., No excess emission is associated with the radio relic.36 The cluster cutission is faint near the radio relic. aud coutinues to weaken at projected cluster-centric distances ereater than that of the relic in the same direction.," The cluster emission is faint near the radio relic, and continues to weaken at projected cluster-centric distances greater than that of the relic in the same direction."37 More interesting structure is seen in the cluster ceuter., More interesting structure is seen in the cluster center.38 Figure 5. shows the siue müage as Fieure |. with a differeut color scale to enhance the variatious iu the inner region of the cluster.," Figure \ref{fig:compcen} shows the same image as Figure \ref{fig:comp}39 with a different color scale to enhance the variations in the inner region of the cluster."40 The knot of ciission iu the upper right is the secoud-brightest elliptical galaxy. F. The brisht vellow knot of cmission is coincident with ealaxy UL. There is a more diffuse bright extension of cluission to the west of galaxy IL. Iu addition. there is other structure iu the cluster center perpendicular to the bright extension.," The knot of emission in the upper right is the second-brightest elliptical galaxy, F. The bright yellow knot of emission is coincident with galaxy H. There is a more diffuse bright extension of emission to the west of galaxy H. In addition, there is other structure in the cluster center perpendicular to the bright extension."41 Iuterestingly. the emission exteusiou to the west of galaxy II secuis to poiut in the same direction as oue night expect the radio relic to extend if it were projected back to galaxy II. Enhanced X-ray enmdsson has been seen in the large elliptical galaxies iu the Coma cluster (Vikliliuiuctal. 2001).," Interestingly, the emission extension to the west of galaxy H seems to point in the same direction as one might expect the radio relic to extend if it were projected back to galaxy H. Enhanced X-ray emission has been seen in the large elliptical galaxies in the Coma cluster \citep{vmf+01}."42. We compared the euhauced euissiou of galaxy F with the results from Coma., We compared the enhanced emission of galaxy F with the results from Coma.43 First. we determined the radial profile of the cuhancement to determine its size and the appropriate source region for galaxy F. The chussion was clearly more extended than the point spread function.," First, we determined the radial profile of the enhancement to determine its size and the appropriate source region for galaxy F. The emission was clearly more extended than the point spread function."44 We fouud that the enhancement mncasured z1755 Pain radius. or roughly 3 kpe at the distance of Abell 13.," We found that the enhancement measured $\approx$ in radius, or roughly 3 kpc at the distance of Abell 13."45 Vikhliuinetal.(2001) found that the compact cores associated with the brightest galaxies in the Coma cluster had a plivsical size of 3 kpc iu radius., \citet{vmf+01} found that the compact cores associated with the brightest galaxies in the Coma cluster had a physical size of 3 kpc in radius.46" We attempted to extract a spectrum of the euliauced enissiou associated with galaxy F frou a circular region with radius of2"".", We attempted to extract a spectrum of the enhanced emission associated with galaxy F from a circular region with radius of.47. To accouut for both the cosmic N-rav backeround and the cluster enmüsson. we used a local background extracted from an aunulus ceutered ou," To account for both the cosmic X-ray background and the cluster emission, we used a local background extracted from an annulus centered on"48bands that are the hallinarks of the T spectral class.,bands that are the hallmarks of the T spectral class.49 We assigned spectral (vpes by measuring the five flux indices delined by ? and then applying the polynomial relations of ?.., We assigned spectral types by measuring the five flux indices defined by \citet{Burgasser2006} and then applying the polynomial relations of \citet{Burgasser2007}.50 We also visually determined spectral (vpes by comparing wilh INTE/Spex prism spectra of the T clwarls spectral standards chosen by ?.., We also visually determined spectral types by comparing with IRTF/Spex prism spectra of the T dwarfs spectral standards chosen by \cite{Burgasser2006}.51" For each object. the depth of the HO and CI, absorption bands were examined. normalizing (he spectra of the objects and the standards to their peak [Iuxesin the J. //. and A bands individually."," For each object, the depth of the $_2$ O and $_4$ absorption bands were examined, normalizing the spectra of the objects and the standards to their peak fluxesin the $J$, $H$, and $K$ bands individually."52 The (vpes from the flux indices agreed well with those from visual examination., The types from the flux indices agreed well with those from visual examination.53 For PSO J226.2590—28.3959. the (wo results disagreed by 1 subclass. therefore we tookthe average {ο assign the final spectral (vpe (Table 4)).," For PSO $-$ 28.8959, the two results disagreed by 1 subclass, therefore we tookthe average to assign the final spectral type (Table \ref{spexclass}) )."54 We calculated photometric distance estimates for our objects based on spectral tvpe vs. absolute magnitude relations., We calculated photometric distance estimates for our objects based on spectral type vs. absolute magnitude relations.55 For PSOJ247.32134-03.59327.. PSO J246.42224-15.4608 and PSO J226.2599-28.8959 we used our UINXIRT photometry and the relations of ?..," For PSO, PSO J246.4222+15.4698 and PSO J226.2599-28.8959 we used our UKIRT photometry and the relations of \cite{Liu2006}."56 This paper «quotes both a bright relation which includes suspected. and confined binaries and a [nl relation which excludes (hese., This paper quotes both a bright relation which includes suspected and confirmed binaries and a faint relation which excludes these.57 The disagreement between these relations is most pronounced in the mid-Ts at approximately one magnitude., The disagreement between these relations is most pronounced in the mid-Ts at approximately one magnitude.58 As the binaritv of our sample of four objects is unknown. we use an average of the bright ancl faint relations to estimate the intrinsic brightness of our objects.," As the binarity of our sample of four objects is unknown, we use an average of the bright and faint relations to estimate the intrinsic brightness of our objects."59 For PSO J201.0320+19.1072 we used 2MAÀSS photometry and the relations of 2? converted into the 2ATASS svstem using the relations of ?.., For PSO J201.0320+19.1072 we used 2MASS photometry and the relations of \cite{Liu2006} converted into the 2MASS system using the relations of \cite{Stephens2004}.60 The distance estimates for each filler were (hen averaged to produce the values shown in Table 52010.0.., The distance estimates for each filter were then averaged to produce the values shown in Table \ref{objphotometry}.61 Three of our objects PSO J201.03204-19.1012. PSO J247.3273+03.5932 and PSO J226.2599-28.8959 have distance estimates within 25 pe wilh PSO J246.42224-15.4698 Ivine further away at 31.3 pe.," Three of our objects PSO J201.0320+19.1072, PSO J247.3273+03.5932 and PSO J226.2599-28.8959 have distance estimates within 25 pc with PSO J246.4222+15.4698 lying further away at 31.3 pc."62 Our four objects are relatively bright and thus could have been detected by previous work such as studies of the 2\LASS survey (? and references therein) the Sloan Digital Sky survey (5D55: ?.. ?.. ancl relerences therein) and the UIXKIRT Deep Sky Survey (UKNIDS5.," Our four objects are relatively bright and thus could have been detected by previous work such as studies of the 2MASS survey \citealt{Burgasser2004} and references therein) the Sloan Digital Sky Survey (SDSS; \citealt{York2000}, , \citealt{Chiu2006}, , and references therein) and the UKIRT Deep Sky Survey (UKIDSS,"63?..,\citet{juett}.64 The residuals show that besides the higher than solar abundance of neon there is a broad emission feature present in the data with a peak at the position of the O VIII Ίνα line., The residuals show that besides the higher than solar abundance of neon there is a broad emission feature present in the data with a peak at the position of the O VIII $\alpha$ line.65 observed 4U 06144091 on 2001 March 13 starting at 9:35 UT for ~ || Ksee and immediately afterwards at 12:32 UT for ~ 17 ksec., observed 4U 0614+091 on 2001 March 13 starting at 9:35 UT for $\sim$ 11 ksec and immediately afterwards at 12:32 UT for $\sim$ 17 ksec.66 In both parts of the observation two Reflection Grating Spectrometers (RGS) were collecting data., In both parts of the observation two Reflection Grating Spectrometers (RGS) were collecting data.67 RGS is a high-resolution spectrometer covering the energy range from 7 to 38 (0.3—2.1 keV) with spectral resolution of Ε/ΔΕ —100 to 500 (2).., RGS is a high-resolution spectrometer covering the energy range from 7 to 38 $0.3-2.1$ keV) with spectral resolution of $E/\Delta E=$ 100 to 500 \citep{denHerder}.68 In the second part of the observation the European Photon Imaging Cameras: MOSI. MOS? and pn were collecting data simultaneously with RGS.," In the second part of the observation the European Photon Imaging Cameras: MOS1, MOS2 and pn were collecting data simultaneously with RGS."69 MOSI. ΜΩΡΟ and pn provide broad band spectral coverage with modest resolution CE/AE| —20 to 50) over the energy range of 0.3 to 10 keV (2)..," MOS1, MOS2 and pn provide broad band spectral coverage with modest resolution $E/\Delta E=$ 20 to 50) over the energy range of 0.3 to 10 keV \citep{turner}."70 The MOSI and pn camera were operated in the Timing mode. while MOS? was operated in the Full Frame (Imaging) We reduce the data using the Data Analysis software SAS version 9.0.," The MOS1 and pn camera were operated in the Timing mode, while MOS2 was operated in the Full Frame (Imaging) We reduce the data using the Data Analysis software SAS version 9.0."71 The light curve shows no signiticant contamination from soft protons., The light curve shows no significant contamination from soft protons.72 We extract source photons with pixel pattern equal to O from the MOSI camera with from 300 to 320 and background with from 260 to 280 (beyond he source PSF)., We extract source photons with pixel pattern equal to 0 from the MOS1 camera with from 300 to 320 and background with from 260 to 280 (beyond the source PSF).73 The net source count rate is 58.54+0.06 c/s (full bandpass)., The net source count rate is $58.54\pm0.06$ c/s (full bandpass).74 We extract the MOS? observation with pixel pattern below or equal to 12., We extract the MOS2 observation with pixel pattern below or equal to 12.75 The MOS? observation sutters from pile-up., The MOS2 observation suffers from pile-up.76 In order to reduce the etfect of pile-up on the source spectrum. we exclude events from within the circle with a radius of 24 arcsec centered on the source position.," In order to reduce the effect of pile-up on the source spectrum, we exclude events from within the circle with a radius of 24 arcsec centered on the source position."77 The background spectrum for he MOS? observation is extracted from CCD-3., The background spectrum for the MOS2 observation is extracted from CCD-3.78 The net source count rate is [3.01+0.03 c/s. We extract source photons with pixel yattern less than 5 from the pn camera withRAwx from 33 to 45 and background with from 15 to 25., The net source count rate is $13.01\pm0.03$ c/s. We extract source photons with pixel pattern less than 5 from the pn camera with from 33 to 45 and background with from 15 to 25.79 The net source count rate is 238.74€0.14 c/s. The data collected by RGS are reduced using standard software pipeline which generates source and background spectra as well as response files., The net source count rate is $238.74\pm0.14$ c/s. The data collected by RGS are reduced using standard software pipeline which generates source and background spectra as well as response files.80 We fit the data using the (2)., We fit the data using the \citep{kaastra1996}.81 Errors on the fit parameters reported throughout this Manuscript correspond to a confidence level for each calculated parameter (Ay?=1)., Errors on the fit parameters reported throughout this Manuscript correspond to a confidence level for each calculated parameter $\Delta\chi^2=1$ ).82 The typical continuum model for X-ray binaries contains two components: a power-law. which is important in the island state and a black body. which is important in the banana state of the source.," The typical continuum model for X-ray binaries contains two components: a power-law, which is important in the island state and a black body, which is important in the banana state of the source."83 In our case adding a soft thermal component (black body) does not improve the fit. therefore we use only a power-law to fit the data.," In our case adding a soft thermal component (black body) does not improve the fit, therefore we use only a power-law to fit the data."84 During the observation the source was in the island state (?).. hence it is not surprising that the soft thermal component is not present.," During the observation the source was in the island state \citep{mendez2002}, hence it is not surprising that the soft thermal component is not present."85 The X-ray spectrum of the source is atfected by the interstellar medium., The X-ray spectrum of the source is affected by the interstellar medium.86 Most of the gas in the line of sight is neutral (?).. therefore as to a first approximation we use only a neutral absorption model(Hor model with very low temperature-neutral gas limit).," Most of the gas in the line of sight is neutral \citep{ferriere}, therefore as to a first approximation we use only a neutral absorption model model with very low temperature-neutral gas limit)."87 This model calculates the transmission of a plasma in collisional ionization equilibrium., This model calculates the transmission of a plasma in collisional ionization equilibrium.88 For the reference abundances of elements in the neutral gas we choose the proto-solar abundances of ?.., For the reference abundances of elements in the neutral gas we choose the proto-solar abundances of \citet{lodders}.89 The data from the MOSI. MOS? and pn cameras do not agree in the soft part of the spectrum.," The data from the MOS1, MOS2 and pn cameras do not agree in the soft part of the spectrum."90 It was reported that the timing mode calibration of MOS has larger uncertainties than the imaging mode calibration (2)., It was reported that the timing mode calibration of MOS has larger uncertainties than the imaging mode calibration \citep{intZand}.91 The differences between MOS? and pn spectra below | keV may indicate also uncertainties in the pn timing mode calibration., The differences between MOS2 and pn spectra below 1 keV may indicate also uncertainties in the pn timing mode calibration.92 Therefore we focus on the MOS? data only., Therefore we focus on the MOS2 data only.93 In the case of the TOS] and pn data we report only the tit in the range |—10 keV. The fit gives a slope of the power-law of 2.17340.004 with y;=2.6 or 242 d.o.f., In the case of the MOS1 and pn data we report only the fit in the range $1-10$ keV. The fit gives a slope of the power-law of $2.173\pm0.004$ with $\chi^2_{\nu}=2.6$ for 242 d.o.f.94 for MOSI and 2.216+0.001 with yz=2.1 for 735 d.o.f., for MOS1 and $2.216\pm0.001$ with $\chi^2_{\nu}=2.1$ for 735 d.o.f.95 for the pn., for the pn.96 We tit MOS? in the range 0.4—10 keV assuming an interstellar medium (SM) with proto-solar abundances., We fit MOS2 in the range $0.4-10$ keV assuming an interstellar medium (ISM) with proto-solar abundances.97 The best-fit model leaves a large residual in the soft X-ray part of the spectrum (Fig., The best-fit model leaves a large residual in the soft X-ray part of the spectrum (Fig.98 | upper left panel).The possible solution proposed by ? is to allow the oxygen and neon abundances. whose absorption," 1 upper left panel).The possible solution proposed by \citet{juett} is to allow the oxygen and neon abundances, whose absorption"99"region before any further run-time refinement are given by Mamret=(1—Q%/OQm)/64c3.53Mo and Megas=(Q¢/Qm)Mam/64~0.72Mo, respectively.","region before any further run-time refinement are given by $M_{\rm dm, ref}=(1-\Omega_b/\Omega_m)M_{\rm dm}/64\simeq 3.53\,{\rm M}_\odot$ and $M_{\rm gas}=(\Omega_b/\Omega_m)M_{\rm dm}/64\simeq 0.72\,{\rm M}_\odot$, respectively."100 We use a comoving gravitational softening length of 17pc for the refined DM component.," We use a comoving gravitational softening length of $17\,{\rm pc}$ for the refined DM component."101 We follow the collapse of the gas in the refined minihalos with the cosmological moving mesh code (Springel, We follow the collapse of the gas in the refined minihalos with the cosmological moving mesh code \citep{springel10a}.102 is a second-order accurate finite volume 2010a)..method that solves the Euler equations based on a piece-wise linear reconstruction and the calculation of hydrodynamical fluxes at every cell face with an exact Riemann solver., is a second-order accurate finite volume method that solves the Euler equations based on a piece-wise linear reconstruction and the calculation of hydrodynamical fluxes at every cell face with an exact Riemann solver.103 The principle difference of compared to Eulerian mesh codes is that its computational mesh is constructed as the Voronoi tessellation of a set of mesh-generating points., The principle difference of compared to Eulerian mesh codes is that its computational mesh is constructed as the Voronoi tessellation of a set of mesh-generating points.104" These points can be moved with the flow velocity itself, making the mesh automatically adaptive in a Lagrangian fashion."," These points can be moved with the flow velocity itself, making the mesh automatically adaptive in a Lagrangian fashion."105" This technique greatly reduces the numerical diffusivity of mesh-based hydrodynamics, especially when large bulk flows are present."," This technique greatly reduces the numerical diffusivity of mesh-based hydrodynamics, especially when large bulk flows are present."106" In fact, the results of the code become fully Galilean-invariant, whereas the truncation error of ordinary Eulerian mesh codes depends on the bulk velocity of the system."," In fact, the results of the code become fully Galilean-invariant, whereas the truncation error of ordinary Eulerian mesh codes depends on the bulk velocity of the system."107" In addition, the unstructured Voronoi mesh of avoids the introduction of preferred directions, which are present in Cartesian meshes."," In addition, the unstructured Voronoi mesh of avoids the introduction of preferred directions, which are present in Cartesian meshes."108 The novelAREPO scheme hence combines the accuracy of mesh-based hydrodynamics with the natural adaptivity and translational invariance usually only provided by the smoothed particle hydrodynamics (SPH) technique., The novel scheme hence combines the accuracy of mesh-based hydrodynamics with the natural adaptivity and translational invariance usually only provided by the smoothed particle hydrodynamics (SPH) technique.109" In terms of hydrodynamical accuracy, the grid-based approach of alleviates a number of shortcomings encountered with SPH (Monaghan 2005)."," In terms of hydrodynamical accuracy, the grid-based approach of alleviates a number of shortcomings encountered with SPH \citep{monaghan05}."110". Among these are the inherent noise of the kernel estimates, the artificial viscosity, and the slow convergence rate of SPH in three dimensions (Springel"," Among these are the inherent noise of the kernel estimates, the artificial viscosity, and the slow convergence rate of SPH in three dimensions \citep{springel10b}."111" Further important improvements ofAREPO lie in the 2010b)..more accurate treatment of shocks and turbulence,"," Further important improvements of lie in the more accurate treatment of shocks and turbulence,"112background. is desirable. but was. prevented. by erowcding.,"background is desirable, but was prevented by crowding."113 Source and background. lighteurves were constructed. fron the pn ancl MOS data in 0.32.5. 2.510 and 0.310 keV bands.," Source and background lightcurves were constructed from the pn and MOS data in 0.3–2.5, 2.5–10 and 0.3–10 keV bands."114 The lighteurves were background subtracted., The lightcurves were background subtracted.115 Source and background. spectra were then created. along with corresponding response matrices and ancillary response files.," Source and background spectra were then created, along with corresponding response matrices and ancillary response files."116 The pn image vielded. 5542 net source photons. while the data from the two MOS cameras were combined. vielding 4893 net source photons.," The pn image yielded 5542 net source photons, while the data from the two MOS cameras were combined, yielding 4893 net source photons."117 The pn and. combined MOS. spectra were modeled simultaneously. using NSPEC yer 11.3.," The pn and combined MOS spectra were modeled simultaneously, using XSPEC ver 11.3."118 Each model consisted of an emission model. with line-of-sight absorption and a normalisation constant that accounts for cillerences in the pn and MOS. callibrations. after. setting. the pn normalisation to 1.," Each model consisted of an emission model, with line-of-sight absorption and a normalisation constant that accounts for differences in the pn and MOS callibrations, after setting the pn normalisation to 1."119 The parameters of the emission models were forced to be the same for the pn and. MOS spectra. but free to vary.," The parameters of the emission models were forced to be the same for the pn and MOS spectra, but free to vary."120 Phree emission. models were initially chosen: blackbody (BB). bremsstrahlung (BR). and power law (PL).," Three emission models were initially chosen: blackbody (BB), bremsstrahlung (BR), and power law (PL)."121 Ehe blackbocds and bremsstrahlung models were characterised by KZ. where 7 is the temperature ancl k is the Boltzmann constant. while power law emission is characterised by photon index. E.," The blackbody and bremsstrahlung models were characterised by $T$, where $T$ is the temperature and k is the Boltzmann constant, while power law emission is characterised by photon index, $\Gamma$."122 Absorption is expressed in terms of Ng. the equivalent absorption by neutral hydrogen.," Absorption is expressed in terms of $N_{\rm H}$, the equivalent absorption by neutral hydrogen."123 We first. modelled the XBo 144 spectra with single component emission models: DD. BR and PL.," We first modelled the XBo 144 spectra with single component emission models: BB, BR and PL."124 We found the pn and. MOS spectra to be best fitted by the PL model. with Ny = Δ.Ε 107 atom 7 and P= 1.48+0.04: all spectral models are provided. in Table 1..," We found the pn and MOS spectra to be best fitted by the PL model, with $N_{\rm H}$ = $\pm$ $\times$ $^{20}$ atom $^{-2}$ and $\Gamma$ = $\pm$ 0.04; all spectral models are provided in Table \ref{specmod}."125 The best fit model is as expected for à LAND in its low state: however. we do not expect such behaviour at X-ray bIuminosities above —0.1 Legg in the 0.01.1000 keV. band (2)...," The best fit model is as expected for a LMXB in its low state; however, we do not expect such behaviour at X-ray luminosities above $\sim$ 0.1 $L_{\rm Edd}$ in the 0.01–1000 keV band \citep{bsh08}."126" Phe 0.3.10 keV luminosity for our best fit PL model is 5.3340.03. 10°"" erg towhich is 0.29. νι for a 1.4 AL. neutron star. and 0.20 Leag for the most massive observed neutron star Al. seee.g.7). Vhis is sullicentlv bright for NBo 144 to be a plausible black hole candidato."," The 0.3–10 keV luminosity for our best fit PL model is $\pm$ $\times$ $^{37}$ erg $^{-1}$, which is 0.29 $L_{\rm Edd}$ for a 1.4 $M_{\odot}$ neutron star, and 0.20 $L_{\rm Edd}$ for the most massive observed neutron star \citep[2.1 M$_{\odot}$ , see e.g. This is sufficently bright for XBo 144 to be a plausible black hole candidate."127 We then modelled the NBo 144. spectra with a wo component (BB|PL) emission. model., We then modelled the XBo 144 spectra with a two component (BB+PL) emission model.128 Phe favoured rlackbods temperature is 0005200010 keV: this tells us hat there is no blackbody component in the 0.310 keV xiuid., The favoured blackbody temperature is $\pm$ 0.0016 keV; this tells us that there is no blackbody component in the 0.3–10 keV band.129 In Figs., In Figs.130 P. and 2.. we compare the PL anc BB|PL emission models with the 0.3.10 keV. pn spectrum. unfolded rom the instrumental responses: we see that the blackbocdy component in Fig.," \ref{plspec} and \ref{bbplspec}, we compare the PL and BB+PL emission models with the 0.3–10 keV pn spectrum, unfolded from the instrumental responses; we see that the blackbody component in Fig."131 2 is set to have minimal contribution to he emission., \ref{bbplspec} is set to have minimal contribution to the emission.132 We also modelled the pn and combined MOS spectra separately with a PL emission. model. to see. if. the instruments are in agreement.," We also modelled the pn and combined MOS spectra separately with a PL emission model, to see if the instruments are in agreement."133 The best fit models are also shown in Table 1: the pn ancl MOS spectra are in excellent agreement., The best fit models are also shown in Table \ref{specmod}; the pn and MOS spectra are in excellent agreement.134 We are therefore confident in our interpretation ol the X-rav emission from Χο 144., We are therefore confident in our interpretation of the X-ray emission from XBo 144.135 The combined pn]MOS. background-subtracted. 0.3.10 keV lighteurve of NBoldd is shown in Fig. 3..," The combined pn+MOS, background-subtracted, 0.3–10 keV lightcurve of XBo144 is shown in Fig. \ref{lc},"136 with 400 s inning., with 400 s binning.137 The intensity appears to vary in the manner of low state LAINBs: however. the uncertainties are large. and the vest fit line of constant intensity vields v col = 630/601. ic. the lighteurve is consistent with being constant.," The intensity appears to vary in the manner of low state LMXBs; however, the uncertainties are large, and the best fit line of constant intensity yields $\chi^2$ /dof = 630/601, i.e. the lightcurve is consistent with being constant."138 The pactional rms., The fractional r.m.s.139 variability is GEL% on time-scales longer han 100 s: this is consistent with a LAIND in the low state. but acds no extra strength to our argument.," variability is $\pm$ on time-scales longer than 100 s; this is consistent with a LMXB in the low state, but adds no extra strength to our argument."140 Without he definite variability exhibited by NBo 45 (?).. the case or Χο 144 being a black hole X-ray binary is somewhat weaker. and rests on its association with the elobular cluster Do 144.," Without the definite variability exhibited by XBo 45 \citep{bsh08}, the case for XBo 144 being a black hole X-ray binary is somewhat weaker, and rests on its association with the globular cluster Bo 144."141 The X-ray source associated with the M31 GC Bo 144. XBo 144. exhibits an emission spectrum that is characteristic ofà LAINB in the low state: the 0.3.10 keV emission is described by à pure power law. with P= I.48:x0.04.," The X-ray source associated with the M31 GC Bo 144, XBo 144, exhibits an emission spectrum that is characteristic of a LMXB in the low state: the 0.3–10 keV emission is described by a pure power law, with $\Gamma$ = $\pm$ 0.04."142 ?. showed that neutron star LMXDs only exhibit this behaviour at, \citet{glad07} showed that neutron star LMXBs only exhibit this behaviour at143refsec:consistency.. we find evidence of a temporally drifting component of the ground signal on I- to 8-hr time scales. subtle but present for all baselines and noticeably stronger for short baselines.,", we find evidence of a temporally drifting component of the ground signal on 1- to 8-hr time scales, subtle but present for all baselines and noticeably stronger for short baselines."144 We therefore apply a linear drift constraint to all visibilities. and a quadratic constraint for [u|<40.," We therefore apply a linear drift constraint to all visibilities, and a quadratic constraint for ${\left| \mathbf{u} \right|} < 40$."145 The additional constraints have little effect on the power spectrum. which makes us confident that sensitivity to ground signal is effectively eliminated.," The additional constraints have little effect on the power spectrum, which makes us confident that sensitivity to ground signal is effectively eliminated."146 As predicted for our experimental configuration ((Tegmark Efstathiou 1996). point sources are the dominant foreground in the DASI data.," As predicted for our experimental configuration \markcite{tegmark96}( (Tegmark Efstathiou 1996), point sources are the dominant foreground in the DASI data."147 To remove point source flux contributions using the constraint matrix formalism above. we require only the positions of the sources.densities.," To remove point source flux contributions using the constraint matrix formalism above, we require only the positions of the sources,."148 We constrain 28 point sources detected in the DASI data itself. in which we can detect an 40 mJy source at beam center with >4.50 significance.," We constrain 28 point sources detected in the DASI data itself, in which we can detect an 40 mJy source at beam center with $>1494.5\,\sigma$ significance."150 The estimated point source flux densities range from 80 mJy to 7.0 Jy., The estimated point source flux densities range from 80 mJy to 7.0 Jy.151 We also constrain point sources from the PMN southern (PMNS) catalog (Wright 1994) with 4.85 GHz flux densities. ὡς. which exceed 50 mJy when multiplied by the DASI primary beam.," We also constrain point sources from the PMN southern (PMNS) catalog \markcite{wright94}( 1994) with 4.85 GHz flux densities, $S_5$, which exceed 50 mJy when multiplied by the DASI primary beam."152 We use this flux density limit for the constrained. point sources because the loss of degrees of freedom resulting from the inclusion of all point sources in the PMNS catalog would be prohibitively large., We use this flux density limit for the constrained point sources because the loss of degrees of freedom resulting from the inclusion of all point sources in the PMNS catalog would be prohibitively large.153 We have tested for the effect of possible absolute pointing error by displacing the point source position templates., We have tested for the effect of possible absolute pointing error by displacing the point source position templates.154 A uniform displacement of the PMNS catalog coordinates by less than our estimated pointing error of 2/ (see Paper D) does not have a significant effect on the angular power spectrum. except in the three highest-/ bins where the effect is ~10%.," A uniform displacement of the PMNS catalog coordinates by less than our estimated pointing error of $2\arcmin$ (see Paper I) does not have a significant effect on the angular power spectrum, except in the three $l$ bins where the effect is $\sim 10\%$."155 For the brightest point sources. positions accurate to «I are required.," For the brightest point sources, positions accurate to $<1\arcmin$ are required."156 We can extract positions to the necessary accuracy from the DASI data (see Paper D., We can extract positions to the necessary accuracy from the DASI data (see Paper I).157 In addition to the point sources constrained above. we make a statistical correction for residual point sources which are too faint to be detected by DASI or included in our PMN source table.," In addition to the point sources constrained above, we make a statistical correction for residual point sources which are too faint to be detected by DASI or included in our PMN source table."158" To do this. we estimate the point source number count per unit flux density at 4.85 GHz. dN/dSs. derived from the PMNS catalog. and the distribution of 31 GHz to 4.85 GHz flux density ratios. S3,/Ss. derived from new observations for this purpose with the OVRO 40 m telescope in Ka band (paper in preparation)."," To do this, we estimate the point source number count per unit flux density at 4.85 GHz, $dN/dS_5$, derived from the PMNS catalog, and the distribution of 31 GHz to 4.85 GHz flux density ratios, $S_{31}/S_5$, derived from new observations for this purpose with the OVRO 40 m telescope in Ka band (paper in preparation)."159" We proceed to calculate the statistical correction for unconstrained residual point sources $3,7| mJy using Monte Carlo techniques: we generate random point source distributions at 4.85 GHz using dN/dSs and statistically extrapolate the flux density of each source to each of our ten frequency channels using $3;/Ss.", We proceed to calculate the statistical correction for unconstrained residual point sources $S_{31} > 1$ mJy using Monte Carlo techniques; we generate random point source distributions at 4.85 GHz using $dN/dS_5$ and statistically extrapolate the flux density of each source to each of our ten frequency channels using $S_{31}/S_5$.160 These simulated point sources are superimposed on CMB temperature fluctuations and observed with DASI simulation software: a power spectrum is then generated with the analysis software., These simulated point sources are superimposed on CMB temperature fluctuations and observed with DASI simulation software; a power spectrum is then generated with the analysis software.161 The resulting mean amplitudes and uncertainties of the residual point source contribution to the nine band powers are [20+70. 70+80. 90+70. 180+70. 240+80. 330+100. 400+100. 500 170. 4304 170] piK.," The resulting mean amplitudes and uncertainties of the residual point source contribution to the nine band powers are $20 \pm 70$, $70 \pm 80$, $90 \pm 70$, $180 \pm 70$, $240 \pm 80$, $330 \pm 100$, $400 \pm 100$, $500 \pm 170$ , $430 \pm 170$ ] $\mu K^2$."162" The reported uncertainties are due to sky sample variance of the point source population in. the simulations. uncertainty in dN/dSs. and uncertainty in $3,/Ss."," The reported uncertainties are due to sky sample variance of the point source population in the simulations, uncertainty in $dN/dS_5$, and uncertainty in $S_{31}/S_5$."163 The residual point source contribution diminishes in the ninth band since that band power is dominated by visibilities from the highest frequency channels where the average point source flux density is lower relative to its mean flux density across all ten frequency channels., The residual point source contribution diminishes in the ninth band since that band power is dominated by visibilities from the highest frequency channels where the average point source flux density is lower relative to its mean flux density across all ten frequency channels.164 We use these statistically estimated amplitudes and uncertainties to adjust our CMB band-power estimates and uncertainties reported below., We use these statistically estimated amplitudes and uncertainties to adjust our CMB band-power estimates and uncertainties reported below.165 The CMB angular power spectrum from the first season of DASI data is shown in Figure 1.. with maximum likelihood estimates of nine band powers. piecewise flat π /(/+DC;/(2x). spanning the range /2100-900.," The CMB angular power spectrum from the first season of DASI data is shown in Figure \ref{fig:ps}, with maximum likelihood estimates of nine band powers, piecewise flat in $l(l+1)C_l/(2\pi)$ , spanning the range $l = $ 100–900."166 Adjacent bands are anticorrelated at the level., Adjacent bands are anticorrelated at the level.167 In addition. we show an alternate analysis of the same data. for nine bands shifted to the right with respect to the original band edges. in order to demonstrate the robustness of the analysis against possible effects due to the anticorrelation of adjacent bands.," In addition, we show an alternate analysis of the same data, for nine bands shifted to the right with respect to the original band edges, in order to demonstrate the robustness of the analysis against possible effects due to the anticorrelation of adjacent bands."168 Note that these two analyses use the same data to estimate band powers in two different piecewise-flat theoretical power spectra; only the first nine-band analysis (filled circles) is used for the cosmological parameter estimation described in Paper III., Note that these two analyses use the same data to estimate band powers in two different piecewise-flat theoretical power spectra; only the first nine-band analysis (filled circles) is used for the cosmological parameter estimation described in Paper III.169 While increasing the number of bands above nine may in principle provide more information about the underlying power spectrum. we have found that this does not significantly improve our ability to constrain cosmological parameters (see Paper III.," While increasing the number of bands above nine may in principle provide more information about the underlying power spectrum, we have found that this does not significantly improve our ability to constrain cosmological parameters (see Paper III)."170" In à separate analysis. we fit for the maximum likelihood value of an additional parameter. the temperature spectral index of the fluctuations. οὐ, where T1""."," In a separate analysis, we fit for the maximum likelihood value of an additional parameter, the temperature spectral index of the fluctuations, $\beta$, where $T \propto \nu^{\beta}$."171 Fitting a single spectral index for all nine bands. we find ./2—0.1+0.2 (10). while fitting a separate spectral index for /«500 and />500 yields }2—0.23:0.3 and 0.0+0.4 respectively. indicating the fluctuation power is consistent with CMB.," Fitting a single spectral index for all nine bands, we find $ \beta = -0.1172\pm 0.2$ $\sigma$ ), while fitting a separate spectral index for $l < 500$ and $l > 500$ yields $\beta = -0.2 \pm 0.3$ and $0.0 \pm 0.4$ respectively, indicating the fluctuation power is consistent with CMB."173 Values and marginal uncertainties for the angular power spectrum m the primary nine bands are given in Table [.., Values and marginal uncertainties for the angular power spectrum in the primary nine bands are given in Table \ref{tab:ps}.174 The center and &7'7 widths of the bands are calculated using band-power window functions adapted from (1999) which are plotted in Paper III., The center and $e^{-1/2}$ widths of the bands are calculated using band-power window functions adapted from \markcite{knox99}{ (1999) which are plotted in Paper III.175 These are the relevant window functions for calculating the expectation value of the band power given a theoretical power spectrum., These are the relevant window functions for calculating the expectation value of the band power given a theoretical power spectrum.176 We give the ratio of the uncertainty due to sky sample variance to the uncertainty due to noise. σςση. estimated using the offset log-normal formalism of Bond.Jaffe. (2000).," We give the ratio of the uncertainty due to sky sample variance to the uncertainty due to noise, $\sigma_{\mathrm{s}} / \sigma_{\mathrm{n}}$, estimated using the offset log-normal formalism of \markcite{bond98}{, (2000)."177 In their notation. a{oy is given by Cp/xp. where Cp is the band power estimate expressed as /(/41)C)/(27) and vg is proportional to the instrument noise contribution to the band-power uncertainty.," In their notation, $\sigma_{\mathrm{s}} / \sigma_{\mathrm{n}}$ is given by $\mathcal{C}_B/x_B$, where $\mathcal{C}_B$ is the band power estimate expressed as $l(l+1)C_l/(2\pi)$ and $x_B$ is proportional to the instrument noise contribution to the band-power uncertainty."178 These values may be used to estimate the non-Gaussianity in the band-power marginal likelihood distributions for parameter estimation caleulations — asymmetric uncertainties due to non-Gaussianity are negligible for most of our band powers and we do not plot them here., These values may be used to estimate the non-Gaussianity in the band-power marginal likelihood distributions for parameter estimation calculations — asymmetric uncertainties due to non-Gaussianity are negligible for most of our band powers and we do not plot them here.179 We also tabulate the band-power correlation matrix (Table 2))., We also tabulate the band-power correlation matrix (Table \ref{tab:corr}) ).180 All of the data products necessary for performing cosmological parameter estimation from this data are available at our., All of the data products necessary for performing cosmological parameter estimation from this data are available at our.181. We perform three types of tests to check the consistency of the data: i) 47 tests on the difference between two visibility data vectors constructed from observations of the same fields on the sky. 11) construction of à nine-band power spectrum of the epoch-differenced visibility data vector. to test for significant deviation from zero power. and iii) 47 tests on the difference between two power spectra constructed from independent fields on the sky.," We perform three types of tests to check the consistency of the data: i) $\chi^2$ tests on the difference between two visibility data vectors constructed from observations of the same fields on the sky, ii) construction of a nine-band power spectrum of the epoch-differenced visibility data vector, to test for significant deviation from zero power, and iii) $\chi^2$ tests on the difference between two power spectra constructed from independent fields on the sky."182 In the second and third types of test. we increase," In the second and third types of test, we increase"183In this section we establish the accuracy with which metallicities can be measured. using the Ca Ll Ix line EW.,In this section we establish the accuracy with which metallicities can be measured using the Ca II K line EW.184 We then show that the estimated metallicities are useful for identifving interlopers in the DIID samples., We then show that the estimated metallicities are useful for identifying interlopers in the BHB samples.185 IxXSIx plotted the Ca LE Ix line EW against (23)12o for stars for which independent accurate metallicitics are known. to determine empirically curves of constant metallicity in this parameter space.," KSK plotted the Ca II K line EW against $(B-V)_0$ for stars for which independent accurate metallicities are known, to determine empirically curves of constant metallicity in this parameter space."186 Ensteacd we have used the theoretical curves of Wilhelm et al. (, Instead we have used the theoretical curves of Wilhelm et al. (1871999a) measured from synthetic spectra.,1999a) measured from synthetic spectra.188 In Fig., In Fig.189 9. we plot their isoabundance contours for. Fe/1] = |. 2 and 8 (the solar-metallicity line is explained below).," \ref{cak_plot} we plot their isoabundance contours for [Fe/H] = –1, –2 and –3 (the solar-metallicity line is explained below)."190 Metallicities can be determined [rom this plot using measurements of ολο and (3Vo. by interpolation.," Metallicities can be determined from this plot using measurements of $_{Ca}$ and $(B-V)_0$, by interpolation."191 To assess the accuracy of this plot we have used it to measure metallicities for all the DIID stars in the INSI total sample for which independent accurate metallicities are known., To assess the accuracy of this plot we have used it to measure metallicities for all the BHB stars in the KSK total sample for which independent accurate metallicities are known.192 The stars used. listed in Table 4 and plotted in Fig. 9..," The stars used, listed in Table 4 and plotted in Fig. \ref{cak_plot},"193 are the BLIB stars in M92 (Fe/II] = 2.2). in M3 (Fe/1] = 1.5). ancl a sample of nearby field DIED stars (individually labeled with the value of Fe/LH]).," are the BHB stars in M92 ([Fe/H] = –2.2), in M3 ([Fe/H] = –1.5), and a sample of nearby field BHB stars (individually labeled with the value of [Fe/H])."194 The accurate metallicitics are taken from WKS ancl are Listed in Table 4 in the column headed ον., The accurate metallicities are taken from KSK and are listed in Table 4 in the column headed $_K$.195 Our interpolated measured: metallicities are listed. in the last column of Table 4. labeled Fe/lje.," Our interpolated measured metallicities are listed in the last column of Table 4, labeled $_C$."196 There is quite good agreement between our interpolatecl estimates and the accurate determinations., There is quite good agreement between our interpolated estimates and the accurate determinations.197 For the sample of 18 stars we measure a mean cillerence Fe/I1l]e ον=0.10 with stancarcl deviation zE0.24., For the sample of 18 stars we measure a mean difference $_C$ $_K=0.10$ with standard deviation $\pm0.24$.198 For the three sub-samiples M92. ALB.and the Ποια BIB stars the results are 0.21+0.20 (S stars). Q.10+4O17 (3 stars). 0.06+0.27. (7 stars) respectively.," For the three sub-samples M92, M3,and the field BHB stars the results are $0.21\pm0.20$ (8 stars), $-0.10\pm0.17$ (3 stars), $0.06\pm0.27$ (7 stars) respectively."199 “Phe scatter of the points could. arise. from a variety of sources including the clifliculty in measuring he continuum. contamination of the line by interstellar absorption. and uncertainties in the accurate estimates henmselves (see IXSIX for à discussion).," The scatter of the points could arise from a variety of sources including the difficulty in measuring the continuum, contamination of the line by interstellar absorption, and uncertainties in the accurate estimates themselves (see KSK for a discussion)."200 Pherefore in. using his plot to measure metallicities of other stars we acid an error of 0.3 dex in quadrature to the random: error. xwed on the measured. scatter.," Therefore in using this plot to measure metallicities of other stars we add an error of 0.3 dex in quadrature to the random error, based on the measured scatter."201 Additional scatter may be introduced. by the spectral peculiarities of the Am and Ap stars (discussed in INSIX. Wilhelm et al.," Additional scatter may be introduced by the spectral peculiarities of the Am and Ap stars (discussed in KSK, Wilhelm et al."202 1999a. 1999b).," 1999a, 1999b)."203 To measure the metallicities of other Atype stars over the full range of metallicities we have delined also a solarmetallicity contour by fitting a straight line to the data for mainsequence A stars in the Pleiaces and Coma clusters. plotted as open triangles.," To measure the metallicities of other A–type stars over the full range of metallicities we have defined also a solar–metallicity contour by fitting a straight line to the data for main–sequence A stars in the Pleiades and Coma clusters, plotted as open triangles."204 Because the lines are converging towards bluer colours we consider this plot reliable for estimating metallicities only for colours reclder than (D.Vyz0.05., Because the lines are converging towards bluer colours we consider this plot reliable for estimating metallicities only for colours redder than $(B-V)_0>0.05$.205 We return now to the issue of interlopers ic. the Low halo stars that are classified by our methods as BILD but classified w ISI as A/DS., We return now to the issue of interlopers i.e. the few halo stars that are classified by our methods as BHB but classified by KSK as A/BS.206 We have discovered that most of these interlopers have anomalously high metallicity., We have discovered that most of these interlopers have anomalously high metallicity.207 If we consider first theColour method. then for he IXSIx halo sample. shown in Fig.," If we consider first the method, then for the KSK halo sample, shown in Fig."208 5— (IIIS). there is one interloper amongst the 33 stars below the classification x»indary. the star numbered IU-70 by ISI.," \ref{colour_width} (RHS), there is one interloper amongst the 33 stars below the classification boundary, the star numbered RR7-70 by KSK."209 This star is the only star below the line with estimated metallicity Ημ0.5., This star is the only star below the line with estimated metallicity $>-0.5$.210 In the KSI total sample. Fig.," In the KSK total sample, Fig."211 5 (LIES). here are 4 interlopers amongst the 51 stars below the classification boundary. and similarly they are the only stars with metallicity Fe/1l]z—0.5.," \ref{colour_width} (LHS), there are 4 interlopers amongst the 51 stars below the classification boundary, and similarly they are the only stars with metallicity $>-0.5$."212 ‘Turning to the method. in the halo sample. Fig.," Turning to the method, in the halo sample, Fig."213 7 (IRIS). thereis one interloper amongst the 33 stars below the classification boundary. the star 117-70.yo- the same interloper in Fig.," \ref{scale_power} (RHS), thereis one interloper amongst the 33 stars below the classification boundary, the star RR7-70, the same interloper in Fig."214 5. (IIIS)., \ref{colour_width} (RHS).215 This star is the only star below the line with estimated metallicity ος—0.5., This star is the only star below the line with estimated metallicity $>-0.5$.216 In the total sample. there are 7 interlopers amongst the 58 stars below the classification boundary.," In the total sample, there are 7 interlopers amongst the 58 stars below the classification boundary."217 For three of these we are unable to estimate metallieities reliably as they have colours (BVyu« 0.05., For three of these we are unable to estimate metallicities reliably as they have colours $(B-V)_0<0.05$ .218 Three of the remaining four are 1ο only stars with metallicity ο] 0.5., Three of the remaining four are the only stars with metallicity $>-0.5$ .219In the literature. usually only the first condition ts meant when generalised three-body binary formation rate is described (e.g...22)..,"In the literature, usually only the first condition is meant when generalised three-body binary formation rate is described \citep[e.g., ][]{bt, bf_nata}."220 There are however limitations that make such treatments applicable to our case., There are however limitations that make such treatments inapplicable to our case.221 The formation rate in ? is derived for equal masses only and ts derived to estimate only the formation of binaries which have the hardness ratio 1) of their binding energy to the kinetic energy of an average object in the core to be |., The formation rate in \cite{bt} is derived for equal masses only and is derived to estimate only the formation of binaries which have the hardness ratio $\eta$ of their binding energy to the kinetic energy of an average object in the core to be 1.222 In ?.. the formation rate 1s derived to take into account unequal masses and different energies: however. the treatment is applicable only when the resulting binaries are near the hard-soft boundary or a bit harder.," In \cite{bf_nata}, the formation rate is derived to take into account unequal masses and different energies; however, the treatment is applicable only when the resulting binaries are near the hard-soft boundary or a bit harder."223 In their derivation of the formation rates of binaries with different energies. ?.. approximate that the orbital separation in the formed binary is the same as the size of the vicinity dy Where the three objects meet.," In their derivation of the formation rates of binaries with different energies, \cite{bf_nata}, approximate that the orbital separation in the formed binary is the same as the size of the vicinity $a_{\rm v}$ where the three objects meet."224 As shown by numerical experiments in ?.. this assumption is approximately satisfied at least for equal masses. where a(1+e)/ay~1.," As shown by numerical experiments in \cite{aar76}, this assumption is approximately satisfied at least for equal masses, where $a(1+e)/a_{\rm v}\approx 1$."225 In our study. we are most interested in the formation rates where objects have a fairly large mass ratio and have high energies ()Z 100).," In our study, we are most interested in the formation rates where objects have a fairly large mass ratio and have high energies $\eta\ga100$ )."226 Compared to +=| binaries. the decrease in the hard binary formation rate due to physical collistons and tidal effects could be significant.," Compared to $\eta=1$ binaries, the decrease in the hard binary formation rate due to physical collisions and tidal effects could be significant."227 We therefore limit ourselves to the consideration of only degenerate objects meeting each other., We therefore limit ourselves to the consideration of only degenerate objects meeting each other.228 On the other hand. the limitation of the considered cases to only very hard binaries eliminates the necessity to consider the second condition: showed that the probability to form à binary is strongly increasing?. as vicinity is decreasing. Px17>. and. even in the case of jj=2. of three-body encounters resulted in binary formation.," On the other hand, the limitation of the considered cases to only very hard binaries eliminates the necessity to consider the second condition: \cite{aar76} showed that the probability to form a binary is strongly increasing as vicinity is decreasing, $P\propto \eta^{-2}$, and, even in the case of $\eta = 2$, of three-body encounters resulted in binary formation."229 They have also shown that the average eccentricity of formed binaries ts a bit higher than in the thermal distribution: «e>=0.77 for i)2 cases. though this was the largest energy considered and the average values were slowly decreasing with hardness.," They have also shown that the average eccentricity of formed binaries is a bit higher than in the thermal distribution: $<e>=0.77$ for $\eta = 2$ cases, though this was the largest energy considered and the average values were slowly decreasing with hardness."230 We may expect that the average eccentricities could go as low as that of the thermal distribution., We may expect that the average eccentricities could go as low as that of the thermal distribution.231 The rate at which three objects meet in the same vicinity ay is the product of the rate Το(αν.nn.n2) for two objects to meet in this neighborhood and the probability P3 that during this event a third object will pass by.," The rate at which three objects meet in the same vicinity $a_{\rm v}$ is the product of the rate $\Gamma_2 (a_{\rm v}, m_1, m_2)$ for two objects to meet in this neighborhood and the probability $P_3$ that during this event a third object will pass by."232" As previously. the two-body encounter rate Τσαν.1.725) is standardly derived às a combination of geometric cross-section and. gravitational focusing: where ""-2GUnims)αν."," As previously, the two-body encounter rate $\Gamma_2 (a_{\rm v}, m_1, m_2)$ is standardly derived as a combination of geometric cross-section and gravitational focusing: where $v_{\rm p}^2=2G(m_1+m_2)/a_{\rm v}$."233" Two objects spend in. this vicinity about 7,z2a,Vp.", Two objects spend in this vicinity about $\tau_{\rm v}\approx 2 a_{\rm v} / v_{\rm p}$.234 The probability that a third object will be within the same vicinity. Ps. is then usually found only assuming that a third object will geometrically sweep a certain volume during τι (??)..," The probability that a third object will be within the same vicinity, $P_3$, is then usually found only assuming that a third object will geometrically sweep a certain volume during $\tau_{\rm v}$ \citep{bt, bf_nata}."235" A ""geometric"" cross-section however is good only for cases of e, large enough so Zv4.", A “geometric” cross-section however is good only for cases of $a_{\rm v}$ large enough so $v_{\rm p} \la v_{\infty}$.236 In the case when we are interested m the encounters voccurring within a small vicinity only. gravitational focusing must be taken into account.," In the case when we are interested in the encounters occurring within a small vicinity only, gravitational focusing must be taken into account."237 In some sense. this implies that we are looking for a probability that. at the same time. two objects will be gravitationally focused by a BH.," In some sense, this implies that we are looking for a probability that, at the same time, two objects will be gravitationally focused by a BH."238" The probability that a third object will pass within e, during 7, is then For very hard binaries. this rate is significantly larger than the corresponding rate derived in ?. and can be written as Therefore. even in a relatively dense stellar cluster. a BH of 15 M.. will have only about 2«10? three-body encounters within its 100 ΔΑ... vicinity during | Gyr."," The probability that a third object will pass within $a_{\rm v}$ during $\tau_{\rm v}$ is then The three-body formation rate then is For very hard binaries, this rate is significantly larger than the corresponding rate derived in \cite{bf_nata} and can be written as Therefore, even in a relatively dense stellar cluster, a BH of 15 $M_\odot$ will have only about $2\times10^{-9}$ three-body encounters within its 100 $R_\odot$ vicinity during 1 Gyr."239 The formation rate necessary to explain BH-WD X-ray binaries formation is consistent only with encounters within ~5000R... or jj=10.," The formation rate necessary to explain BH-WD X-ray binaries formation is consistent only with encounters within $\sim 5000 R_\odot$, or $\eta=10$."240" We also note that in this case v,2»v4. the assumption used for the derivation. is no longer valid."," We also note that in this case $v_{\rm p}\gg v_\infty $, the assumption used for the derivation, is no longer valid."241 Even m a denser cluster (i.7 10°pe) with a somewhat smaller velocity dispersion. the rate given by the above estimate is still about an order of magnitude less than the observed formation rate.," Even in a denser cluster $n_{\rm c}\sim 10^6 {\rm pc}$ ) with a somewhat smaller velocity dispersion, the rate given by the above estimate is still about an order of magnitude less than the observed formation rate."242 We conclude that unless a significant fraction of small-vicinity three-body encounters lead to a formation of binaries much smaller than the vicinity where they met (10-50 times smaller). three-body binary formation can not explain the observed rates of BH-WD X-ray binary formation.," We conclude that unless a significant fraction of small-vicinity three-body encounters lead to a formation of binaries much smaller than the vicinity where they met (10-50 times smaller), three-body binary formation can not explain the observed rates of BH-WD X-ray binary formation."243 We have discussed which BH-WD binaries can become X-ray binaries. how these binaries can be formed and at what rate.," We have discussed which BH-WD binaries can become X-ray binaries, how these binaries can be formed and at what rate."244 We have considered the following formation channels: All of the channels require that at least a fraction of BHs interacts strongly with other stars m the cluster., We have considered the following formation channels: All of the channels require that at least a fraction of BHs interacts strongly with other stars in the cluster.245 We find that the most important mechanism to make a BH-WD X-ray binary from an initially dynamically formed, We find that the most important mechanism to make a BH-WD X-ray binary from an initially dynamically formed246"such that for some cube Q,,.",such that for some cube $Q_{r_{1}}$.247 Iudeed. estimates on all such cubes Q77*.>sl are already controlled by (3.1ZEE1)).," Indeed, estimates on all such cubes $Q^{trans{'}}_{r\sqrt{2}}$ are already controlled by \ref{moa3})."2480. In this case we repeat the same arguments as in Step 2 of Lemma 3.5.., In this case we repeat the same arguments as in Step 2 of Lemma \ref{Hgibi2}.249) Indeed. in the present case. it is even simpler since the function ?t is symunetric with respect to {/=Tj.," Indeed, in the present case, it is even simpler since the function $\tilde{u}$ is symmetric with respect to $\{t=T\}$."250" W We now show Low to prove estimate (3.5)). The parabolic BAO, norm (3.9)) of Và could be estimated taking the supremum of f,Wa—(Varo, |. MC7. over small parabolic cubes (6) with r€ro/2). aud big parabolic cubes (Q, with r> r3/2)."," $\hfill{\blacksquare}$ We now show how to prove estimate \ref{5s5_2}) The parabolic $BMO_p$ norm\ref{eq_bmo_nor}) ) of $\Psi\tilde{u}$ could be estimated taking the supremum of $\inm_{Q_{r}}|\Psi\tilde{u}-(\Psi\ti{u})_{Q_{r}}|$ , $Q_{r}\subseteq \R^{2}$, over small parabolic cubes $Q_{r}$ with $r\leq r_{2}/2$ ), and big parabolic cubes $Q_{r}$ with $r>r_{2}/2$ )."251" The proofis then divided intotwo (Analysis ou big parabolic cubes Q,. r>r2/2). We compute. using the fact that V=0 on NRe\Zo. aud V<1 on Re (see (3.6))): (Analysis ou sinall parabolic cubes Qp.r€72/2).From the (3.2 [)) ro. aul the coustruction (3.6)) of the function V. we deduce that ifQ, intersects Zo then delinitionlorcecdly »C CO."," The proof is then divided intotwo (Analysis on big parabolic cubes $Q_{r}$ , $r>r_{2}/2$ We compute, using the fact that $\Psi=0$ on $\R^{2}\setminus \mathcal{Z}_{2}$, and $\Psi\leq2521$ on $\R^{2}$ (see \ref{cut_off}) )): (Analysis on small parabolic cubes $Q_{r}$, $r\leq r_{2}/2$From the definition \ref{mabaa0}) ) of $r_{2}$, and the construction \ref{cut_off}) ) of the function $\Psi$ , we deduce that if$Q_{r}$ intersects $\mathcal{Z}_{2}$ then forcedly $Q_{r}\subseteq253\widetilde{\Omega}_{T}$ ."254" Itnot. Le. GQ.C25=0 then V—0 on Q,. aud therefore: Then we have ouly to consider Q,COy. "," Ifnot, i.e. $Q_{r}\cap \mathcal{Z}_{2}=\emptyset$ then $\Psi=0$ on $Q_{r}$ , and therefore: Then we have only to consider $Q_{r}\subseteq255\widetilde{\Omega}_{T}$ "256members with masses in the 0.8—1.2Mo range. in this field.,"members with masses in the $0.8-1.2~M_{\sun}$ range, in this field."257 The median luminosity value for this sample is 1.3x10-SereIu the quantiles bars also refer to this sample.," The median luminosity value for this sample is $1.3 \times25810^{28}$; the quantiles bars also refer to this sample."259 Including the two candidate members with M.—0.8M5 (Ids 136 and 246 — see reftab:Ix)) does not change the median value: including the one ddetected (solar-mass) member outside the FFOV does not change this value either. although it increases the size of the quantile bar.," Including the two candidate members with $M \sim 0.8 M_{\sun}$ (Ids 136 and 246 – see \\ref{tab:lx}) ) does not change the median value; including the one detected (solar-mass) member outside the FOV does not change this value either, although it increases the size of the quantile bar."260 The error on the distance of the cluster (430+20 pe) is relatively small and does not impact significantly on the position of the median luminosity of NGC 752 in the plot: for the distance range of 410—450 pe. the value of the median X-ray luminosity stays within the two quantile vertical bars.," The error on the distance of the cluster $430\pm 20$ pc) is relatively small and does not impact significantly on the position of the median luminosity of NGC 752 in the plot; for the distance range of $410-450$ pc, the value of the median X-ray luminosity stays within the two quantile vertical bars."261 The X-ray luminosities of the seven points in NGC 752 have a spread of more than one order of magnitude. confirming the large spread in luminosity of stars with the same age and," The X-ray luminosities of the seven points in NGC 752 have a spread of more than one order of magnitude, confirming the large spread in luminosity of stars with the same age and"262,".,"263" The quantities p, and ce. can be expressedinterms of p, and e, by requiring p and e be continuous.", where $r_c$ is the “capture” or Bondi accretion redius r_c = The quantities $\rho_c$ and $v_c$ can be expressedinterms of $\rho_s$ and $v_s$ by requiring $\rho$ and $v$ be continuous.264 Consider a cross section per unit nass gpa; Which variesas some power of velocity where op. ey. and @ are determined by (he fundamental physies of the interaction.," Consider a cross section per unit mass $\sigma_{DM}$, which variesas some power of velocity ) = where $\sigma_0$, $v_0$, and $a$ are determined by the fundamental physics of the interaction."265 Then the optical depth is ," Then the optical depth is ) . ,"266where, where _s _s ; _c _c267The authors wish to thauk Tom herr and the UIKIRT Service Observing Proeranune for obtaining preliminary spectra near 1500 uuxng COSL Ryan Porter and Cary Ferlaud for supplying unpublished values of the recombination cocficicnts and for suggesting the possibility of electrou Hupact excitation into the metastable state. and the anonviuous referee for helpful sugeestions aud οΟΙΤΗ.,"The authors wish to thank Tom Kerr and the UKIRT Service Observing Programme for obtaining preliminary spectra near 10830 using CGS4, Ryan Porter and Gary Ferland for supplying unpublished values of the recombination coefficients and for suggesting the possibility of electron impact excitation into the metastable state, and the anonymous referee for helpful suggestions and comments."268 NI aud DJM have been supported by NSF eraut PIIY 05-55I86., NI and BJM have been supported by NSF grant PHY 05-55486.269 Based on observations obtained. at the Coin Observatory. which is operated by the Association of Universities for Research in Astronomy. Πιο. uuder à cooperative agrecinent with the NSF on behalf of the Gemini partnership: the National Scieuce. Foundatiou (United States). the Science aud Technology. Facilities Council (United. Kiuedomi. the National Research Council (Canada). CONICYT (Chile). the Australian Research Council (Australia). \linistérrio da Cienncia ο Tecnologia (Brazil) and Ministerio de Ciencia. Tecnologiaa e Dunovaciónn. Productiva (Arecutina).," Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Ministérrio da Ciênncia e Tecnologia (Brazil) and Ministerio de Ciencia, Tecnologíaa e Innovaciónn Productiva (Argentina)."270 The Comini/Phoeuix spectra were obtained through program CS-2008A-O-11., The Gemini/Phoenix spectra were obtained through program GS-2008A-Q-14.271 The observations were obtained with the Phoenix infrared spectrograph. which was developed bv the National Optical Astronomy Observatory.," The observations were obtained with the Phoenix infrared spectrograph, which was developed by the National Optical Astronomy Observatory."272consider images located near the critical curves. when the contribution to the total magnification due to any other images becomes important (see below),"consider images located near the critical curves, when the contribution to the total magnification due to any other images becomes important (see below)."273 The RBAL was developed for applications in the weak lensing limit. and should be used with caution for strong lensing cases.," The RBM was developed for applications in the weak lensing limit, and should be used with caution for strong lensing cases."274 We can investigate the validity of the RBAL by comparison with the various analytic solutions which exist for the Schwarzschilel lens (sce for a summary)., We can investigate the validity of the RBM by comparison with the various analytic solutions which exist for the Schwarzschild lens (see for a summary).275" Consider first a circular source of racius 2, with centre atq,=(μις.tee)."," Consider first a circular source of radius $R_{\rm s}$ with centre at $\bvec{y}_{\rm c} = (y_{1, \rm c}, y_{2, \rm c})$."276 Phe circumference of the source is then described by the set of vectors y=(yi.ye) with where 0<6<2s.," The circumference of the source is then described by the set of vectors $\bvec{y} = (y_1, y_2)$ with where $0 \leq \phi < 2 \pi$."277 For each g. we can solve for the two solutions. a. withCX3).," For each $\bvec{y}$, we can solve for the two solutions, $\bvec{x}_{\pm}$, with."278. In this case we are using the GLE to map rom the source plane to the image plane., In this case we are using the GLE to map from the source plane to the image plane.279 A source far [rom he lens axis produces one highly demagnified image Cpu) ocated near the lens axis (at api)., A source far from the lens axis produces one highly demagnified image $\mu_{\rm faint}$ ) located near the lens axis (at $\bvec{x}_{\rm faint}$ ).280 Phe second image will rave a magnification /haisniο 1. and an angular position near the source (al amisi ).," The second image will have a magnification $\mu_{\rm bright} \geq 1$ , and an angular position near the source (at $\bvec{x}_{\rm bright}$ )."281" As the source is moved towarcWu. he lens axis. the imagesare stretched in the tangential and become comparable in brightness (πμνfis| las gy, 0)."," As the source is moved towards the lens axis, the imagesare stretched in the tangential and become comparable in brightness $\vert \mu_{\rm faint}\vert 282\sim \vert \mu_{\rm bright} \vert \gg 1$ as $\bvec{y}_c \rightarrow 0$ )."283 When y=0. the two images merge into a highly magnilied ring (the Einstein ring) with otal magnification given byCXIO).," When $\bvec{y}_c = 0$, the two images merge into a highly magnified ring (the Einstein ring) with total magnification given by."284. Now consider a circular image of radius 2; centred on he location of the bright image. a=anisn(y.). with circumferential points a=(μήνο) This time. the GLE maps in the opposite direction from the image plane to the source. plane.," Now consider a circular image of radius $R_{\rm i}$ centred on the location of the bright image, $\bvec{x}_{\rm c} = \bvec{x}_{\rm bright} (\bvec{y}_{\rm c})$, with circumferential points $\bvec{x} = (x_1, x_2)$ This time, the GLE maps in the opposite direction -- from the image plane to the source plane."285 Phe source shape we obtain is stretched along the radial direction. and dillerentially compressed in the tangential direction.," The source shape we obtain is stretched along the radial direction, and differentially compressed in the tangential direction."286 demonstrates. the dillerences. between the shape and locations of a circular source (solid. line). for which the cireumferential points may be determined for both images. and a circular image and its corresponding source (short dashed line).," demonstrates the differences between the shape and locations of a circular source (solid line), for which the circumferential points may be determined for both images, and a circular image and its corresponding source (short dashed line)."287 In this example. we have considered a source which is near the Einstein radius where strong lensing elfects dominate. and there may be a significant contribution o the llux from the second image.," In this example, we have considered a source which is near the Einstein radius where strong lensing effects dominate, and there may be a significant contribution to the flux from the second image."288 In all cases where we will apply the RBAL. the image is chosen to be well away from he Einstein raclius (eritical curve). and so the Dux lost from he second image is not important. as we now show.," In all cases where we will apply the RBM, the image is chosen to be well away from the Einstein radius (critical curve), and so the flux lost from the second image is not important, as we now show."289 We now look at how accurately the RBAL approximates the total magnification. even though it includes the contribution of only one image.," We now look at how accurately the RBM approximates the total magnification, even though it includes the contribution of only one image."290 HE we set 2.=Ay. as we expect only small changes to the shape and hence radius of the image in the weak lensing limit (|a.| 1). then this is à two parameter problem. (GR. Ny).," If we set $R_{\rm s} = R_{\rm i}$, as we expect only small changes to the shape and hence radius of the image in the weak lensing limit $\vert \bvec{x}_{\rm c} \vert \gg 1$ ), then this is a two parameter problem $R_{\rm i}, N_{\rm ray}$ )."291 Defining the IUDM magnification in terms of the rav bundle image and source areas Glipist loan) as and the true (total) magnification as where «μμ slain are the areas of the two images. then the relative error in f/ppyt is Due to the circular svmmetry of the Sebwarzschild lens model we need only determine the radius. roa. within which he RBAL produces a relative error ausu>p per cent.," Defining the RBM magnification in terms of the ray bundle image and source areas $A_{\rm i, RBM}, A_{\rm s, RBM}$ ) as and the true (total) magnification as where $A_{\rm faint}$, $A_{\rm bright}$ are the areas of the two images, then the relative error in $\mu_{\rm RBM}$ is Due to the circular symmetry of the Schwarzschild lens model, we need only determine the radius, $x_{\rm cut}$, within which the RBM produces a relative error $\frac{\Delta \mu_{\rm RBM}}{\mu_{\rm RBM}} > p$ per cent."292 Dv using. Nia. favs in. the image. PARISancML source bundles. we are approximating the shape of a circular image/source w à polveon with Nya sides.," By using $N_{\rm ray}$ rays in the image and source bundles, we are approximating the shape of a circular image/source by a polygon with $N_{\rm ray}$ sides."293 Clearly. when Nya.X1. we will have a reasonable approximation to the truc shape of he image/source.," Clearly, when $N_{\rm ray} \gg 1$, we will have a reasonable approximation to the true shape of the image/source."294 However. to improve the speed of the ray xindle method. (at the cost of a small error). we ideally want to select a small value of [NS (2 90).," However, to improve the speed of the ray bundle method (at the cost of a small error), we ideally want to select a small value of $N_{\rm ray}$ $\lesim 20$ )."295" The areas are calculated. as a sum of triangular components within the imagefsouree polygon. where cach triangle has a common vertex at gj, orae (ie."," The areas are calculated as a sum of triangular components within the image/source polygon, where each triangle has a common vertex at $\bvec{y}_{\rm c}$ or$\bvec{x}_{\rm c}$ (ie."296 the null &eocdesic), the null geodesic).297 We need to first check that the caleulated μις is not significantly in error using a particular value of Nye., We need to first check that the calculated $\mu_{\rm true}$ is not significantly in error using a particular value of $N_{\rm ray}$ .298 This, This299is seen at the location of this radio source. although recdshifts are not available for either of them.,"is seen at the location of this radio source, although redshifts are not available for either of them."300 The 23rd magnitude galaxy which lies within half an arcsecond of the position of this weak racio source has no measured redshift., The 23rd magnitude galaxy which lies within half an arcsecond of the position of this weak radio source has no measured redshift.301 This unresolved radio source is unambiguously identified with an early type galaxy at he redshift of the cluster., This unresolved radio source is unambiguously identified with an early type galaxy at the redshift of the cluster.302 This radio source corresponds to a 22nd magnitude Sab galaxy of unknown redshift., This radio source corresponds to a 22nd magnitude Sab galaxy of unknown redshift.303 Phe brighter ealaxy a few aresee to the northeast of the radio source jàs à recdshift of z=0.17. and is unlikely to be related.," The brighter galaxy a few arcsec to the north–east of the radio source has a redshift of $z=0.17$, and is unlikely to be related."304 This radio source is associated with a diffuse optical galaxy at the cluster redshift., This radio source is associated with a diffuse optical galaxy at the cluster redshift.305 This weak extended radio emission appears ο be associated with a bright spiral galaxy at recishift 0.25., This weak extended radio emission appears to be associated with a bright spiral galaxy at redshift 0.25.306 The measured flux density corresponds to a radio luminosity of 1.2107 HIIz.+. similar to that of MS2.," The measured flux density corresponds to a radio luminosity of $1.2 \times30710^{22}$ $^{-1}$, similar to that of M82."308 No optical counterpart brighter than /=26 is seen associated with this unresolved radio source.," No optical counterpart brighter than $I309\approx 26$ is seen associated with this unresolved radio source."310 At the location of this radio source there is only a marginal detection of an extremely. faint galaxy (1~ 26)., At the location of this radio source there is only a marginal detection of an extremely faint galaxy $I \sim 26$ ).311 Vhis luminous extended racio source. which," This luminous extended radio source, which"312"Using the ITO4 sample we showed that the neelect of the stellar absorption on the helium emission lines causes a large svstematic effect on Y, and dY7dZ.",Using the IT04 sample we showed that the neglect of the stellar absorption on the helium emission lines causes a large systematic effect on $Y_p$ and $dY/dZ$.313 This is due to the fact that the stellar absorption equivalent width shows a trend (hat increases towards metal poor svstenis., This is due to the fact that the stellar absorption equivalent width shows a trend that increases towards metal poor systems.314 The inclusion of helium stellar absorption improves enormously the acceptance ol fit to the He I emission data. especially when the «quality. was bad without the inclusion of stellar absorption.," The inclusion of helium stellar absorption improves enormously the acceptance of fit to the He I emission data, especially when the quality was bad without the inclusion of stellar absorption."315 The resulüng magnitude of the absorption equivalent width of He I ÀJAT1 line (and also of the IH I Balmer lines) is on the order that is expected in a population svithesis calculation (Gonzállez Delgado οἱ al., The resulting magnitude of the absorption equivalent width of He I $\lambda4471$ line (and also of the H I Balmer lines) is on the order that is expected in a population synthesis calculation (Gonzállez Delgado et al.316 1999) for the nebula phase., 1999) for the nebula phase.317" With the inclusion of stellar absorption we obtained (the primordial helium abundance increased from y,=0.2342:0.004 to 0.2502£0.004.", With the inclusion of stellar absorption we obtained the primordial helium abundance increased from $y_p=0.234\pm 0.004$ to $0.250\pm 0.004$.318 We do not claim that the latter is the true value. but it is much preferred to the former on the ground of much smaller 47? [or the fit to individual HII regions. while the 6-Iine fit without stellar absorption is barely acceptable.," We do not claim that the latter is the true value, but it is much preferred to the former on the ground of much smaller $\chi^2$ for the fit to individual HII regions, while the 6-line fit without stellar absorption is barely acceptable."319" Or. most conservatively, one can claim (hat we cannot obtain the primordial helium abundance to the error of OY),20.004 or less unless underlying stellar absorption is properly understood."," Or, most conservatively, one can claim that we cannot obtain the primordial helium abundance to the error of $\delta Y_p\approx 0.004$ or less unless underlying stellar absorption is properly understood."320 In our analvsis we noted that the minimisation of the sum of the 4? over individual lines. ignoring (he correlation among lines. is likely to underestimate the error.," In our analysis we noted that the minimisation of the sum of the $\chi^2$ over individual lines, ignoring the correlation among lines, is likely to underestimate the error."321" If we accept our higher helium abundance. we find np/n,=7.91x10.1!"" with the aid ol the standard Dig Bang nucleosvnthesis caleulation (e.g.. Olive et al."," If we accept our higher helium abundance, we find $n_B/n_\gamma=3227.9^{+4.0}_{-2.4}\times 10^{-10}$ with the aid of the standard Big Bang nucleosynthesis calculation (e.g., Olive et al."323 2001)., 2001).324 This barvon abundance is consistent wilh that inferred from cosmic microwave background. anisotropies (Spereel et al., This baryon abundance is consistent with that inferred from cosmic microwave background anisotropies (Spergel et al.325 2003)., 2003).326 We also found that dY/dZ is largely allectecd upon the inclusion of stellar absorption: the original value of αγ/dZ£z4—5 decreases to 1+1., We also found that $dY/dZ$ is largely affected upon the inclusion of stellar absorption: the original value of $dY/dZ\approx 4-5$ decreases to $1\pm 1$.327 This smaller number is consistent with the derivative inferred form the standard solar model (Baheall et al., This smaller number is consistent with the derivative inferred form the standard solar model (Bahcall et al.328 2001) AY/AZ= 1-4. and also with 2.10.4 from the dwarf star atmosphere (Jimenez et al.," 2001) $\Delta Y/\Delta Z=(Y_{\rm initial}-Y_p)/Z_{\rm initial}=1.4$ , and also with $2.1\pm 0.4$ from the dwarf star atmosphere (Jimenez et al."329 2003) using the method introduced by Pagel and Portinari (1993) who presented 342 for this derivative., 2003) using the method introduced by Pagel and Portinari (1998) who presented $3\pm2$ for this derivative.330 This work is supported in part by Grants in Aid of the Ministry of Exlucation of Japan al Kashiwa., This work is supported in part by Grants in Aid of the Ministry of Education of Japan at Kashiwa.331 ME received support from the Monell Foundation at Princeton., MF received support from the Monell Foundation at Princeton.332SALARTS observations were obtained with a cadence designed to ensure that the final combined frames in each filter ave referenced to the same mid-exposure time.,SMARTS observations were obtained with a cadence designed to ensure that the final combined frames in each filter are referenced to the same mid-exposure time.333 The images obtained during the first SMAIIS epoch have a common mid-exposure time of 2.2 hr post-burst., The images obtained during the first SMARTS epoch have a common mid-exposure time of 2.2 hr post-burst.334 The OAG was brightest during this epoch. so we use these data to build the spectral energy. distribution (SED) of the afterglow in order (ο evaluate whether (here is significant extinction along the line-ol-sight through the GRB host galaxy.," The OAG was brightest during this epoch, so we use these data to build the spectral energy distribution (SED) of the afterglow in order to evaluate whether there is significant extinction along the line-of-sight through the GRB host galaxy."335 To help constrain the host reddening. we extend the SED blueward of the S\IARTS DB--band filter using Swift--UVOT ultraviolet observations.," To help constrain the host reddening, we extend the SED blueward of the SMARTS -band filter using -UVOT ultraviolet observations."336" The UVOT light eiurve was best sampled in the U filter. and a power-law Π to the data vield a decay slope of (lor f,x/"". where f, is the (ransient’s [lux density and / is the time since the burst trigger). consistent with the decay rates inferred. [rom SAIARTS (Cobb2009) and SkvNet/PROMPT (Laislipetal.2009) observations at similar (imes."," The UVOT light curve was best sampled in the $U$ filter, and a power-law fit to the data yield a decay slope of $\alpha=0.55\pm0.05$ (for $f_\nu\propto t^{-\alpha}$, where $f_\nu$ is the transient's flux density and $t$ is the time since the burst trigger), consistent with the decay rates inferred from SMARTS \citep{gcn10244} and SkyNet/PROMPT \citep{gcn10219} observations at similar times."337 Assuming {his decay rate. UVOT magnitudes were extrapolated to thecommon time of 2.2 hr post-burst (U= mag. UVIV1=1022010 mag. CVA2=16.080.10 mag) in order to match the SMARTS epoch.," Assuming this decay rate, UVOT magnitudes were extrapolated to thecommon time of 2.2 hr post-burst $U=16.45\pm0.10$ mag, $UVW1=16.22\pm0.10$ mag, $UVM2=16.08\pm0.10$ mag) in order to match the SMARTS epoch."338 Figure 3 shows the SMARTS/UVOT SED of the OAG of GRD 091127., Figure 3 shows the SMARTS/UVOT SED of the OAG of GRB 091127.339 After correction or Galactic extinction. the observed UV-optical-II SED was fit assuming an intrinsic power aw allected by an extinction screen at the host redshilt of 0.49.," After correction for Galactic extinction, the observed UV-optical-IR SED was fit assuming an intrinsic power law affected by an extinction screen at the host redshift of 0.49."340 We examined several models. including Milkv. Was-like extinction. SAIC-like extinction. and LMC-like extinction using the parameterization of Fitzpatrick(1999).," We examined several models, including Milky Way-like extinction, SMC-like extinction, and LMC-like extinction using the parameterization of \citet{Fitzpatrick99}."341. The data were also fit to an unextinguished power aw., The data were also fit to an unextinguished power law.342 Fits wil a small amount of host-galaxy. extinction were statistically acceptable. as was the fit to the unextinguished case.," Fits with a small amount of host-galaxy extinction were statistically acceptable, as was the fit to the unextinguished case."343 Given (hat the addition of host-galaxy. extinction does 100 significantlv inprove the moclel fit. we suggest there is little to no significant extinction along the line-of-sight to GRD 091127.," Given that the addition of host-galaxy extinction does not significantly improve the model fit, we suggest there is little to no significant extinction along the line-of-sight to GRB 091127."344 The 30 upper limit on host-galaxv extinction is Ay\<O.5 mag., The $3\sigma$ upper limit on host-galaxy extinction is $A_V<0.5$ mag.345oS If a small amount of extinction is present in the host ogalaxy. we will slightly5 underestimate the peak brightness of the SN associated with GRB 091127.," If a small amount of extinction is present in the host galaxy, we will slightly underestimate the peak brightness of the SN associated with GRB 091127."346 The decav of the OAG of GRB 091127 is modeled by a broken power law., The decay of the OAG of GRB 091127 is modeled by a broken power law.347 During the first three epochs. the power-law decay index ealeulated from the SMARTS /-band observations is à=0.54zE0.02.," During the first three epochs, the power-law decay index calculated from the SMARTS $I$ -band observations is $\alpha = 0.54\pm0.02$."348 At ~0.3 das post-burst. the power law steepens to a=1.29+0.03.," At $\sim 0.3$ days post-burst, the power law steepens to $\alpha =3491.29\pm0.03$."350 This decay index is similar to that reported in other optical filters (llaislipοἱal. 2009).., This decay index is similar to that reported in other optical filters \citep{gcn10219}. .351" Observations taken up to 6 davs post-burst are dominated by the ΟΛ,", Observations taken up to 6 days post-burst are dominated by the OAG.352 The transient’s, The transient's353a conversion of 4.5 per cent of stellar mass into IGM metals with a single generation of stars.,a conversion of 4.5 per cent of stellar mass into IGM metals with a single generation of stars.354" ln Fig. ον,"," In Fig. \ref{fig:sfrd_lar},"355" we show the level of global star-formation that between >=15 and z=2, would enrich the IGM to one-tenth solar metallicity.", we show the level of global star-formation that between $z=15$ and $z=z_r$ would enrich the IGM to one-tenth solar metallicity.356 This is already considerably higher than measurements such as those of Aguirre.Schave&‘Phe-uns(2002). and Bouehéctal.(2007).. which suggest a median enrichment of x;107Z. at redshift 2. so our bound is quite conservative.," This is already considerably higher than measurements such as those of \citet{aguirre02} and \citet{bouche07}, which suggest a median enrichment of $\lesssim 10^{-2}\: Z_\odot$ at redshift 2, so our bound is quite conservative."357 There are several ways in which this metal production could be suppressed., There are several ways in which this metal production could be suppressed.358 Hf the IME were modified. [rom the Larson function assumed here. so as to reduce the number of stars forming in the mass range. metal production. would be greatly diminished.," If the IMF were modified from the Larson function assumed here, so as to reduce the number of stars forming in the mass range, metal production would be greatly diminished."359 Stars at masses higher than ~260M. are generally believed to follow a dillerent evolutionary. path than lighter stars and end with most of their metals inside remnant black holes (Lleger&Woosley2002)., Stars at masses higher than $\sim 260$ are generally believed to follow a different evolutionary path than lighter stars and end with most of their metals inside remnant black holes \citep{heger&woosley02}.360. Shifting high-mass star production to either above or immediately below the mass range (ic. 50 t0 140M... or 4; 260 M.) would reduce IGM enrichment. while having little impact on the UV spectra produced per stellar mass.," Shifting high-mass star production to either above or immediately below the mass range (i.e. 50 to 140, or $\gtrsim$ 260 ) would reduce IGM enrichment, while having little impact on the UV spectra produced per stellar mass."361 Lt is also possible that our assumption of all metals escaping to the IGM after an event is incorrect., It is also possible that our assumption of all metals escaping to the IGM after an event is incorrect.362 With a Salpeter EME. the mass fraction of barvons allotted to the mass range is much smaller. ~0.6 per cent. and the bound would thus be a factor of 15 higher.," With a Salpeter IMF, the mass fraction of baryons allotted to the mass range is much smaller, $\sim$ 0.6 per cent, and the bound would thus be a factor of 15 higher."363 The number of barvons available at a given reclshilt to form stars puts another constraint on any carly star-formation model., The number of baryons available at a given redshift to form stars puts another constraint on any early star-formation model.364 The utilization of gas in. collapsed structures can be quantified by the. star-formation cllicieney parameter. f;., The utilization of gas in collapsed structures can be quantified by the star-formation efficiency parameter $f_*$.365 We take the total mass in all structures above the molecular cooling mass scale as a function of redshift (rom fig., We take the total mass in all structures above the molecular cooling mass scale as a function of redshift from fig.366 1 of MS05. and use eq. (," 1 of MS05, and use eq. ("3676) of thatpaper to compute the total mass conversion into pop-LLE stars for a given value fi: 0.1 is used. here.,6) of thatpaper to compute the total mass conversion into pop-III stars for a given value $f_*$; 0.1 is used here.368 The correction factor g in this formula is set to 0.65., The correction factor $g$ in this formula is set to 0.65.369 While slightlv higher masses of f. might be allowable. NS05 argue that f.20.3 is probably implausibly. high.," While slightly higher masses of $f_*$ might be allowable, MS05 argue that $f_* \gtrsim 0.3$ is probably implausibly high."370 Our results indicate that the GeV sources seen byΓΗ ab 21.5 dislavor a scenario in which pop-LLl stars with a strongly top-heavy LAI are. formed. in. copious numbers (ic. SERD £0.2 to O04 ve+ )in the late stages of reionization. 6<z«SN.," Our results indicate that the GeV sources seen by at $z > 1.5$ disfavor a scenario in which pop-III stars with a strongly top-heavy IMF are formed in copious numbers (i.e. SFRD $\gtrsim 0.2$ to 0.4 $^{-1}$ $^{-3}$ ) in the late stages of reionization, $6 < z < 8$."371 At higher redshift. very high. star-formation rate densities {κ1) are disfavored by our result. ancl global limits on star-formation ellicienev in. proto-galaxies and limits on IGM metallicity impose a constraint at approximately the same level.," At higher redshift, very high star-formation rate densities $\gtrsim 1$ ) are disfavored by our result, and global limits on star-formation efficiency in proto-galaxies and limits on IGM metallicity impose a constraint at approximately the same level."372 Switching to a more moderate LME. Le. a Salpeter IME with a cutolf at 5M... raises our limits by about a factor of 1.7.," Switching to a more moderate IMF, i.e. a Salpeter IMF with a cutoff at 5, raises our limits by about a factor of 1.7."373 The results in Figs., The results in Figs.374 3. and 40 can be compared to those calculated in RISALOO for zero-metallicity stars that are based on low-recshilt blazar limits on the local background. (their fis., \ref{fig:sfrd_lar} and \ref{fig:sfrd_sal} can be compared to those calculated in RKM09 for zero-metallicity stars that are based on low-redshift blazar limits on the local background (their fig.375 9: the model with a=10 and j=0 most closely. resembles our simple step function applied for our pop-HI SETUD function)., 9; the model with $\alpha =10$ and $\beta = 0$ most closely resembles our simple step function applied for our pop-III SFRD function).376 We place slightly stronger constraints on the SERD at redshifts 7 and & than this work., We place slightly stronger constraints on the SFRD at redshifts 7 and 8 than this work.377 However. the significance of our strongest claims are only marginal. due to the very limited number of high-redshift photons that are relevant to our calculation.," However, the significance of our strongest claims are only marginal, due to the very limited number of high-redshift photons that are relevant to our calculation."378 The RINALOO result. is based on an upper limit to the local background light of 5 nW ? sr+ at 2qun., The RKM09 result is based on an upper limit to the local background light of 5 nW $^{-2}$ $^{-1}$ at 2.379 As shown in Fig. 1..," As shown in Fig. \ref{fig:ebllims},"380 our results most strongly constrain the EBL below ~1ji. and do not bound the EBL at co-moving wavelengths longer than 2 (mi.," our results most strongly constrain the EBL below $\sim 1$, and do not bound the EBL at co-moving wavelengths longer than 2 ."381.. For similar reasons. high redshift gamma-ray observations have little hope of providing meaningful constraints on a hypothetical contribution to the Ht background. [rom dark-matter burning stars," For similar reasons, high redshift gamma-ray observations have little hope of providing meaningful constraints on a hypothetical contribution to the IR background from dark-matter burning stars"382Thus. πο large excess of cawart galaxies exists around NGC 1399.,"Thus, no large excess of dwarf galaxies exists around NGC 1399."383 Iu Fig., In Fig.384 1l we show the distribution of point sources around NGC 1399., 11 we show the distribution of point sources around NGC 1399.385 We selected all objects with classificr values larger than 0.86. magnitudes between 20.5 22.0. aud colors in the range 0.6<<(VI)1.6. in total 218 objects.," We selected all objects with classifier values larger than 0.86, magnitudes between $20.5 < V < 22.0$ , and colors in the range $0.6 < (V-I) < 3861.6$, in total 248 objects."387 The isodeusitv contours have levels between L1 aud 6.8 objects per arci., The isodensity contours have levels between 1.1 and 6.8 objects per $^2$.388 The coutribution of background objects was determined by counting poit ποιαος. ni he local backeround field D2., The contribution of background objects was determined by counting point sources in the local background field B2.389 With the same selection criteria 0.13 objects poer arch9 lave Όσοι counted., With the same selection criteria 0.13 objects per $^2$ have been counted.390 Thus. about of the poiut sources should belong to the GCS of NGC |399 or the background galaxy cluster.," Thus, about of the point sources should belong to the GCS of NGC 1399 or the background galaxy cluster."391 Again. the couuts in all fields except F3 are coniplete (see also EKohle e al. 1996)).," Again, the counts in all fields except F3 are complete (see also Kohle et al. \cite{kohl}) )."392 The peak of the distribution is located about 173 cast of NGC 1399., The peak of the distribution is located about $1\farcm3$ east of NGC 1399.393 This result has to be considered with caution. since he counts in the ceutral density pixels as well as in the SW field are not complete.," This result has to be considered with caution, since the counts in the central density pixels as well as in the SW field are not complete."394 However. we foie that this effect also (even. nore pronounced) occurs when taking brighter siuuples of point sources. which should rot be affected by incompleteness effects;," However, we found that this effect also (even more pronounced) occurs when taking brighter samples of point sources, which should not be affected by incompleteness effects."395 Iu previous investigations of the GCS of NGC 1399. the properties of GCs in the NE (F2) and NW (Fl) feld have been examined (Iissler-Patig et al. 1997)).," In previous investigations of the GCS of NGC 1399, the properties of GCs in the NE (F2) and NW (F4) field have been examined (Kissler-Patig et al. \cite{kiss97}) )."396 We fouud that the radial surface density profile of the GCs (centered on NGC 1399) is shallower iu the NE field than iu the NW field., We found that the radial surface density profile of the GCs (centered on NGC 1399) is shallower in the NE field than in the NW field.397 Furthermore Forbes et al. (1998)), Furthermore Forbes et al. \cite{forb}) )398" investigated the augular distribution of GCs within 100""(25 ""deusitv pixels” ina TST WEPC? image that covers the NE part of the center of NGC 1399).", investigated the angular distribution of GCs within $= 5$ “density pixels” in a HST WFPC2 image that covers the NE part of the center of NGC 1399).399 They found a peak in the augular distribution iu the east direction., They found a peak in the angular distribution in the east direction.400 Both results are consistent with our fiuding of an excess of poiut sources east of NGC 1399., Both results are consistent with our finding of an excess of point sources east of NGC 1399.401 Note that iu this direction also the deusitv of resolved objects is very hieh., Note that in this direction also the density of resolved objects is very high.402 There are basically three possible explanations for this displaceiieut., There are basically three possible explanations for this displacement.403 First. a significant amount of uuresolved ealaxies in the background cluster was counted together with the globular clusters.," First, a significant amount of unresolved galaxies in the background cluster was counted together with the globular clusters."404 Since the background cluster ies cast of NCC 1399 the peak of the distribution of all poiut sources would then be shifted to the cast., Since the background cluster lies east of NGC 1399 the peak of the distribution of all point sources would then be shifted to the east.405 Iu his case. the galaxies would have absolute magunitucle:μα vetween δικMae16.5 mae and half light radii sanaller than about 1.5 kpc assunüug a distance of 1950 Mpc to the background cluster.," In this case, the galaxies would have absolute magnitudes between $-18.0 < M_V < -16.5$ mag and half light radii smaller than about 1.5 kpc assuming a distance of 480 Mpc to the background cluster."406 Such properties cau only ο explained by cEs. which are believed to represeut oulv a neeligible fraction of the galaxy population 1- he local universe.," Such properties can only be explained by cEs, which are believed to represent only a negligible fraction of the galaxy population in the local universe."407 Iu the Fornax cluster. for example. Drinkwater et al. (1997..," In the Fornax cluster, for example, Drinkwater et al. \cite{drin97},"408 soe also Drinkwater Creee 1998)) investigated by radial velocity measurements that all galaxies in their sample that are classified as M32 type compact ellipticals in the FCC are background ealaxies., see also Drinkwater Gregg \cite{drin98}) ) investigated by radial velocity measurements that all galaxies in their sample that are classified as M32 type compact ellipticals in the FCC are background galaxies.409 Further. in the UST counts CC's should be clearly distinguishable from galaxies even at 2=0.11.," Further, in the HST counts GCs should be clearly distinguishable from galaxies even at $z = 0.11$."410 Secoud. the GCS is really displaced with respect to the bulge of NGC 1399.," Second, the GCS is really displaced with respect to the bulge of NGC 1399."411 This would be a lint that the CC'S follows another potential than the stellar Πο and may belong rather to the cluster as a whole than to NGC 1399 itsclt. as supported by the velocity dispersion iieasurenments iu Cailliuzür et al. (1991))," This would be a hint that the GCS follows another potential than the stellar light and may belong rather to the cluster as a whole than to NGC 1399 itself, as supported by the velocity dispersion measurements in Grillmair et al. \cite{gril}) )"412 ancl Wissler-Patig et al. (1998a))., and Kissler-Patig et al. \cite{kiss98a}) ).413 Iu this respect. it is worthwhile noting that the ceuter of the eas distribution detected by N-ray observations Is also displaced. to the north-cast of NGC 1399 (Isehe et al. 1996...," In this respect, it is worthwhile noting that the center of the gas distribution detected by X-ray observations is also displaced, to the north-east of NGC 1399 (Ikebe et al. \cite{ikeb},"414 Jones et al. 1997))., Jones et al. \cite{jone}) ).415" Finally. solmewhat simular to the latter poiut. the istribution of GCs melt be a temporary displaceimoeut of a ""normal GCS centered on NGC. 1399."," Finally, somewhat similar to the latter point, the distribution of GCs might be a temporary displacement of a “normal” GCS centered on NGC 1399."416 A scenario that supports this possibility is related to the enriclineut of the ceutral GCS by the accretion of GCs from other galaxies., A scenario that supports this possibility is related to the enrichment of the central GCS by the accretion of GCs from other galaxies.417 Iissler-Patig et i (1998b)), Kissler-Patig et al. \cite{kiss98b}) )418 sugsest that tidal tails of GCs from the last passage of a Fornax galaxy might still be visible: these could mimic a skewed distribution of CC's around NGC 1399., suggest that tidal tails of GCs from the last passage of a Fornax galaxy might still be visible; these could mimic a skewed distribution of GCs around NGC 1399.419 Further iuvestisatiouns have to show if this riddle eiui be solved., Further investigations have to show if this riddle can be solved.420 Iu selected. fields in the Fornax cluster. located near ellipticals. nore than 870 ealaxies were identified down to a V maguitude of 22 mae.," In selected fields in the Fornax cluster, located near ellipticals, more than 870 galaxies were identified down to a $V$ magnitude of 22 mag."421 Photometric properties. such as total V. maguitude. (VI) color. peak aud effective surface briehtuess. and effective radius of these galaxies are compiled in a catalog.," Photometric properties, such as total $V$ magnitude, $(V - I)$ color, peak and effective surface brightness, and effective radius of these galaxies are compiled in a catalog."422 For a bright subsample of the galaxies we determined he seemg corrected true central surface briehtnesses., For a bright subsample of the galaxies we determined the seeing corrected true central surface brightnesses.423 Expoucutial and/or de Vaucouleur profiles were fitted to hese data aud their parameters are elven in a further catalog (Appendix B)., Exponential and/or de Vaucouleur profiles were fitted to these data and their parameters are given in a further catalog (Appendix B).424 Ouly few galaxies could clearly ο indeuti&Bed as dwiurf galaxies by their location iu he surface brightness magnitude diagram. where thev ollow the expected sequence for dwarf ellipticals.," Only few galaxies could clearly be indentified as dwarf galaxies by their location in the surface brightness magnitude diagram, where they follow the expected sequence for dwarf ellipticals."425 Most of them are already listed in the Fornax Cluster Catalog (Fereuson 1989))., Most of them are already listed in the Fornax Cluster Catalog (Ferguson \cite{ferg}) ).426 ILoxcever. our survey limit of about 21 nag 7 in peak surface brightness is too bright o detect cawarts like the faintest Local Croup dwarf spheroidals. if they were located at the Fornax distance.," However, our survey limit of about 24 mag $^{-2}$ in peak surface brightness is too bright to detect dwarfs like the faintest Local Group dwarf spheroidals, if they were located at the Fornax distance."427 Ou the other haud. among the hieh surface briehtuess objects compact dwarfs might be hidden. as shown by two uucleus-lke objects that are Fornax mcmbers as derived from their radial velocities (sce also Paper II).," On the other hand, among the high surface brightness objects compact dwarfs might be hidden, as shown by two nucleus-like objects that are Fornax members as derived from their radial velocities (see also Paper II)."428" The ""u parameter of the Sérrsic profile fits is clearly correlated to the absolute liuinesityv of the dwarf ellipticals in Foruax.", The “n” parameter of the Sérrsic profile fits is clearly correlated to the absolute luminosity of the dwarf ellipticals in Fornax.429" The nbhuuinositv relation secs to continue in the region of the ""normal ellipticals.", The n–luminosity relation seems to continue in the region of the “normal” ellipticals.430 Tn the color maguitude diagriuu the dwarfs tend to follow a color maenitude relation iu the sense that fainter galaxies are bluer., In the color magnitude diagram the dwarfs tend to follow a color – magnitude relation in the sense that fainter galaxies are bluer.431 The common explanation for this relation is a decreasing metallicity with decreasing huninosity., The common explanation for this relation is a decreasing metallicity with decreasing luminosity.432 However. for some dwarts the blue colors might also be a lint for au existing voune or interiiedciate-age stellar population. as it is seen iu several Local Croup dwarf spheroidals.," However, for some dwarfs the blue colors might also be a hint for an existing young or intermediate-age stellar population, as it is seen in several Local Group dwarf spheroidals."433 Evident sigus for recent and ougoine, Evident signs for recent and ongoing434in [imite-difference codes (hat results in an artificial (ransler of power [rom trailing slwaves into leading shwaves as the shwave is lost at the grid scale.,in finite-difference codes that results in an artificial transfer of power from trailing shwaves into leading shwaves as the shwave is lost at the grid scale.435" One expects aliasing {ο occur approximatelv at intervals οἱ = where n, is Che azimuthal shwave number.", One expects aliasing to occur approximately at intervals of = where $n_y$ is the azimuthal shwave number.436" This interval corresponds to δν)=2z/fel. where dr=L,/N, is the radial grid scale."," This interval corresponds to $\Delta437\tilde{k}_x(t) = 2\pi/dx$ , where $dx = L_x/N_x$ is the radial grid scale."438 Dased upon expression (??)). aliasing ellects should be more pronounced al lower numerical resolution because (he code has less time to evolve a shwave before the wavelength of the shwave becomes smaller than the grid scale.," Based upon expression \ref{ALIAS}) ), aliasing effects should be more pronounced at lower numerical resolution because the code has less time to evolve a shwave before the wavelength of the shwave becomes smaller than the grid scale."439 It can be seen trom the far-right data point in Figure 11. (Iun 7 in Table 1)) that the measured erowth ratedecreases wilh increasing resolution., It can be seen from the far-right data point in Figure \ref{f11} (Run 7 in Table \ref{pap4t1}) ) that the measured growth rate with increasing resolution.440 The evolution of v; for this run is shown in Figure 12.., The evolution of $v_t$ for this run is shown in Figure \ref{f12}.441 The effects of aliasing can be seen explicitly bv evolving a single shwave. as was clone for our linear theory test (Figure 1)).," The effects of aliasing can be seen explicitly by evolving a single shwave, as was done for our linear theory test (Figure \ref{f1}) )."442" Figure 13. shows the evolution of the density perturbation [or a single shwave using the same parameters that were used for Run 7: L,=Ly,12. E "," Figure \ref{f13} shows the evolution of the density perturbation for a single shwave using the same parameters that were used for Run 7: $L_x = L_y = 12$, $N_{x,min}^2 = -0.01$ and $q = 0.2$."443"""T −−∙ ↜↴⊺∐↕⊳∖⇁≺∢∪↕⋅↕⋅≼↲⊳∖⊽↕↽≻∪∐≼⇂⋝∖⊽↥∪∣∣∙∩∶∔⋅≀↧↴∐≺⇂⊔∐↲≼↲⇀↸↕↽≻"," The initial shwave vector used was $(k_{x0}, k_y) = (-8\pi/L_x, 8\pi/L_y)$."444≼↲≺∢∥↲≼⇂≀↧↴∐≀↧↪∖⊽↕∐≸↽↔↴↕∐∥↲↕⋅∖↽≀↧↴↥≼⋝∊∙↗∊∙↗↕⋝⇄⊔⊳∖⇁⊔∐↲↕⋅≼↲↓⋟∪↕⋅≼↲∆∣⊑∶↼∖⊽⊽∣∣∕∕∕∔⋅ ∐∏∐⋟∖⊽≀↕↴↥⊔∐⋅≼↲≼↲," This corresponds to $n_y = 4$, and the expected aliasing interval \ref{ALIAS}) ) is therefore $\Delta \tilde{\tau} = N_x/4$."445∐∏∐∐↲↕⋅↕≺∢≀↕↴↥↕⋅≼↲⋟∖⊽∪↥∏∐∪∐⋟∖⇁≀↧↴↕⋅≼↲↕↽≻↥∪⊔≼↲≺⇂↕∐⊡≸≟⋯⋅≼↲↓⊑↽⊰⋅⋅≀↧↴∐≼⇂⊔∐↲≀↧↴∐≀↧⊔∖⊽↕∐≸≟↕∐∩↲↕⋅∖↽≀↧↴↥≀↧↴↥ each resolution is consistent with expression (??7)).," Runs at three numerical resolutions are plotted in Figure \ref{f13}, , and the aliasing interval at each resolution is consistent with expression \ref{ALIAS}) )."446 It is clear from Figure 123. (hat. a lower resolution results in a larger overall growth at the end of the run., It is clear from Figure \ref{f13} that a lower resolution results in a larger overall growth at the end of the run.447 It also appears that the erowlh seen in Figure 13. requires a negative entropy gradient., It also appears that the growth seen in Figure \ref{f13} requires a negative entropy gradient.448 We have performed (his same test with V2>0. and while aliasing oceurs at the same interval. (here is no overall growth in the perturbations.," We have performed this same test with $N_x^2 > 0$, and while aliasing occurs at the same interval, there is no overall growth in the perturbations."449 This may be due to the fact that the perturbationsdecay asvaptotically for V2>0 (see expression [??]])., This may be due to the fact that the perturbationsdecay asymptotically for $N_x^2 > 0$ (see expression \ref{ASYMP}] ]).450 Fieure 14. summarizes (he parameter space we have surveved. indicating (hat there is instability only for d~0 and NV?<0.," Figure \ref{f14} summarizes the parameter space we have surveyed, indicating that there is instability only for $\qe \simeq 0$ and $N_x^2 < 0$."451 The numerical resolution in all of these runs is 512x512., The numerical resolution in all of these runs is $512 \times 512$.452" Figure 15. shows the evolution of the radial velocity in Run 10. a run with realistic parameters for a disk with a nearlv-Ixeplerian rotation profile ancl racial gradients on the order of the disk radius: ¢=1.5 and NZ,,,,=—0.01 (corresponding to Rie—0.004)."," Figure \ref{f15} shows the evolution of the radial velocity in Run 10, a run with realistic parameters for a disk with a nearly-Keplerian rotation profile and radial gradients on the order of the disk radius: $q = 1.5$ and $N_{x,min}^2 =453-0.01$ (corresponding to ${\rm Ri} \simeq -0.004$ )."454 Clearly no instability is occurring on a dynamical timescale., Clearly no instability is occurring on a dynamical timescale.455 This plot is typical of all runs for which the evolution was stable., This plot is typical of all runs for which the evolution was stable.456 To give a sense for the minimum growth rate that we are able to measure. we have also plotted in Figure 15. (he results [rom several unstable runs with g=0 and a boost equivalent to the velocity at the minimum in .V2 for Run 10.," To give a sense for the minimum growth rate that we are able to measure, we have also plotted in Figure \ref{f15}457 the results from several unstable runs with $q = 0$ and a boost equivalent to the velocity at the minimum in $N_x^2$ for Run 10."458 lt is difficult tomeasure a growth rate for the smallest value of N> Penne but it is clear that there is activity present in (his run which does not occur inthe stable run.," It is difficult tomeasure a growth rate for the smallest value of $N_{x,min}^2$ , but it is clear that there is activity present in this run which does not occur inthe stable run."459 Based upon, Based upon46010st ealaxy is stronglv elliptical. with a (projected) eccentricity of e~0.6.,"host galaxy is strongly elliptical, with a (projected) eccentricity of $e\sim 0.6$."461" Furthermore. the wines of (he ""X are closely aligned with the minor axis of (he gas distribution. supporting a nodel in which the wings correspond to a colliding backflow that has ""blown"" out of the ISM along the direction of least resistance."," Furthermore, the wings of the “X” are closely aligned with the minor axis of the gas distribution, supporting a model in which the wings correspond to a colliding backflow that has “blown” out of the ISM along the direction of least resistance."462 Although 3C 403 is the only X-shaped radio galaxy or which high-resolution. X-ray maps of the hot ISAT ave available. Capettietal.(2002) wave noted that a number of N-shaped sources have wings that are oriented along the minor axis of theoplical host galaxy.," Although 3C 403 is the only X-shaped radio galaxy for which high-resolution X-ray maps of the hot ISM are available, \citet{capetti02} have noted that a number of X-shaped sources have wings that are oriented along the minor axis of the host galaxy."463 This suggests that the conclusion of Kraftetal.(2005). [or 3C 403 may be more generally (rue., This suggests that the conclusion of \citet{kraft05} for 3C 403 may be more generally true.464 There is one particular svstem. 0402+379. (hat might provide a direct. view of spin alignment in a binary black hole svstem.," There is one particular system, $+$ 379, that might provide a direct view of spin alignment in a binary black hole system."465 Very Long Baseline Array (VLBA) imaging of this radio galaxy by. Manessοἱal.(2004) cliscoverecl (wo compact flat spectrum radio cores. ancl follow-up VLBA observations presented by Rodriguezοἱal.(2006). showed the cores to be stationary.," Very Long Baseline Array (VLBA) imaging of this radio galaxy by \citet{maness04} discovered two compact flat spectrum radio cores, and follow-up VLBA observations presented by \citet{rodriguez06} showed the cores to be stationary."466 A binary supermassive black hole is the most satisfactory explanation for this source. with the projected clistance between the two black holes being only ppc.," A binary supermassive black hole is the most satisfactory explanation for this source, with the projected distance between the two black holes being only pc."467 Within the context of our gas-rich merger scenario. these (wo black holes already have aligned spins.," Within the context of our gas-rich merger scenario, these two black holes already have aligned spins."468 Existing VLBA data only show a jet associated with one of the radio cores., Existing VLBA data only show a jet associated with one of the radio cores.469 We predict that. ifa jet is eventually Found associated with the other radio core. it will have the same position angle as the existing jet.," We predict that, if a jet is eventually found associated with the other radio core, it will have the same position angle as the existing jet."470 Evidence already exists for alignment of a coalescence remnant with its galactic scale eas disk., Evidence already exists for alignment of a coalescence remnant with its galactic scale gas disk.471 Perlmanetal.(2001) imaged the host galaxies of three compact svnuuelric objects and discovered nuclear gas disks approximately normal (o the jet axis., \citet{perlman01} imaged the host galaxies of three compact symmetric objects and discovered nuclear gas disks approximately normal to the jet axis.472 The presence of such a nuclear gas disk as well as disturbances in the outer isophotes of all three host galaxies suggests (hat these galaxies had indeed suffered. major gas-cich mergers within the past 10? vvr., The presence of such a nuclear gas disk as well as disturbances in the outer isophotes of all three host galaxies suggests that these galaxies had indeed suffered major gas-rich mergers within the past $10^8$ yr.473 In conclusion. we propose Chat in the majority of ealactic mergers. torques from gas accretion alien the spins of supermassive black holes aud their orbital axis with large-scale eas disks.," In conclusion, we propose that in the majority of galactic mergers, torques from gas accretion align the spins of supermassive black holes and their orbital axis with large-scale gas disks."474 This scenario helps explain the ubiquity of black holes in galaxies despite the potentially large kicks Ilrom gravitational radiation recoil., This scenario helps explain the ubiquity of black holes in galaxies despite the potentially large kicks from gravitational radiation recoil.475 Further observations. particularly of galaxv mergers (hat do not involve significant. amounts of gas. will test our. predictions and may point to a class of large galaxies without central black holes.," Further observations, particularly of galaxy mergers that do not involve significant amounts of gas, will test our predictions and may point to a class of large galaxies without central black holes."476 We thank Doug Hamilton and Eve Ostriker for insightful discussions., We thank Doug Hamilton and Eve Ostriker for insightful discussions.477 TB thanks the UMCP-Astronomy Center for Theory and Computation Prize Fellowship program for support., TB thanks the UMCP-Astronomy Center for Theory and Computation Prize Fellowship program for support.478 CSR and MCM gratefully acknowledge support from the National Science Foundation uuder erants AST0205990 (CSR) and ASTOGOTA28 (CSR ancl MCA., CSR and MCM gratefully acknowledge support from the National Science Foundation under grants AST0205990 (CSR) and AST0607428 (CSR and MCM).479"efficiency case, MBH seeds form in the regions of earliest star formation, which tend to be the halos which become the most massive later on in every simulation.","efficiency case, MBH seeds form in the regions of earliest star formation, which tend to be the halos which become the most massive later on in every simulation."480" The most massive halos have also experienced the greatest number of mergers, which further populates them with MBH seeds brought in by satellites."," The most massive halos have also experienced the greatest number of mergers, which further populates them with MBH seeds brought in by satellites."481" Thus, our model predicts that halos with masses greater than Mnyajc10? will be extremely likely to host MBH seeds, even if the formation efficiency of such seeds is small."," Thus, our model predicts that halos with masses greater than $M_{\rm halo} \sim 10^{9}$ will be extremely likely to host MBH seeds, even if the formation efficiency of such seeds is small."482" 'This result is consistent with observations of the local universe, where MBHs are found to occupy halos above a similar mass threshold with high likelihood (Ferrareseetal.2006;Wehner&Harris2006)."," This result is consistent with observations of the local universe, where MBHs are found to occupy halos above a similar mass threshold with high likelihood \citep{Ferrarese06,Wehner06}."483". To estimate the fate of MBHs at later cosmic times, we have run a lower-resolution simulation of galaxy h258 to z—0 and traced the mass evolution of halos."," To estimate the fate of MBHs at later cosmic times, we have run a lower-resolution simulation of galaxy $h258$ to $z = 0$ and traced the mass evolution of halos."484" Halos with mass between 10"" and 10!? at z=5 can remain in the same mass range at z—0, be stripped, or become incorporated in the main halo."," Halos with mass between $10^7$ and $10^{10}$ at $z = 5$ can remain in the same mass range at $z=0$, be stripped, or become incorporated in the main halo."485" In this particular case, ~ of halos with Mpalo>10? - that host MBHs already at z—5 — are stripped during their evolution, ending up as halos with z=0 masses of 105— "," In this particular case, $\sim$ of halos with $M_{\rm halo} > 10^{9}$ – that host MBHs already at $z = 5$ – are stripped during their evolution, ending up as halos with $z = 0$ masses of $10^8 - 10^{10}$ ."486"While one zoomed-in simulation cannot give us broad Μο...statistics on the evolution of all z=5 halos, we can deduce that while a large number of high redshift MBH hosts undergo hierarchical merging and settle in massive galaxies, a non-negligible fraction have more quiescent merger histories and do not grow substantially, and make up today's population of low-mass galaxies which may harbor MBHs."," While one zoomed-in simulation cannot give us broad statistics on the evolution of all $z = 5$ halos, we can deduce that while a large number of high redshift MBH hosts undergo hierarchical merging and settle in massive galaxies, a non-negligible fraction have more quiescent merger histories and do not grow substantially, and make up today's population of low-mass galaxies which may harbor MBHs."487" On the other hand, the occupation fraction of halos with masses less than Mhalo~10° is sensitive to the physics of MBH formation, via xsceq."," On the other hand, the occupation fraction of halos with masses less than $M_{\rm halo} \sim 10^{9}$ is sensitive to the physics of MBH formation, via $\chi_{\rm seed}$ ."488" For example, in galaxies h258 and hz1, halos with mass M~10? are nearly likely to host a MBH seed if yseeq=0.5, but only likely when Yseeq=0.05."," For example, in galaxies $h258$ and $hz1$, halos with mass $M \sim 10^9$ are nearly likely to host a MBH seed if $\chi_{\rm seed} = 0.5$, but only likely when $\chi_{\rm seed} = 0.05$."489" The dependence of occupation fraction with yseeq becomes weaker in galaxy h603, however."," The dependence of occupation fraction with $\chi_{\rm seed}$ becomes weaker in galaxy $h603$, however."490" We suggest that in the larger halos, metals diffuse throughout the simulation more efficiently (due to increased star formation and possible ejection of metals from galaxies) which suppresses MBH formation in their satellites."," We suggest that in the larger halos, metals diffuse throughout the simulation more efficiently (due to increased star formation and possible ejection of metals from galaxies) which suppresses MBH formation in their satellites."491" In the isolated low-mass disk galaxy, however, the relative lack of metal diffusion allows MBHs to continue forming in these halos at a time when MBH formation would have been truncated in a larger overdensity."," In the isolated low-mass disk galaxy, however, the relative lack of metal diffusion allows MBHs to continue forming in these halos at a time when MBH formation would have been truncated in a larger overdensity."492" In this study we have aimed for high resolution with a sample of field galaxies; for a simulated uniform volume, including voids and high density regions more completely samples halos in the universe) our (whichresults may be slightly different."," In this study we have aimed for high resolution with a sample of field galaxies; for a simulated uniform volume, including voids and high density regions (which more completely samples halos in the universe) our results may be slightly different."493" However, a test study using a small (6 Mpc) uniform volume gives results consistent with our findings here."," However, a test study using a small (6 Mpc) uniform volume gives results consistent with our findings here."494" We have undertaken a study of theformation and evolution of seed MBHs using cosmological simulations,"," We have undertaken a study of theformation and evolution of seed MBHs using cosmological simulations,"495Ultra Luminous Infrared Galaxies (hereafter ULIRGs) are an enigmatic class of sources which emit most of their energy in the far-infrared (FIR. 8-100047: domain (Sanders&Mirabel. 1996).. with luminosities above ~10! 1.. (Le.. comparable to QSO luminosities).,"Ultra Luminous Infrared Galaxies (hereafter ULIRGs) are an enigmatic class of sources which emit most of their energy in the far-infrared (FIR, $\mu496m$ ) domain \citep{Sanders}, with luminosities above $\sim 10^{12}$ $_\odot$, (i.e., comparable to QSO luminosities)."497 The importance of understanding the physical processes at work in ULIRGs ts strengthened by the observational evidence that they are generally advanced mergers of gas-rich galaxies: these events are now considered to be at the origin of some of the massive elliptical and SO galaxies (Hopkinsetal..2005.2006;Springel2005) and the QSO stage could be a phase during the evolution of these systems.," The importance of understanding the physical processes at work in ULIRGs is strengthened by the observational evidence that they are generally advanced mergers of gas-rich galaxies; these events are now considered to be at the origin of some of the massive elliptical and S0 galaxies \citep{hopkins2005, hopkins2006, springel05} and the QSO stage could be a phase during the evolution of these systems."498 However. understanding their physical nature is complicated by the large amount of obscuration from dust present in these sources. which makes it difficult to directly observe the nuclear X-ray observations of ULIRGs performed with XMM-Newton (Franceschinietal..2003:Braitoal...2003) and Chandra (Ptaketal.2003:Tengal..2005) and recently Suzaku (Tengetal..2009) have proved to be a fundamental tool to disentangle the contribution of starburst and AGN activity and to investigate the presence of hidden AGNs in these sources.," However, understanding their physical nature is complicated by the large amount of obscuration from dust present in these sources, which makes it difficult to directly observe the nuclear X-ray observations of ULIRGs performed with XMM-Newton \citep{France, Braito} and Chandra \citep{Ptak,Teng05} and recently Suzaku \citep{Teng} have proved to be a fundamental tool to disentangle the contribution of starburst and AGN activity and to investigate the presence of hidden AGNs in these sources."499 These observations have shown that ULIRGs are intrinsically faint X-ray sources. their observed X-ray luminosities being typically L(2-10 10?—0οςs7'..," These observations have shown that ULIRGs are intrinsically faint X-ray sources, their observed X-ray luminosities being typically L(2–10 $10^{42}-10^{43}$."500 The X-ray spectra of ULIRGs are complex and present the signatures of both the starburst eand the AGN activity. confirming the composite. nature. of ULIRGs.," The X-ray spectra of ULIRGs are complex and present the signatures of both the starburst and the AGN activity, confirming the composite nature of ULIRGs."501" These studies have also shown that more than half of the local brightest AGN-ULIRGs (5/8) host an obscured AGN. with three being Comption Thick (Nj,> 107*em7?; GC60240. Vignatietal.1999;; Mrk 231. Braitoetal.2004 and UGC35101. Imanishietal. 2003))."," These studies have also shown that more than half of the local brightest AGN-ULIRGs (5/8) host an obscured AGN, with three being Comption Thick $_H >50210^{24}$ $^{-2}$; NGC6240, \citealt{6240}; Mrk 231, \citealt{231} and UGC5101, \citealt{5101}) )."503 Observations above 10 keV are thus fundamental for measuring the intrinsic X-ray luminosity of obscured AGN hosted in ULIRGs and its contribution to their high observed FIR IRAS 19254-7245 (also known as the SuperAntennae) belongs to a flux limited sample at 60m composed of the 15 brightest nearby ULIRGs (Genzeletal..1998)., Observations above 10 keV are thus fundamental for measuring the intrinsic X-ray luminosity of obscured AGN hosted in ULIRGs and its contribution to their high observed FIR IRAS 19254-7245 (also known as the ) belongs to a flux limited sample at $60\mu$ m composed of the 15 brightest nearby ULIRGs \citep{Genzel}.504". Located at 2=0.062. hhas an infrared luminosity of Lsioo;—1.1x10""L corresponding to a bolometric luminosity of L5,~4x10y?erg s!."," Located at $z=0.062$, has an infrared luminosity of $L_\mathrm{8-1000\mum}=1.1\times 10^{12}L_\odot$ corresponding to a bolometric luminosity of $L_{bol}\sim 4\times 10^{45}$ erg $^{-1}$."505 Like most of the ULIRGs. IRAS 19254-7245 is a merger system of two gas-rich spiral galaxies.," Like most of the ULIRGs, IRAS 19254-7245 is a merger system of two gas-rich spiral galaxies."506 The southern nucleus. optically classified as a Seyfert 2. 1s one of the brightest nearby ULIRGs which has proved to host both a powerful starburst and an obscured AGN. while there is no evidence of AGN activity in the northern A previous X-ray observation of performed with suggested that this ULIRG harbors a heavily obscured and high-luminosity AGN.," The southern nucleus, optically classified as a Seyfert 2, is one of the brightest nearby ULIRGs which has proved to host both a powerful starburst and an obscured AGN, while there is no evidence of AGN activity in the northern A previous X-ray observation of performed with suggested that this ULIRG harbors a heavily obscured and high-luminosity AGN."507 Indeed the hard power-law continuum above 2 keV (photon index I= 1.3) and the detection of a strong Fe-K« emission line at 6.5+0.1 keV (EW~1.4 keV) were highly indicative of a Compton-thick source al.. 2003)..," Indeed the hard power-law continuum above 2 keV (photon index $\Gamma508=1.3$ ) and the detection of a strong $\alpha$ emission line at $6.5\pm 0.1$ keV $EW \sim 1.4$ keV) were highly indicative of a Compton-thick source \citep{Braito}."509 As the two nuclei are located ~ 9 aresee apart from each other. ccould not resolve them: however the centroid of the hard X-ray emission was spatially coincident with the southern nucleus.," As the two nuclei are located $\sim$ 9 arcsec apart from each other, could not resolve them; however the centroid of the hard X-ray emission was spatially coincident with the southern nucleus."510 Chandra observation which would settle or solve this issue has not been performed yet., Chandra observation which would settle or solve this issue has not been performed yet.511 The best fit model obtained for the 0.5-10 keV X-ray emission detected with wwas composed by a strong soft thermal component. associated with the starburst emission and à hard X-ray component associated with the AGN activity.," The best fit model obtained for the 0.5-10 keV X-ray emission detected with was composed by a strong soft thermal component, associated with the starburst emission and a hard X-ray component associated with the AGN activity."512 This AGN component was parametrized with a Compton thick AGN model and was composed of a pure Compton-reflected continuum (with D~ 1.8). a scattered power law component and a strong Fe emission line.," This AGN component was parametrized with a Compton thick AGN model and was composed of a pure Compton-reflected continuum (with $\Gamma\sim5131.8$ ), a scattered power law component and a strong Fe emission line."514 The observed 2-10 keV, The observed 2–10 keV515(see Fie. 8)).,(see Fig. \ref{ircameradrawing}) ).516 The surfaces of the optical hatie can and other components adjaceut to the optical bean. such as the filter wheel aud array housing. are anodized black to suppress internal reflections.," The surfaces of the optical baffle can and other components adjacent to the optical beam, such as the filter wheel and array housing, are anodized black to suppress internal reflections."517 The outer surface of the optical baffle can is wrapped with a sinele laver of aluuinized niwlar o chhance its performance as a radiation shield., The outer surface of the optical baffle can is wrapped with a single layer of aluminized mylar to enhance its performance as a radiation shield.518 The optical baffle can ataches to the base of the uiror cell. aud thereby provides the aligumeut and mechanical support or the secondary. nuinmror.," The optical baffle can attaches to the base of the mirror cell, and thereby provides the alignment and mechanical support for the secondary mirror."519 One of the two Linch flaILCs Is Used as all access ort for the model 21 CTT Ue retriecrator. while he other carries a vaciimn valve aud pumping nanifold.," One of the two 4-inch flanges is used as an access port for the model 21 CTI He refrigerator, while the other carries a vacuum valve and pumping manifold."520 During final assembly this port also oovides access to the interior of the camera., During final assembly this port also provides access to the interior of the camera.521 The camera is equipped with a cold filter wheel hat can hold up to six. onc-inch filters., The camera is equipped with a cold filter wheel that can hold up to six one-inch filters.522 The uonment of inertia of the filter wheel is low aud he required precision is undemanudiug so the filter wheel is driven directly by a stepper motor with 10 gear reduction., The moment of inertia of the filter wheel is low and the required precision is undemanding so the filter wheel is driven directly by a stepper motor with no gear reduction.523 Cryogenic mechanisms are frequently the most unreliable part of infrared instruments., Cryogenic mechanisms are frequently the most unreliable part of infrared instruments.524 These uechanisnis are usually actuated by a wari external motor via ferro-fliidic fecd-througls., These mechanisms are usually actuated by a warm external motor via ferro-fluidic feed-throughs.525 A cryogenic motor simplifies design and assenblv, A cryogenic motor simplifies design and assembly526the distance at which the star would have V=7.3.,the distance at which the star would have $V=7.3$.527 Actually. we are not directly interested in the LF at M=4. but rather in the integral of LY7R7 over the subpopulations that make up the AM.=4 bin of the LF.," Actually, we are not directly interested in the LF at $M_V=4$, but rather in the integral of $L^{3/2} R^{-7/2}$ over the subpopulations that make up the $M_V = 4$ bin of the LF."528" ence The stellar radii are determined from IHipparcos/Tvcho (D4.V) photometry and the color/surlace-brightness relation of Gould&Morgan(2003).. My,/5. ultimately derived from vanBelle(1999)."," Hence The stellar radii are determined from Hipparcos/Tycho $(B_T,V_T)$ photometry and the color/surface-brightness relation of \citet{gould03}, $\log {R/R_\odot} = 0.597 + 0.536(B_T - V_T) - {M_{V_T}/5}$ , ultimately derived from \citet{belle99}."529". To normalize F. we adopt the values associated with A4:=5 stus: Ly0.501... Ry=0.97 R..m,=0.0025pe*."," To normalize $F$ , we adopt the values associated with $M_V = 5$ stars: $L_0 = 0.86\,L_\odot$, $R_0 = 0.97\,R_\odot$, $n_0 = 0.0025\,{\rm pc}^{-3}$."530 The resulting Iunction (Fig. 1)), The resulting function (Fig. \ref{figF}) )531" shows that the majority of stars that are probed will be F and G tvpe (2<M,S 6.5).", shows that the majority of stars that are probed will be F and G type $2 \la M_{V} \la 6.5$ ).532 To check how this distribution depends on the size of the volume sampled. we recaleulate the distribution function for the case in which the observed stars cover a much larger volume — one which would be better described bv a thin disk. rather than a spherically uniform distribution.," To check how this distribution depends on the size of the volume sampled, we recalculate the distribution function for the case in which the observed stars cover a much larger volume – one which would be better described by a thin disk, rather than a spherically uniform distribution."533 For this case we find that the scalings shown in equation (7)) are replaced by FCU)xnhLR?. where fr is the scale height of each population.," For this case we find that the scalings shown in equation \ref{equF}) ) are replaced by $F(M_{V}) \propto n h L R^{-2}$, where $h$ is the scale height of each population."534 The distribution function is still dominated by F and G stars. but there are more Ix. and early to mid M stars (0MyS 12).," The distribution function is still dominated by F and G stars, but there are more K and early to mid M stars $7 \la M_{V} \la 12$ )."535 Of course. it is common knowledge that magnitude-limited (FoxnL*?) samples of main-sequence stars will be dominated by F and G stars.," Of course, it is common knowledge that magnitude-limited $(F\propto n L^{3/2})$ samples of main-sequence stars will be dominated by F and G stars."536 The interesting feature of Figure 1 is Chat this result does not qualitatively change despite the addition of the factor 2? in equation (7)). which very strongly. favors laler-twpe stars.," The interesting feature of Figure \ref{figF} is that this result does not qualitatively change despite the addition of the factor $R^{-7/2}$ in equation \ref{equF}) ), which very strongly favors later-type stars."537 Equation (6)) describes what kinds of NSP systems can be detected by a certain survey. given a photometric detection limit.," Equation \ref{equfin}) ) describes what kinds of XSP systems can be detected by a certain survey, given a photometric detection limit."538" The threshold A\2,,, is determined by taking account of the [act Chat the data stream [rom an ASP search must be analvzed for any. combination of the parameters 2. a. and r within reasonable ranges."," The threshold $\Delta\chi^2_{\rm min}$ is determined by taking account of the fact that the data stream from an XSP search must be analyzed for any combination of the parameters $R$, $a$, and $r$ within reasonable ranges."539 Such analvsis will. however. vield a number of false-positive detections due to random noise.," Such analysis will, however, yield a number of false-positive detections due to random noise."540" The threshold value of Ay, must be chosen to vield a manageable nunber of candidate svstems for follow-up observations.", The threshold value of $\Delta \chi_{\rm min}^{2}$ must be chosen to yield a manageable number of candidate systems for follow-up observations.541 To determine Az. we generate 1000 independent simulated streams of photometric observations of a single svstem with a host star of onesolar mass and a circular planetary orbit.," To determine $\Delta \chi_{\rm min}^{2}$, we generate 1000 independent simulated streams of photometric observations of a single system with a host star of onesolar mass and a circular planetary orbit."542 We attempt to simulate a schedule that would be characteristic of an all-skv survey, We attempt to simulate a schedule that would be characteristic of an all-sky survey543 P- and Rox-abuuc vices determinations will be discussed.,$P$ - and $R_{23}$ -abundances determinations will be discussed.544 In order to obtali a more reliable determination of the gradieut. a siugle value of the abundauces is obtained for each region by averaging the values from the VISTA data-set and the one.," In order to obtain a more reliable determination of the gradient, a single value of the abundances is obtained for each region by averaging the values from the VISTA data-set and the ALICE-MIDAS one."545 This €au be clone because the abuauces between tliese two data-set are very similar., This can be done because the abundances between these two data-set are very similar.546 The gradients determined in tliis way are showrin column 3 of Table 1 ancl i Figure 3.., The gradients determined in this way are shown in column 3 of Table 1 and in Figure \ref{fig3}.547 The valie rou the average al»indauce is —0.6 dex/kpc. which is quite unrealistic as weLas hose determined rom the 223 and DP methods which are eve steeper.," The value from the average abundance is $-0.6$ dex/kpc, which is quite unrealistic as well as those determined from the $R{23}$ and $P$ methods which are even steeper."548" The reasou for suc as eep value is the arge dillereuce in metallicity among the internal and the external regions 1u this galaxy. 1.2 dex. aud the small gala""l'ocentric distauce rauge. 2 spc."," The reason for such a steep value is the large difference in metallicity among the internal and the external regions in this galaxy, $1.2$ dex, and the small galactocentric distance range, $2$ kpc."549 Although a eraclient at fixed distance |as litte sense because it will be cletermined with ouly two regions. it is useful becaise it shows hat tle abundances change crauatically. [rom $8.5 dex at the center to 7.7 dex at tle o utskirts.," Although a gradient at fixed distance has little sense because it will be determined with only two regions, it is useful because it shows that the abundances change dramatically, from $8.5$ dex at the center to $7.7$ dex at the o utskirts."550 T1e slope derived usie the more/less metallic regions is —0.6 dex/kpe while the valle using e ou most/inmost regions is —0.3 dex/kpc., The slope derived using the more/less metallic regions is $-0.6$ dex/kpc while the value using the out most/inmost regions is $-0.3$ dex/kpc.551 These values are similar to those determined with tle uatlieiatical fittiug., These values are similar to those determined with the mathematical fitting.552 Tus is becatse U06a2 is probably over euriclied as discussed in pape: . Such richinent does not force the steep [n]gradient. because a value of —0.52 dex/kpc is obtained witlot is region.," This is because U96a2 is probably over enriched as discussed in paper I. Such enrichment does not force the steep gradient, because a value of $-0.52$ dex/kpc is obtained without this region."553 One might thirk that 'eelon U»56al. with its very low metallicity. is lorci1g such steep acien.," One might think that region U56a1, with its very low metallicity, is forcing such steep gradient."554 Actually. the ox aud P eracients do uot chiauge when this regiou is uot c«οιed.," Actually, the $R_{23}$ and $P$ gradients do not change when this region is not considered."555 Ou e corary. the average 1jetallicity gradient does shallow mainly due to the higlier abtdauces or U96aI.," On the contrary, the average metallicity gradient does shallow mainly due to the higher abundances for U96a4."556 Values as large as —0.6 dex/spe are not acceptable because they are quite uu'ealistic., Values as large as $-0.6$ dex/kpc are not acceptable because they are quite unrealistic.557 The value of the gradient adopted Lor this galaxy is —0.[ dex/kpe. as 5ated in Table 1.," The value of the gradient adopted for this galaxy is $-0.4$ dex/kpc, as stated in Table 1."558 The clis)ersioln is about 0.1., The dispersion is about $0.1$.559 Such eradient is large but suniar to the oue o. UGC' 6205., Such gradient is large but similar to the one of UGC 6205.560 From the 'esults of the averaged. distauces determination discussed above. it can be COLLCuced there is a real. aud large. Change iu the abundauces of this galaxy.," From the results of the averaged distances determination discussed above, it can be concluded there is a real, and large, change in the abundances of this galaxy."561 Sucl differences cauιοί be explained by differeuces iu the SER (Réyyes-Pérrez 2009)., Such differences cannot be explained by differences in the SFR (Réyyes-Pérrez 2009).562" Finally. i tis important to point out that a more accurate cetermiuation ol the eracdient will be very difficul. because there is ouls ""nOie more region in the galaxy."," Finally, it is important to point out that a more accurate determination of the gradient will be very difficult, because there is only one more region in the galaxy."563 Iu aly case. more high quality data are needed in order to confin these results.," In any case, more high quality data are needed in order to confirm these results."564 The number o“HAUL regious is also small in this barre ealaxy. but for the same reasous as lor UGC 5296. a gradient will be determined.," The number of regions is also small in this barred galaxy, but for the same reasons as for UGC 5296, a gradient will be determined."565 Moreover. as belig the ouly barred galaxy in our saimple. it will be very inte‘esting to kuow if the gradient is siinilar o tlie other barred late-type galaxies.," Moreover, as being the only barred galaxy in our sample, it will be very interesting to know if the gradient is similar to the other barred late-type galaxies."566 The oxygen abuncleuices can be determiued for ouly five regeious., The oxygen abundances can be determined for only five regions.567 Also. as for UGC 5296. there are two set of data for each region.," Also, as for UGC 5296, there are two set of data for each region."568 Therelore. a total of eigit different. estimations of the gracdieut cali be obtained for this galaxy.," Therefore, a total of eight different estimations of the gradient can be obtained for this galaxy."569 Those values determijecl wih the ALICE-MIDAS abundances are very similar. but 11ose obtaiied. from the VISTA abuia1‘es are not.," Those values determined with the ALICE-MIDAS abundances are very similar, but those obtained from the VISTA abundances are not."570 Two of them are positive., Two of them are positive.571 The best way to proceed is. jist as in UCC 5296. to cetermiue an average abundance for each," The best way to proceed is, just as in UGC 5296, to determine an average abundance for each"572Faleke. Kórrding Markoff (2004: see also Koerding. Faleke Corbel 2006): scaling of fast variability properties with mass and accretion rate are reported in McHardy et al. (,"Falcke, Körrding Markoff (2004; see also Koerding, Falcke Corbel 2006); scaling of fast variability properties with mass and accretion rate are reported in McHardy et al. ("5732006) and Kórrding et al. (,2006) and Körrding et al. (5742007): more qualitative similarities between XRB and AGN üccretion are noted in Kórrding. Jester Fender (2006) and also discussed in Marscher et al. (,"2007); more qualitative similarities between XRB and AGN accretion are noted in Körrding, Jester Fender (2006) and also discussed in Marscher et al. ("5752002) and Chatterjee et al. (,2002) and Chatterjee et al. (5762009).,2009).577 The temporal evolution of XRB jets. relatively rapid compared to AGN. has allowed many estimates of the power (e.g. Fender 2001: Gallo et al.," The temporal evolution of XRB jets, relatively rapid compared to AGN, has allowed many estimates of the power (e.g. Fender 2001; Gallo et al."578" 2005: Kórrding. Fender Migliari 2006) and speed (e.g. Mirabel Rodriguez 1994: Tiller-Jones. Fender Nakar 2006) of the jets and their connection © accretion ""state"" as characterized by the X-ray emission (e.g. Fender et al."," 2005; Körrding, Fender Migliari 2006) and speed (e.g. Mirabel Rodriguez 1994; Miller-Jones, Fender Nakar 2006) of the jets and their connection to accretion `state' as characterized by the X-ray emission (e.g. Fender et al."579 1999: Fender. Belloni Gallo 2004. [hereafter FBGO4]: Corbel et al.," 1999; Fender, Belloni Gallo 2004 [hereafter FBG04]; Corbel et al."580 2004)., 2004).581 Importantly these studies have shown hat the jet power of a black hole XRB. as well as the radiative efficiency of the accretion flow. can change dramatically in the same source at the same overall radiative luminosity on timescales ur shorter than those associated with significantly changing mass or angular momentum.," Importantly these studies have shown that the jet power of a black hole XRB, as well as the radiative efficiency of the accretion flow, can change dramatically in the same source at the same overall radiative luminosity on timescales far shorter than those associated with significantly changing mass or angular momentum."582 In very brief summary. in black hole XRBs the coupling of radio emission (and hence jets) to X-ray state and luminosity is as follows: at Eddington ratios (in terms of X-ray luminosity) below about 0.01. sources seem to be exclusively in the “hard” X-ray state in which the X-ray emission is dominated by a componen extending to ~100 keV. widely (but not universally) accepted to arise via thermal Comptonisation of seed photons by a ho flow / corona.," In very brief summary, in black hole XRBs the coupling of radio emission (and hence jets) to X-ray state and luminosity is as follows: at Eddington ratios (in terms of X-ray luminosity) below about 0.01, sources seem to be exclusively in the `hard' X-ray state in which the X-ray emission is dominated by a component extending to $\sim 100$ keV, widely (but not universally) accepted to arise via thermal Comptonisation of seed photons by a hot flow / corona."583 In this state there is strong aperiodic variability and a steady. powerful. flat-spectrum jet.," In this state there is strong aperiodic variability and a steady, powerful, flat-spectrum jet."584 The luminosities of the two components scales roughly as LyxLy?, The luminosities of the two components scales roughly as $L_{\rm radio} \propto L_{X}^{0.6-0.7}$.585"UA higher Eddington ratios. reached generally by transient outbursting systems. sources can switch into ""softer states in which the X-ray spectrum is dominated by a cooler (~1 keV) componen with a near-blackbody spectrum. generally interpreted as the inner accretion dise."," At higher Eddington ratios, reached generally by transient outbursting systems, sources can switch into `softer' states in which the X-ray spectrum is dominated by a cooler $\sim 1$ keV) component with a near-blackbody spectrum, generally interpreted as the inner accretion disc."586" In this state the radio emission is either dramatically suppressed by a factor =50 or evolves to a fading. optically thin. state. both scenarions suggesting the ""quenching of the core jet (possibly with some remnant extended emission)."," In this state the radio emission is either dramatically suppressed by a factor $\geq 50$ or evolves to a fading, optically thin, state, both scenarions suggesting the `quenching' of the core jet (possibly with some remnant extended emission)."587 In transitions from hard to soft states major radio flares. often resolved as discrete. powerful. ejection events. are commonly observed.," In transitions from hard to soft states major radio flares, often resolved as discrete, powerful, ejection events, are commonly observed."588 Sources generally fade in the soft state until they are once again at à few Eddington (in. Lx Y. and then make a transition back to the hard state in which mode they fade further.," Sources generally fade in the soft state until they are once again at a few Eddington (in $L_{X}$ ), and then make a transition back to the hard state in which mode they fade further."589 The initial hard - Soft state transition is usually at a higher luminosity than the soft = hard return branch. i.e. hysteresis when spectral hardness is compared to luminosity.," The initial hard $\rightarrow$ soft state transition is usually at a higher luminosity than the soft $\rightarrow$ hard return branch, i.e. hysteresis when spectral hardness is compared to luminosity."590 Note that the same source has been observed to make both hard + soft and soft» hard transitions at different luminosities in different outbursts: note further that some sources e.g. Cyg X-I never drop below the Eddington threshold and remain “persistent and variable’., Note that the same source has been observed to make both hard $\rightarrow$ soft and soft $\rightarrow$ hard transitions at different luminosities in different outbursts; note further that some sources e.g. Cyg X-1 never drop below the Eddington threshold and remain 'persistent and variable'.591 For comprehensive reviews on these phenomena. see FRGO4: Remillard MeClintock 2006: Done. Kubota Gierlinski 2007: Fender. Homan Belloni 2009: Belloni 2009.," For comprehensive reviews on these phenomena, see FBG04; Remillard McClintock 2006; Done, Kubota Gierlinski 2007; Fender, Homan Belloni 2009; Belloni 2009."592 The most comprehensive compilation of X-ray data on black hole binaries is presented in Dunn et al. (, The most comprehensive compilation of X-ray data on black hole binaries is presented in Dunn et al. (5932010).,2010).594 Note that X-ray binary systems with comparable properties but hosting a neutron star instead of a black hole (candidate) also show jets. but with a lower ratio of Lus to Ly (Pender Kuulkers 2001: Mighari Fender 2006).," Note that X-ray binary systems with comparable properties but hosting a neutron star instead of a black hole (candidate) also show jets, but with a lower ratio of $L_{\rm radio}$ to $L_{X}$ (Fender Kuulkers 2001; Migliari Fender 2006)."595 In parallel with these advances in the study of black hole jet power. speed and relation to accretion state. there has been a rapid recent growth in the number of estimates of spin of black holes in XRB systems (the spin is generally discussed in terms of the dimensionless spin parameter αν—04/67A7 ," In parallel with these advances in the study of black hole jet power, speed and relation to accretion state, there has been a rapid recent growth in the number of estimates of spin of black holes in XRB systems (the spin is generally discussed in terms of the dimensionless spin parameter $a_* = cJ / GM^2$ "596The effect on LDB ages for the choice of bolometric correction is shown in the bottom panel of Figure 9..,The effect on LDB ages for the choice of bolometric correction is shown in the bottom panel of Figure \ref{bcerr}.597 The bolometric correction differences as a fanction ol effective temperature are mapped to LDB Iuminositv using the effective temperatures at lithium depletion given in Table 1.., The bolometric correction differences as a function of effective temperature are mapped to LDB luminosity using the effective temperatures at lithium depletion given in Table \ref{reftab}.598 For the 1-6 error in LDB ages due to bolometric corrections. we take the average absolute value of the percentage age differences shown in the bottom panel of Figure 9 as 1-0 deviations.," For the $\sigma$ error in LDB ages due to bolometric corrections, we take the average absolute value of the percentage age differences shown in the bottom panel of Figure \ref{bcerr} as $\sigma$ deviations."599 Unlike for the input plivsies variations. it 15 not clear which bolometric correction is the best choice.," Unlike for the input physics variations, it is not clear which bolometric correction is the best choice."600 Thus. the bolometric correction error budget is more conservative than the theoretical error budget.," Thus, the bolometric correction error budget is more conservative than the theoretical error budget."601 Adcditionallv. there are (wo sources of svstematic error in the bolometric correction when applied to the LDB age technique.," Additionally, there are two sources of systematic error in the bolometric correction when applied to the LDB age technique."602 Lithium depletion occurs during the contraction phase., Lithium depletion occurs during the pre-main-sequence contraction phase.603 Thus. the stars al the LDB have smaller log(g) values than main sequence stars.," Thus, the stars at the LDB have smaller $\log$ (g) values than main sequence stars."604 Since empirical bolometric corrections are derived [rom main sequence stars. the bolometric corrections are svstematically different.," Since empirical bolometric corrections are derived from main sequence stars, the bolometric corrections are systematically different."605 Theoretical models used by Lejeune.Chuisinier.&Buser(1997). show the bolometric corrections roughlyvary by 0.05- magnitudes per 0.5 dex change in log(g).," Theoretical models used by \citet{lej97}606 show the bolometric corrections roughlyvary by 0.05 magnitudes per 0.5 dex change in $\log$ (g)."607 The reference calculations have 70.4 dex larger log(g) values after 1 Gyr., The reference calculations have $\sim$ 0.4 dex larger $\log$ (g) values after 1 Gyr.608 Using bolometric corrections of log(g)—5.0 for stars that actually have log(g)—4.5 resulis in a overestimate of the LDD ages across all luminosities., Using bolometric corrections of $\log$ (g)=5.0 for stars that actually have $\log$ (g)=4.5 results in a overestimate of the LDB ages across all luminosities.609 The second source of svstematie error in the bolometric correction results [rom the potential for the stellar cluster (o differ in metallicity [rom the empirical stellar templates enmploved in determining (he bolometric correction., The second source of systematic error in the bolometric correction results from the potential for the stellar cluster to differ in metallicity from the empirical stellar templates employed in determining the bolometric correction.610 Bolometric corrections vary as a function of metallicity especially for temperatures « 3000 Ix. The theoretical models used by (1997). predict a 0.1 dex change in ΕΟΤ) results in a difference in the resulting LDD ages lor the lowest luminosities., Bolometric corrections vary as a function of metallicity especially for temperatures $<$ 3000 K. The theoretical models used by \citet{lej97} predict a 0.1 dex change in [Fe/H] results in a difference in the resulting LDB ages for the lowest luminosities.611" The elfect is less than toward higher luminosities,", The effect is less than toward higher luminosities.612 These (wo sources of systematic uncertainty are added in (quadrature with the dominant error source based on intercomparisons between alternative bolometric correction determinations as discussed in (he preceding paragraph., These two sources of systematic uncertainty are added in quadrature with the dominant error source based on intercomparisons between alternative bolometric correction determinations as discussed in the preceding paragraph.613 The resulüng 1-0 error is shown as ihe long-dashe line in Figuree 8.., The resulting $\sigma$ error is shown as the long-dash line in Figure \ref{errfig}.614 A comparison of theoretical isochrone fits to the observed color-magnitude diagram of the Pleiades is an additional check (o the svstematic errors that may be present in the bolometrie corrections., A comparison of theoretical isochrone fits to the observed color-magnitude diagram of the Pleiades is an additional check to the systematic errors that may be present in the I-band bolometric corrections.615 For the coolest observed Pleiades members. Tell~3400 Ix. the isochrone fit is too blue by 0.30 magnitudes in (V-I) when using the Hauschildt..(1999) bolometric corrections.," For the coolest observed Pleiades members, $\sim$ 3400 K, the isochrone fit is too blue by 0.30 magnitudes in (V-I) when using the \citet{hau99} bolometric corrections."616 The details of this isochrone fit to the Pleiacles using the YREC stellar evolution code is given in Pinsonneaultetal.—(1998)., The details of this isochrone fit to the Pleiades using the YREC stellar evolution code is given in \citet{pin98}.617. Similar sized discrepancies are generic to all isochrone fits to open cluster coloramagnitude diagrams (Grocholski&Sarajedini 2003).., Similar sized discrepancies are generic to all isochrone fits to open cluster color-magnitude diagrams \citep{gro03}. .618 If the I-band bolometric correction is solely responsible for the discrepancy in (he coloranagnitude diagram. revisions to the bolometric corrections of a," If the I-band bolometric correction is solely responsible for the discrepancy in the color-magnitude diagram, revisions to the bolometric corrections of a"619programs. suffer the observational bias of detecting more easily the short-pertod and massive planets -the rarest ones-. which may be the reason why only of planets in the solar neighbourhood show the signature of a giant planet.,"programs, suffer the observational bias of detecting more easily the short-period and massive planets -the rarest ones-, which may be the reason why only of planets in the solar neighbourhood show the signature of a giant planet."620 Note that this proportion of stars with planetary systems greatly increases when planets in the mass range of Neptune or below are discovered(?).., Note that this proportion of stars with planetary systems greatly increases when planets in the mass range of Neptune or below are discovered\citep{may2008}.621 Extending the planet sample. especially in well-defined volume-limited samples of main-sequence stars as monitored by HARPS. is one of the new challenges of this scientific field. to help understanding the mechanisms which form and maintain planets around other stars.," Extending the planet sample, especially in well-defined volume-limited samples of main-sequence stars as monitored by HARPS, is one of the new challenges of this scientific field, to help understanding the mechanisms which form and maintain planets around other stars."622nol dominate (he opacity of (he F ring.,not dominate the opacity of the F ring.623 This is consistent with the lack of a pronounced slope in the transmission spectrum outside of the dip at 2.87 jam. IH sub-micron grains were common in (the ring. then far [rom the strong absorption band we would expect Onnox(S/AY. which would produce a steep slope in the transmission spectrum which is not seen either here or in earlier ground-based oceultation data (Boshefaf.2002).," This is consistent with the lack of a pronounced slope in the transmission spectrum outside of the dip at 2.87 $\mu$ m. If sub-micron grains were common in the ring, then far from the strong absorption band we would expect $Q_{ext} \propto (s/\lambda)^2,$ which would produce a steep slope in the transmission spectrum which is not seen either here or in earlier ground-based occultation data \citep{Bosh02}."624. The above considerations suggest that only a rather limited set of particle size distributions will be able to reproduce the observed (ransuiission spectrum., The above considerations suggest that only a rather limited set of particle size distributions will be able to reproduce the observed transmission spectrum.625" To test (his supposition. we computed the predicted transmission spectra for various populations of particles using a Mie scattering code. assuming that all the particles are composed of pure crvstalline water ice al 50 Ix. Initial investigations indicated that. neither simple power-laws nor narrow Ilansen-Hovenier distributions could reproduce the observed transmission spectra,"," To test this supposition, we computed the predicted transmission spectra for various populations of particles using a Mie scattering code, assuming that all the particles are composed of pure crystalline water ice at 80 K. Initial investigations indicated that neither simple power-laws nor narrow Hansen-Hovenier distributions could reproduce the observed transmission spectra."626 We therefore considered slightly more complex size distributions. ancl found (hat certain broken power laws could reproduce many of the salient features of the observed. transmission spectra.," We therefore considered slightly more complex size distributions, and found that certain broken power laws could reproduce many of the salient features of the observed transmission spectra."627" In these models the particle size distribution follows a power law n(s)xsSee! up to a critical size 55,444. above which (he size distribution changes to a different (sleeper) power law n(s)xs“i (see Figure 4))."," In these models the particle size distribution follows a power law $n(s) \propto s^{-q_{small}}$ up to a critical size $s_{break}$, above which the size distribution changes to a different (steeper) power law $n(s) \propto s^{-q_{big}}$ (see Figure \ref{pldiag}) )."628" Figure 5. shows the calculated extinction spectra for a range of models with different values of qj, and q;;5. but the same value of ρω= 10jm and an assumed maximunm particle size of 1 mm."," Figure \ref{plmod} shows the calculated extinction spectra for a range of models with different values of $q_{big}$ and $q_{small}$, but the same value of $s_{break}=10 \mu$ m and an assumed maximum particle size of 1 mm."629 As expected. increasing ρω which increases the fraction of micron-sized ancl smaller particles in (he ring— causes the extinction spectrum (o develop a prominent slope al short wavelengths and a peak αἱ 3.1yan. Since these features are not seen in the F-ring data. (μυ needs to be rather low (around 2 assuming s$5;;;» 10jm).," As expected, increasing $q_{small}$ —which increases the fraction of micron-sized and smaller particles in the ring— causes the extinction spectrum to develop a prominent slope at short wavelengths and a peak at $\mu$ m. Since these features are not seen in the F-ring data, $q_{small}$ needs to be rather low (around 2 assuming $s_{break}\sim 10 \mu$ m)."630" On the other hand. decreasing qui; which increases the fraction of very large particles in (he ring— tends to dilute the 2,874 opacity dip."," On the other hand, decreasing $q_{big}$ —which increases the fraction of very large particles in the ring— tends to dilute the $\mu$ m opacity dip."631" Thus to produce an opacity dip of the appropriate magnitude (big Must be fairly high (around3.5 assuming 554,457 10m).", Thus to produce an opacity dip of the appropriate magnitude $q_{big}$ must be fairly high (around3.5 assuming $s_{break}\sim 10 \mu$ m).632" Together. these findings indicate that the size distribution must have a rather sharp break (wilh 4,45~2 aud (big~ 3.5) to reproduce the observed transmission spectra."," Together, these findings indicate that the size distribution must have a rather sharp break (with $q_{small} \sim 2$ and $q_{big} \sim 3.5$ ) to reproduce the observed transmission spectra."633 Such a break is consistent with other spectral data for the F ring (Vahidiniaefa£.2011)., Such a break is consistent with other spectral data for the F ring \citep{Vahidinia11}.634. Ol course. such transmission spectra alone cannot uniquely determine the particle size distributions of these rings.," Of course, such transmission spectra alone cannot uniquely determine the particle size distributions of these rings."635" For example. different. assumed. values of 55,444 lead to somewhat different preferred values of qj."," For example, different assumed values of $s_{break}$ lead to somewhat different preferred values of $q_{big}$."636 Indeed. these data can only place limits on the fraction of the ring particles larger than LO jam across. and do not strongly constrain the (vpical size of these larger particles.," Indeed, these data can only place limits on the fraction of the ring particles larger than 10 $\mu$ m across, and do not strongly constrain the typical size of these larger particles."637 Furthermore. (he ring particles are nol spheres of pure water ice. and (he detailed microstructure of the grains can alter the depth and location of the dip and complicate efforts to quantitatively constrain (he particle size distribution (Vahliidiniaefa£.2011).," Furthermore, the ring particles are not spheres of pure water ice, and the detailed microstructure of the grains can alter the depth and location of the dip and complicate efforts to quantitatively constrain the particle size distribution \citep{Vahidinia11}."638. Indeed. the observed dip occurs at slightly longer wavelengths than predicted by (the simple models used here.," Indeed, the observed dip occurs at slightly longer wavelengths than predicted by the simple models used here."639 Nevertheless. (hese transmission spectra still provide a unique resource [or exploring i the particle size distributions of these dusty rings.," Nevertheless, these transmission spectra still provide a unique resource for exploring in the particle size distributions of these dusty rings."640estimates of the expected velocity variability of a given star from the Doppler analysis pipeline. activity iudices. and comparison stars.,"estimates of the expected velocity variability of a given star from the Doppler analysis pipeline, activity indices, and comparison stars."641 For 77921. the measurement uncertainties use)) aud jitter usej) predict the level of velocity ius use)) in the absence of a planetary companion.," For 7924, the measurement uncertainties ) and jitter ) predict the level of velocity rms ) in the absence of a planetary companion."642 Thus. when interpreting this FAP. we should discount it bv the xobabilitv that the observed velocity variability is due ο LOINC.," Thus, when interpreting this FAP, we should discount it by the probability that the observed velocity variability is due to noise."643 Ou balance. we reject the null hypothesis and with confidence attribute the observed velocity variation to a alanet orbiting 77921 with a period of dd. Motivated by the success of our single-Ieplerian fit and the higher than expected velocity rms to that fit 11s0)). we performed a search for a double-I&eplerau fit.," On balance, we reject the null hypothesis and with confidence attribute the observed velocity variation to a planet orbiting 7924 with a period of d. Motivated by the success of our single-Keplerian fit and the higher than expected velocity rms to that fit ), we performed a search for a double-Keplerian fit."644 The Lonmb-Scargle periodogram of the residuals to the sinele-planct fit reffe:pergraun jshouwsseceral short periodnarrowpeakswithpsi) modest power(atl23.16.7.3 periodpeakwithbroadpower(atllr.5dd).," The Lomb-Scargle periodogram of the residuals to the single-planet fit \\ref{fig:pergram_1p_resid}) ) shows several short-period narrow peaks with modest power (at 12.3, 16.7, d) and an intermediate-period peak with broad power (at d)."645Eachofthesepeaksicgbeseewddnoniss bhdearomaiea, Each of these peaks represent a possible second planet in the system.646 We seeded. the Ieplerian search with several trial two-planet solutions., We seeded the Keplerian search with several trial two-planet solutions.647 The first plauet was always seeded with the orbital parameters for the best-fit sinele-planct model iu refseciorbital., The first planet was always seeded with the orbital parameters for the best-fit single-planet model in \\ref{sec:orbital}.648 The secoucd planet was seceded with the largest peaks in the periodoeram of single-planet residuals reffiie:pereramyposi) jandawidecarietyofothertrialperiods, The second planet was seeded with the largest peaks in the periodogram of single-planet residuals \\ref{fig:pergram_1p_resid}) ) and a wide variety of other trial periods.649Al lettlle patalic μμdde d'édieet gati it hoyMél process(2)., All of the parameters were allowed to vary during the fitting process \citep{Wright08}.650 fitschangedthcorbitalparametcrs for 7792 lbs fia, None of these two-planet fits changed the orbital parameters for 7924b substantially.651 planetsolutionshadbest fitsecond periodsofl2.3.3 {δι 16.7. andll.5 dd. andbestmagnitudes fiteccentriciticsof0.72. 0.76.0.LL 83. respect," The top four two-planet solutions had best-fit second-planet periods of 12.3, 34.8, 16.7, and d, and best-fit eccentricities of 0.72, 0.76, 0.44, 0.83, respectively."652 Phehigh το αμ feat att eopleelurtonoiscana , The high eccentricities raise our suspicion that the signals may be due to noise and increase the need for a robust tests of the null hypothesis.653We calculated FAPs for cach of these two-plauctcelu models. building ou the methods described iu refseciorbital..," We calculated FAPs for each of these two-planet models, building on the methods described in \\ref{sec:orbital}."654 Svuthetic data sets were constructed: by drawing with replacement from the residuals to the best-fit sinele-planct solution and adding the coherent. best-fit solution. back to the scrambled residuals.," Synthetic data sets were constructed by drawing with replacement from the residuals to the best-fit single-planet solution and adding the coherent, best-fit solution back to the scrambled residuals."655 We performed ai thorough search for the best-fit two-↕↘⊽↸∖↻↕↸∖∏⋜⋯⋯∪≼∐∖↕⋖⋜↧↴∖↴⋜∏⋝∪↖↽↸∖⋝∪∐↸∖⋜↧↸⊳∐↴∖↴⋅↖↽∐↑↕∐∖, We performed a thorough search for the best-fit two-Keplerian model (as above) on each synthetic data set.656↑↕↸⊳≺↧⋜↧↑⋜↕↴∖↴↸∖↑∙ ⋮⋜↧↕↴∖↴↸∖⋜↧↕⋜∐⋅↕⊔↴∖↴↖↖↽↸∖↥⋅↸∖⊓⋅↕∩⊾∩⊾↸∖↥⋅↸∖≼↧↖↖⇁∐↸∖∐⋜↧↴∖↴⋅↖↽∐↑∐↸∖↑↕↸⊳≼↧⋜↧↑⋜↧↴∖↴↸∖↑ 2tOD," False alarms were triggered when a synthetic data set had a lower value of than the original, unscrambled data."657 had a lower value of thau the original.," All of the two-planet models had an $>$, even when we restricted the scrambled trials to low-eccentricity solutions."658 unseraibled data.," Thus, we do not consider any of these signals to be viable planet candidates at the present time."659 Al of the two-plauct inodels had an 7.220%..cvenehbeueeresteictedtheseraimbledteialstolow," Nevertheless, we will continue to hunt for additional planets orbiting 7924 by making additional observations and by refining our Doppler analysis algorithms to reanalyze our extant Doppler spectra more precisely."660 We acquired photometric observations of 77921 with the T12 0.5nuu automated plotometric telescope (APT) at Fairborn observatory ju southern Arizoua., We acquired photometric observations of 7924 with the T12 m automated photometric telescope (APT) at Fairborn observatory in southern Arizona.661 Our brightness. measurements were made between 2006 December aud 2008 October and cover the last part of the 200607. the complete 200705. and the first part of the 200809 observing seasons.," Our brightness measurements were made between 2006 December and 2008 October and cover the last part of the 2006–07, the complete 2007–08, and the first part of the 2008–09 observing seasons."662 The T12 APT and its precision plotometer are very similar to the Ts APT described in ?.., The T12 APT and its precision photometer are very similar to the T8 APT described in \citet{Henry1999}.663 The precision photometer uses two telpcrature-stabilized EMT 912LOB photomultiplicr tubes to measure photon couut rates sinultaucouslv through Strouuneren b aud y filters., The precision photometer uses two temperature-stabilized EMI 9124QB photomultiplier tubes to measure photon count rates simultaneously through Strömmgren $b$ and $y$ filters.664" The telescope was progranuned to gather observations of 77921 with respect to three nearby comparison stars du the following sequence: DARK. A. D. C. D. A. SKY4. D. SkYp. C. SkKYc. D. SkKYp. A. D. C. D. The coluparison stars A. B. aud C are 11295 (V.=6.39. BPale F3V). HID110971(V —691 DB.V= 19. F5), and IID110259 (V=6.59.B V= 1.0L C5). respectively. while star Dis IID77921(V =7.17.DB.V S83. INO)."," The telescope was programmed to gather observations of 7924 with respect to three nearby comparison stars in the following sequence: DARK, A, B, C, D, A, $_{\rm A}$ , B, $_{\rm B}$, C, $_{\rm C}$, D, $_{\rm D}$, A, B, C, D. The comparison stars A, B, and C are 4295 $V=6.39$, $B-V=0.42$, F3 V), 10971 $V=6.94$, $B-V=0.49$ , F5), and 10259 $V=6.59, B-V=1.04$ , G5), respectively, while star D is 7924 $V=7.17, 665B-V=0.83$ , K0)."666 Each complete sequence. referred to as a group observation. was reduced to form 3 cut ucasures of cach of the 6 differcutial magnitudesindepen A. DBD ο. A.C D. and A. The differential nagnitudes were corrected for differential extinction with wieltly extinction cocfiicents and. trausformedt o the standard Strónuugren system with vearly nean transformation cocficicuts.," Each complete sequence, referred to as a group observation, was reduced to form 3 independent measures of each of the 6 differential magnitudes $-$ A, $-$ B, $-$ C, $-$ A, $-$ B, and $-$ A. The differential magnitudes were corrected for differential extinction with nightly extinction coefficents and transformed to the standard Strömmgren system with yearly mean transformation coefficients."667 The three indepeudeut neasures of cach differeutial magnitude were combined. eiving one mean data point per complete sequence for Laddandaclintdér thelétecliffercutial magnitudes.," The three independent measures of each differential magnitude were combined, giving one mean data point per complete sequence for each of the 6 differential magnitudes."668 To filter any meburfirotostetric conditious. an entire eroup observation was discarded if the standard deviation of auy of the six mean differential magnitudes exceeded 0.01 mae.," To filter any observations taken under non-photometric conditions, an entire group observation was discarded if the standard deviation of any of the six mean differential magnitudes exceeded 0.01 mag."669 Finally. we combined the Stronuneren b aud y differential magnitudes iuto a single (b|y)/2 passband to improve the precision.," Finally, we combined the Strömmgren $b$ and $y$ differential magnitudes into a single $(b+y)/2$ passband to improve the precision."670 Our complete data set consists of 192. 135. aud 115 good eroup observations from the 200607. 2007Os. and 200809 observing seasons. respectively. for πο ," Our complete data set consists of 192, 435, and 145 good group observations from the 2006–07, 2007--08, and 2008–09 observing seasons, respectively, for a total 772."671ubstautiahhyclcado three comparison stars. we averaged the three A. D. aud C differential of77921 withincach eroup iuto asingle πμοπα al brwbuidss 77921 aud the mean of the three comparison stars: D(A|B ο.," To minimize the effect of any low-level intrinsic variation in the three comparison stars, we averaged the three $-$ A, $-$ B, and $-$ C differential magnitudes of7924 withineach group into asingle value, representing the difference in brightness between 7924 and the mean of the three comparison stars: $D-(A+B+C)/3$ ."672 The standard deviations of these eusenible differential maguitudes for cach of the three observing seasons are 1.77. 1.85. and l.1?11nuuag.," The standard deviations of these ensemble differential magnitudes for each of the three observing seasons are 1.77, 1.85, and mmag,"673vary over a timescale of 3 vears: the Fe: vi]]A3759 bv in intensity over 3 years. the Fe A5159 line by τος... and HenA4686 by65'4.,"vary over a timescale of 3 years: the [Fe $\lambda$ 3759 by in intensity over 3 years, the [Fe $\lambda$ 5159 line by , and $\lambda$ 4686 by."674. In contrast. the same lines in QI1131116 did not vary significant over à similar 3 wear timescale.," In contrast, the same lines in Q1131+16 did not vary significant over a similar 3 year timescale."675 Moreover. the Balmer lines in QI131116 lack the variable. multi-peakec profiles observed in SDSS.J095209.56|214313.5. (Ixomossactal.2008. Ixomossaetal. 2009)).," Moreover, the Balmer lines in Q1131+16 lack the variable, multi-peaked profiles observed in SDSSJ095209.56+214313.3 \citealt{komossa1b} \citealt{komossa2}) )."676 Rather than an illumination event. an explanation for the unusual strength. of the ΕΠΗ may by. provided: by the inner torus wall being viewed atς a specific orientation. at which the inner wall on the far side of the torus can be seen by the observer. but the quasar itself. remains hidden. because the torus is being viewed at a relatively large angle to its axis (see f 4.3).," Rather than an illumination event, an explanation for the unusual strength of the FHILs may by provided by the inner torus wall being viewed at a specific orientation, at which the inner wall on the far side of the torus can be seen by the observer, but the quasar itself remains hidden, because the torus is being viewed at a relatively large angle to its axis (see $\oint$ 4.3)."677 In this case. the visible area of the torus wall is high enough that the ΕΤΗ) are strong. but the direct. quasar continuum emission does not dilute the torus line emission (see Figure 22).," In this case, the visible area of the torus wall is high enough that the FHILs are strong, but the direct quasar continuum emission does not dilute the torus line emission (see Figure 22)."678 Therefore the ΕΤΗ] are observed with Large equivalent widths., Therefore the FHILs are observed with large equivalent widths.679 Viewed at smaller angles to the axis of the torus. the luminous οπαρα. nuclear emission would become directly visible. substantially reducing the equivalent widths of the FLILLSs.," Viewed at smaller angles to the axis of the torus, the luminous quasar nuclear emission would become directly visible, substantially reducing the equivalent widths of the FHILs."680 In this contest we note that. despite the unusually Large equivalent widths of its ΕΕ. the measured ratios of the FLILs to lower ionization lines (e.g. Fe ΙΟ iu]]. Fe vus uu]. ο O i] and. Fe x]]/O i] for Q1131|16 fall well within the ranges measured for typical broad. line AGN (Sv1. ον1.5. NLSv1) in the sample compiled by Nagaoetal. (2000).," In this context we note that, despite the unusually large equivalent widths of its FHILs, the measured ratios of the FHILs to lower ionization lines (e.g. [Fe ]/[O ], [Fe ]/[S ], [Fe ]/[O ] and [Fe ]/[O ]) for Q1131+16 fall well within the ranges measured for typical broad line AGN (Sy1, Sy1.5, NLSy1) in the sample compiled by \citet{nagao2}."681.. Pherefore. viewed from a direction closer to the torus axis. we would. expect QI131|16 to display a normal quasar/Sevfert 1. spectrum. but. the multitude. of fainter FLL would be difficult to detect. because of the dilution by the bright. quasar continuum. ancl broad. line emission.," Therefore, viewed from a direction closer to the torus axis, we would expect Q1131+16 to display a normal quasar/Seyfert 1 spectrum, but the multitude of fainter FHIL would be difficult to detect because of the dilution by the bright quasar continuum and broad line emission."682 However. it is more challenging to explain why the ELILL are significantly stronger relative to the low ionization lines in QI1131116 than in typical Tvpe 2 AGN (Nagaoctal. 2000).," However, it is more challenging to explain why the FHIL are significantly stronger relative to the low ionization lines in Q1131+16 than in typical Type 2 AGN \citep{nagao2}."683. Given the geometry in Figure 22. we would expect to observe significant ELLEL emission over a range of viewing angles. even if the quasar nucleus itself is οποιος: rather than a sharp cut-oll. for this simple geometry. the strength of the FLL would gradually cüminish as the orientation of the axis of the torus to the line of sight increased (ie. the torus became more edge-on).," Given the geometry in Figure 22, we would expect to observe significant FHIL emission over a range of viewing angles, even if the quasar nucleus itself is obscured; rather than a sharp cut-off, for this simple geometry, the strength of the FHIL would gradually diminish as the orientation of the axis of the torus to the line of sight increased (i.e. the torus became more edge-on)."684 Therefore. although the particular ecometry shown in Figure 22 would maximise the FIL emission. it is unlikely that a specific geometry. alone can explain the rarity of objects like QLIS1|16.," Therefore, although the particular geometry shown in Figure 22 would maximise the FHIL emission, it is unlikely that a specific geometry alone can explain the rarity of objects like Q1131+16."685 We now cliseuss three further factors that might lead to the FULL appearing relatively strong compared. with the &eneral. population of Type 2 ACN., We now discuss three further factors that might lead to the FHIL appearing relatively strong compared with the general population of Type 2 AGN.686 Overall. while a specific viewing angle is likely to provide an important part of the explanation for the unusual spectrum of Q1131]116. one or more of the other factors discussed above may also be significant.," Overall, while a specific viewing angle is likely to provide an important part of the explanation for the unusual spectrum of Q1131+16, one or more of the other factors discussed above may also be significant."687 As mentioned. previously. there are a few similar objects in which the ΕΠΗ are unusually strong (LLL Zw τα. Tololo 0109-383 and ESO 138 GI).," As mentioned previously, there are a few similar objects in which the FHILs are unusually strong (III Zw 77, Tololo 0109-383 and ESO 138 G1)."688 Phe kev properties of these objects. as derived. from their ΕΠΗ5. are. presented. in ‘Table 4.," The key properties of these objects, as derived from their FHILs, are presented in Table 4."689 In particular. we note that the densities derived from the FLULLs are consistently. high.," In particular, we note that the densities derived from the FHILs are consistently high."690 Also all the objects show a particularly low ο H1](5007/4363) line ratio when compared to “typical Sevfert galaxies. (although in. this respect QI131116 is the most extreme)., Also all the objects show a particularly low [O III](5007/4363) line ratio when compared to `typical' Seyfert galaxies (although in this respect Q1131+16 is the most extreme).691" Finally. all the objects show relatively nioclest ΕΠΗ, line widths (ENLEM). with the exception of ESO138 G1 whose line widths are larger."," Finally, all the objects show relatively modest FHIL line widths (FWHM), with the exception of ESO138 G1 whose line widths are larger."692laws (hereafter ScMIL disks). when submitted to (he Newtonian potentials. reveal to have secular centrifugal ecuilibrium. maintaining their densitv. profiles for eigavears of simulated time (Springel&White1999:Springel.DiMatteoIlernquist.2005).,"laws (hereafter SdMH disks), when submitted to the Newtonian potentials, reveal to have secular centrifugal equilibrium, maintaining their density profiles for gigayears of simulated time \citep{sw1999,sdmh2005}."693. Llowever. under (he DM's picture. we do not know how long the DM potential can maintain galactic disks in dynamical equilibrium. and in fact if it really can.," However, under the BM's picture, we do not know how long the BM potential can maintain galactic disks in dynamical equilibrium, and in fact if it really can."694 This is the main aim of the present paper., This is the main aim of the present paper.695 In previous works (Drandao&deAraujo2010a.b.c).. we developed an N-body method io study alternative theories of egravitv al egalactic seales.," In previous works \citep{ca2009a, ca2009b,ca2009c}, we developed an $N$ -body method to study alternative theories of gravity at galactic scales."696 In particular. we have studied an Yukawian exgravitational potential and have shown that this potential is viable only if the Yukawian parameter is such that the Yukawian potential is nearly Newtonian.," In particular, we have studied an Yukawian gravitational potential and have shown that this potential is viable only if the Yukawian parameter is such that the Yukawian potential is nearly Newtonian."697 In the present work. we also use N-bocly simulations to construct aid evolve spiral ealaxies now submitted to the potential andacceleration given by Equations (3)) aud (4)). respectively.," In the present work, we also use $N$ -body simulations to construct and evolve spiral galaxies now submitted to the potential andacceleration given by Equations \ref{moffat_equation2}) ) and \ref{moffat_aceleration2}) ), respectively."698 We adopted. in (hese simulations. the same set of fitted. parameters (Af)=96x10A. and ry=13.96 kpe) as by DM.," We adopted, in these simulations, the same set of fitted parameters $M_0 = 96 \times 10^{10} M_{\odot}$ and $r_0 = 13.96 $ kpc) as by BM."699 Our aim is to probe if the DM model is dvuamucally consistent. as it should be in order to be considered a realistic model.," Our aim is to probe if the BM model is dynamically consistent, as it should be in order to be considered a realistic model."700 In this way. we are able to test the DMs claim aud verily (he reliabilitv of (his kind of law of gravity. using the concepts aud techniques of the galactic dvnamies studies (Dinnev&Tremaine2003).," In this way, we are able to test the BM's claim and verify the reliability of this kind of law of gravity, using the concepts and techniques of the galactic dynamics studies \citep{bt2008}."701. In Section 2.. we present the N-body code used in (his work and the numerical techniques used to model spiral galaxies will the DM potential: in Section 3.. (he results of (he numerical simulation ave present: and finally. in Section 4.. we discuss our main findings and conclusions.," In Section \ref{numtech}, we present the $N$ -body code used in this work and the numerical techniques used to model spiral galaxies with the BM potential; in Section \ref{simulresul}, the results of the numerical simulation are present; and finally, in Section \ref{conclusao}, , we discuss our main findings and conclusions."702the resulting excess spectrum from the inner region. binned to a minimum of 20 excess counts per bin. together with an absorbed power-law model fit.,"the resulting excess spectrum from the inner region, binned to a minimum of 20 excess counts per bin, together with an absorbed power-law model fit."703 The spectral parameters are collected in Table 1. region)., The spectral parameters are collected in Table \ref{tab-specparams} region).704 Fitting a thermal plasma (MEKAL. y; = 1.4(50)) or a thermal bremsstrahlung model (BREMSS. y; = 1.4(50)) to the spectrum gave temperatures kT>17 keV and kT>25 keV. respectively.," Fitting a thermal plasma (MEKAL, $\chi^{\scriptscriptstyle{2}}_{\scriptscriptstyle{\nu}}$ = 1.4(50)) or a thermal bremsstrahlung model (BREMSS, $\chi^{\scriptscriptstyle{2}}_{\scriptscriptstyle{\nu}}$ = 1.4(50)) to the spectrum gave temperatures $>$ 17 keV and $>$ 25 keV, respectively."705 The most prominent feature of the Galactic thermal emission observed from the outer region is an emission line. centered Οἱ ~1.8 keV. Introducing a Gaussian line at that energy or an additional thermal component to the excess spectrum from the inner annulus does not significantly improve the fit (y; 1.2).," The most prominent feature of the Galactic thermal emission observed from the outer region is an emission line, centered on $\sim$ 1.8 keV. Introducing a Gaussian line at that energy or an additional thermal component to the excess spectrum from the inner annulus does not significantly improve the fit $\chi^{\scriptscriptstyle{2}}_{\scriptscriptstyle{\nu}}$ = 1.2)."706 Therefore. we conclude that the spectrum from the outer region describes the Galactic. diffuse background component sufficiently.," Therefore, we conclude that the spectrum from the outer region describes the Galactic diffuse background component sufficiently."707 The total unabsorbed diffuse excess flux in the 1-7 keV band measured from the innerregion above the Galactic background is Fx = (5.520.8)xs7!., The total unabsorbed diffuse excess flux in the 1–7 keV band measured from the innerregion above the Galactic background is $_{\rm X}$ = $\pm$.708. Assuming a cistance of 5.5 kpe the intrinsic luminosity is Lx = (2.0x0.3.., Assuming a distance of 5.5 kpc the intrinsic luminosity is $_{\rm X}$ = $\pm$.709 We found no indication of a variation in the spectral index with increasing radius. when subdividing the inner region into two or more sub-regions.," We found no indication of a variation in the spectral index with increasing radius, when subdividing the inner region into two or more sub-regions."710 Furthermore. there is no evidence of a directional variation in. the index and flux.," Furthermore, there is no evidence of a directional variation in the index and flux."711 Table | regions) illustrates the result for directional dependence of spectra extracted from pie-shaped regions towards the north. east. south and west with respect to the cluster center (Fig. 1)).," Table \ref{tab-specparams} regions) illustrates the result for directional dependence of spectra extracted from pie-shaped regions towards the north, east, south and west with respect to the cluster center (Fig. \ref{fig-map}) )."712 Their inner and outer radii are 607 and 1207. respectively.," Their inner and outer radii are 60"" and 120"", respectively."713 The latter value was chosen such that the southern region is not truncated by the FoV. As background we again used the spectrum from the outer annulus., The latter value was chosen such that the southern region is not truncated by the FoV. As background we again used the spectrum from the outer annulus.714 A similar result was achieved when the regions were rotated by 45 degrees., A similar result was achieved when the regions were rotated by 45 degrees.715 The present results indicate GC-centered diffuse hard X-ray excess emission above Galactic background. which extends significantly beyond ry.," The present results indicate GC-centered diffuse hard X-ray excess emission above Galactic background, which extends significantly beyond $_{\rm h}$."716 In this section we briefly discuss standard thermal and non-thermal emission scenarios. leaving out more exotic possibilities. as described by. eg. ?..," In this section we briefly discuss standard thermal and non-thermal emission scenarios, leaving out more exotic possibilities, as described by, e.g., \citet{2005A&A...444L..33D}."717 Throughout this section we use a distance to oof 3.5 kpe., Throughout this section we use a distance to of 5.5 kpc.718 The larger distance estimate (8.7 kpe) would increase the energy requirements for the models by a factor of 2.5, The larger distance estimate (8.7 kpc) would increase the energy requirements for the models by a factor of 2.5.719 The lummosity of unresolved point sources inside ἡ has been estimated by 2. tos7'.. , The luminosity of unresolved point sources inside $_{\rm h}$ has been estimated by \citet{2006ApJ...651.1098H} to.720They furthermore determined the spatial surface distribution of X-ray sources in tto SC)eexpectfrofindtHe? with qe143., They furthermore determined the spatial surface distribution of X-ray sources in to $S(r)\propto(1 + (r/r_c)^2)^{(1-3q)/2}$ with q=1.43.721 From this distribution we to in the 1—3 aremin annulus only of the luminosity of unresolved point sources within rj., From this distribution we expect to find in the 1–3 arcmin annulus only of the luminosity of unresolved point sources within $_{\rm h}$.722 The expected is much lower than the measured emission (2.00.3) 7l. so we conclude that the contribution from unresolved point sources is negligible.," The expected is much lower than the measured emission $\pm$ , so we conclude that the contribution from unresolved point sources is negligible."723 , 724observed. periods may therefore be associated. with the bulk rotation of the white cwarl — which is certainly a short period clock available in the svstem.,observed periods may therefore be associated with the bulk rotation of the white dwarf – which is certainly a short period clock available in the system.725 This suggests that some mass accretion onto two magnetic poles on the primary is at times taking place (as well as onto the belt)., This suggests that some mass accretion onto two magnetic poles on the primary is at times taking place (as well as onto the belt).726 Unlike an LP. in a low field CY the angular momentum accreting at the equator will be distributed in latitude as well as racially: the outer lavers of the primary will therefore rotate dillerentially and it follows that the accretion regions may not always have exactly the same angular velocity.," Unlike an IP, in a low field CV the angular momentum accreting at the equator will be distributed in latitude as well as radially; the outer layers of the primary will therefore rotate differentially and it follows that the accretion regions may not always have exactly the same angular velocity."727 To investigate these IpDNOs. more thoroughly we have computed. LPs for subsets of the data ancl have also produced. phase/amplitude diagrams., To investigate these lpDNOs more thoroughly we have computed FTs for subsets of the data and have also produced phase/amplitude diagrams.728 In such diagrams the formal errors. derived from least squares fits. are sometimes very large.," In such diagrams the formal errors, derived from least squares fits, are sometimes very large."729 This is particularly the case where there are large amplitude QPOs or Uiekering which has not been filtered. out., This is particularly the case where there are large amplitude QPOs or flickering which has not been filtered out.730 We start by describing the behaviour of WW Ην in run 86133., We start by describing the behaviour of VW Hyi in run S6133.731 The FT in the vicinity of 100 s for the entire run. shown in Fig. 10..," The FT in the vicinity of 100 s for the entire run, shown in Fig. \ref{ft6133},"732 has three strong poriocdicities., has three strong periodicities.733 Using prewhitening and multiple sinusoidal non-linear least squares fitting we [ind these are best represented. by oscillations at S444 s and 85.70 s. with an amplitude of 1.7 mamag and 1.5 mmag. respectively. ancl 94.13 s with amplitude 1.6 mmag.," Using prewhitening and multiple sinusoidal non-linear least squares fitting we find these are best represented by oscillations at 84.44 s and 85.70 s, with an amplitude of 1.7 mmag and 1.5 mmag, respectively, and 94.13 s with amplitude 1.6 mmag."734 The doublet at S5 s is strong only in he first half of the run. which is not long enough to resolve it fully: in the second half of the run it appears as a single »ealkk at 85.3 s. “Phe separation of the doublet is therefore not well determined: although formally it corresponds to a oat period. of 5740 s. we Found. periods up to 6200 s while analvsing the data in various ways.," The doublet at 85 s is strong only in the first half of the run, which is not long enough to resolve it fully; in the second half of the run it appears as a single peak at 85.3 s. The separation of the doublet is therefore not well determined; although formally it corresponds to a beat period of 5740 s, we found periods up to 6200 s while analysing the data in various ways."735 “Phe orbital period. of VA DINI is 6173 s. so we suggest. in analogy. with IPs. that he S5 s doublet consists of a rotating beam with a period of δΕ s and an orbital sideband produced by reprocessing ov the secondary. or by the vertical thickening of the disc where the stream inipacts.," The orbital period of VW Hyi is 6173 s, so we suggest, in analogy with IPs, that the 85 s doublet consists of a rotating beam with a period of 84.4 s and an orbital sideband produced by reprocessing by the secondary or by the vertical thickening of the disc where the stream impacts."736 Η the 94 s modulation is also a reprocessed signal then. » analogy with the model for double DNOs. it implies a »ograde travelling wave with a period of 820 s. There is no »eak in the EE at this period. but shadowing of the disc and shadowing in our direction can be very dillerent for a system that is not of high inclination.," If the 94 s modulation is also a reprocessed signal then, by analogy with the model for double DNOs, it implies a prograde travelling wave with a period of 820 s. There is no peak in the FT at this period, but shadowing of the disc and shadowing in our direction can be very different for a system that is not of high inclination."737 Of the other runs with IpDNOs in them all but. one have narrow spikes in their Fs., Of the other runs with lpDNOs in them all but one have narrow spikes in their FTs.738 The 93 s peak in the ΕΙ) of run 50129 is too broad to be a single period. and its phase/amplitude diagram shows a sinusoidal modulation suggesting the presence. of two unresolved periodicities beating together — see Fig. Ll.," The 93 s peak in the FT of run S0129 is too broad to be a single period, and its phase/amplitude diagram shows a sinusoidal modulation suggesting the presence of two unresolved periodicities beating together – see Fig. \ref{omc0129}."739 Phe sine wave fitted. to the amplitude modulation has a period ~ 5100 s. but. is uncertain because only one evcle is observed.," The sine wave fitted to the amplitude modulation has a period $\sim$ 5100 s, but is uncertain because only one cycle is observed."740 This is of a similar time scale to the doublet separation in run 86133., This is of a similar time scale to the doublet separation in run S6133.741 The occasional. presence of unresolved doublets in the KP could account for. some of the spread in the IpDNO periods seen in VW. Lyi by Llaefner ct al., The occasional presence of unresolved doublets in the FT could account for some of the spread in the lpDNO periods seen in VW Hyi by Haefner et al.742 In Section 3.2 we gave the first. evidence that the ~ 23 s oscillations are DNOs., In Section 3.2 we gave the first evidence that the $\sim$ 23 s oscillations are DNOs.743 To study this further we have computed. amplitude/phase diagrams for all of the runs. some of which are shown in Fig. 12.," To study this further we have computed amplitude/phase diagrams for all of the runs, some of which are shown in Fig. \ref{omc6551}."744. The variations seen (for example the change of slope in the O€ diagram from negative to positive in run S655) at 1LJD 2452552.335 (bottom left panel of Fig. 12)).," The variations seen (for example the change of slope in the O–C diagram from negative to positive in run S6551 at HJD 2452552.335 (bottom left panel of Fig. \ref{omc6551}) ),"745 which is caused by a change in period from 22.66 s to 22.74 s) are characteristic of DNOs (see Paper D., which is caused by a change in period from 22.66 s to 22.74 s) are characteristic of DNOs (see Paper I).746 Phere is no svstematie variation of DNO phase with orbital phase of the kind seen in DO Ler (ODonoghue 1985)., There is no systematic variation of DNO phase with orbital phase of the kind seen in DQ Her (O'Donoghue 1985).747 Phase changes in the DNO (as illustrated in Fie. 12)), Phase changes in the DNO (as illustrated in Fig. \ref{omc6551}) )748We use all galaxies with 23 <Re 26 that are detected at significance 2106. as sources.,We use all galaxies with 23 $ < R_{C} <$ 26 that are detected at significance $ \gtrsim 10 \sigma$ as sources.749 On average there are 35 sources 7: a total of 254.010 sources meet our cits.," On average there are 35 sources $^{-2}$; a total of 254,010 sources meet our cuts."750 To derive the ellipticities of the sources. we use the routinesgetshapes.efi. ancl in the package (Ixaiser. Squires Droadhlurst 1999: KSB).," To derive the ellipticities of the sources, we use the routines, and in the package (Kaiser, Squires Broadhurst 1999: KSB)."751 We measure the object shears usingσας (IND)., We measure the object shears using (KSB).752 From (he shears. we construct lensing maps using a modifiel version of (he KSB algorithm (Lloekstva et al.," From the shears, we construct lensing maps using a modified version of the KSB algorithm (Hoekstra et al."753 1998)., 1998).754 First. we measure (he size ancl ellipticity of all the stars identified in the stacked images.," First, we measure the size and ellipticity of all the stars identified in the stacked images."755 We then use a second order polynomial function in pixel e and y coordinates to fit for (he PSF size and ellipticity in each subfield image., We then use a second order polynomial function in pixel $x$ and $y$ coordinates to fit for the PSF size and ellipticity in each subfield image.756" This procedure vields an estimate of the PSF size aud shape αἱ each galaxy. position. ancl. therefore. the stellar shear and smear polarizabilities. 2%, and P7,. respectively."," This procedure yields an estimate of the PSF size and shape at each galaxy position, and, therefore, the stellar shear and smear polarizabilities, $P_{sh}^*$ and $P_{sm}^*$, respectively."757 To caleulate the PSF shear and smear polarizabilitv appropriate for each galaxys position and size. we first caleulate the polarizability tensor for a set of Gaussian kernels wilh FWIIM ranging from {to ((in increments of 0.8)).," To calculate the PSF shear and smear polarizability appropriate for each galaxy's position and size, we first calculate the polarizability tensor for a set of Gaussian kernels with FWHM ranging from to (in increments of )."758" At each position. we calculate the appropriate polarizabilities Ly, and £2, by linear interpolation of the polarizability components for the kernels (hat bracket the galaxy size."," At each position, we calculate the appropriate polarizabilities $P_{sh}$ and $P_{sm}$ by linear interpolation of the polarizability components for the kernels that bracket the galaxy size."759 We can thus caleulate the induced polarizability of the galaxy as (Iloekstra el al., We can thus calculate the induced polarizability of the galaxy as (Hoekstra et al.760 1998: NSB) Because { can be affected. by uncertainty in the size determination of individual galaxies. our corrected shear measure for each galaxy is where © is the measured. ellipticity and orientation of the galaxy and (74) is the average polarizability of the 20 nearest neighbors to each galaxy. we consider.," 1998; KSB) Because $P_\gamma$ can be affected by uncertainty in the size determination of individual galaxies, our corrected shear measure for each galaxy is where $\vec{e}$ is the measured ellipticity and orientation of the galaxy and $\left<P_\gamma\right>$ is the average polarizability of the 20 nearest neighbors to each galaxy we consider."761 Once we have produced a & map from the shear distribution. we estimate a noise map bv taking (he variance over LOO realizations of the & map.," Once we have produced a $\kappa$ map from the shear distribution, we estimate a noise map by taking the variance over 100 realizations of the $\kappa$ map."762 In each realization. we randomize the orientation of the galaxies (leaving their positions fixed) and regenerate (he map.," In each realization, we randomize the orientation of the galaxies (leaving their positions fixed) and regenerate the map."763 We show the significance map (the & map divided by (the variance map) in Figure 8.., We show the significance map (the $\kappa$ map divided by the variance map) in Figure \ref{fig:GTOkappawopeak.ps}.764 We use Sextiractor to identilv peaks in (he &-S/N map., We use Sextractor to identify peaks in the $\kappa$ -S/N map.765" The pixels in the map are 0.3x0.3"" and we use a smoothing kernel with a 1.5 FWIIM."," The pixels in the map are $0.3^{\prime} \times 0.3^{\prime}$, and we use a smoothing kernel with a $1.5^{\prime}$ FWHM."766" A peak must have at least 3 connected pixels significant al the 20,s;x level or more to be selected: 6,sy is the variance in the & -S/N map.", A peak must have at least 3 connected pixels significant at the $\sigma_{\kappa-S/N}$ level or more to be selected; $\sigma_{\kappa -S/N}$ is the variance in the $\kappa$ -S/N map.767Here we demonstrate the ability of our algorithm to realistically simulate Galactic synchrotron and free-free emission both in total and polarized intensity.,Here we demonstrate the ability of our algorithm to realistically simulate Galactic synchrotron and free-free emission both in total and polarized intensity.768 We also quantitatively compare results of our simulations with the observations., We also quantitatively compare results of our simulations with the observations.769 The simulated data cubes are presented for four cases of the Galactic emission., The simulated data cubes are presented for four cases of the Galactic emission.770" First three models show ability of our algorithm to simulate different cases of Faraday rotation and depolarization, while the last model is tailored to be in agreement with ? observations of the Fan region."," First three models show ability of our algorithm to simulate different cases of Faraday rotation and depolarization, while the last model is tailored to be in agreement with \citet{bernardi09a} observations of the Fan region."771 Cartoons of these four models are given in Fig. 2.., Cartoons of these four models are given in Fig. \ref{fig:models}.772" The Galactic emission maps are obtained in the frequency range from 115to 180MHz, with 0.5MHz step."," The Galactic emission maps are obtained in the frequency range from $115$to $180~{\rm MHz}$, with $0.5~{\rm MHz}$ step."773" In the first three models the synchrotron emission originates from the same CR electron distribution with uniform B, component of the Galactic magnetic field.", In the first three models the synchrotron emission originates from the same CR electron distribution with uniform $B_{\perp}$ component of the Galactic magnetic field.774" Therefore, the brightness temperature maps of the total and polarized synchrotron emission are equivalent in all three models."," Therefore, the brightness temperature maps of the total and polarized synchrotron emission are equivalent in all three models."775" The same is valid for the free-free emission, i.e., Τε is normalized to the same value of EM."," The same is valid for the free-free emission, i.e., $n_e$ is normalized to the same value of $EM$."776 The resulting 10°x maps at 150MHz are shown in Fig. 3.," The resulting $10^\circ\times10^\circ$ maps at $150~{\rm777MHz}$ are shown in Fig. \ref{fig:synff}."778".a, b d, while the mean and rms of the maps are given in Table 1.."," .a, b d, while the $mean$ and $rms$ of the maps are given in Table \ref{tab:maps}."779 A map of the brightness temperature spectral index 3 of simulated total intensity synchrotron emission is shown in Fig. 3., A map of the brightness temperature spectral index $\beta$ of simulated total intensity synchrotron emission is shown in Fig. \ref{fig:synff}.780.c., .c.781 Simulated polarized emission maps of the first three Galactic synchrotron emission models are shown in Fig. 4., Simulated polarized emission maps of the first three Galactic synchrotron emission models are shown in Fig. \ref{fig:PImodels}.782".a, b c. Their mean and rms values together with the degree of polarization and depolarization are listed in Table 1.."," .a, b c. Their $mean$ and $rms$ values together with the degree of polarization and depolarization are listed in Table \ref{tab:maps}."783 A random line through synchrotron total and polarized intensity frequency data cubes of first three models are given in Fig. 6.., A random line through synchrotron total and polarized intensity frequency data cubes of first three models are given in Fig. \ref{fig:los}.784 Simulated total and polarized emission images of the Fan region (model D) are given at 150 MHz in Fig., Simulated total and polarized emission images of the Fan region (model D) are given at 150 MHz in Fig.785" 5 a b. For easier comparison with the observation an RM image, obtained by applying rotation measure synthesis technique (?) on the simulated data, is shown in Fig."," \ref{fig:Fan} a b. For easier comparison with the observation an RM image, obtained by applying rotation measure synthesis technique \citep{brentjens05} on the simulated data, is shown in Fig."786" 5 c. First test of the algorithm is to estimate the degree of intrinsic polarization, II, of the simulated map (Fig. 3."," \ref{fig:Fan} c. First test of the algorithm is to estimate the degree of intrinsic polarization, $\Pi$, of the simulated map (Fig. \ref{fig:synff}."787.b)., .b).788" By dividing the intrinsic polarized emission map and the total intensity map, we obtain II—0.69 (see Table 1))."," By dividing the intrinsic polarized emission map and the total intensity map, we obtain $\Pi=0.69$ (see Table \ref{tab:maps}) )."789 This value is in a good agreement with the expected theoretical value II5—2=9/13 (see Eq. 6))., This value is in a good agreement with the expected theoretical value $\Pi_{p=2}=9/13$ (see Eq. \ref{eq:intpol}) ).790 Further we explore morphological characteristics of the simulated polarized emission by comparing the image of intrinsically polarized emission (Fig. 3., Further we explore morphological characteristics of the simulated polarized emission by comparing the image of intrinsically polarized emission (Fig. \ref{fig:synff}.791.b) with the images of polarized emission of the first three models (see Fig. 4))., .b) with the images of polarized emission of the first three models (see Fig. \ref{fig:PImodels}) ).792 Note that degree of polarized and depolarized emission is given in Table 1.., Note that degree of polarized and depolarized emission is given in Table \ref{tab:maps}. .793 Model A assumes that there is no region in which the plasma (thermal electron cloud) is mixed with CR electrons., Model A assumes that there is no region in which the plasma (thermal electron cloud) is mixed with CR electrons.794" Therefore, polarization angle of synchrotron emission along the LOS are"," Therefore, polarization angle of synchrotron emission along the LOS are"795 (e.c.Fujita2001:Corteseetal.2006: Owenctal.(2005):Miller&(2003):Ferrarietal.(2005) found evideuce for triggeredao star formation imn cluster mergers based ou au euhauced fraction of star-Kktune radio salaxies aud a preferential distribution Q[o cndssion line galaxies between imereiug subchlps.," \citep{dressler1980,gomez2003} \citep{zabludoff1998,kodama2001}, \cite[e.g.][]{fujita2004,cortese2006,koyama2008,berrier2009}, \citet{owen2005,miller2003,ferrari2005} found evidence for triggered star formation in cluster mergers based on an enhanced fraction of star-forming radio galaxies and a preferential distribution of emission line galaxies between merging subclumps."796 Iu contrast. Pogoejutietal.(2MOL) found tha the post-starburst pomulation of dwarf ealaxies in the €‘olla cluster lie near the edees of two mereiue substrucures. sugecstine that he Inerecr queiches star formation. though these could be relics of starbursts induced dwine an earlier phase «Xf the merger (Alahajanetal.201n..," In contrast, \citet{poggianti2004} found that the post-starburst population of dwarf galaxies in the Coma cluster lie near the edges of two merging substructures, suggesting that the merger quenches star formation, though these could be relics of starbursts induced during an earlier phase of the merger \citep{mahajan2010}."797 There has aSO bec nuci effort place| in nuderstanudine the specific]physical mechaΠΡΙΝ responsible for raustornius ealaxy propertics such as norplology and star formaion., There has also been much effort placed in understanding the specific physical mechanisms responsible for transforming galaxy properties such as morphology and star formation.798 Although there is sole observatiolal aid/or theoretical support for various plivsical processes colsered to occur iji the cluster environnrent. such as ra pressure (Cauπι&ot 1972).. straugulatio1 (IiaaMulchaev.2tKs... and galaxy hareissnucnt (Mooreetal.1996). it Is still uukuown wlich of hese play a dominant role in transforming a stir-forwine ealaxy into a quiescent one.," Although there is some observational and/or theoretical support for various physical processes considered to occur in the cluster environment, such as ram pressure \citep{gunn1972}, strangulation \citep{kawata2008}, and galaxy harassment \citep{moore1996}, it is still unknown which of these play a dominant role in transforming a star-forming galaxy into a quiescent one."799 This is the sec‘ond in a series of papers to examine the star formatioli proyerties of 1E0657-56. also kuown as the Bullet Cluster.," This is the second in a series of papers to examine the star formation properties of 1E0657-56, also known as the Bullet Cluster."800 The Bullet Cluster is a galaxy cluster at 2.=0.207 τιvderegoing a major niereer event. with a collision betwoee1i fhe main cluster aud subcluster occurug close Τε| the plane of the sky with /cow. (Markevitchetal.200D..," The Bullet Cluster is a galaxy cluster at $z=0.297$ undergoing a major merger event, with a collision between the main cluster and subcluster occuring close to the plane of the sky with $i<8^{\circ}$ \citep{markevitch2004}."801 A welbdefiued bow shock frout has been coufirmed by Markevitchetal. (2002).. ancl is propagating through the X-ray easin the subcluster," A well-defined bow shock front has been confirmed by \citet{markevitch2002}, , and is propagating through the X-ray gasin the subcluster"802Soft excesses are the steep. upturn in the X-ray spectra of active galactic nuclei (AGN) below —2kkeV.— (e.g. Alushotzkyv.Done&Pounds 1993)).,Soft excesses are the steep upturn in the X-ray spectra of active galactic nuclei (AGN) below $\sim$ keV (e.g. \ncite{MUS93}) ).803 They were first detected in Sevlert ealaxies. (Arnaucletal.1985)... where they are quite common (Lurner&Pounds1989: Alushotzky ct al.," They were first detected in Seyfert galaxies \cite{ARN85}, where they are quite common \ncite{TAP89}; Mushotzky et al."804 1993)., 1993).805 They can be modelled by a steep power law (a 72). or by à cool thermal component (P< 150eeV). and are thought to be the high energy tail of the accretion dise radiation (e.g. Alushotzky et al.," They can be modelled by a steep power law $\alpha\!>$ 2), or by a cool thermal component $<\!150$ eV), and are thought to be the high energy tail of the accretion disc radiation (e.g. Mushotzky et al."806 1993)., 1993).807 Soft excesses are reasonably common in quasars (Urryal.1989:Masnouet1992:Saxton 1993)... ancl there appears to be a trend for the more distant objects to have Hatter spectra (Scharteletal.1992).. as expected if a soft component were being redshifted: below the soft. X-ray energy range.," Soft excesses are reasonably common in quasars \cite{URR89,MAS92,SAX93}, and there appears to be a trend for the more distant objects to have flatter spectra \cite{SCH92}, as expected if a soft component were being redshifted below the soft X-ray energy range."808 Initial detections were in raclio-quict objects (Comastriοἱal.1992).. but more recent work (Bühleretal.1995:Prieto1906:Scharteletal.1996). suggests that a soft excess may also be a common feature in racdio-oud objects.," Initial detections were in radio-quiet objects \cite{COM92}, but more recent work \cite{BUH95,PRI96,SCH96} suggests that a soft excess may also be a common feature in radio-loud objects."809 One of these with a consistently detected soft. excess is 3€ 273 Clurneretal.1990:Staubert1992:Leach.M'LHlardy&Papadakis 1995).," One of these with a consistently detected soft excess is 3C 273 \cite{TUR90,STA92B,LEA95}."810. It is also one object for which the spectral parameters of the soft excess itself are reasonably well known (Leach et al., It is also one object for which the spectral parameters of the soft excess itself are reasonably well known (Leach et al.811 1995)., 1995).812 The main instrument on (Bradt.Rotbschilel&Swank1903) is the proportional counter array (PCA)., The main instrument on \cite{BRA93} is the proportional counter array (PCA).813 With five xenon filled counters (PCU's). it has an energy range of GÜOkkeV. an energy resolution of 18 per cent at GkkeV. a large (0.727) elective area. and a circular. Geld of view (ΑΔΗΝΟ 17).," With five xenon filled counters (PCU's), it has an energy range of keV, an energy resolution of 18 per cent at keV, a large $\sim$ $^2$ ) effective area, and a circular field of view (FWHM: $^{\rm o}$ )."814 Each. PCU has three separate lavers., Each PCU has three separate layers.815 Only data from laver 1 has been used. as for observations of Lain sources most of the counts in lavers 2 and 3 are cause by the background.," Only data from layer 1 has been used, as for observations of faint sources most of the counts in layers 2 and 3 are caused by the background."816 Problems with two of the PCUs mean they were switehed olf for some of the observations (note in Table 1))., Problems with two of the PCUs meant they were switched off for some of the observations (noted in Table \ref{tab-xteobs}) ).817 TheRN observations (see Table 1)) were made on a roughly daily basis between 1996 June 27 to July 14 and have an average exposure of about 650 seconds., The observations (see Table \ref{tab-xteobs}) ) were made on a roughly daily basis between 1996 June 27 to July 14 and have an average exposure of about 650 seconds.818 Data reduction was done usingANE specilic programs., Data reduction was done using specific programs.819" The Standard2 data were filtered. using internally appliec good times (c.g. times of SAA passage). as well as the criteria that elevation above the horizon of greater than 10"" ane olfsource angle of less than 0.17. and then reduced to ligh curves and spectra usingTROY."," The Standard2 data were filtered using internally applied good times (e.g. times of SAA passage), as well as the criteria that elevation above the horizon of greater than $^{\rm o}$ and offsource angle of less than $^{\rm820o}$, and then reduced to light curves and spectra using."821 The PCA cannot measure the background. during an observation and so it must. be estimated., The PCA cannot measure the background during an observation and so it must be estimated.822 Phe accuracy of this estimation is very important for faint sources such as 3€ 270. (, The accuracy of this estimation is very important for faint sources such as 3C 279. (823vl4te with the οὐ model) was used to produce models of the background due to particles and cosmic X-rays. and spectra and light curves were then extracted [rom these models in the same way as for the source observations.,"v1.4g with the q6 model) was used to produce models of the background due to particles and cosmic X-rays, and spectra and light curves were then extracted from these models in the same way as for the source observations."824 In order to check the accuracy of the PCA background. model. ancl to. provide a systematic error for the PCA datapoints. we have analysed 11. slew observations with individual exposures of greater than 400 seconds. (average exposure is 671 seconds).," In order to check the accuracy of the PCA background model, and to provide a systematic error for the PCA datapoints, we have analysed 11 slew observations with individual exposures of greater than 400 seconds (average exposure is 671 seconds)."825 The mean residual count rate after subtraction of the background in the 2 kkeV. band is -0.25 cts/sec (ie. the mocoel is an over estimate of the data on average). with an intrinsic scatter of 0.44 cts/sec.," The mean residual count rate after subtraction of the background in the $\sim$ keV band is -0.25 cts/sec (i.e. the model is an over estimate of the data on average), with an intrinsic scatter of 0.44 cts/sec."826"density contributions for all these mono-abundance bins in aand ((with [a/Fe]widths 0.1 in aand 0.05 in [a/Fe])) versus the scale height of those sub-population fromB11,, color-coded by their eenhancement.","density contributions for all these mono-abundance bins in and (with widths 0.1 in and 0.05 in ) versus the scale height of those sub-population from, color-coded by their enhancement."827" We then sum the surface-mass[a/Fe] contributions (hz|[Fe/H],[a/Fe])_ of sub-populations ([Fe/H],,[a/Fe]))Up, into bins in scale height."," We then sum the surface-mass contributions $\Sigma_{R_0}(h_z|\feh,\afe)$ of sub-populations ) into bins in scale height."828" This results in the thick black histogram, which represents Xp,(hz), or more simply p(h;), the surface-mass weighted distribution of vertical scale heights in the Solar neighborhood (which we will refer to in the remainder simply as the 'scale-height distribution"")."," This results in the thick black histogram, which represents $\Sigma_{R_0}(h_z)$, or more simply $p(h_z)$, the surface-mass weighted distribution of vertical scale heights in the Solar neighborhood (which we will refer to in the remainder simply as the `scale-height distribution')."829" That is, for any random stellar-mass element this function gives the probability density for the scale height of the structural component to which it belongs."," That is, for any random stellar-mass element this function gives the probability density for the scale height of the structural component to which it belongs."830 This is the function we set out to construct in order to examine whether it makes sense to think of distinct thin and thick disk components in the Milky Way., This is the function we set out to construct in order to examine whether it makes sense to think of distinct thin and thick disk components in the Milky Way.831" Remarkably, we find that the scale-height distribution simply decreases quite smoothly towards larger scale heights, with an approximately exponential relation between surface- density and scale height H(Ro)οexp(—h,)."," Remarkably, we find that the scale-height distribution simply decreases quite smoothly towards larger scale heights, with an approximately exponential relation between surface-mass density and scale height $\Sigma(R_0) \propto \exp(-h_z)$."832" 'The scale height distribution does not show any gaps, excesses, or hints of bi-modality, beyond this simple relation."," The scale height distribution does not show any gaps, excesses, or hints of bi-modality, beyond this simple relation."833 By combining all of the stellar surface-mass density estimates we can precisely measure the total visible stellar surface-mass density at the Solar radius., By combining all of the stellar surface-mass density estimates we can precisely measure the total visible stellar surface-mass density at the Solar radius.834" We find Lp,*=3041Mo Ρο”.", We find $\Sigma^{^*}_{R_0} = 30 \pm 1\ M_\odot$ $^{-2}$ .835" This is similar to the estimate of Flynnetal.(2006),, who report Yn=29Mo pc-?."," This is similar to the estimate of \citet{Flynn06a}, who report $\Sigma^{^*}_{R_0} = 29\ M_\odot$ $^{-2}$."836 This estimate depends slightly on the assummed IMF., This estimate depends slightly on the assummed IMF.837" Using the exponential IMF (IMF3) of Chabrier(2001) gives Zp,=29.5Mo pc?; the IMF from Chabrier(2003) gives Yn.=29Μα pc?; and a Kroupa(2003) IMF gives Up,=32Μο Ρο”."," Using the exponential IMF (IMF3) of \citet{Chabrier01a} gives $\Sigma^{^*}_{R_0} = 29.5\ M_\odot$ $^{-2}$; the IMF from \citet{Chabrier03a} gives $\Sigma^{^*}_{R_0} =83829\ M_\odot$ $^{-2}$; and a \citet{Kroupa03a} IMF gives $\Sigma^{^*}_{R_0} = 32\ M_\odot$ $^{-2}$ ."839 1 shows that properly correcting for the spectroscopic sampling of the underlying stellar populations is crucial in assessing the elemental-abundance distribution at Ro of spectroscopically selected samples of stars., \ref{fig:mass_afe_feh} shows that properly correcting for the spectroscopic sampling of the underlying stellar sub-populations is crucial in assessing the elemental-abundance distribution at $R_0$ of spectroscopically selected samples of stars.840 The abundance distribution without this correction is heavily influenced by the survey-specific spatial and mass sampling of the underlying stellar population—both of which act to make the metal-poor and a-enhanced sub-populations more prominent in the high-latitude and color-selected ssample—which leads to a spuriously enhanced modality in elemental-abundance space., The abundance distribution without this correction is heavily influenced by the survey-specific spatial and mass sampling of the underlying stellar population—both of which act to make the metal-poor and $\alpha$ -enhanced sub-populations more prominent in the high-latitude and color-selected sample—which leads to a spuriously enhanced bi-modality in elemental-abundance space.841 The mass-weighted metallicity distribution in 1 has no bi-modality., The mass-weighted metallicity distribution in \ref{fig:mass_afe_feh} has no bi-modality.842" The mass-weighted [a/Fe]--distribution in the same Figure has only a hint of a bi-modality, as is expected in standard smooth star formation and enrichment scenarios (Schónrich&Binney2009a),, and reflects enrichment physics and not galaxy evolution."," The mass-weighted -distribution in the same Figure has only a hint of a bi-modality, as is expected in standard smooth star formation and enrichment scenarios \citep{Schoenrich08a}, and reflects enrichment physics and not galaxy evolution."843 Improper weighting of the distribution of structural parameters can again lead to spurious bi-modal signatures., Improper weighting of the distribution of structural parameters can again lead to spurious bi-modal signatures.844" 5 in sshows the location of chemically-defined in the space of the structural parameters (radial scale length, vertical scale height)."," 5 in shows the location of chemically-defined in the space of the structural parameters (radial scale length, vertical scale height)."845" This scatterplot, based on equal-area bins in ([Fe/H],,|o//Fe])), gives undue prominence to thelow-[Fe/H],, ((with respect to solar abundances) bins, as many of them only contribute a negligible amount to the total stellar mass."," This scatterplot, based on equal-area bins in ), gives undue prominence to the, (with respect to solar abundances) bins, as many of them only contribute a negligible amount to the total stellar mass."846" Here, 2 shows that the proper mass-weighting of the ([Fe/H],,[a/Fe])) sub-populations gives a vertical-scale-height distribution that is smooth and monotonically declining between thinner disk component scale heights of 200 pc and thicker components' scale heights of αγ1200 pc, with no gaps or excesses beyond the smooth, approximately exponential distribution."," Here, \ref{fig:mass_hz_afe} shows that the proper mass-weighting of the ) sub-populations gives a vertical-scale-height distribution that is smooth and monotonically declining between thinner disk component scale heights of 200 pc and thicker components' scale heights of $\approx 1200$ pc, with no gaps or excesses beyond the smooth, approximately exponential distribution."847" ffound that each elemental-abundance bin was preferentially fit by a single exponential rather than by two disk components, such that the smooth scale-height distribution in 2 is merely the result of the smoothing out of an intrinsically bi-modal distribution by elemental-abundance errors (which are small for the ssample, see ??)) or overlapping abundance distributions of distinct “thin” and “thick” disks; if either of these were the case B11 should have detected two components in the abundance bins with single-exponential scale heights in the range of approximately 400 to 600 pc."," found that each elemental-abundance bin was preferentially fit by a single exponential rather than by two disk components, such that the smooth scale-height distribution in \ref{fig:mass_hz_afe} is merely the result of the smoothing out of an intrinsically bi-modal distribution by elemental-abundance errors (which are small for the sample, see \ref{sec:main}) ) or overlapping abundance distributions of distinct “thin” and “thick” disks; if either of these were the case B11 should have detected two components in the abundance bins with single-exponential scale heights in the range of approximately 400 to 600 pc."848" The large uncertainties on the radial scale lengths of the chemically-defined mono-abundance sub-populations in ccomplicate a similar assessment of the radial structure of the disk, but this has no bearing on the analysis of the vertical structure."," The large uncertainties on the radial scale lengths of the chemically-defined mono-abundance sub-populations in complicate a similar assessment of the radial structure of the disk, but this has no bearing on the analysis of the vertical structure."849 Upcoming surveys such as (Eisensteinetal.2012) or (Freemanetal.2010) that willHERMES/GALAH sample stellar populations in the plane of the Milky Way will be able to study the radial structure in more detail., Upcoming surveys such as \citep{Eisenstein12a} or \citep{Freeman10a} that will sample stellar populations in the plane of the Milky Way will be able to study the radial structure in more detail.850" Thus, stars in the Solar neighborhood have a smoothly decreasing probability of belonging to structural components with increasing scale heights."," Thus, stars in the Solar neighborhood have a smoothly decreasing probability of belonging to structural components with increasing scale heights."851 This implies that the thicker disk component in the Milky Way is simply the tail of a continuous and monotonic scale-height distribution., This implies that the thicker disk component in the Milky Way is simply the tail of a continuous and monotonic scale-height distribution.852" This has been suggested before (e.g.,Norris1987;Schónrich&Binney 2009b), but never directly measured as we do here."," This has been suggested before \citep[\eg,][]{Norris87a,Schoenrich09a}, but never directly measured as we do here."853" As such, there is no distinct thick-disk component in our Galaxy."," As such, there is no distinct thick-disk component in our Galaxy."854" Together with the findings in tthat the thicker and older components of the Galactic disk have a shorter radial scale lengths than the thinner and younger components this qualitatively points toward a continuous internal mechanism such as radial migration or turbulent disk evolution being predominantly responsible for the thickening of the disk, rather than an external merger or heating event."," Together with the findings in that the thicker and older components of the Galactic disk have a shorter radial scale lengths than the thinner and younger components this qualitatively points toward a continuous internal mechanism such as radial migration or turbulent disk evolution being predominantly responsible for the thickening of the disk, rather than an external merger or heating event."855" However, a rigorous comparison with models, where thick stellar-disk components arise from one or a few distinct events triggered,e.g.,, by satellite infall, is needed to see whether the present data are indeed inconsistent with the results presented here."," However, a rigorous comparison with models, where thick stellar-disk components arise from one or a few distinct events triggered, by satellite infall, is needed to see whether the present data are indeed inconsistent with the results presented here."856 Formation age of stars may serve as a sensible marker in simulations to associate individual stars with a ‘parentsub-population’ whose scale height can be determined., Formation age of stars may serve as a sensible marker in simulations to associate individual stars with a `parentsub-population' whose scale height can be determined.857" Inthis context, it is worth noting that combined data from and will soon provide ages for a large number of stars through asteroseismology (e.g.,Gillilandetal. 2010),,"," Inthis context, it is worth noting that combined data from and will soon provide ages for a large number of stars through asteroseismology \citep[\eg,][]{Gilliland10a}, ,"858δἱ xov.L(TT) δουν αμ..,"P v_x, _x and v_y t v_x."859".loo This leads to one of our tain Couclusious: the density αμα y-velocity perturbations will grow οH forH a>4 1. ke.H RiHxNZD«0. For4 small. RichardsonH munber. however (así isH expected for a Ixeplerian disk with mocest radial gracdieuts). a1—3141 and the asymptotic erowth is extremely slow: de,M"," This leads to one of our main conclusions: the density and $y$ -velocity perturbations will grow asymptotically for $\alpha > 1$ , i.e. ${\rm{Ri}} \propto N_x^2 < 0$ For small Richardson number, however (as is expected for a Keplerian disk with modest radial gradients), $\alpha \sim 1 - 2{\rm Ri}$ and the asymptotic growth is extremely slow: v_y."86080) Iu the stratified shearing sheet. the right-haud side of equation (22)) governing the evolution of the perturbed potential vorticity is uo longer zero.," In the stratified shearing sheet, the right-hand side of equation \ref{PVEVLIN}) ) governing the evolution of the perturbed potential vorticity is no longer zero."861 The form of this equation for the incompressive slwaves is Sea oS, The form of this equation for the incompressive shwaves is = ) =.862"NL) The asymptotic time dependeuce of the perturbed potential vorticity can be obtained by iutegratiug equation (??)):B82)"" for ?5l1 and |Ri|«1.", The asymptotic time dependence of the perturbed potential vorticity can be obtained by integrating equation \ref{DPVEV}) ): for $\tilde{\tau} \gg 1$ and $|{\rm Ri}| \ll 1$.863 As uoted in 82. an entropy gradient is not required to generate vorticity.," As noted in 2, an entropy gradient is not required to generate vorticity."864H For4 NzD=0. a=41 aud the perturbed potentialH vorticityH grows linearlyH withH time.," For $N_x^2 = 0$, $\alpha = 1$ and the perturbed potential vorticity grows linearly with time."865"H The unstratified shearing sheet is recovered in thelimit of zero stratification (1/Lp— 0). sincein this limit equation (2?)) reduces to €= constant,"," The unstratified shearing sheet is recovered in thelimit of zero stratification $1/L_P \rightarrow 8660$ ), sincein this limit equation \ref{DPVEV}) ) reduces to $\xi = constant$ ."8672006).,.868". Each pointing consisted of three 800s exposures giving a total integration time of 2400s for the extended emission, and 4800s in the overlapping region centered on the disk."," Each pointing consisted of three 800s exposures giving a total integration time of 2400s for the extended emission, and 4800s in the overlapping region centered on the disk."869 The seeing was ~1.2” during the course of our exposures., The seeing was $\sim1.2$ during the course of our exposures.870 Sky subtraction was carried out using a sky spectrum drawn from relatively clean sky regions at the ends of each data cube., Sky subtraction was carried out using a sky spectrum drawn from relatively clean sky regions at the ends of each data cube.871" The data were reduced and flux calibrated using the WiFeS pipeline, briefly described in Dopitaetal.(2010),, which uses IRAF routines adapted primarily from the Gemini NIFS data reduction package."," The data were reduced and flux calibrated using the WiFeS pipeline, briefly described in \citet{Dopita10}, which uses IRAF routines adapted primarily from the Gemini NIFS data reduction package."872 The bias subtraction is somewhat complicated by the use of quad-readout to decrease chip read-out time as well as due to a slight slope and instability in the bias across each region of the chip., The bias subtraction is somewhat complicated by the use of quad-readout to decrease chip read-out time as well as due to a slight slope and instability in the bias across each region of the chip.873 Bias frames are taken immediately before and after each set of observations and a two-dimensional fit of the surface is subtracted from the temporally nearest object data in order to avoid additional noise to the data., Bias frames are taken immediately before and after each set of observations and a two-dimensional fit of the surface is subtracted from the temporally nearest object data in order to avoid adding additional noise to the data.874 Any resulting residual is addingaccounted for with a fit to unexposed regions of the detector., Any resulting residual is accounted for with a fit to unexposed regions of the detector.875 Quartz lamp flats are used to account for the response curve of the and twilight flats are used to correct for illumination chipvariation along skyeach of the slitlets., Quartz lamp flats are used to account for the response curve of the chip and twilight sky flats are used to correct for illumination variation along each of the slitlets.876 Spatial calibration is carried out by placing a thin wire in the filter wheel and illuminating the slitlet array with a continuum lamp., Spatial calibration is carried out by placing a thin wire in the filter wheel and illuminating the slitlet array with a continuum lamp.877 This procedure defines the center of each slitlet., This procedure defines the center of each slitlet.878 The individual spectra have no spatial distortion because the camera corrects the small amount of distortion introduced by the spectrograph., The individual spectra have no spatial distortion because the camera corrects the small amount of distortion introduced by the spectrograph.879 Thus only low-order mapping of the slitlets is required., Thus only low-order mapping of the slitlets is required.880 Wavelength calibration is performed using a CuAr arc lamp to provide sufficient lines in both the blue and red arms of the camera., Wavelength calibration is performed using a CuAr arc lamp to provide sufficient lines in both the blue and red arms of the camera.881 Arc lamp data were taken in between sets of object exposures., Arc lamp data were taken in between sets of object exposures.882 Each of the 25 slitlets is then rectified by the pipeline into full data cube (one for each arm) sampled on a common wavelengtha scale., Each of the 25 slitlets is then rectified by the pipeline into a full data cube (one for each arm) sampled on a common wavelength scale.883 The resulting data cubes from each exposure were flux and telluric calibrated using observations of the white dwarf Feige 110 and corrected for the effect of atmospheric dispersion., The resulting data cubes from each exposure were flux and telluric calibrated using observations of the white dwarf Feige 110 and corrected for the effect of atmospheric dispersion.884" Cosmic-ray removal was performed with the ""dcr"" routine (Pych2004)."," Cosmic-ray removal was performed with the ""dcr"" routine \citep{Pych04}."885". The final reduced and flux-calibrated data cubes were binned by 2 pixels along the slit in order to increase signal-to-noise ratio and to produce square spatial elements 1""x1"",, corresponding to ~260 pc at the distance of the galaxy."," The final reduced and flux-calibrated data cubes were binned by 2 pixels along the slit in order to increase signal-to-noise ratio and to produce square spatial elements $1\arcsec \times 1$, corresponding to $\sim260$ pc at the distance of the galaxy."886" We analyzed each spectrum using an automated fitting routine written in IDL, UHSPECFIT, which is based on the code created Zahidetal.(2010) and is also by Rupkeetal.by(2010)."," We analyzed each spectrum using an automated fitting routine written in IDL, UHSPECFIT, which is based on the code created by \citet{Zahid10} and is also employed by \citet{Rupke10}."887. Our routine fits and subtracts a employedstellar continuum from each spectrum using population synthesis models from GonzálezDelgadoetal.(2005) and an IDL routine which fits a linear combination of stellar templates to a galaxy spectrum using the method of Moustakas&Kennicutt (2006)., Our routine fits and subtracts a stellar continuum from each spectrum using population synthesis models from \citet{Gonzalez05} and an IDL routine which fits a linear combination of stellar templates to a galaxy spectrum using the method of \citet{Moustakas06}.888". Although our population synthesis models include stellar absorption features, it was necessary to fit excess absorption in or Hain some spectra to ensure a reliable emission line flux estimates."," Although our population synthesis models include stellar absorption features, it was necessary to fit excess absorption in or in some spectra to ensure a reliable emission line flux estimates."889 We believe this is due to the inadequacy of the stellar models., We believe this is due to the inadequacy of the stellar models.890" Lines in the resulting emission are fit a one or two-component Gaussian, dependingspectra on the usinggoodness of fit determined by the routine."," Lines in the resulting emission spectra are fit using a one or two-component Gaussian, depending on the goodness of fit determined by the routine."891 All of the emission lines are fit simultaneously using the same Gaussian component or components., All of the emission lines are fit simultaneously using the same Gaussian component or components.892 Resulting fits were given a cursory inspection by eye to ensure the routine did not fail., Resulting fits were given a cursory inspection by eye to ensure the routine did not fail.893" Both continuum and emission lines were fit using the MPFIT package, which performs a least- analysis using the Levenberg-Marquardt algorithm squares(Markwardt2009)."," Both continuum and emission lines were fit using the MPFIT package, which performs a least-squares analysis using the Levenberg–Marquardt algorithm \citep{Markwardt09}."894". The fitting routine provides a redshift, flux, and width for each Gaussian component in each spatial pixel."," The fitting routine provides a redshift, flux, and width for each Gaussian component in each spatial pixel."895" The blue and red spectra are fit simultaneously, while for slight variations in wavelength zero point andaccounting instrumental resolution between the two arms as parameters in the fitting routine."," The blue and red spectra are fit simultaneously, while accounting for slight variations in wavelength zero point and instrumental resolution between the two arms as parameters in the fitting routine."896" The,,55007A,,66300A,, 66583A,, ,, and emission lines were fit reliably across the disk and in the extended emission."," The, and emission lines were fit reliably across the disk and in the extended emission."897" Where there are two components, we consistently apply the ""principal"" label to the component with the strongest Ha emission line flux."," Where there are two components, we consistently apply the ""principal"" label to the component with the strongest Ha emission line flux."898 We apply a signal-to-noise cutoff of 5o for accurate measurement of emission line fluxes., We apply a signal-to-noise cutoff of $5\sigma$ for accurate measurement of emission line fluxes.899" This sets an effective flux sensitivity of ~2.5x10-7! ""erg s! em? inI], ΤΠ. aand aand ~1.5x107! ""erg s! cm""? in aand."," This sets an effective flux sensitivity of $\sim2.5 \times 10^{-17}$ erg $^{-1}$ $^{-2}$ in, , and and $\sim1.5 \times 10^{-17}$ erg $^{-1}$ $^{-2}$ in and."900". In all, some 300 spatial pixels (spaxels) lie above the >5c cutoff in the strongest line oor )) out of a total of ~1500 possible spatial elements in the final binned and mosaicked cube."," In all, some 300 spatial pixels (spaxels) lie above the $>5\sigma$ cutoff in the strongest line or ) out of a total of $\sim1500$ possible spatial elements in the final binned and mosaicked cube."901 An example of one of our fits is shown in Figure 2., An example of one of our fits is shown in Figure 2.902" In this example, the line splitting is clearly evident, as well as the underlying stellar continuum."," In this example, the line splitting is clearly evident, as well as the underlying stellar continuum."903" In this example, all of the emission lines shown in Region 1 are above the signal-to-noise cutoff mentioned above, while in Region 2 only aand aare above the cutoff, the rest of the emission line fluxes are thrown out for the purposes of ouranalysis."," In this example, all of the emission lines shown in Region 1 are above the signal-to-noise cutoff mentioned above, while in Region 2 only and are above the cutoff, the rest of the emission line fluxes are thrown out for the purposes of ouranalysis."904 The cutoff is applied to each component in order to avoid contamination, The cutoff is applied to each component in order to avoid contamination905the right panel in refAurlFS as a dotted line.,the right panel in \\ref{Aur1FS} as a dotted line.906 One finds that the general trend of the slope values is retained., One finds that the general trend of the slope values is retained.907" However. individual 71, values in the rebinned images can differ by almost the one sigma uncertainties of the values obtained in the maps determined at the respective spatial resolution."," However, individual $\gamma_{\rm low}$ values in the rebinned images can differ by almost the one sigma uncertainties of the values obtained in the maps determined at the respective spatial resolution."908" See refapp2 and A? for the effects the rebinning has on the 21, and ni&n Values.", See \\ref{app2} and \ref{app3} for the effects the rebinning has on the $\gamma_{\rm low}$ and $\gamma_{\rm high}$ values.909 In general the differences are well below the one sigma uncertainties., In general the differences are well below the one sigma uncertainties.910 However. in a few cases. such as the above quoted example of |I. larger differences can be found.," However, in a few cases, such as the above quoted example of 1, larger differences can be found."911 As a consequence of this result we chose to perform option iii) to determine our cloud structure parameters for all clouds at ppe resolution., As a consequence of this result we chose to perform option iii) to determine our cloud structure parameters for all clouds at pc resolution.912 We calculate all parameters at each of the four spatial resolutions available to us. and then interpolate to obtain the values at ppe.," We calculate all parameters at each of the four spatial resolutions available to us, and then interpolate to obtain the values at pc."913 All subsequent analysis is performed this way., All subsequent analysis is performed this way.914 Using the technique deseribed in reffittwo we obtained fit parameters for each cloud we analysed with reflom!2.., Using the technique described in \\ref{fittwo} we obtained fit parameters for each cloud we analysed with \\ref{lom12}.915 In refbintab we summarise the fit parameters and the root mean square deviation (7725). values obtained for the various spatial resolutions and histogram bin sizes for the 11 cloud as an example., In \\ref{bintab} we summarise the fit parameters and the root mean square deviation $rms$ ) values obtained for the various spatial resolutions and histogram bin sizes for the 1 cloud as an example.916 There is a general trend visible for all parameters., There is a general trend visible for all parameters.917 In particular the width of the distribution increases with spatial scale., In particular the width of the distribution increases with spatial scale.918 This is expected. since more andmore small scale high extinction cores are not detected anymore at these coarse resolutions.," This is expected, since more andmore small scale high extinction cores are not detected anymore at these coarse resolutions."919 In refAurlPBSFI we show the normalised column density distribution for the H1 cloud as a solid line (shown is the data for the spatial resolution closest to ppc., In \\ref{Aur1PBSF1} we show the normalised column density distribution for the 1 cloud as a solid line (shown is the data for the spatial resolution closest to pc).920 Overplotted is a fit with parameters scaled to ppe spatial scale., Overplotted is a fit with parameters scaled to pc spatial scale.921 Similar plots for all individual clouds can be seen in the refappd.., Similar plots for all individual clouds can be seen in the \\ref{app4}.922 We list the fit parameters and rijs values (scaled to ppe resolution) for all clouds in refanatab.., We list the fit parameters and $rms$ values (scaled to pc resolution) for all clouds in \\ref{anatab}.923 As described in reffitone we calculate gradients ~ for each cloud in our sample., As described in \\ref{fitone} we calculate gradients $\gamma$ for each cloud in our sample.924 As discussed. there are usually at least two distinct regions with different slopes.," As discussed, there are usually at least two distinct regions with different slopes."925" One region. at low extinction values (51,4) characterises the general turbulence of the cloud."," One region, at low extinction values $\gamma_{\rm low}$ ) characterises the general turbulence of the cloud."926 At higher extinction values (shies) gravity becomes important and changes the column density away from a log-normal distribution., At higher extinction values $\gamma_{\rm high}$ ) gravity becomes important and changes the column density away from a log-normal distribution.927 As an example we show the log(.N) vs Ay for LI in PBSFI.., As an example we show the $\log(N)$ vs $A_V$ for 1 in \\ref{Aur1PBSF1}. .928 The plots for the other clouds are shown in, The plots for the other clouds are shown in929of either interacting or merging galaxies.,of either interacting or merging galaxies.930 However. there is a large uncertainty in estimating merger fraction from these observations due to their shallow survey depth.," However, there is a large uncertainty in estimating merger fraction from these observations due to their shallow survey depth."931 In Figure 6. we compare this result. with our moclel predictions.," In Figure 6, we compare this result with our model predictions."932 Ehe downward arrow indicates the upper limit on the merger fraction at 2=5.7., The downward arrow indicates the upper limit on the merger fraction at $z=5.7$.933 A this recshift. our modcl predicted major merger fraction is about 50%.," At this redshift, our model predicted major merger fraction is about $50\%$."934 While comparing our mocdel predictions with the observations we have accounted for the observed limiting. ILaminosity at corresponcding recshilts., While comparing our model predictions with the observations we have accounted for the observed limiting luminosity at corresponding redshifts.935 At slightly lower redshift z25. (Pirzkalctal.2007) studied morphologics of nine. eemitters in HUDE.," At slightly lower redshift $z\approx 5$, \citep{pir07}936 studied morphologies of nine emitters in HUDF."937 To quantify. the morphologies. they used. concentralion(C). and. asgmmelrg(GY) parameters (Conseliceetal.2000). and found that nearly 44% have clumpy or complex structures.," To quantify the morphologies, they used $concentration (C)$, and $asymmetry (A)$ parameters \citep{con00}938 and found that nearly $44\%$ have clumpy or complex structures."939 Visually. about 3314. sources look morphologically disturbed or as ongoing mergers. while nearly LOY of the sources can not be reliably identified. as mergers.," Visually, about $33\%$ sources look morphologically disturbed or as ongoing mergers, while nearly $10\%$ of the sources can not be reliably identified as mergers."940 This observed merger fraction is nearly same às our nmiocel prediction., This observed merger fraction is nearly same as our model prediction.941 Bondetal.(2009). studied morphological properties of about 120 z— 3.1 cemitting galaxies in the rest-L[rame ultra-violet band., \citet{bon09} studied morphological properties of about 120 $z\sim$ 3.1 emitting galaxies in the rest-frame ultra-violet band.942 They found that at least 174 of the total eemitters contain multiple components which might indicate that these are either individual star-forming regions within a single galaxy. a merged system or ongoing mergers.," They found that at least $17\%$ of the total emitters contain multiple components which might indicate that these are either individual star-forming regions within a single galaxy, a merged system or ongoing mergers."943 Since ib ds very dillicult. due to their compact sizes (e.g. Malhotra et al 2011. Bond et al 2011). to definitely conclude whether the multi-component cemitters are the remnants of mergers or if these are individual star-forming regions in a single system. we have shown the merger fraction by upward and. downward arrows indicating uncertainties in both directions.," Since it is very difficult, due to their compact sizes (e.g. Malhotra et al 2011, Bond et al 2011), to definitely conclude whether the multi-component emitters are the remnants of mergers or if these are individual star-forming regions in a single system, we have shown the merger fraction by upward and downward arrows indicating uncertainties in both directions."944 Their classification is based on counting the number of components in a fixed aperture., Their classification is based on counting the number of components in a fixed aperture.945 At z=3.1. our model predicted major merger fraction is about 35%. much higher than the observed fraction.," At $z=3.1$, our model predicted major merger fraction is about $35\%$, much higher than the observed fraction."946 Llowever. in Section 4 we show that. some of this dillerence between model predictions ancl observations can be attributed to the definition of major merger mass ratio.," However, in Section 4 we show that, some of this difference between model predictions and observations can be attributed to the definition of major merger mass ratio."947 At lower redshift. οz0.3. Cowieetal.(2010). studied morphologies of cemitters in the GROTIOO and SIRLEELOO fields.," At lower redshift, $z\approx 0.3$, \citet{cow10} studied morphologies of emitters in the GROTH00 and SIRTFFL00 fields."948 They found that z30% of the cemitters show signs of ongoing mergers., They found that $>30\%$ of the emitters show signs of ongoing mergers.949 Dased on the above comparisons. there is some discrepancy. between our model prediction. ancl observations.," Based on the above comparisons, there is some discrepancy between our model prediction and observations."950 In the following section. we investigate the uncertainty due to the mass ratio used to define a merger.," In the following section, we investigate the uncertainty due to the mass ratio used to define a merger."951" We now vary the progenitor mass ratio criteria from 1:3 to 1:2. Le. major mergers are now defined as those for which 0.5 mefmy1l. while for 0.1x:mom,«0.5. they are Classified as minor mergers."," We now vary the progenitor mass ratio criteria from 1:3 to 1:2, $i.e.$ major mergers are now defined as those for which $0.5 \leq m_{2}/m_{1} \leq 1$, while for $0.1 \leq m_{2}/m_{1} < 0.5$, they are classified as minor mergers."952 All other. eemitters with their progenitor halo mass ratio moy«0.1 ave defined as smooth accreting., All other emitters with their progenitor halo mass ratio $m_{2}/m_{1} < 0.1$ are defined as smooth accreting.953 In figure 7 (see also Figure 2) we show the predicted merger fraction dependence on the progenitor mass ratio definition., In figure 7 (see also Figure 2) we show the predicted merger fraction dependence on the progenitor mass ratio definition.954 Dv changing the progenitor mass ratio from 1: 3 to 1: 2. the major (minor) merger fraction drops. (increases) bv about 15%.," By changing the progenitor mass ratio from 1: 3 to 1: 2, the major (minor) merger fraction drops (increases) by about $15\%$."955 Thus. it is clear that the predicted. merger fraction of comitters depends. on the the progenitor mass ratio definition. anc one needs to be careful when comparing model predicted. merger. fractions with the observations since the merecr defining mass ratio inlluences the merger fractions.," Thus, it is clear that the predicted merger fraction of emitters depends on the the progenitor mass ratio definition, and one needs to be careful when comparing model predicted merger fractions with the observations since the merger defining mass ratio influences the merger fractions."956 1n addition to the above uncertainty. the observed. merger," In addition to the above uncertainty, the observed merger"957the overall slope of the continuum of the phase 0.0 A-band spectrum is consistent will a very late-tvpe M cw (see Fig.,the overall slope of the continuum of the phase 0.0 -band spectrum is consistent with a very late-type M dwarf (see Fig.958 6). the ZI-band spectrum bears no resemblence to such an object.," 6), the -band spectrum bears no resemblence to such an object."959 The departure becomes more pronounced when we compare the spectrum of EF Evi to those of L-dwarls., The departure becomes more pronounced when we compare the spectrum of EF Eri to those of L-dwarfs.960 While brown clwarls around L5 show a similar spectral slope al 1.55—1.6 jm. there is no feature comparable to the strong flix peak at yan. While the absence of Nall in the A-band is consistent wilh a spectral (wpe later than LO (McLean et al.," While brown dwarfs around L5 show a similar spectral slope at $-$ $\mu$ m, there is no feature comparable to the strong flux peak at $\mu$ m. While the absence of I in the $K$ -band is consistent with a spectral type later than L0 (McLean et al."961 2003). the water vapor absorption bands of such an object are much deeper ancl clepress the blue side of the A-band spectrum much farther to the red (2 jm) (han is seen in EF Evi.," 2003), the water vapor absorption bands of such an object are much deeper and depress the blue side of the $K$ -band spectrum much further to the red $>$ $\mu$ m) than is seen in EF Eri."962 Also. the first overtone CO bandhead at jim remains strong in brown dwarls until CII; takes over in the late L types.," Also, the first overtone CO bandhead at $\mu$ m remains strong in brown dwarfs until $_4$ takes over in the late L types."963 In addition to the difference in //- ancl A-band spectral features. the 77/7A photometry is almost impossible to reconcile with anv known dwarf.," In addition to the difference in $H$ - and $K$ -band spectral features, the $JHK$ photometry is almost impossible to reconcile with any known dwarf."964 While the7/ to A-band fIux ratios for EF Exi are more-or-less consistent with those of late (wpe objects. giving rise to (I4—iy) = 0.42 (zz MSV). the observed (J— A) color (> 2.5) exceeds even the reddest IR. colors observed in mid-L dwarls (Legget οἱ al.," While the to -band flux ratios for EF Eri are more-or-less consistent with those of late type objects, giving rise to $H - K$ ) = 0.42 $\approx$ M8V), the observed $J - K$ ) color $>$ 2.5) exceeds even the reddest IR colors observed in mid-L dwarfs (Legget et al."965 2002)., 2002).966 As shown in Ilarrison et al. (, As shown in Harrison et al. (9672004). many. CV. secondary stars show evidence [or weak CO absorption features.,"2004), many CV secondary stars show evidence for weak CO absorption features."968 Since the water vapor features in those objects appear to be relatively normal. Harrison et al.," Since the water vapor features in those objects appear to be relatively normal, Harrison et al."969 suspect that Carbon is deficient in those svstems., suspect that Carbon is deficient in those systems.970 This view is supported by analvsis of FUV observations of CVs by Gdusicke et al. (, This view is supported by analysis of FUV observations of CVs by $\ddot{a}$ nsicke et al. (9712003) where large N V/C IV ratios have been found. suggesting Nitrogen enhancements. aud Carbon delicits in the photospheres of the white dwarf primaries. presumably arising [rom matter transferred Irom the secondary star.,"2003) where large N V/C IV ratios have been found, suggesting Nitrogen enhancements, and Carbon deficits in the photospheres of the white dwarf primaries, presumably arising from matter transferred from the secondary star."972 It would therefore not be unexpected for the secondary star in EF Eri to have anomalously weak CO absorption features., It would therefore not be unexpected for the secondary star in EF Eri to have anomalously weak CO absorption features.973 A lowered Carbon, A lowered Carbon974extensive modeling to extract this measure (Lahavetal.1993).,extensive modeling to extract this measure \citep{lah93}.975. Additionally. given the potentially complex residual contamination by diffuse. or unresolved. Galactic emission. and the known selection biases of the Abell catalog. the stacking approach allows for more intuitive modeling (see 833.1 below).," Additionally, given the potentially complex residual contamination by diffuse, or unresolved, Galactic emission, and the known selection biases of the Abell catalog, the stacking approach allows for more intuitive modeling (see 3.1 below)."976" Given an equatorial pixel size of 0.5"" for the EGRET data and the flux enclosure ab 1° radius for the energv weighted 100 MeV. PSF. we choose annular bins of width 1 in ow analvsis."," Given an equatorial pixel size of $0.5^{\circ}$ for the EGRET data and the flux enclosure at $1^{\circ}$ radius for the energy weighted $>100$ MeV PSF, we choose annular bins of width $1^{\circ}$ in our analysis."977" We nole (hat the energv dependent PSF can be described as 3, where ϐ is the energv. dependent radius for flux enclosure (Thompsonetal.1993:Esposito1999)."," We note that the energy dependent PSF can be described as $\theta \leq 5.85^{\circ}978(E_{\gamma}/100 {\rm MeV})^{-0.534}$ , where $\theta$ is the energy dependent radius for flux enclosure \citep{tho93,esp99}."979. The area-noise weighted mean flux. excess (hereafter referred. (ο as the mean. see 822) above the global mean (calenlated [rom all unmmasked. Galactic-corrected pixels) is then evaluated by using all umnasked pixels in each radial bin. out to 20° [or each cluster. ancl then averaging over all clusters used.," The area-noise weighted mean flux excess (hereafter referred to as the mean, see 2) above the global mean (calculated from all unmasked, Galactic-corrected pixels) is then evaluated by using all unmasked pixels in each radial bin, out to $20^{\circ}$ for each cluster, and then averaging over all clusters used."980 Pixels are used if their centers lie within a given annulus. consequently (here is some variation in (he number of pixels counted between different cluster centers. however (his variation is negligible on averaging over many objects.," Pixels are used if their centers lie within a given annulus, consequently there is some variation in the number of pixels counted between different cluster centers, however this variation is negligible on averaging over many objects."981" We refer to this angular [uncetion as i4,(8) or <Af>=</—-—J/> in the results in $3 below.", We refer to this angular function as $w_{c\gamma}(\theta)$ or $<\Delta I>=<I-\bar{I}>$ in the results in 3 below.982 In 83.1 we assess ils significance., In 3.1 we assess its significance.983 In Figure 4 we plot our principal results., In Figure 4 we plot our principal results.984 The curve for all 2469 Abell elusters with |b]>45°- peaks αἱ a value of 6.2x10* phs tem 7? ! in the 1 bin. and gently declines with 9.," The curve for all 2469 Abell clusters with $|b|>45^{\circ}$ peaks at a value of $6.2\times10^{-7}$ ph $^{-1}$ $^{-2}$ $^{-1}$ in the $^{\circ}$ bin, and gently declines with $\theta$."985 This is generically the behavior expected [or a positive correlation between the two datasets., This is generically the behavior expected for a positive correlation between the two datasets.986 Upper and lower heavy curves are for (he richest (22> 2) ancl poorest (RH.< 2) Abell subsets., Upper and lower heavy curves are for the richest $R\geq 2$ ) and poorest $R<2$ ) Abell subsets.987 We note that. based on the theoretical predictions. more massive cluster svslenms are expected (o be the sites of higher egamma-ray eniission (Loeb&Waxman2000).," We note that, based on the theoretical predictions, more massive cluster systems are expected to be the sites of higher gamma-ray emission \citep{loe00}."988. Consequently. the increased amplitude of (he rich cluster ic relative to Chat of the poorer clusters goes in the expected sense if clusters are indeed difDuse gammimna-ray sources.," Consequently, the increased amplitude of the rich cluster $w_{c\gamma}$ relative to that of the poorer clusters goes in the expected sense if clusters are indeed diffuse gamma-ray sources."989 The peak amplitude for the rich cluster subset. in the1° bin. is 1.19xI0 phs !em ?2sr |...," The peak amplitude for the rich cluster subset, in the$1^{\circ}$ bin, is $1.19\times 10^{-6}$ ph $^{-1}$ $^{-2}$ $^{-1}$."990 We have also investigated the effect of using subsets of the Abell catalog divided using the distance parameter D. which is based on the brighter galaxy magnitudes in a cluster.," 	 We have also investigated the effect of using subsets of the Abell catalog divided using the distance parameter $D$, which is based on the brighter galaxy magnitudes in a cluster."991 The increased noise of using smaller subsets of clusters makes (hese measurements less significant., The increased noise of using smaller subsets of clusters makes these measurements less significant.992 Using only the more distant. (2> 4) subset. which is still large. has little effect on the results presented here. with variations well within (he noise.," Using only the more distant $D> 4$ ) subset, which is still large, has little effect on the results presented here, with variations well within the noise."993 In addition. since N-rav-selected clusters are considered to be more robust in terms of being real. eravilationally relaxed svstems. we have run our cross-correlation using the 304 clusters of the BCS (Ebelingetal.1998. 2000)..," In addition, since X-ray-selected clusters are considered to be more robust in terms of being real, gravitationally relaxed systems, we have run our cross-correlation using the 304 clusters of the BCS \citep{ebe98,ebe00}. ."994 However. only 159 svstems remain alter the," However, only 159 systems remain after the"995GR was often quoted and used as a significance).,$R$ was often quoted and used as a significance).996 A conservative threshold 2=3 means that we consider only QPOs for which we can measure the power of the Lorentzian with an accuracy of 30 or more., A conservative threshold $R=3$ means that we consider only QPOs for which we can measure the power of the Lorentzian with an accuracy of $3\sigma$ or more.997 Such a threshold corresponds (oa ~Go excess power in the PDS for a single5 (rial. equivalent to ~4o if we account for the number of trials of the scanning5 procedure (vanderlis1989)..," Such a threshold corresponds to a $\sim 6\sigma$ excess power in the PDS for a single trial, equivalent to $\sim9984\sigma$ if we account for the number of trials of the scanning procedure \citep{van-der-Klis:1989kn}."999 The integrated. power of the Lorentzian is (hen converted into a root mean square (RAIS). expressed as a fraction of the total source count rate.," The integrated power of the Lorentzian is then converted into a root mean square (RMS), expressed as a fraction of the total source count rate."1000 We detect a significant QPO in twelve segments of observations. within the same ObsIDs as in Sannaοἱal.(2010)..," We detect a significant QPO in twelve segments of observations, within the same ObsIDs as in \citet{sanna10mnras}."1001 The reduced number of active PCUs on the PCA. and the relative [aintness of the source in its atoll phase (count rate between ~40 and ~120 counts/s/PCU). means (that special care must be taken to correct lor the Irequeney dift of the lower κκ QPO.," The reduced number of active PCUs on the PCA, and the relative faintness of the source in its atoll phase (count rate between $\sim 40$ and $\sim 120$ counts/s/PCU), means that special care must be taken to correct for the frequency drift of the lower kHz QPO."1002 We apply the very same technique as in Barretοἱal.(2006)..., We apply the very same technique as in \citet{barret06mnras}.1003 IU is an iterative procedure. which enables us to bound the QPO frequencies on shorter aud shorter integration (ime. will narrower and narrower frequeney intervals.," It is an iterative procedure, which enables us to bound the QPO frequencies on shorter and shorter integration time, with narrower and narrower frequency intervals."1004 As the integration time decreases. the significance threshold set to the scanning technique (Boirinetal.2000) is adjusted (e.g the threshold is decreased [rom 4o. [or an interval of 50 Lz over a few thousanud seconds. to 36 [or an interval of 10 Lz over a few hundreds of seconds).," As the integration time decreases, the significance threshold set to the scanning technique \citep{Boirin:2000jt} is adjusted (e.g the threshold is decreased from $4\sigma$ for an interval of 50 Hz over a few thousand seconds, to $3 \sigma$ for an interval of 10 Hz over a few hundreds of seconds)."1005 The QPO path is then recovered through a linear interpolation between the most significant detections with the shortest integration time., The QPO path is then recovered through a linear interpolation between the most significant detections with the shortest integration time.1006 An example of the application of the analvsis to (vo segments of real observations is shown in Figure 1.., An example of the application of the analysis to two segments of real observations is shown in Figure \ref{fig1}.1007 such a technique. which we have used extensively on real and simulated data. is able to correct precisely for the Irequency drift of QPOs of the strength reported here.," Such a technique, which we have used extensively on real and simulated data, is able to correct precisely for the frequency drift of QPOs of the strength reported here."1008 In (those simulations. (he QPO frequency evolution is modeled by a random walk of a given step (e.g. 0.2-0.5 IIz/second).," In those simulations, the QPO frequency evolution is modeled by a random walk of a given step (e.g. 0.2-0.5 Hz/second)."1009 We follow Timmer&Ixoenig(1995) ancl generate svnthetic 1 second PDS. lor which the underlying model consists of a constant (2) to account [or the Poisson noise plus a lorentzian to model the οΡΟ profile. with parameters (Q and RAIS) appropriate for each frequency (estimated from the random walk).," We follow \cite{timmer95aa} and generate synthetic 1 second PDS, for which the underlying model consists of a constant (2) to account for the Poisson noise plus a lorentzian to model the QPO profile, with parameters (Q and RMS) appropriate for each frequency (estimated from the random walk)."1010 The simulated PDS are then combined on (16 second timescales) ancl scanned as the real data to recover the time evolution of the QPO Irequency. and to determine Q and RAIS. lor comparison with the parameters injected in the simulations.," The simulated PDS are then combined on (16 second timescales) and scanned as the real data to recover the time evolution of the QPO frequency, and to determine Q and RMS, for comparison with the parameters injected in the simulations."1011 The difference between (he reconstructed and original Q ancl RAIS parameters can (hen be evaluated., The difference between the reconstructed and original Q and RMS parameters can then be evaluated.1012 Those simulations have shown (hat lor QPOs of comparable strength to the one of, Those simulations have shown that for QPOs of comparable strength to the one of1013Applving this to the second jump condition (31). we obtain a solution for the eigenvalue ωτσ--ο. where jo is the density contrast parameter defined in equation (29).,"Applying this to the second jump condition ), we obtain a solution for the eigenvalue $\omega = \sigma + m\Omega$, where $\mu$ is the density contrast parameter defined in equation )."1014 We are interested in unstable modes. ie.. modes with complex w. where the erowth rate of the instability is given by à.," We are interested in unstable modes, i.e., modes with complex $\omega$, where the growth rate of the instability is given by $\omega_{\rm I}$."1015 From equation (36). we see that unstable modes exist whenever μι3+ur2.9G7«0. Le.te il Thus. the instability exists only if ji is positive. Le. if the density on the outside is ereater (han that on the inside.," From equation ), we see that unstable modes exist whenever $-\mu m \Omega_{\rm eff}^2 +\mu^2 \zeta^2 <0$, i.e. if Thus, the instability exists only if $\mu$ is positive, i.e., if the density on the outside is greater than that on the inside."1016 This is perfectly natural for the Ravleigh-Tavlor instability., This is perfectly natural for the Rayleigh-Taylor instability.1017 Surprisinglv. if the density contrast is (oo large. ie. if 4627yey. the instability shuts off.," Surprisingly, if the density contrast is too large, i.e., if $\mu>\mu_1$, the instability shuts off."1018 This is Clearly the result of rotation. or more specilically vorticitv.," This is clearly the result of rotation, or more specifically vorticity."1019 When q=2 aud (he vorticity C—0. then jn—x. and the instability is present [or any positive value of µ.," When $q=2$ and the vorticity $\zeta=0$, then $\mu_1\to\infty$, and the instability is present for any positive value of $\mu$."1020 However. for q—0. we have €=Ὁ and jn is finite.," However, for $q=0$, we have $\zeta=\Omega$ and $\mu_1$ is finite."1021 In this case. if O2; small (i.e.. the effective gravity is weak) and if we consider a low value of the azimuthal waveniunber i. the vorticity is able io eliminate the instability.," In this case, if $\Omega_{\rm eff}^2$ small (i.e., the effective gravity is weak) and if we consider a low value of the azimuthal wavenumber $m$, the vorticity is able to eliminate the instability."1022 The real part of w gives the observed oscillation frequency. of the mode., The real part of $\omega$ gives the observed oscillation frequency of the mode.1023" When q=2. (his is simply equal to mQ,,."," When $q=2$, this is simply equal to $m\Omega_m$."1024" In (his case. the mode is stationary in (he frame of the [Iuid al (he boundary. aid what one observes is simply the Keplerian [requency Q,,TH of the gas al (he inner edge of the disk."," In this case, the mode is stationary in the frame of the fluid at the boundary, and what one observes is simply the Keplerian frequency $\Omega_m$ of the gas at the inner edge of the disk."1025" However. when q=0. the observed mode frequency is not equal to mQ,,. but is equal to (i+p), which differs from the orbital frequency."," However, when $q=0$, the observed mode frequency is not equal to $m\Omega_m$, but is equal to $(m+\mu)\Omega_m$, which differs from the orbital frequency."1026 In this case. because of vorticity. we have a traveling mode in the fIuid frame.," In this case, because of vorticity, we have a traveling mode in the fluid frame."1027 For a moregeneral disk. with 40 and/or qz0.2. we have to solve the wave equation and the boundary. conditions numerically.," For a moregeneral disk, with $\gamma\ne0$ and/or $q\ne 0,2$, we have to solve the wave equation and the boundary conditions numerically."1028 Although we derived in a differential equation as our basic wave equation (eq. [19]).," Although we derived in a second-order differential equation as our basic wave equation (eq. ]),"1029 for numerical purposes a pair of equivalent first-orcer cdilferential equations is more convenient., for numerical purposes a pair of equivalent first-order differential equations is more convenient.1030 Defining V—re. we," Defining $V = r v$, we"1031‘This is why the ∖− ∖⇁⋜↧↓⋯⊾⊳∖⊀↓⊔↾↓∖⋜↧∣⋡↓⋖⋅⇉⇂⋅∪↓⋅↿↓∐⋅∣⇜∩∆∶∪⋅↓⋅↱≻⊔↓⋯⇂∢⊾↓ ⋜⊔⋅⋖⋅↓∪∖∖⊽∢⊾↓⋅↿↓⋯⊔↿↓↕∪⊳∖⋖⊾∪⇂⋅↿↓⊔⋅∣∽↙⋎∆∶∪⊔↓⋯⇂⋖⋅↓⋜∐∐⇀∖⋖⋅∠⇂∣⋯∣∟⇘⊳ ↾↓∖↓↥∢⊾∢⊾∐⋅⋖⋅≼∼↿⊳∖∪⇂⋅⋏∙≟↓⋅⋜↧∖⋰∐⋜,"This is why the $\hat \chi^2$ values in Table \ref{tab1}1032 for the $\omega_K = 0.15$ model are lower than those of the $\omega_K=0$ model at fixed $\ell_{\rm max}$."1033∐⊲↓∪↓⋯↓⇂∢⊾⊔⊳∖⊀↓⊔⋏∙≟∪⊔↿↓↥∢⋅↓≻⋖⋟↓⋜↧↓⋅↕≻⋜↧↿↕⋖≱↓↕ ≻∪∖∖⊽∢⊾↓⋅⊳∖↓≻⋖⊾≼∼↓⋅⊔⊔↓⋜⊔⋅∢⊾∢⋅∖⇁∢⊾⊔⊔↓∪↓⋅⋖⋅⊳∖⊲↓⋏∙≟⊔⊀↓∐≼⇍∥⊔↿↿↓⋯⊔⇂⋅∪↓⋅ he temperature power spectrum. producing significant distortions at £o1000.," The effects of gravitational lensing on the polarisation power spectrum are even more significant than for the temperature power spectrum, producing significant distortions at $\ell \sim 1000$."1034 At first sight the high values of (7 for »olarisation might seem surprising., At first sight the high values of $\hat \chi^2$ for polarisation might seem surprising.1035 They arise because the »olarisation power spectrum has sharper minima ancl peaks compared to the temperature power spectrum Figures 2w.b) and is therefore. more sensitive to the gravitational ensingconvolution (equation 5).," They arise because the polarisation power spectrum has sharper minima and peaks compared to the temperature power spectrum Figures 2a,b) and is therefore more sensitive to the gravitational lensingconvolution (equation 5)."1036 This build up of X7 with multipole for polarisation is illustrated in Figure 3., This build up of $\hat \chi^2$ with multipole for polarisation is illustrated in Figure 3.1037 Llowever. as the amplitude of the polarisation power spectrum is almost two orders of magnitude lower than that of the temperature Duetuations. an experiment. with high sensitivity to polarisation. in addition to high angular resolution. is required to detect the lensing contribution.," However, as the amplitude of the polarisation power spectrum is almost two orders of magnitude lower than that of the temperature fluctuations, an experiment with high sensitivity to polarisation, in addition to high angular resolution, is required to detect the lensing contribution."1038 Indeed. for a ALAP type mission. the results in Table 2 show that gravitational lensing is cillicult to detect in either temperature or polarisation spectra.," Indeed, for a MAP type mission, the results in Table \ref{tab1} show that gravitational lensing is difficult to detect in either temperature or polarisation spectra."1039 Llowever. for Planck the lensing ellects might be significant in both. power spectra and. despite the lower amplitude. are usually more. easily detected: in. polarisation than in the temperature signal.," However, for Planck the lensing effects might be significant in both power spectra and, despite the lower amplitude, are usually more easily detected in polarisation than in the temperature signal."1040 In a Planck-like (or cosmic variance limited) experiment. the lensing contribution is measurable also in the cross-correlation power spectrum. though less easily than in the polarisation spectrum.," In a Planck-like (or cosmic variance limited) experiment, the lensing contribution is measurable also in the cross-correlation power spectrum, though less easily than in the polarisation spectrum."1041 The dependence of the gravitational lensing contribution on cosmological parameters can be quantified by computing the derivatives of the lensed power spectra with respect to the cosmological paramoeter of interest., The dependence of the gravitational lensing contribution on cosmological parameters can be quantified by computing the derivatives of the lensed power spectra with respect to the cosmological parameter of interest.1042 However. this presents a difficult numerical problem. because the lensing contribution is small and hence the derivatives can be easily swamped by numerical errors.," However, this presents a difficult numerical problem, because the lensing contribution is small and hence the derivatives can be easily swamped by numerical errors."1043 In particular. the direct finite dillerencing scheme used to compute linear power spectrum derivatives (Bond 1997. Zaldarriaga LOOT: Eisenstein 1998. see also Section 4.2). cannot beapplied to the lensed case given the typical numerical errors of ~1% in CMD 3oltzmann codes.," In particular, the direct finite differencing scheme used to compute linear power spectrum derivatives (Bond 1997, Zaldarriaga 1997; Eisenstein 1998, see also Section 4.2), cannot be applied to the lensed case given the typical numerical errors of $\sim 1 \%$ in CMB Boltzmann codes."1044 Instead. we have applied two clillerent semi-analvtical approaches basec on numerical derivatives of the matter power spectrum rather than those of ἐς , Instead we have applied two different semi-analytical approaches based on numerical derivatives of the matter power spectrum rather than those of $\wtilde C_\ell$ .1045"The derivative with respect to a cosmological parameter 5; can be expressed as (4=1. ie""). where computing derivatives of the window functions MV. Moinvolves computations. of⋅ the derivatives. of the photon path clispersions o(8) and σο(06)."," The derivative with respect to a cosmological parameter $s_i$ can be expressed as $I=T,E,C$ ), where computing derivatives of the window functions ${\cal W}^I$ involves computations of the derivatives of the photon path dispersions $\sigma\l(\theta\r)$ and $\sigma_2\l(\theta\r)$."1046 In our first approach. derivatives of a7(6). 83(8) zux C; are computed numericallyby finite dilferencing.," In our first approach, derivatives of $\sigma^2\l(\theta\r)$, $\sigma_2^2\l(\theta\r)$ and $C_\ell$ are computed numericallyby finite differencing."1047 In the second. method. only the derivative of C5 is compute bv finite dillerencing. and the derivatives of both o7(8) and a52(0) are calculated: using numerically precomputec derivatives of the gravitational potential power spectrum D.(ag).," In the second method, only the derivative of $C_\ell$ is computed by finite differencing, and the derivatives of both $\sigma^2\l(\theta\r)$ and $\sigma_2^2\l(\theta\r)$ are calculated using numerically precomputed derivatives of the gravitational potential power spectrum $P_{\phi}\l(k,\eta\r)$."1048 For example. derivatives of σ(6) are. obtaine through the formula. derived. from equation (9)).," For example, derivatives of $\sigma^2\l(\theta\r)$ are obtained through the formula, derived from equation \ref{sigmath}) )."1049 An analogous expression can be written for 03(6)., An analogous expression can be written for $\sigma_2^2\l(\theta\r)$.1050" To compute the derivative of I, at any given value of the scale [actor a. we assumed that £2,grows. according to Linear theory, {δα= (haps). where Di.eps) is the linear erowth [actor ancl depends only the scale factor and. cosmological"," To compute the derivative of $P_{\phi}$ at any given value of the scale factor $a$ , we assumed that $P_{\phi}$grows according to linear theory, $P_{\phi}(k,a)=D^2(a,a_{LS})P_{\phi}(k,a_{LS})$ , where $D(a, a_{LS})$ is the linear growth factor and depends only the scale factor and cosmological"1051 , 1052The shape of our space is one of the greatest and the oldest questions in human history.,The shape of our space is one of the greatest and the oldest questions in human history.1053 Ancient people tried to answer the question mythologically or philosophically. but in the present era. We are ready to answer it scientifically.," Ancient people tried to answer the question mythologically or philosophically, but in the present era, we are ready to answer it scientifically."1054 When we consider the universe as a 4-manifold consisting of 3-space and |-time. answering the question is translated into determining the local geometry and the global topology of ourUniverse.," When we consider the universe as a 4-manifold consisting of 3-space and 1-time, answering the question is translated into determining the local geometry and the global topology of our."1055 Local geometry Is described by Einstein's General Relativity (GR)., Local geometry is described by Einstein's General Relativity (GR).1056 The metric of space-time plays a fundamental role. and the assumption of the cosmological principle. which states that our is (locally) homogeneous and isotropic. leads us to a famous Friedman-Lemaittre-Robertson-Walker (FLRW) metric. where & is the normalized curvature of space. such as k=+1 (spherical geometry). 0 (Euclidean geometry). or —1 (hyperbolic geometry).," The metric of space-time plays a fundamental role, and the assumption of the cosmological principle, which states that our is (locally) homogeneous and isotropic, leads us to a famous Friedman-Lemaîttre-Robertson-Walker (FLRW) metric, where $k$ is the normalized curvature of space, such as $k= +1$ (spherical geometry), 0 (Euclidean geometry), or $-1$ (hyperbolic geometry)."1057 Together with Einstein's equations. curvature is related to the average energy density of the universe and then to other physical quantities.," Together with Einstein's equations, curvature is related to the average energy density of the universe and then to other physical quantities."1058 Recent observations favor a A-CDM universe with curvature k=0 (e.g. Ow=1.0050700001 from WMAP++BAO+SN data. by Hinshaw et al.," Recent observations favor a $\Lambda$ -CDM universe with curvature $k\simeq0$ (e.g. $\Omega_{\mathrm{tot}}=1.0050^{+0.0060}_{-0.0061}$ from +BAO+SN data, by Hinshaw et al."1059 2009). suggesting that our ts one of the flat spaces with Euclidean geometry.," 2009), suggesting that our is one of the flat spaces with Euclidean geometry."1060 Global topology. on the other hand. has no reliable physical theories to describe it. so. constrain it mathematically through direct observations.," Global topology, on the other hand, has no reliable physical theories to describe it, so constrain it mathematically through direct observations."1061 This situation is analogous to the case of Carl Gauss who had to measure the angles of the large triangle formed by three peaks of mountains to know which geometry describes our space. since he did not know GR àd modern cosmology.," This situation is analogous to the case of Carl Gauss who had to measure the angles of the large triangle formed by three peaks of mountains to know which geometry describes our space, since he did not know GR and modern cosmology."1062 Moreover. it is true that the global topology of the universe has been a relatively less popular concept than the local geometry.," Moreover, it is true that the global topology of the universe has been a relatively less popular concept than the local geometry."1063 À pioneering work on cosmic topology was done by Ellis (1971)., A pioneering work on cosmic topology was done by Ellis (1971).1064 Other early works include Sokolov and Shvartsman (1974). Fang and Sato (1985). Gott (1980). Fangundes (1983). and so on.," Other early works include Sokolov and Shvartsman (1974), Fang and Sato (1985), Gott (1980), Fangundes (1983), and so on."1065 This theoretical or observational research has not due to the lack of observational data., This theoretical or observational research has not due to the lack of observational data.1066 In the past two decades. however. we have seen tremendous progress in this field. along with progress in observational techniques.," In the past two decades, however, we have seen tremendous progress in this field, along with progress in observational techniques."1067 Specifically. a possibility of the Poincaré dodecahedral space topology suggested by Luminet et al. (," Specifically, a possibility of the Poincaré dodecahedral space topology suggested by Luminet et al. ("10682003) was a breakthrough in cosmic topology.,2003) was a breakthrough in cosmic topology.1069 This field recieves more and more interest these days. not only from theorists but also from observational astronomers.," This field recieves more and more interest these days, not only from theorists but also from observational astronomers."1070 The overall topology of the is how becoming one of the major concerns in astronomy and cosmology., The overall topology of the is now becoming one of the major concerns in astronomy and cosmology.1071 In modern cosmology with an FLRW= metric. already described above. we have three geometriessign.," In modern cosmology with an FLRW metric, already described above, we have three geometries."1072" For each geometry. there i5 only one type of space with a simply connected topology. namely. 3-sphere 22. 3-Euclidean space 2°. and 3-hyperbolic space IE"". for k=«1.0. and -I. respectively."," For each geometry, there is only one type of space with a simply connected topology, namely, 3-sphere $\mathbb{S}^3$ , 3-Euclidean space $\mathbb{E}^3$, and 3-hyperbolic space $\mathbb{H}^3$, for $k=+1, 0,$ and $-1$, respectively."1073 It is possible to construct a space that is locally indistinguishable from the simply connected one. re.. having the same It is not simply connected. but multiconnected.," It is possible to construct a space that is locally indistinguishable from the simply connected one, i.e., having the same It is not simply connected, but multiconnected."1074 We give a brief review of multiconnected spaces below., We give a brief review of multiconnected spaces below.1075 Detailed treatments are found in various reviews (e.g. Lachiézze-Rey and Luminet 1995)., Detailed treatments are found in various reviews (e.g. Lachièzze-Rey and Luminet 1995).1076 Amulticonnected space M is a quotient space of the simply connected space U with the same geometry. by a holonomy group I. where a holonomy is an isometry on UÜ without any fixed points (except for the identity).," Amulticonnected space $M$ is a quotient space of the simply connected space $U$ with the same geometry, by a holonomy group $\Gamma$, where a holonomy is an isometry on $U$ without any fixed points (except for the identity)."1077 Hence. M.=U/T can be as LU tiled by polyhedra identified by holonomies y€I. possessing structures.," Hence, $M=U/\Gamma$ can be as $U$ tiled by polyhedra identified by holonomies $\gamma \in \Gamma$, possessing structures."1078 This polyhedron. called a fundamental cell. ts a 2K- polyhedron whose K pairs of faces are glued mathematically by holonomies.," This polyhedron, called a fundamental cell, is a $2K$ -polyhedron whose $K$ pairs of faces are glued mathematically by holonomies."1079 The definition of fundamental cell is not unique. and in this paper by fundamental cell we mean the Dirichlet domain seen from the “center” of the universe. D(Xj)=λεU|vy€T.|X-Xol[V— γα].," The definition of fundamental cell is not unique, and in this paper by fundamental cell we mean the Dirichlet domain seen from the “center"" of the universe, $D(\vec x_0)=\{\vec x \in U \ | \ \forall \gamma \in \Gamma, |\vec x - \vec x_0| \leq |\vec x -\gamma \vec x _0| \}$ ."1080 In some space that is globally inhomogeneous. we can define a center. and the fundamental cells for 17 multiconnected flat spaces are found in Figure |.," In some space that is globally inhomogeneous, we can define a center, and the fundamental cells for 17 multiconnected flat spaces are found in Figure 1."1081 No two points in a Dirichlet domain be linked by any holonomies: in this sense a Dirichlet domain represents the whole universe., No two points in a Dirichlet domain be linked by any holonomies; in this sense a Dirichlet domain represents the whole universe.1082 In a globally homogeneous space. the shape of the Dirichlet domain is independent of the observer's position. since it is always the same as that of the fundamental cell.," In a globally homogeneous space, the shape of the Dirichlet domain is independent of the observer's position, since it is always the same as that of the fundamental cell."1083 In a globally inhomogeneous space. on the other hand. in general we do not stand atthe center Yo. so the observed Dirichlet domain D(Xop.) can vary from the fundamental cell DC).," In a globally inhomogeneous space, on the other hand, in general we do not stand atthe center $\vec x_0 $ , so the observed Dirichlet domain $D(\vec x_{\mathrm{obs}})$ can vary from the fundamental cell $D(\vec x_{\mathrm{0}})$ ."1084Dust is an important constituent of interstellar medium (ISAD. molecular clouds iud accretion disks (see Whüttet 2003. Draine 2009).,"Dust is an important constituent of interstellar medium (ISM), molecular clouds and accretion disks (see Whittet 2003, Draine 2009)."1085 It ects involved im many kev processes. for imstance. it controls heating and cooling of the ΤΟΝΤ (see. Draine 2003. Tieliens 2005). reveals magnetic fields through eram aligunmieut (see Lazariau 2007 for a review) aud interferes with the attempts to measure properties of CAIB radiation (see Lazarian Fiukbeiuer 2003. Fraisse et al.," It gets involved in many key processes, for instance, it controls heating and cooling of the ISM (see Draine 2003, Tieliens 2005), reveals magnetic fields through grain alignment (see Lazarian 2007 for a review) and interferes with the attempts to measure properties of CMB radiation (see Lazarian Finkbeiner 2003, Fraisse et al."1086 2009)., 2009).1087 Small. ie. less than ~10? cm erains. are an iurportaut component of the interstellar dust population. with a notable fraction of very small erains Polveyclie Avomatic Uvdrocarborn (PAID) particles. =which are esscutially large (sce Leger Puget 1981).," Small, i.e., less than $\sim10^{-5}$ cm grains, are an important component of the interstellar dust population, with a notable fraction of very small grains – Polycyclic Aromatic Hydrocarborn (PAH) particles, which are essentially large (see Leger Puget 1984)."1088 Iu what follows. due to the reasons that are explained iu 85. we do not directly address PAIT particles. but our approach to the siall erain acceleration may be exteuded to this iuportaut population of eraius.," In what follows, due to the reasons that are explained in 5, we do not directly address PAH particles, but our approach to the small grain acceleration may be extended to this important population of grains."1089 Most properties of grains. ποιαπιο light extinction. electron pliotoenmission. aud chemical activity depend not oulv on erain chemical composition. but also on their sizes.," Most properties of grains, including light extinction, electron photoemission, and chemical activity depend not only on grain chemical composition, but also on their sizes."1090 In astroplivsical media. these sizes are affected by erain-erain collisions.," In astrophysical media, these sizes are affected by grain-grain collisions."1091 The minimal velocities of erains are determined by their Brownian motion corresponding to the temperature of the ambicut eas., The minimal velocities of grains are determined by their Brownian motion corresponding to the temperature of the ambient gas.1092 Large-scale lvdvodvuamic motions associated with turbulence cau make eraius move faster (sce Draine 1985)., Large-scale hydrodynamic motions associated with turbulence can make grains move faster (see Draine 1985).1093 Since most astrophysical inedia. are magnetized and eraius are charged the hywdrodyuanüc tfreatineut of acceleration ds frequently uot adequate., Since most astrophysical media are magnetized and grains are charged the hydrodynamic treatment of acceleration is frequently not adequate.1094 A proper treatment of the erain acceleration through the interaction of charged erains with maguetolivdrodyuianic (ANID) turbulence las heen developed recently (Lazavian Yan 2002. Yan Lazarian 2003.—- Yan. Lazariau Draine 2001. Yan 2009).," A proper treatment of the grain acceleration through the interaction of charged grains with magnetohydrodynamic (MHD) turbulence has been developed recently (Lazarian Yan 2002, Yan Lazarian 2003, Yan, Lazarian Draine 2004, Yan 2009)."1095 This troatineut makes extensive use of the adwauces of compressible ATID turbulence (see Cho Lazarian 2002. 2003) aud provides the mathematical formalisia of the secoud-order Ferma acceleration of charged eraius interacting with MIID turbulence.," This treatment makes extensive use of the advances of compressible MHD turbulence (see Cho Lazarian 2002, 2003) and provides the mathematical formalism of the second-order Fermi acceleration of charged grains interacting with MHD turbulence."1096 However. the acceleration mechanisuis based on the AIIID interaction of turbulence and. charged. graius exhibit acceleration rates that decrease as dust particles ect smaller.," However, the acceleration mechanisms based on the MHD interaction of turbulence and charged grains exhibit acceleration rates that decrease as dust particles get smaller."1097 This decrease arises from the fact that the Larinor radius of charged erains becomes sinaller witli the decrease of erain nass and. correspondingly. erains have to interact with smaller. 10. less powerful. turbulent fluctuations.," This decrease arises from the fact that the Larmor radius of charged grains becomes smaller with the decrease of grain mass and, correspondingly, grains have to interact with smaller, i.e., less powerful, turbulent fluctuations."1098 lu addition. compressible fluctuations. boc. fast dgnodes. which were identified i Yan Lazarian (2003) with the most cficient acceleration. ect suppressed at the small scales due to plasima damipine. while the Alfvenic mode gets inefficient for acceleration at sinall scale due to anisotropy (sce Yan Lazarian 2003 for more discussion).," In addition, compressible fluctuations, i.e., fast modes, which were identified in Yan Lazarian (2003) with the most efficient acceleration, get suppressed at the small scales due to plasma damping, while the Alfvenic mode gets inefficient for acceleration at small scale due to anisotropy (see Yan Lazarian 2003 for more discussion)."1099 The acceleration oferaius with sizes less than ~10.7 cin becomes rather inefficient for most media discussed in Yau Lazarian (2003). Yan. Lazarian Draine (2001. henceforth YLDO1).," The acceleration of grains with sizes less than $\sim10^{-5}$ cm becomes rather inefficient for most media discussed in Yan Lazarian (2003), Yan, Lazarian Draine (2004, henceforth YLD04)."1100 Ave there other mechanisius ofdust acceleration which dominate for grains snialler than 10 cn?, Are there other mechanisms of dust acceleration which dominate for grains smaller than $10^{-5}$ cm?1101 One can expect that dust-plasima iuteractious may be important for such erains., One can expect that dust-plasma interactions may be important for such grains.1102 This paper prescuts a novel promising acceleration mechauisi based on erain-grain C'ouloiib collisions in- the preseuce of erai charge fluctuations., This paper presents a novel promising acceleration mechanism based on grain-grain Coulomb collisions in the presence of grain charge fluctuations.1103 This iiechauisi utilizes intrinsic non-equilibriuni nature of the dusty astrophvsical plasmas., This mechanism utilizes intrinsic non-equilibrium nature of the dusty astrophysical plasmas.1104 Iu additiou to acting ou the small erains. the mechanism may drive the acceleration of larger erains whenever the MIID acceleration mechanisia is suppressed. (e.g.. when MIID turbulence is camped).," In addition to acting on the small grains, the mechanism may drive the acceleration of larger grains whenever the MHD acceleration mechanism is suppressed (e.g., when MHD turbulence is damped)."1105 Iu what follows. we introduce major timescales characterizing dynamics of charged interstellar eraius in 822. present the mechanisia of sinall eram acceleration associated with charge fluctuations during their Coulonib collisions iu 8233. calculate the acceleration for graius of a given size distribution in ll. analyze iuplicatious of the effect for interstellar phases im 855. discuss the iuportauce of our results in 866. aud sunuuanze them im," In what follows, we introduce major timescales characterizing dynamics of charged interstellar grains in 2, present the mechanism of small grain acceleration associated with charge fluctuations during their Coulomb collisions in 3, calculate the acceleration for grains of a given size distribution in 4, analyze implications of the effect for interstellar phases in 5, discuss the importance of our results in 6, and summarize them in"1106We determine the plasma pressure at every point. by calculating the path of the field line through that point and solving for isothermal. hydrostatic equilibrium along that »»h.,"We determine the plasma pressure at every point by calculating the path of the field line through that point and solving for isothermal, hydrostatic equilibrium along that path."1107 In this case. the gas pressure is p=prattes where m is the mean particle mass. Ae is DBolzmann's constant. T is the temperature and p=po is the eas oessure at the base of the field line.," In this case, the gas pressure is $p=p_{0}e^{\frac{m}{k_B T}\int g_{s}ds}$ where $m$ is the mean particle mass, $k_B$ is Bolzmann's constant, T is the temperature and $p=p_{0}$ is the gas pressure at the base of the field line."1108 We note that the integral in this expression is performed. along the path of he field line and that ος=(g.B)/|B| is the component of eravity along the field., We note that the integral in this expression is performed along the path of the field line and that $g_{s} =( {\bf g.B})/|{\bf B}|$ is the component of gravity along the field.1109 We note that the plasma pressure is set to zero at any point where the field. line. through hat point experiences a plasma pressure greater than the magnetic pressure somewhere along its length., We note that the plasma pressure is set to zero at any point where the field line through that point experiences a plasma pressure greater than the magnetic pressure somewhere along its length.1110 In this case. we assume that this field line should have been forced open o» the pressure of the plasma.," In this case, we assume that this field line should have been forced open by the pressure of the plasma."1111 The gas pressure at the ootpoint of the field line po is a free parameter of this model., The gas pressure at the footpoint of the field line $p_0$ is a free parameter of this model.1112 Following Jardineetal.(2002a.b) we choose to scale po o the magnetic pressure at the base of the field line. such hat po=ADS where A is a constant that is the same on every field. line.," Following \citet{JardineM:2002a,JardineM:2002b} we choose to scale $p_{0}$ to the magnetic pressure at the base of the field line, such that $p_{0}=K B^{2}_{0}$ where $K$ is a constant that is the same on every field line."1113 By sealing A up or down we can scale he overall level of the coronal gas pressure ancl hence the density and emission measure., By scaling $K$ up or down we can scale the overall level of the coronal gas pressure and hence the density and emission measure.1114 Our model therefore has two xwameters: the radius at which the field lines are opened up (the surface) and the constant Av which determines he gas pressure po at the base of each field line., Our model therefore has two parameters: the radius at which the field lines are opened up (the ) and the constant $K$ which determines the gas pressure $p_0$ at the base of each field line.1115 These two xwameters determine the magnetic field. structure ane X-rav emission measure of the closed-field regions of the stellar corona., These two parameters determine the magnetic field structure and X-ray emission measure of the closed-field regions of the stellar corona.1116 These are given in Fig., These are given in Fig.1117 3. for the three epochs., \ref{cor_struct} for the three epochs.1118 The Hao line in active solar-tvpe stars is often used as an activity indicator with the line being lled-in in more active stars (Le.Soderblometal.1993)., The $\alpha$ line in active solar-type stars is often used as an activity indicator with the line being “filled-in” in more active stars \citep[i.e.][]{SoderblomDR:1993}.1119.. In. addition. prominence activity in the stellar chromosphere can also be mapped on such stars that show emission in the Ho line (ie.Cameron&Robinson1989:Donatietal. 2000).," In addition, prominence activity in the stellar chromosphere can also be mapped on such stars that show emission in the $\alpha$ line \citep[i.e.][]{CameronAC:1989, DonatiJF:2000}."1120. We rave analvsed the Ho. line of HD. 141943 in March/April 2007 ancl Alarch/April 2010 as the most complete datasets ο look for possible variations that may be attributable to prominence activity in the stellar chromosphere., We have analysed the $\alpha$ line of HD 141943 in March/April 2007 and March/April 2010 as the most complete datasets to look for possible variations that may be attributable to prominence activity in the stellar chromosphere.1121 As shown in Fig. 4..," As shown in Fig. \ref{Fig_ha},"1122 LID 141943 is a more active star than the Sun with a variable level of activity in 2007. but it cloes not have he overt Ho. emission of a voung T Tauri star.," HD 141943 is a more active star than the Sun with a variable level of activity in 2007, but it does not have the overt $\alpha$ emission of a young T Tauri star."1123 In order to look for prominence activity on LID 141943 we have divided cach of the Ho. profiles in the MarchApril 2007 dataset by the mean La profile from the dataset., In order to look for prominence activity on HD 141943 we have divided each of the $\alpha$ profiles in the March/April 2007 dataset by the mean $\alpha$ profile from the dataset.1124 We did the same for the March/Xpril 2010 dataset. clivicling by the average Ho. profile from the 2010 dataset.," We did the same for the March/April 2010 dataset, dividing by the average $\alpha$ profile from the 2010 dataset."1125 Dynamic spectra of these are displayed in Fig., Dynamic spectra of these are displayed in Fig.1126 5. (2007 on the left and 2010 on the right) with darker areas showing regions of lower activity and lighter areas higher activity., \ref{Fig_dynha} (2007 on the left and 2010 on the right) with darker areas showing regions of lower activity and lighter areas higher activity.1127 As can be seen the Πα profile of LID 121943. in 2007 is variable in its activity level with regions of higher activity located. between phases 20.1 to ~0.6 while the other phases have lower activity than average., As can be seen the $\alpha$ profile of HD 141943 in 2007 is variable in its activity level with regions of higher activity located between phases $\sim$ 0.1 to $\sim$ 0.6 while the other phases have lower activity than average.1128 There would appear to be some evolution in the velocity of these active/non-active regions with most of the motion restricted to within the rotational velocity of the star. although there is some enhancement/reduction of the Πα emission outside these velocities.," There would appear to be some evolution in the velocity of these active/non-active regions with most of the motion restricted to within the rotational velocity of the star, although there is some enhancement/reduction of the $\alpha$ emission outside these velocities."1129 Given the inclination angle of LID 1431943. wominences located: at the co-rotation radius of the star are likely to be seen as absorption features quickly crossing he range of the star.," Given the inclination angle of HD 141943, prominences located at the co-rotation radius of the star are likely to be seen as absorption features quickly crossing the range of the star."1130 Phe absorption features seen in Fig., The absorption features seen in Fig.1131 5. around. phase «0.50 appear to travel across the stellar profile. but. not rapidly.," \ref{Fig_dynha} around phase $\sim$ 0.85 appear to travel across the stellar profile, but not rapidly."1132 A sine wave fitted to the vcak/troughs of the Le emission shows the amplitude of he sine wave to be restricted. to within the of LD 141943., A sine wave fitted to the peak/troughs of the $\alpha$ emission shows the amplitude of the sine wave to be restricted to within the of HD 141943.1133 Thus we believe that this region is located fairly close to the surface of the star and appears to be an inactive (in Hla) low-to-mid latitude feature on. or near. the stellar surface.," Thus we believe that this region is located fairly close to the surface of the star and appears to be an inactive (in $\alpha$ ) low-to-mid latitude feature on, or near, the stellar surface."1134 Comparing this to the X-ray. emission image right image in Fie. 3)), Comparing this to the X-ray emission image (top-right image in Fig. \ref{cor_struct}) )1135 we see that for phases around. 0.0 here is little N-rav emission from the star., we see that for phases around 0.0 there is little X-ray emission from the star.1136 For 2010 there appears to be little or no variation in he Lla emission of LID 1431943., For 2010 there appears to be little or no variation in the $\alpha$ emission of HD 141943.1137 Ehe level of Lla emission at all phases in 2010 is similar to the highest level of emission in 2007., The level of $\alpha$ emission at all phases in 2010 is similar to the highest level of emission in 2007.1138 Thus it would appear that the Lla feature seen in 2007 is not present on the star in 2010., Thus it would appear that the $\alpha$ feature seen in 2007 is not present on the star in 2010.1139 As mentioned in Paper lL. HD. 141943 is only the second. (or2010) voung early-G star for which the large-scale magnetic topology has been determined. the other being HD. 171488 (Marsdenetal.2006:Jellers&Donati2008:οἱ 2010).," As mentioned in Paper I, HD 141943 is only the second \citep[or third including the results for HD 106506 by][]{WaiteIA:2010} young early-G star for which the large-scale magnetic topology has been determined, the other being HD 171488 \citep{MarsdenSC:2006, JeffersSV:2008, JeffersSV:2010}."1140. In total there have been five voung carly-C stars for which cülferential rotation measures have been determined., In total there have been five young early-G stars for which differential rotation measures have been determined.1141 The five stars are HD. 141943. LID 171488. LID 106506. R5s (Marsdenοἱal.2005a.b) and LO Lup (Donatietal.2000).," The five stars are HD 141943, HD 171488, HD 106506, R58 \citep{MarsdenSC:2005a, MarsdenSC:2005b} and LQ Lup \citep{DonatiJF:2000}."1142. The stellar parameters for all five stars have been given in Paper L. However. for ease of comparison we have given these in Table 3. and included the cillerential rotation results of the five stars.," The stellar parameters for all five stars have been given in Paper I. However, for ease of comparison we have given these in Table \ref{Tab_gstars} and included the differential rotation results of the five stars."1143 The differential rotation rate found for HD. 14143 is one of he largest vet found using the Doppler imaging method and is similar to that of the other voung earlv-CG star HD 171488 (Marsdenetal.2006:Jeffers&Donati2008: 2010).," The differential rotation rate found for HD 14143 is one of the largest yet found using the Doppler imaging method and is similar to that of the other young early-G star HD 171488 \citep{MarsdenSC:2006, JeffersSV:2008, JeffersSV:2010}."1144. Lt is also similar in level to that of the more mature xanet-hosting late-F star Tau Boo (Donatietal.2008:Faresetal.2009). and is in agreement with the findings of high evels of dillerential rotation on inactive E stars by Reiners)06) using the line-profile method to measure differential rotation.," It is also similar in level to that of the more mature planet-hosting late-F star Tau Boo \citep{DonatiJF:2008, FaresR:2009} and is in agreement with the findings of high levels of differential rotation on inactive F stars by \citet{ReinersA:2006} using the line-profile method to measure differential rotation."1145 However. the level of differential rotation for LED 141943 and LID 171488 are significantly above that of the οier voung earlv-Ci stars studied using the Doppler imaging method (see Table 3)).," However, the level of differential rotation for HD 141943 and HD 171488 are significantly above that of the other young early-G stars studied using the Doppler imaging method (see Table \ref{Tab_gstars}) )."1146 There appears to be little differences between the stars to explain the dilferences in. differential rotation with the possible exception of their convective zone depth., There appears to be little differences between the stars to explain the differences in differential rotation with the possible exception of their convective zone depth.1147 Barnesetal.(2005) have shown that generally (LEferential rotation increases with stellar elfective temperature. however. their dataset. does not have any μαtees with a cillerential rotation greater than dQ — 0.2 rad 5," \citet{BarnesJR:2005} have shown that generally differential rotation increases with stellar effective temperature, however, their dataset does not have any stars with a differential rotation greater than $d\Omega$ $\sim$ 0.2 rad ."1148 As an extension of the work of Barnesetal.(2005) we have plotted the differential rotation measured: using the (Zeeman) Doppler imaging technique for these stars plus the aclelitional stars listed in Table 3 (11D141943.," As an extension of the work of \citet{BarnesJR:2005} we have plotted the differential rotation measured using the (Zeeman) Doppler imaging technique for these stars plus the additional stars listed in Table \ref{Tab_gstars} (HD141943,"1149Because complex plasma interactions among CRs. waves. and the underlving gas flow are not fully understood vet. it is not possible to make a precise quantitative prediction for the injection. process from first principles (e.g..Malkov&Drury2001).,"Because complex plasma interactions among CRs, waves, and the underlying gas flow are not fully understood yet, it is not possible to make a precise quantitative prediction for the injection process from first principles \citep[e.g.,][]{maldru01}."1150. Here. we adopt a phenomenological injection model that can emulate the thermal leakage process. through which particles above a certain injection monientum pij cross the shock and get injected to the CR population (Ixangetal.2002:INang&Raw2010).," Here, we adopt a phenomenological injection model that can emulate the thermal leakage process, through which particles above a certain injection momentum $p_{\rm inj}$ cross the shock and get injected to the CR population \citep{kjg02, kr10}."1151. Then. the CHR. distribution [unction at pij is anchored to the downstream Maxwellian distribution as where n3 is the downstream proton number density.," Then, the CR distribution function at $p_{\rm inj}$ is anchored to the downstream Maxwellian distribution as where $n_2$ is the downstream proton number density."1152 Here. pi; aid Qing ave delined as where pin—/2mykpT» is the thermal peak momentum of the downstream gas with temperature 75 and fy isthe Boltzmann constant.," Here, $p_{\rm inj}$ and $Q_{\rm inj}$ are defined as where $p_{\rm th}= \sqrt{2m_p k_B T_2}$ is the thermal peak momentum of the downstream gas with temperature $T_2$ and $k_B$ isthe Boltzmann constant."1153" We note that the functional form οἱ Qu;cinj was adopted to represent an ""effective"" injection momentiun. since particles in the ⊳∖⊽∏↕↽≻↕⋅≀↧↴⊔∐↲↕⋅∐↓≀↧↴↥↥≀↧↴∐≺∢≀↧↴∐≺∢↕⋅∪⊳∖⇁⊳∖⊽⊔∐↲⋝∖⊽∐∪≺∢↳↽∖∖↽↕⊔↥≀↧↴⊳∖⇁∐⋯∪↥⊔⋡∖↽−∖↽≀↧↴↕⋅∡∖⇁↕∐≸≟↕↽≻↕⋅∪∣↽≻≀↧↴∣↽≻∐∐∡"," We note that the functional form of $Q_{\rm inj}$ was adopted to represent an “effective” injection momentum, since particles in the suprathermal tail can cross the shock with a smoothly-varying probability distribution \citep[see][]{kjg02}. ."1154∖↽≼∐⊳∖⊽∏⋅↕∣↽≻∏∐∪∐≼⋝⊳∖⊽≼↲≼↲ ↕∖⊽≀⋯∩≼↲↥≀↧↴↥⋅⊋∪∩⊋⋝⋝⋅⋅≼↽≽∐≼↲∐⋅≼↲≼↲↽≻≀↧↴↕⋅≀↧↴∐∐↲∩↲↕⋅⊔⋯↥≺∢∪∐∏⋅∪↥⊳∖⇁⊔∐↲↥≼↲≀↧↓≶≀↧↴∩≼↲↽≻↕⋅∪≺∢≼↲⊳∖⊽⊳∖⇁↕⊳∖⇁⊔∐↲↕∐∎≼↲≺∢∐∪∐  ↕↽≻≀↧↴↕⋅≀↧↴∐∐↲∩↲↕⋅⋅←∐∶∐∣∣∕∕∕∐≓⋅∖∖↽∐↥≺∢∐↕⊳∖⊽⊔∐↲↕⋅≀↧↴∐∪∪↓⋟⊔∐↲≸↽↔↴≼↲∐≼↲↕⋅≀↧↴↥∐↓≀↧↴≸↽↔↴∐≼↲∐≺∢∐≼↲↥≺⇂≀↧↴↥∪∐≸≟⊔∐↲⊳∖⇁∐∪≺∢↥≶ ∐∪↕⋅∐⋯↴↥⋅∐∣∣⋅↥∪⊔∐↲≀↧↴∐↓↕↽≻∐↥⋯⇂≼↲∪↓⋟⊔∐↲≼⇂∪∖∖⇁∐⋟∖⊽⊔⋅≼↲≀↧↴↕⊔⋅∐⋯↴↖≺↽↔↴∐≼↲↥∪∐∡∖↽≼∐⋅⋯⇂⋡∖⇁∐≀↧↴∐↓↕≺∢⋖⋡↼∖∐∐⊃↕⋟∖∖⇁≀↧↴∖↽≼↲⊓∐⋅∣↽≻∏↥≼↲∐≺∢≼↲⋅ DB.," One free parameter that controls the leakage process is the injection parameter, $\epsilon_B = B_0/B_{\perp}$ which is the ratio of the general magnetic field along the shock normal, $B_0$, to the amplitude of the downstream, magnetohydrodynamic (MHD) wave turbulence, $B_{\perp}$."1155 Although5 plasma hybrid simulations aud theories both suggestedo5 that 0.25eoXegtB0.35 (Alalkov&Volk1998).. the physical range of this parameter remains to be rather uncertain due to lack of full understanding of relevant plasma interactions.," Although plasma hybrid simulations and theories both suggested that $0.25 \la \epsilon_B \la 0.35$ \citep{mv98}, the physical range of this parameter remains to be rather uncertain due to lack of full understanding of relevant plasma interactions."1156 The second term in Equation (6)) is fixed bv q. pi. and finj.," The second term in Equation \ref{drury}) ) is fixed by $q$, $p_{\rm inj}$, and $f_{\rm inj}$ ."1157 The fraction of particles injected into the CR population can be estimated analvlically as well: which is fixed only byQj and q., The fraction of particles injected into the CR population can be estimated analytically as well: which is fixed only by$Q_{\rm inj}$ and $q$ .1158 The injection fraction depends strongly on ej (through (inj) lor weak shocks with AZ5 (seealsoWang&Ryu 2010).., The injection fraction depends strongly on $\epsilon_B$ (through $Q_{\rm inj}$ ) for weak shocks with $M \la 5$ \citep[see also][]{kr10}. .1159 Forexample.it variesfrom 5x10? to 10? forey=0.25—0.3 for shocks with M= 3. K," Forexample,it variesfrom $5\times10^{-5}$ to $10^{-3}$ for$\epsilon_B=0.25-0.3$ for shocks with $M=3$ ."1160ane&Ryu(2010)demonstrated that the /me-dependent. test-particle solutions of the," \citet{kr10} demonstrated that the , test-particle solutions of the"1161The nature of the SW turbulence below the ion scale (typically p;100 kan. corresponding to an observed frequency in the spacecraft frame of f0.5 IIz) has attracted considerable iutercst in the space and astrophysical comiuuitics in recent vears.,"The nature of the SW turbulence below the ion scale (typically $\rho_i \sim 100$ km, corresponding to an observed frequency in the spacecraft frame of $f_{sc} \sim 0.5$ Hz) has attracted considerable interest in the space and astrophysical communities in recent years."1162 This has Όσσα encouraged in particular by recent observations frou. the Cluster iissiou that provided the most complete and detailed picture of the SW turbulence cascade from maguctolivdrodvuamic (MIID) scales (2>> p;} to electrou scales (EL~p.1 kin) (Salraouietal.2009:2011:Salraouietal. 2010a).," This has been encouraged in particular by recent observations from the Cluster mission that provided the most complete and detailed picture of the SW turbulence cascade from magnetohydrodynamic (MHD) scales $L>>\rho_i$ ) to electron scales $L\sim\rho_e\sim 1$ km) \citep{sahraoui09,kiyani09,alexandrova09,chen10,sahraoui10a}."1163. Determining the nature aud properties (o.g.. scaling. anisotropy) of the turbulence at sanall scales is indeed a crucial point to uuderstaudiug the problems of cherey dissipation aud heating. particle acceleration. and magnetic reconnection dn space and astrophysical plasmas (Schekochihinetal.2009).," Determining the nature and properties (e.g., scaling, anisotropy) of the turbulence at small scales is indeed a crucial point to understanding the problems of energy dissipation and heating, particle acceleration, and magnetic reconnection in space and astrophysical plasmas \citep{schekochihin09}."1164. Recent Cluster observations provided clear evidence that SW turbulence cascades below the ion scale p; down to the electron scale pi where dissipation becomes iaportaut aud the spectra stecpen to f. “witha2 L(Sahraouietal.2009. 2010a).," Recent Cluster observations provided clear evidence that SW turbulence cascades below the ion scale $\rho_i$ down to the electron scale $\rho_e$ where dissipation becomes important and the spectra steepen to $\sim {f_{sc}}^{-\alpha}$, with $\alpha \gtrsim 4$ \citep{sahraoui09,sahraoui10a}."1165. The spectrum thus formed was termed the while the range of scales between p; aud p. las been termed the iu reference to the dispersive uature of the plasma modes at those scales (Stawickietal.2001)., The spectrum thus formed was termed the while the range of scales between $\rho_i$ and $\rho_e$ has been termed the in reference to the dispersive nature of the plasma modes at those scales \citep{stawicki01}.1166.. In a case study. Saliraouietal.(20108). performed a detailed analysis of the enuergv cascade from ΑΠΟ to. sub-on scales ming a nmldtipoiut imeasurenient techuique called ᾖ-filtering (Pingou&Lefeuvre1991:Saliraouietal.20032.2010b:Naritaetal.2010a.b:Tjuliux 2005).," In a case study, \cite{sahraoui10a} performed a detailed analysis of the energy cascade from MHD to sub-ion scales using a multipoint measurement technique called }-filtering \citep{pincon91,sahraoui03a, sahraoui10b, narita10a, narita10b, tjulin05}."1167". The results showed clearly that: 1) The maguetic turbulence is 19strongly anisotropic (ej>> hy) down to the observed scale. ερ~2: un) Tre cascade d8consistent wd[η KAW turbuleuce as pro]οσο in Towesetal.(20084) and Schekochihinetal.(2009) with frequencies in f Plasma rest frame ο=Ole; although the frequencies in je spacecraft frame reached 20w,.;: 3i) The turbulence uudergocs arange near the ion scale p; characterized by a steepeniug of the spectrum frou &9 ο kok which has been interpreted as due to Laudau damping of maenetic energy iuto ion heating (owesot. 2008bj."," The results showed clearly that: i) The magnetic turbulence is strongly anisotropic $k_\perp>>k_\parallel$ ) down to the observed scale $k_\perp\rho_i\sim 2$; ii) The cascade isconsistent with KAW turbulence as proposed in \cite{howes08a} and \cite{schekochihin09} with frequencies in the plasma rest frame $\omega \lesssim0.1\omega_{ci}$ although the frequencies in the spacecraft frame reached $20\omega_{ci}$; iii) The turbulence undergoes a near the ion scale $\rho_i$ characterized by a steepening of the spectrum from $ k_\perp^{-1.6}$ to $k_\perp^{-4.5}$, which has been interpreted as due to Landau damping of magnetic energy into ion heating \citep{howes08b}."1168. The remaining cucrev cascades following a power-law ~£277 down to the electrou scale where the jergv was sugeested to dissipate iuto electron heating., The remaining energy cascades following a power-law $\sim f_{sc}^{-2.8}$ down to the electron scale where the energy was suggested to dissipate into electron heating.1169 The wave απλα spectra at those frequencies could rot have been directly measured using the A-filteriug echuique due to the limitation iniposed by the satellites separations. which are larger than 100 ka," The wave number spectra at those frequencies could not have been measured using the -filtering technique due to the limitation imposed by the satellites separations, which are larger than $100$ km"1170 The wave απλα spectra at those frequencies could rot have been directly measured using the A-filteriug echuique due to the limitation iniposed by the satellites separations. which are larger than 100 kaw," The wave number spectra at those frequencies could not have been measured using the -filtering technique due to the limitation imposed by the satellites separations, which are larger than $100$ km"1171in the asviuptotic giant brauch. (ACB).,in the asymptotic giant branch (AGB).1172 However. the discrepancy becomes larecr for more massive stars where the ratio can reach —100 in the AGB phase (Clirbounel 1995: Weiss et al.," However, the discrepancy becomes larger for more massive stars where the ratio can reach $\sim$ 100 in the AGB phase (Charbonnel 1995; Weiss et al."1173 1996: Forestini Charbonnel 1997. hereafter FC: van den Toek Grocneweecn 1997. lerafter HIC: Alarigo 1998: Doothrovd Sackinaun 1999. hereafter DS).," 1996; Forestini Charbonnel 1997, hereafter FC; van den Hoek Groenewegen 1997, herafter HG; Marigo 1998; Boothroyd Sackmann 1999, hereafter BS)."1174 From an observational viewpoiut. it is important to obtain accurate measuremients of the isotopic ratio in those PNe where the Πο abundance has Όσοι determined.," From an observational viewpoint, it is important to obtain accurate measurements of the isotopic ratio in those PNe where the $^3$ He abundance has been determined."1175" Should+ these objects. show a high- value of. οςP ALBο, then no modifications to the standard stellar models would be required."," Should these objects show a high value of $^{12}$ $^{13}$ C, then no modifications to the standard stellar models would be required."1176 Otherwise. one has to invoke another selective process (nmüxius. diffusion etc.)," Otherwise, one has to invoke another selective process (mixing, diffusion etc.)"1177" that operates on some isotopes bu not ou ""Πο.", that operates on some isotopes but not on $^3$ He.1178 However. the παρ of PNe with Πο measurements is small (e.g. Balser et al.," However, the number of PNe with $^3$ He measurements is small (e.g. Balser et al."1179 1999). whereas the suggestedOO plavsical processes should be quite eeneral aud should affect the uucleosvuthetie vields of all stars of mass less than ~2M...," 1999), whereas the suggested physical processes should be quite general and should affect the nucleosynthetic yields of all stars of mass less than $\sim 2$."1180 Hence. it is critical to measure the carbon isotopic ratio in a sample of PNe as large as possible.," Hence, it is critical to measure the carbon isotopic ratio in a sample of PNe as large as possible."1181 The molecular cuvelopes of PNe have been studied extensively at near mfrared aud millimeter wavelengths (see e.c. [xastuer et al., The molecular envelopes of PNe have been studied extensively at near infrared and millimeter wavelengths (see e.g. Kastner et al.1182 1996. Ilugeis οἳ al.," 1996, Huggins et al."1183 1996. Bachiller et al.," 1996, Bachiller et al."1184 1997)., 1997).1185 These observations have shown that massive cuvelopes (Z10D7 M.) coutainug a rich variety of molecular species are comuiuonlv found around PNe., These observations have shown that massive envelopes $\simgreat 10^{-2}$ ) containing a rich variety of molecular species are commonly found around PNe.1186 CO is the most widely observed species. aud the isotopic ratio has been measured toward several PNe (Bachiller et al.," CO is the most widely observed species, and the isotopic ratio has been measured toward several PNe (Bachiller et al."1187 1989. 1997: Cox et al.," 1989, 1997; Cox et al."1188 1992)., 1992).1189 These initial studies have shown that the ratio is iu the rauee 1020., These initial studies have shown that the ratio is in the range 10–20.1190 Oi project consists of two parts., Our project consists of two parts.1191" In the first one. we rave caaricd out high quality observations of °CO aud ""CO insix PNe that have been searched for ""Ie emission."," In the first one, we have carried out high quality observations of $^{12}$ CO and $^{13}$ CO in six PNe that have been searched for $^3$ He emission."1192 Iu the second run. a larger sample of uchulae with strong ?(O0 line enission lias been observed iu the CO lines in order to determine the isotopic ratio iu PNe Πο neasurenients.," In the second run, a larger sample of nebulae with strong $^{12}$ CO line emission has been observed in the $^{13}$ CO lines in order to determine the isotopic ratio in PNe $^3$ He measurements."1193 Galli et al. (, Galli et al. (1194"1997. hereafter GSTP) have argued that extra-uixiung processes must be at work iu nore than of low-mass stars CM2 M.) in order o reconcile the predictions of the Galactic evolution of Πο with the obscrvationa constraints,","1997, hereafter GSTP) have argued that extra-mixing processes must be at work in more than of low-mass stars $M\simless 2$ ) in order to reconcile the predictions of the Galactic evolution of $^3$ He with the observational constraints."1195 We set out this experiaueut to determine the isotopic ratio in a relatively aree saluple of PNe., We set out this experiment to determine the isotopic ratio in a relatively large sample of PNe.1196 The observations were carried out with the IRAM 3041 telescope at Pico Veleta (near Canada. Spain) duriug two observing rus in November 1996 aud May. 1997.," The observations were carried out with the IRAM 30-m telescope at Pico Veleta (near Granada, Spain) during two observing runs in November 1996 and May 1997."1197 In. the first one we observed the six PNe studied by Balser et al. (, In the first one we observed the six PNe studied by Balser et al. (11981997).,1997).1199 The observations were mace simultaneously in the J—2 Laud 10 lines of CO., The observations were made simultaneously in the $J = 2$ –1 and 1–0 lines of $^{12}$ CO.1200 The strongest cuitters were then observed in the CO lines., The strongest emitters were then observed in the $^{13}$ CO lines.

Showing the first 1,200 of 10445 lines. Download the file for the rest.