ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 In Section 6 we discuss the properties of the overdensities. and the colours of the galaxies within the 3 protocluster candidates.," In Section \ref{discussion} we discuss the properties of the overdensities, and the colours of the galaxies within the 3 protocluster candidates."3 Fluxes are calibrated using the Vega magnitude scale unless noted otherwise., Fluxes are calibrated using the Vega magnitude scale unless noted otherwise.4 A flat ACDM cosmology is assumed throughout. with Όλι=0.3. Q4—0.7 and Hy=70kms |.," A flat $\Lambda$ CDM cosmology is assumed throughout, with $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$ and $_0$ $^{-1}$."5" The distance scale at z=24 isKKpe/"".. or ~ eco-movingMpc/""."," The distance scale at $z=2.4$ is, or $\sim$ co-moving."6. Distances are given in co-moving units. unless stated otherwise.," Distances are given in co-moving units, unless stated otherwise."7 Six HzRGs with redshifts between 2.28 and 2.55 were selected from the compendium of ?.., Six HzRGs with redshifts between 2.28 and 2.55 were selected from the compendium of \citet{MileydeBreuck2008}.8 The targets were chosen based on their distributions in right ascension. declination. and redshift. and their bright radio luminosity «ΩΜΗ;L089 . and not on any prior information about surrounding galaxy overdensities.," The targets were chosen based on their distributions in right ascension, declination, and redshift, and their bright radio luminosity $_{500 \rm{MHz}}>10^{28.5}$ $^{-1}$ ), and not on any prior information about surrounding galaxy overdensities."9 Co-ordinates of the selected targets and the control field are given in Table Ι.., Co-ordinates of the selected targets and the control field are given in Table \ref{tab:obs}.10 The 6 HzRG fields and a blank control field were observed in service mode using the High Acuity Wide field K-band Imager (HAWK-I: 23) on the ESO Very Large Telescope (VLT) YEPUN telescope in Paranal. Chile during the period April-September 2008.," The 6 HzRG fields and a blank control field were observed in service mode using the High Acuity Wide field K-band Imager (HAWK-I; \citealt{Kissler-Patig2008}) ) on the ESO Very Large Telescope (VLT) UT4-YEPUN telescope in Paranal, Chile during the period April--September 2008."11" HAWRK-I is a near-infrared camera comprising of four Hawaii-2 2048x2048 pixel detectors separated by a gap of 7-5"".", HAWK-I is a near-infrared camera comprising of four Hawaii-2 2048x2048 pixel detectors separated by a gap of $\sim$.12 The camera spans 7.5x aaremin with a pixel scale of 0.106 aresee per pixel., The camera spans $7.5\times7.5$ arcmin with a pixel scale of 0.106 arcsec per pixel.13 The telescope pointing was optimised so that the HZRG were placed close to the centre of the HAWK-I field of view. but at least aaremin away from the chip gaps. and bright stars did not fall within the field of view.," The telescope pointing was optimised so that the HzRG were placed close to the centre of the HAWK-I field of view, but at least arcmin away from the chip gaps, and bright stars did not fall within the field of view."14 Each target was observed through theJ.. aand ‘filters and total exposure times for all targets are provided in reftab:obs..," Each target was observed through the, and filters and total exposure times for all targets are provided in \\ref{tab:obs}."15 The telescope dithered every mmins which resulted in a I aaremin-wide cross-shaped region on the final mosaic where he image depth is shallower., The telescope dithered every mins which resulted in a $\sim$ arcmin-wide cross-shaped region on the final mosaic where the image depth is shallower.16 During the observing period the anti-reflection coating of he Dewar window of HAWK-I was damaged resulting in small cross-shaped patterns on the exposures., During the observing period the anti-reflection coating of the Dewar window of HAWK-I was damaged resulting in small cross-shaped patterns on the exposures.17 These patterns. caused by the spider of the secondary mirror. rotated as the telescope racked the targets as HAWK-I is situated on the Nasmyth focus.," These patterns, caused by the spider of the secondary mirror, rotated as the telescope tracked the targets as HAWK-I is situated on the Nasmyth focus."18 These patterns were not adequately removed by the background subtraction step unless the telescope had not moved far between successive exposures., These patterns were not adequately removed by the background subtraction step unless the telescope had not moved far between successive exposures.19 Thus the integration time between dithers was reduced from mmins to mmin for all images obtained after May 2008., Thus the integration time between dithers was reduced from mins to min for all images obtained after May 2008.20 The data were reduced using the ESO/MVM 01 data reduction pipeline optimised for the reduction of our HAWK-I data., The data were reduced using the ESO/MVM \citep{Vandame2004} data reduction pipeline optimised for the reduction of our HAWK-I data.21 The usual near-infrared reduction steps were taken. including dark subtraction. flat-field removal. harmonising the gain of the four detectors. removal of fringing. sky-subtraction. creation of bad pixel and weight maps. calculating the relative astrometry between the chips and the absolute astrometry.," The usual near-infrared reduction steps were taken, including dark subtraction, flat-field removal, harmonising the gain of the four detectors, removal of fringing, sky-subtraction, creation of bad pixel and weight maps, calculating the relative astrometry between the chips and the absolute astrometry."22 USNO-BI catalogues (2) were used to calculate the guess astrometric solutions for the images., USNO-B1 catalogues \citep{Monet2003} were used to calculate the first-guess astrometric solutions for the images.23 A catalogue was then compiled from the objects detected in the iimages and used to calculate the astrometric solution of the / and H images., A catalogue was then compiled from the objects detected in the images and used to calculate the astrometric solution of the $J$ and $H$ images.24 The accuracy of the relative astrometry between the 3 images of each target is typically within ppixel. but the absolute astrometry is limited by the USNO-BI catalogue which has an accuracy of aaresec. equivalent to 2 HAWK-I pixels.," The accuracy of the relative astrometry between the 3 images of each target is typically within pixel, but the absolute astrometry is limited by the USNO-B1 catalogue which has an accuracy of arcsec, equivalent to 2 HAWK-I pixels."25 After the sky background was removed from each exposure by MVM in a two step process. low level large-scale variations were still seen across each quadrant and most notably between the quadrants of the reduced images.," After the sky background was removed from each exposure by MVM in a two step process, low level large-scale variations were still seen across each quadrant and most notably between the quadrants of the reduced images."26 Therefore the sky background of the tinal images was measured and subtracted using a local background estimator with SExtractor LOCAL)., Therefore the sky background of the final images was measured and subtracted using a local background estimator with SExtractor ).27 The data were flux calibrated with 2MASS catalogues. using 13. I5.mmag stars within the target fields.," The data were flux calibrated with 2MASS catalogues, using $13-15.5$ mag stars within the target fields."28 The calibration was checked5 using 13[4+ mmag stars within standard star fields taken within hhours of the observations., The calibration was checked using $13-14$ mag stars within standard star fields taken within hours of the observations.29 Typically the zero-points determined from both methods agreed within the uncertainties (~0.05 mmag)., Typically the zero-points determined from both methods agreed within the uncertainties $\sim0.05$ mag).30 The end products of this reduction process are a science image containing the reduced data of the target. and an effective exposure-time Image. which is an exposure-time map that has been normalised to account for differences in sensitivity between the four detectors of HAWK-I and the chip gaps.," The end products of this reduction process are a science image containing the reduced data of the target, and an effective exposure-time image, which is an exposure-time map that has been normalised to account for differences in sensitivity between the four detectors of HAWK-I and the chip gaps."31 The 3 colour images of a target were convolved to match the lowest resolution image using the packagePSFMATCH., The 3 colour images of a target were convolved to match the lowest resolution image using the package.32 The convolution kernels were optimised so that the stellar growth curves (created by median combining at least 20 bright and unsaturated stars in each image) converged to within at and beyond a 1” radius., The convolution kernels were optimised so that the stellar growth curves (created by median combining at least 20 bright and unsaturated stars in each image) converged to within at and beyond a $\arcsec$ radius.33" The 56 image depths given in reftab:obs were measured by placing 2"" diameter apertures at multiple random positions,", The $\sigma$ image depths given in \\ref{tab:obs} were measured by placing $\arcsec$ diameter apertures at multiple random positions.34 Each of the four HAWK-I detectors contain 32 amplifiers., Each of the four HAWK-I detectors contain 32 amplifiers.35 Cross-talk between the amplifiers produce a series of artefacts arranged horizontally with respect to each star in the field at regular 64 pixel intervals (2).., Cross-talk between the amplifiers produce a series of artefacts arranged horizontally with respect to each star in the field at regular 64 pixel intervals \citep{finger2008}.36 These artefacts appear crater-like and although are produced for every object in the field. they are only detectable above the noise if the star is brighter than approximately /=14.5 mmag (although this is strongly dependent on the seeing).," These artefacts appear crater-like and although are produced for every object in the field, they are only detectable above the noise if the star is brighter than approximately $J=14.5$ mag (although this is strongly dependent on the seeing)."37 The artefacts Were most pronounced in the deep images., The artefacts were most pronounced in the deep images.38" 2""x|"" rectangular regions were masked at 64 pixels intervals from each star brighter than J14.5 across the entiredetector quadrant.", $2\arcsec\times1\arcsec$ rectangular regions were masked at 64 pixels intervals from each star brighter than $J\simeq14.5$ across the entiredetector quadrant.39 Cross-talk was greatly reduced in programs observed after ESO semester 82., Cross-talk was greatly reduced in programs observed after ESO semester 82.40 Bright stars and nearby galaxies were also masked because they cover a significant amount of area in some fields and therefore, Bright stars and nearby galaxies were also masked because they cover a significant amount of area in some fields and therefore41As the waves need more time to reach the eas located παμοι away from the ceuter. the heating rate rises at xogressively later times for more distant aunuli.,"As the waves need more time to reach the gas located further away from the center, the heating rate rises at progressively later times for more distant annuli."42 Ouce the first wave has reached a given distance. viscous beating )ocomies comparable to the cooling rate.," Once the first wave has reached a given distance, viscous heating becomes comparable to the cooling rate."43 This is cousisteut with heating rate predictions made by Fabianotal.(2003a).. also assuniug Spitzer viscosity.," This is consistent with heating rate predictions made by \citet{fab03a}, also assuming Spitzer viscosity."44 We also note hat dissipating waves of ereater initial amplitude in our simulations would eive even more heating to offset cooling., We also note that dissipating waves of greater initial amplitude in our simulations would give even more heating to offset cooling.45 Tuterestinely. the average ratio of heating to cooling secus o be relatively stable as a function of time.," Interestingly, the average ratio of heating to cooling seems to be relatively stable as a function of time."46 We have also computed the vohune-inteerated heating aud cooling rates and found that their ratio couverges to a value of the order of a few., We have also computed the volume-integrated heating and cooling rates and found that their ratio converges to a value of the order of a few.47 However. the balance of heating aud cooling is uot automatic as it depends ou the choice of parameters (0.8. ACN power and deusitv eradieut in the intracluster mediunu) aud here feedback may plav a role.," However, the balance of heating and cooling is not automatic as it depends on the choice of parameters (e.g., AGN power and density gradient in the intracluster medium) and here feedback may play a role."48 Note that the curves display a pronounced periodic behavior., Note that the curves display a pronounced periodic behavior.49 This reflects the imutenuittenev of the ceutral source. with on- and off-states of 1.5«10* vears;," This reflects the intermittency of the central source, with on- and off-states of $1.5\times 10^{7}$ years."50 This is cousisteut with the observational estimates based on observations of ripples iu the Perseus cluster (Fabianetal...2003a.b).," This is consistent with the observational estimates based on observations of ripples in the Perseus cluster \citep{fab03a,fab03b}."51. We performed a series of numerical experiments to mvoestieate whether a single ACN outburst can generate waves for which the dissipation rates could offset local radiative cooling rates., We performed a series of numerical experiments to investigate whether a single AGN outburst can generate waves for which the dissipation rates could offset local radiative cooling rates.52 These simulations demonstrated that. whereas secondary waves ecucrated by the interaction of the rising bubble with the surrounding iutraclustor medi are clearly present. the viscous heating associated with a single outburst is insufficient to balance radiative cooling.," These simulations demonstrated that, whereas secondary waves generated by the interaction of the rising bubble with the surrounding intracluster medium are clearly present, the viscous heating associated with a single outburst is insufficient to balance radiative cooling."53 This suggests that the ripples observed in the Perseus cluster can be interpreted as beiug due to the AGN duty cvcle. ie.. they trace ACN activity The work doue by the expanding cavities on the ambieut medium is limited to a modest fraction of the οποιον injected by the ACN.," This suggests that the ripples observed in the Perseus cluster can be interpreted as being due to the AGN duty cycle, i.e., they trace AGN activity The work done by the expanding cavities on the ambient medium is limited to a modest fraction of the energy injected by the AGN."54 If the cavities are approximately in pressure balance with their surroundings. the iucrease in the enerey of the ambicut eas. dU.=d(PV)/(1). is related to the work done. (IT=PdV. by dU~AVA(> 1) ," If the cavities are approximately in pressure balance with their surroundings, the increase in the energy of the ambient gas, $dU = d(PV)/(\gamma - 1)$, is related to the work done, $dW = P \, dV$, by $dU \simeq dW/ (\gamma -1)$ ."55The first law of thermodynamics then implies that (ITx-- where dQ is the heat injected into the cavity.," The first law of thermodynamics then implies that $dW \simeq {\gamma - 156\over \gamma} dQ$, where $dQ$ is the heat injected into the cavity."57 This meaus that. depending on the effective value of 5 (which can range between 1/3 and 5/3). 25—LO% of the energv input can be trausterred to the züubient medi (see also. c.e.. Chirazov ct al.," This means that, depending on the effective value of $\gamma$ (which can range between 4/3 and 5/3), $25-40\%$ of the energy input can be transferred to the ambient medium (see also, e.g., Churazov et al."58 2001)., 2001).59 The fraction of the iuput power transterred to the ICAL will be laveer if the cavities are overpressured., The fraction of the input power transferred to the ICM will be larger if the cavities are overpressured.60" The fraction of this work that goes iuto acoustic enerev. as opposed to other types of disturbance (οι, e-modes or internal waves). depends on the timescale of pressure fluctuations. as well as detailed structure of the cavityICAL interface."," The fraction of this work that goes into acoustic energy, as opposed to other types of disturbance (e.g., g-modes or internal waves), depends on the timescale of pressure fluctuations, as well as detailed structure of the cavity–ICM interface."61 We expect the production of sound waves to be efficicut wheu the AGN duty evcle is of the same order as the sound crossing time at the cavity radius. or shorter.," We expect the production of sound waves to be efficient when the AGN duty cycle is of the same order as the sound crossing time at the cavity radius, or shorter."62 This condition is satisfied for our choscn duty exele of 3«10* vr., This condition is satisfied for our chosen duty cycle of $3 \times 10^7$ yr.63 In addition to work done by in situ expansion of the cavities. a roughly comparable amount of euergv is transferred to the surrounding medimm. iu the form of kinetic and eravitational potential energv. as the cavities vise through the backgound pressure eradieut.," In addition to work done by in situ expansion of the cavities, a roughly comparable amount of energy is transferred to the surrounding medium, in the form of kinetic and gravitational potential energy, as the cavities rise through the backgound pressure gradient."64 The latter is the generic mechanisin appealed to bv Beechuan (2001) and Ruszkowski Beechuan (2002) iu thei discussion of “effervescent heating”.," The latter is the generic mechanism appealed to by Begelman (2001) and Ruszkowski Begelman (2002) in their discussion of “effervescent heating""."65 Enerey injected in this way can also be converted to heat through viscous dissipation., Energy injected in this way can also be converted to heat through viscous dissipation.66 Our sinulatious map the total viscous dissipation rate. aud do not distinguish between dissipation of sound waves and other kinds of motion.," Our simulations map the total viscous dissipation rate, and do not distinguish between dissipation of sound waves and other kinds of motion."67 Note. however. that sound waves have lareer propagation speeds than other modes. aud therefore should progressively dominate the energeties at radii well outside the Although we have devised a specific model iu which the viscous dissipation rate of sound waves roughly balances local radiative cooling. such a balance may not be a universal property of ACN heating iu cluster cores.," Note, however, that sound waves have larger propagation speeds than other modes, and therefore should progressively dominate the energetics at radii well outside the Although we have devised a specific model in which the viscous dissipation rate of sound waves roughly balances local radiative cooling, such a balance may not be a universal property of AGN heating in cluster cores."68 The distribution of sound cucrey dissipation is larecly determined by the radial structure of the model., The distribution of sound energy dissipation is largely determined by the radial structure of the model.69 The sound dissipation leneth for a fixed waveleneth L(A.1) for the paraincters in our simulations decreases from the center to the edge of the simulated region. mainly due to the decrease in density (Fabian et al.," The sound dissipation length for a fixed wavelength $L(\lambda, r)$ for the parameters in our simulations decreases from the center to the edge of the simulated region, mainly due to the decrease in density (Fabian et al."70 2003)., 2003).71 For the paralcters chosen in our simulation the characteristic dissipation leusth near the outer edge of the exid is of order the size of the simulation region. implying that the dissipation is spread over a volune that far exceeds that of the bubbles. aud ach of the acoustic energy goes iuto heating.," For the parameters chosen in our simulation the characteristic dissipation length near the outer edge of the grid is of order the size of the simulation region, implying that the dissipation is spread over a volume that far exceeds that of the bubbles, and much of the acoustic energy goes into heating."72 The steady rate of production of acoustic energv then leads to a rough balance between heating aud cooling. eiven the adopted deusity aud temperature profile.," The steady rate of production of acoustic energy then leads to a rough balance between heating and cooling, given the adopted density and temperature profile."73 These conditious may not be satisfied in all clusters., These conditions may not be satisfied in all clusters.74 As the sound dissipation leusth is proportional to the square of the period of the sound waves. more frequent outbursts should lead to more ceutrally concentrated dampine.," As the sound dissipation length is proportional to the square of the period of the sound waves, more frequent outbursts should lead to more centrally concentrated damping."75 Uowever. the dissipation rate does uot depend on the ACN intermittency period as such. since the pressure pulses ecuerated by the bubbles are likely to be far from sinusoidal aud will contain a wide rauge of frequencies.," However, the dissipation rate does not depend on the AGN intermittency period as such, since the pressure pulses generated by the bubbles are likely to be far from sinusoidal and will contain a wide range of frequencies."76" The dispersion of the waves as they propagate sugecst that the ""effective"" wavelength will increase with r. au effect that will partially counteract the decrease of L(A.1) with +."," The dispersion of the waves as they propagate suggest that the “effective"" wavelength will increase with $r$, an effect that will partially counteract the decrease of $L(\lambda, r)$ with $r$."77" Where the velocity field has small-scale structure or where the damping rate is much lareer than the Spitzer rate, acoustic waves (as well as eravity aud internal waves) can be dissipated much closer to the sites where they are eecuerated."," Where the velocity field has small-scale structure or where the damping rate is much larger than the Spitzer rate, acoustic waves (as well as gravity and internal waves) can be dissipated much closer to the sites where they are generated."78 Distributed heating would then occur oulv after he bubbles had penetrated most of the cluster., Distributed heating would then occur only after the bubbles had penetrated most of the cluster.79 This is the situation envisaged by Begehuan (2001) and Ruszkowski Beechuan (2002) in the effervescent heating scenario., This is the situation envisaged by Begelman (2001) and Ruszkowski Begelman (2002) in the effervescent heating scenario.80 This orm of heating may be occumiug concurrently with the argc-scale acoustic heating iu Perseus. aud may dominate he heating in other clusters (e.g. those with smaller acoustic energev ecucration due to the intermittency xoperties of the central AGN).," This form of heating may be occurring concurrently with the large-scale acoustic heating in Perseus, and may dominate the heating in other clusters (e.g., those with smaller acoustic energy generation due to the intermittency properties of the central AGN)."81 Note that viscosity may welp the bubbles penetrate to large distances without excessive LUNI., Note that viscosity may help the bubbles penetrate to large distances without excessive mixing.82 We stress that our two-dimensional simulations do not accurately represent the behavior of three-dineusional acoustic heating in several respects., We stress that our two-dimensional simulations do not accurately represent the behavior of three-dimensional acoustic heating in several respects.83 In three dimensions, In three dimensions84small effect.,small effect.85 For example. ife =1. changing Z bv a factor 5 would only change rias by ~20%.," For example, if $\epsilon=-1$, changing $Z$ by a factor 5 would only change $r_{\rm max}$ by $\sim20\%$."86 The existence of a cutolf in the distribution bevond some size is naturally expected in the above picture. given the existence of some minimum density ys in the regions of the galaxies where SNe explode. and in view of the decrease Of Puuas With raciius.," The existence of a cutoff in the distribution beyond some size is naturally expected in the above picture, given the existence of some minimum density $\rho_{\rm min}$ in the regions of the galaxies where SNe explode, and in view of the decrease of $\rho_{\rm max}$ with radius."87 At some radius. gas will equal fiiia and the integral in Eq.," At some radius, $\rho_{\rm max}$ will equal $\rho_{\rm min}$ and the integral in Eq."88 13. will therefore become zero., \ref{dNdr} will therefore become zero.89 In other words. there will be nowhere in the MCSs a region with a density low enough to permit a SNR of that size that is still in its bright Sedov phase.," In other words, there will be nowhere in the MCs a region with a density low enough to permit a SNR of that size that is still in its bright Sedov phase."90 A deficit of SNRs at small racii. as observed in M33 ane the MCs. is also expected in this scenario.," A deficit of SNRs at small radii, as observed in M33 and the MCs, is also expected in this scenario."91 Before the onse of the Sedov stage. faster shock velocities will lead to fewer objects observed. in the bins with the smallest. radii.," Before the onset of the Sedov stage, faster shock velocities will lead to fewer objects observed in the bins with the smallest radii."92" This onset happens at ages (sizes) that depend on the details of the ejecta structure. as well as the ambient density (see87in""Truelove&Melxee 1999).. but for most SNRs it should occur around a few hundred vears (a few pc). which is consisten with the deficits that we have observed in M33 and the MICs."," This onset happens at ages (sizes) that depend on the details of the ejecta structure, as well as the ambient density \citep[see \S~7 in][]{truelove99:adiabatic-SNRs}, but for most SNRs it should occur around a few hundred years (a few pc), which is consistent with the deficits that we have observed in M33 and the MCs."93 This regime alfects only a small number of objects in the ALCs. and does not impact any of the arguments made above. so we will ignore it for the remainder of the paper.," This regime affects only a small number of objects in the MCs, and does not impact any of the arguments made above, so we will ignore it for the remainder of the paper."94 We have also ignored the deviations from the standard evolutionary picture that can be introduced by the shape of he circumstellar medium excavated by the SN progenitors., We have also ignored the deviations from the standard evolutionary picture that can be introduced by the shape of the circumstellar medium excavated by the SN progenitors.95 Baclenesοἱal.(2007) showed that most Type la SNRs with known ages have sizes that are consistent with an interaction witha uniform ambient medium. but no such study has been done for CC SNRs.," \citet{badenes07:outflows} showed that most Type Ia SNRs with known ages have sizes that are consistent with an interaction with a uniform ambient medium, but no such study has been done for CC SNRs."96 The fast stellar outflows expected from he more massive CC SN progenitors will modify the sizes ofa few individual SNlis at certain stages of their evolution (c.g..Dwarkaclas2005.2007).. but most of the objects that we consider here are too large to be expanding in even the most extreme wind-blown cavities.," The fast stellar outflows expected from the more massive CC SN progenitors will modify the sizes of a few individual SNRs at certain stages of their evolution \citep[e.g.,][]{dwarkadas05:SNR-Bubbles_1D,dwarkadas07:SNRs_Bubbles_WR}, but most of the objects that we consider here are too large to be expanding in even the most extreme wind-blown cavities."97 As long as the bulk of he SNRs in the sample spend most of their lifetimes in the Seclov stage. this should not allect our scenario.," As long as the bulk of the SNRs in the sample spend most of their lifetimes in the Sedov stage, this should not affect our scenario."98 Similarly. he fact that some SNRs evolve inside superbubbles (e.g.AlacLow&AleCray1988) is naturally incorporated. into our picture superbubbles merely become one more of the actors driving the density distribution in the interstellar medium.," Similarly, the fact that some SNRs evolve inside superbubbles \citep[e.g.][]{maclow88:superbubbles} is naturally incorporated into our picture – superbubbles merely become one more of the factors driving the density distribution in the interstellar medium."99 Incidentally. the size distribution of superbubbles also relates to the properties of the interstellar medium. as shown bv Oev&Clarke(1997). [or several nearby. galaxies. including the SAIC.," Incidentally, the size distribution of superbubbles also relates to the properties of the interstellar medium, as shown by \citet{oey97:superbubble_sizes} for several nearby galaxies, including the SMC."100 As we have seen. a uniform SNR. size distribution can be understood as he result of SecOV expansion. οςmbined with a transition to the radiative pdase al an age hat depends on the local cknsbÜv. provided that the clensiv of the gas in the interstelar medium [οlows a clistribution close to a power law with an index. of 1. Pidp~p7.," As we have seen, a uniform SNR size distribution can be understood as the result of Sedov expansion, combined with a transition to the radiative phase at an age that depends on the local density, provided that the density of the gas in the interstellar medium follows a distribution close to a power law with an index of $-1$, $dP/d\rho\sim \rho^{-1}$."101 In this Section. we tes this hypothesis by examining hree indirect tracers of gas censity in the Aagellanie Cloucs: HIE column density: star-formation rate (SER) based on resolved stellar populations: and Ho emission-ine surface briginess.," In this Section, we test this hypothesis by examining three indirect tracers of gas density in the Magellanic Clouds: HI column density; star-formation rate (SFR) based on resolved stellar populations; and $\alpha$ emission-line surface brightness."102 These tracers are wel suited for our goals because they are valid over a wide range of densities. and the necessary data are available [rom public surveys that cover the whole extent of the Clouds. as deseribed in detail below.," These tracers are well suited for our goals because they are valid over a wide range of densities, and the necessary data are available from public surveys that cover the whole extent of the Clouds, as described in detail below."103 We have taken the surface brightness of HE 21 cm line emission in the ALCS from the maps of Iximetal.(2003) and Stanimiroviectal.(1999)... which combine single-dish Parkes and aperture-svnthesis APCA data to. probe both small and laree scales in the LAIC and the SAIC. respectively.," We have taken the surface brightness of HI 21 cm line emission in the MCs from the maps of \citet{kim03:LMC_HI_Parkes_ATCA} and \citet{stanimirovic99MNRAS.302..417S}, which combine single-dish Parkes and aperture-synthesis ATCA data to probe both small and large scales in the LMC and the SMC, respectively."104 “Phe 21 cem emission is optically thin. so the surface brightness is directly. proportional to the HIE column density.," The 21 cm emission is optically thin, so the surface brightness is directly proportional to the HI column density."105 Since the LMC possesses a fairly face-on (inclinationPee357. 2001). wellorderecl LIE clisk. the column density should. in turn. be roughly. proportional to the volume density p.," Since the LMC possesses a fairly face-on \citep[inclination $i\sim 35^{\circ}$ , well-ordered HI disk, the column density should, in turn, be roughly proportional to the volume density $\rho$."106 Iximal.(2007) report that the LIL column density distribution ofet individual “clouds” of neutral hvdrogen in the LMC follows à log-normal form. rather than a power law.," \citet{kim07:HI_Clouds_LMC} report that the HI column density distribution of individual “clouds” of neutral hydrogen in the LMC follows a log-normal form, rather than a power law."107" However. their figure 13 suggests that. above a low cutoll of 2107""em.7. the distribution does behave as a power law of slope l.over at least an order of magnitude."," However, their figure 13 suggests that, above a low cutoff of $2\times 10^{20}~{\rm cm}^{-2}$, the distribution does behave as a power law of slope $\sim -1$ , over at least an order of magnitude."108 ‘To re-examine this. in Figure 5 we show the cillerentia distribution of HIE column density in the individual sized pixels of the Iximetal.(2003). LMC map 2007)..," To re-examine this, in Figure \ref{HIHist} we show the differential distribution of HI column density in the individual beam-sized pixels of the \citet{kim03:LMC_HI_Parkes_ATCA} LMC map \citep[as opposed to109the cumulative plot for ``clouds'' shown in][]{kim07:HI_Clouds_LMC}."110" We see that the LL column in the LMC does follow an index l power Law fairlv well. between a column of 3.107""em ?2and 61073em.7."," We see that the HI column in the LMC does follow an index $-1$ power law fairly well, between a column of $3\times 10^{20}~{\rm cm}^{-2}$ and $6\times 10^{21}~{\rm111 cm}^{-2}$."112 The observed low cutoll in the column density is unavoidable because of the integration through the disk and over the beam size (every line of sigh is basically sampling the densest regions at that point)., The observed low cutoff in the column density is unavoidable because of the integration through the disk and over the beam size (every line of sight is basically sampling the densest regions at that point).113 In regions with low density. the LE atoms might be ionized. as in the swarm ionized” phase of the interstellar medium (Ferriere2001).. so the tracer may become loss reliable there.," In regions with low density, the H atoms might be ionized, as in the “warm ionized” phase of the interstellar medium \citep{ferriere01:ISM}, so the tracer may become less reliable there."114 ]t is quite plausible that. in the regions where SNe actually explode. the underlving distribution of densities also reaches a minimum as we recall. a minimum density. is required in order to reproduce the observed upper cutoll in SNIt size rut this does not necessarily correspond with the lower threshold in the LIE distribution.," It is quite plausible that, in the regions where SNe actually explode, the underlying distribution of densities also reaches a minimum – as we recall, a minimum density is required in order to reproduce the observed upper cutoff in SNR size – but this does not necessarily correspond with the lower threshold in the HI distribution."115" To illustrate this. we have calculated the mean LIE columns in each of the spatial ""cells"" defined by Llarris&Zaritsky(2004) and Harris&Zaritsky(2009) that contain SNRs (see 5.2 below for adescription of the cells). which we display with the horizontal rulers in Figure 5.."," To illustrate this, we have calculated the mean HI columns in each of the spatial “cells” defined by \citet{harris04:SMC_SFH} and \citet{harris09:LMC_SFH} that contain SNRs (see \ref{sec:SFR} below for adescription of the cells), which we display with the horizontal rulers in Figure \ref{HIHist}."116 We note that. in the LMC. these cells have average column values between ο10720em72 (close to. but higher: than the low LIE cutoll) and 61072lem2τι although most SNRs appear clustered around 21073em7.," We note that, in the LMC, these cells have average column values between $5\times10^{20}~{\rm cm}^{-2}$ (close to, but higher than the low HI cutoff) and $6\times 10^{21}~{\rm cm}^{-2}$, although most SNRs appear clustered around $2\times 10^{21}~{\rm cm}^{-2}$."117 As shown in Figure 5..5.. the cistribution of HIE column densities in the SAIC is flat over the same range. although the rise ancl fall from the plateau happen at the same densities as the rise and fall of the powerlaw in the LMC.," As shown in Figure \ref{HIHist}, the distribution of HI column densities in the SMC is flat over the same range, although the rise and fall from the plateau happen at the same densities as the rise and fall of the powerlaw in the LMC."118 While we do not know the reason for this. we speculate that it may. be related to an SAIC geometry. that is elongated along our line of sight. and the integration ellect that results.," While we do not know the reason for this, we speculate that it may be related to an SMC geometry that is elongated along our line of sight, and the integration effect that results."119 The actual “depth” of the SAIC. whether just a few kpe or as much as 20 kpe. is debated. (Llatzicimitriou&Llawkins 2000)... but is likely at least a few times larger than that of the nearly [ace-on. LMC.," The actual “depth” of the SMC, whether just a few kpc or as much as 20 kpc, is debated \citep{hatzidimitriou89:SMC_structure,harris04:SMC_SFH,subramanian09:LMC_SMC_Depth}, , but is likely at least a few times larger than that of the nearly face-on LMC."120 Such an integration elfect would explain why the SAIC SNRsare found at HE column, Such an integration effect would explain why the SMC SNRsare found at HI column121Dynamics auc cor'espouding Laclative signatwes of nou-spherical relativistic shocks remains an linportaut unresolved issues tu studies on Gatuna Ray Bursts (GRBs).,Dynamics and corresponding radiative signatures of non-spherical relativistic shocks remains an important unresolved issues in studies on Gamma Ray Bursts (GRBs).122 Since GRBs produce uarrowlv collimated οullows that evolve lateralN. uuderstandiug the overall dyuamies - both theoretical aud in terns of agreement between cliΠοιοί numerical results - is imperative to the iuterpretation of the 'oadbaud observatious of GRBs ??..," Since GRBs produce narrowly collimated outflows that evolve laterally, understanding the overall dynamics - both theoretical and in terms of agreement between different numerical results - is imperative to the interpretation of the broadband observations of GRBs \cite{Rhoads99,Frail01}."123 Presently. there aο two coimpeti ews ο he lateral evolution of the relativistic outflows.," Presently, there are two competing views on the lateral evolution of the relativistic outflows."124 Theoretically. it is twpically argued t je lateral evolution of the flow proceeds with relativistic velocities ?).(2).. (seealso? )..," Theoretically, it is typically argued that the lateral evolution of the flow proceeds with relativistic velocities \citep{PiranReview}, \citep[see also][]{2011arXiv1102.5618W}."125 This view Οἱαςicted by he results of numerical simulations that show very little lateral evolΠιο in t alivistic reelme ?TTT?..," This view is contradicted by the results of numerical simulations that show very little lateral evolution in the relativistic regime \cite{2004ApJ...601..380C,2009ApJ...698.1261Z,2010A&A...520L...3M,2011arXiv1105.2485V}."126 In this Letter we argue hat thisdisagreemen results L'om the incorrect theoretical assumptions about the lateral evolution of the flow., In this Letter we argue that this disagreement results from the incorrect theoretical assumptions about the lateral evolution of the flow.127 Wjab is linj»ortant for the interpretation of observations is the evolution of a curved stock., What is important for the interpretation of observations is the evolution of a curved shock.128 Previously t aeral evolution of the nor-spherical shocks was incorrectly treated as a f‘ee lateral expansion LO vacuttn (e.g..?.Eq.5) 2s , Previously the lateral evolution of the non-spherical shocks was incorrectly treated as a free lateral expansion into vacuum \citep[\eg][Eq. 5]{2011arXiv1102.5618W}. ..129"The assumption of the lateral expausion witl Iie souud speed results 1 lagramophone-type"" p‘oliles aud. slowing down of the ejeca.", The assumption of the lateral expansion with the sound speed results in a “gramophone-type” profiles and slowing down of the ejecta.130 This has drastic iinplicatious for the uuderlviug light curves (eg??)..," This has drastic implications for the underlying light curves \citep[eg][]{2000ApJ...541L...9K,2003ApJ...592..390P}."131" Iu fact the dynamics of tle nou-sphlierical shocks is uore subtle: the correct treatineut. as we argue below. is cousistent with slow lateral evolution ""eel in uuimerical simulations."," In fact the dynamics of the non-spherical shocks is more subtle; the correct treatment, as we argue below, is consistent with slow lateral evolution seen in numerical simulations."132 Evolution of strong ion-spherical shocks is a well studies problems iu fluid dynanics., Evolution of strong non-spherical shocks is a well studies problems in fluid dynamics.133 The two fuudamental works that have laid the founcatiou for nou-sphlerical (two-dimeusional) shocks.," The two fundamental works that have laid the foundation for non-spherical (two-dimensional) shocks,"134was located in the inter-binary system.,was located in the inter-binary system.135 However. further work (eg Juett Chakrabarty 2005) showed that for individual sources. the Ne/O ratio showed evidence for variability from epoch to epoch. which they attributed. to source variability.," However, further work (eg Juett Chakrabarty 2005) showed that for individual sources, the Ne/O ratio showed evidence for variability from epoch to epoch, which they attributed to source variability."136 This implied that the abundances could not be used to determine the composition of the mass donating star., This implied that the abundances could not be used to determine the composition of the mass donating star.137 There is a clear similarity between the neutron star UCBs described by Juctt et al and RX 1914|24., There is a clear similarity between the neutron star UCBs described by Juett et al and RX J1914+24.138 1n each observation of RN J1914|24. there is clear evidence that the absorption component has an overabundance of neon.," In each observation of RX J1914+24, there is clear evidence that the absorption component has an overabundance of neon."139 For the reasons outlined in 85.1. we rule out AA 1914|24 being an isolated neutron star.," For the reasons outlined in \ref{ins}, we rule out RX J1914+24 being an isolated neutron star."140 Since all the known neutron star UCBs have X-ray emission extending up to many 10s of keV we also rule out an accreting neutron star UCB model., Since all the known neutron star UCBs have X-ray emission extending up to many 10's of keV we also rule out an accreting neutron star UCB model.141 We cannot rule out that a neutron star is in a binary system where a secondary star was not filling its Roche Lobe., We cannot rule out that a neutron star is in a binary system where a secondary star was not filling its Roche Lobe.142 In this scenario. an N-rav. bright svstemi would have to be powered by Ul.," In this scenario, an X-ray bright system would have to be powered by UI."143 Lt is highly. unlikely that the line of sight absorption to RN 121914|24 has a chance enhancement of neon., It is highly unlikely that the line of sight absorption to RX J1914+24 has a chance enhancement of neon.144 LH is much more likely that this overabundance is concentrated in the binary svstem., It is much more likely that this overabundance is concentrated in the binary system.145 Juett Chakrabarty (2005) noted that for some neutron star UCBs the Ne/O abundance varied. from epoch to epoch and hence the observations could not be used to determine the abundance of the secondary. mass-donating star. in the binary svstem.," Juett Chakrabarty (2005) noted that for some neutron star UCBs the Ne/O abundance varied from epoch to epoch and hence the observations could not be used to determine the abundance of the secondary, mass-donating star, in the binary system."146 In the case of RN 1914|24 there is clear evidence for a significant. over-abundance of neon in the absorption Component at each epoch., In the case of RX J1914+24 there is clear evidence for a significant over-abundance of neon in the absorption component at each epoch.147 At this stage it is not clear if this over-abuncance is due to circumbinary materia eft over from a previous stage in the binary formation »Focess or can give us a direct. insight into the chemica composition of the secondary star (if accretion is occurring)., At this stage it is not clear if this over-abundance is due to circumbinary material left over from a previous stage in the binary formation process or can give us a direct insight into the chemical composition of the secondary star (if accretion is occurring).148 What are the implications of our findings reearcing he X-rav luminositv of RA .1914]|24?, What are the implications of our findings regarding the X-ray luminosity of RX J1914+24?1497 Stecehs et a (2006) discuss the extinction and distance estimates to UX J1914]24., Steeghs et al (2006) discuss the extinction and distance estimates to RX J1914+24.150 While the distance is rather uncertain. it is likely that it is greater than 1 kpe.," While the distance is rather uncertain, it is likely that it is greater than $\sim$ 1 kpc."151 We can rule ou he lower estimates (Lx~10%? Hor a distance of kpe) which were derived. using a low temperature thermal plasma model., We can rule out the lower estimates $L_\mathrm{X}\sim10^{33}$ for a distance of 1 kpc) which were derived using a low temperature thermal plasma model.152 Taking the unabsorbed bolometric [uxes derived using the blackhocky with absorption component with variable abundances ancl assuming a distance of 1 kpe we find Ly—2.105.L6.107?," Taking the unabsorbed bolometric fluxes derived using the blackbody with absorption component with variable abundances and assuming a distance of 1 kpc we find $L_\mathrm{X}=2\times10^{34} -1531.6\times10^{35}$."154 DallOsso ct al (2007) made a detailed. investigation of the Ul model in the context of RN J1914|24 and RX JOSOG|15., Dall'Osso et al (2007) made a detailed investigation of the UI model in the context of RX J1914+24 and RX J0806+15.155 They. predicted that for low luminosities. Lx105 the asvnchronism between the orbit. and he magnetic star in RA 1914|24 would have to be a~1.9 0.98. where a=wjfe. andy is the rotation frequency of the primary star and cy is the orbital frequency.," They predicted that for low luminosities, $L_\mathrm{X}\sim10^{33}$ , the asynchronism between the orbit and the magnetic star in RX J1914+24 would have to be $\alpha\sim0.9-$ 0.98, where $\alpha=\omega_{1}/\omega_{o}$, and $\omega_{1}$ is the rotation frequency of the primary star and $\omega_{o}$ is the orbital frequency."156 Unless UN. J1914124 was located at a distance significantly. less han Ikpce. we can rule these low values of asynchronism.," Unless RX J1914+24 was located at a distance significantly less than 1kpc, we can rule these low values of asynchronism."157 bor ο] luminosities (Lx=107 13). Dall'Osso et al (2007) predicted that an asvnchronism of a few was required (o~ 4).," For high luminosities $L_\mathrm{X}=10^{34-35}$ ), Dall'Osso et al (2007) predicted that an asynchronism of a few was required $\alpha\sim$ 4)."158 For an observed period of 569 sec. à=5.10 gives a predicted: period. of 760-300 sec.," For an observed period of 569 sec, $\alpha=2-10$ gives a predicted period of $\sim$ 60-300 sec."159 There is no evidence for power at these periods in the power spectra of the X-ray ight curves (Ramsay ct al 2006)., There is no evidence for power at these periods in the power spectra of the X-ray light curves (Ramsay et al 2006).160 Until now the nature of the emission source that powers the X-ray spectrum of RA J1914124 has been far from clear., Until now the nature of the emission source that powers the X-ray spectrum of RX J1914+24 has been far from clear.161 In this paper we have shown that it can be well modelled using a simple blackbody model. with an absorption component which has non-solar abuncances. in particular. an enhancement of neon.," In this paper we have shown that it can be well modelled using a simple blackbody model with an absorption component which has non-solar abundances, in particular, an enhancement of neon."162 Since the X-ray light curves of RX 1914|24 and RX JOSOG|15 are practically identical. it suggests that their X-ray emission source is the same.," Since the X-ray light curves of RX J1914+24 and RX J0806+15 are practically identical, it suggests that their X-ray emission source is the same."163 The fact wt their X-pav spectra were apparently different (albeit both being soft) was therefore. perplexing., The fact that their X-ray spectra were apparently different (albeit both being soft) was therefore perplexing.164 Our result. showing. that the emission. source is the same for both RN 1914|24 and RX JOSOG|15 is therefore very attractive., Our result showing that the emission source is the same for both RX J1914+24 and RX J0806+15 is therefore very attractive.165 Indeed. their temperatures are virtually identical owe obtain a mean value of AL~67 eV for RN JL914|24 compared to kl65 eV for RX JOSOG|15 Csracl οἱ al 2003)., Indeed their temperatures are virtually identical – we obtain a mean value of $kT\sim67$ eV for RX J1914+24 compared to $kT\sim65$ eV for RX J0806+15 (Israel et al 2003).166 The cillerence between the X-ray spectrum. of RX J1914]24 and RA JOSOG|15 is that the absorption component of IUX J1914|24 has enhanced neon abundance., The difference between the X-ray spectrum of RX J1914+24 and RX J0806+15 is that the absorption component of RX J1914+24 has enhanced neon abundance.167were acceptable.,were acceptable.168 Such spectra contrast with the cool optically thick coronae obtained when modeling high huuinositv neutron star binaries with (e.g.7).," Such spectra contrast with the cool, optically thick coronae obtained when modeling high luminosity neutron star binaries with \citep[e.g.][]{disalvo00}."169 The disk blackbody fit wasacceptable (4? /dof = 110/117): however the inner disk temperature (kta) was 3.20.1 keV. which is too hot for a black hole high state.," The disk blackbody fit wasacceptable $\chi^2$ /dof = 140/147); however the inner disk temperature $T_{\rm in}$ ) was $\pm$ 0.4 keV, which is too hot for a black hole high state."170 Hence. we reject this fit ou physical grounds. even though it is a statistically acceptable fit to the data.," Hence, we reject this fit on physical grounds, even though it is a statistically acceptable fit to the data."171 The two component inodel vielded Αι = I«1075402 and D — ME12. as with. the simple. power aw quodol.," The two component model yielded $N_{\rm H}$ = $\times 10^{21}$ and $\Gamma$ = 1.2, as with the simple power law model."172" Furthermore. NSPEC was unable o estimate the uncertainties in the blackbody xuineters,"," Furthermore, XSPEC was unable to estimate the uncertainties in the blackbody parameters."173 We therefore couchide that αν hermal componcut is too marginal to be detected: rence. NBOs2 is unlikely to be in the high accretion rate state for neutron star LAINBs.," We therefore conclude that any thermal component is too marginal to be detected; hence, XB082 is unlikely to be in the high accretion rate state for neutron star LMXBs."174 We oxeseut the spectra and best fit two componcut nodel iu Fig. 3.., We present the spectrum and best fit two component model in Fig. \ref{4bhspec}.175 We conclude that NBUS2 exhibited nou-thermal enission approximated by a power law., We conclude that XB082 exhibited non-thermal emission approximated by a power law.176 However. active galactic unclei (ACN) also exhibit similar Cluission spectra. hence we calculated the probability of a coincident ACN from the 2.10 keV huuinositv function provided bv ?..," However, active galactic nuclei (AGN) also exhibit similar emission spectra, hence we calculated the probability of a coincident AGN from the 2–10 keV luminosity function provided by \citet{moretti03}."177" The observed 210 keV flux of NDUS2 derived from the best fit power law model was 2.041007 erg en? hk the probability of an ACN this bright existiue within 1"" of one of the Bs GCs is +2 «10. ©.", The observed 2–10 keV flux of XB082 derived from the best fit power law model was $\times 10^{-12}$ erg $^{-2}$ $^{-1}$; the probability of an AGN this bright existing within $''$ of one of the 428 GCs is $\sim$ $\times$ $^{-6}$.178 We therefore conclude that NDOS2 is a candidate black hole LMXND., We therefore conclude that XB082 is a candidate black hole LMXB.179 If we assume P=1.l. then the black hole would require a mass ~30 ML. for a 1550 keV Iuuinositv <0.1 Lepp. consistent with the low state (following?)..," If we assume $\Gamma$ =1.4, then the black hole would require a mass $\sim$ 30 $_{\odot}$ for a 15–50 keV luminosity $\la$ 0.1 $L_{\rm EDD}$, consistent with the low state \citep[following][]{tang10}."180 This is rather more massive than the Galactic stellar mass black holes (?).. but consistent with the priuary in the dyvuunicallv confirmed black hole | WolfRavet binary ICIO N-L (07).," This is rather more massive than the Galactic stellar mass black holes \citep{ozel10}, but consistent with the primary in the dynamically confirmed black hole + Wolf-Rayet binary IC10 X-1 \citep{silverman08}."181 We note that the observed spectrum of NBOs2 (D ~1.2) is more often associated with ligh mass A-rav binaries (ΠΑΛΙΑΟΚ) with Be donors (seec.g.T.forarecent review): e.g. 7. examined NAIANewton observations of 11. Be NBs in the Suiall Magellanic Cloud. aud found the photon iudex distribution to be strongly peaked at D — 1.00. with a standard devition of 0.16.," We note that the observed spectrum of XB082 $\Gamma$ $\sim$ 1.2) is more often associated with high mass X-ray binaries (HMXBs) with Be donors \citep[see e.g.][ for a recent review]{reig11}; e.g. \citet{haberl04} examined XMM-Newton observations of 11 Be XBs in the Small Magellanic Cloud, and found the photon index distribution to be strongly peaked at $\Gamma$ = 1.00, with a standard devition of 0.16."182 Therefore NDUS2 could be a IIMXD superposed on the elobular cluster., Therefore XB082 could be a HMXB superposed on the globular cluster.183 However. we note that ποσα Be NBs tend to be transient. N-rav sources. as they have lone. eccentric orbits and oulv accrete for a short time when the neutron star is closest to the donor: fiuthermore. this cussion spectrum is ecucrally confined to the range ~107+ 107 eve 1 (?.andreferenceswithin)...," However, we note that known Be XBs tend to be transient X-ray sources, as they have long, eccentric orbits and only accrete for a short time when the neutron star is closest to the donor; furthermore, this emission spectrum is generally confined to the range $\sim$ $^{34}$ $^{37}$ erg $^{-1}$ \citep[][and references within]{reig11}."184" Persisteuthy bright Be NBs exist. but tend to have N-ray Imniuosities < 10%"" ere +."," Persistently bright Be XBs exist, but tend to have X-ray luminosities $\la$ $^{35}$ erg $^{-1}$."185 Tf NBOs2 is a neutron star IIMXD. it is unlike any know thus far.," If XB082 is a neutron star HMXB, it is unlike any know thus far."186" NBI53 was observed in 75 ACTS and 15 URC observations. and appears to be persistently bright: its 0.3.10 keV huninosityv varied by a factor ~3up to ~2.1< 10°"" ere st."," XB153 was observed in 75 ACIS and 45 HRC observations, and appears to be persistently bright; its 0.3–10 keV luminosity varied by a factor $\sim$ 3 up to $\sim$ $\times 10^{38}$ erg $^{-1}$."187 It was also observed in the 2002 NNMM-Newton observation with 60 ks eood time., It was also observed in the 2002 XMM-Newton observation with 60 ks good time.188 The best fit power kuv to the NATALNewton pu spectrum vielded Αι = 8.541.0< 1029 atom ?. aud P= 1.6240.01: 42 /dof = 129/tlLaud Los.i10 = LOTEO.OL SIO ore ," The best fit power law to the XMM-Newton pn spectrum yielded $N_{\rm H}$ = $\pm$ $\times 10^{20}$ atom $^{-2}$, and $\Gamma$ = $\pm$ 0.04; $\chi^2$ /dof = 429/444 and $L_{0.3-10}$ = $\pm$ $\times 10^{38}$ erg $^{-1}$."189Fitting a cluission mocel vielded a hot. optically thin corona with uncoustraimed kT. 060 keV aud 7r —1.," Fitting a emission model yielded a hot, optically thin corona with unconstrained $T_{\rm e}$ $\sim$ 60 keV and $\tau$ $\sim$ 1."190 The disk blackbody: model was rejected. with v /dof = 1067/LL.," The disk blackbody model was rejected, with $\chi^2$ /dof = 1067/445."191 Again. Nyy and E for the two component model were consistent with the values for the single power law. and NSPEC was unable to estimate the uncertainties for the blackbody parameters.," Again, $N_{\rm H}$ and $\Gamma$ for the two component model were consistent with the values for the single power law, and XSPEC was unable to estimate the uncertainties for the blackbody parameters."192 We present the spectrum and two component fit in Fie 3.., We present the spectrum and two component fit in Fig \ref{4bhspec}.193 NBI53 also appears to have been m a non-thermal state., XB153 also appears to have been in a non-thermal state.194" The observed 210 keV fux for the best fit power law modelwas «10.1 ere ein ? thence. the probability of finding an ACN of this brightnesso within 1"" of αν of tlhe 128 CC's is 1.0410 7. The ↽∙∶≻∖↽"," The observed 2–10 keV flux for the best fit power law model was $\times10^{-13}$ erg cm $^{-2}$ $^{-1}$; hence, the probability of finding an AGN of this brightness within $''$ of any of the 428 GCs is $\times 10^{-5}$ ."195⊽∙↴∐∖∏⊔↕⋯↴∖↴↕↖⇁⋟∪↥⋅∐∖⊸⊽⋀⋀≓⋀⊽↸∖↖↖↽∪∐I tvfor the NMM-Newt observation was a factor ~2 lower than the peak., The luminosity for the XMM-Newton observation was a factor $\sim$ 2 lower than the peak.196 XD 163 was observed in 10. ACTS aud 29 TRC observations., XB 163 was observed in 10 ACIS and 29 HRC observations.197 NB163 exhibited at least 5 outbursts over the ~ lo00 day viewing period. aud was brightest during ACTS observation slsl (2007 February 11) the exposure was 5 ks.," XB163 exhibited at least 5 outbursts over the $\sim$ 4000 day viewing period, and was brightest during ACIS observation 8184 (2007 February 14); the exposure was 5 ks."198 The 0.3 keV huuinositv of NBIG3 varied by a factor ~650. lence it cannot be an ACN.," The 0.3--10 keV luminosity of XB163 varied by a factor $\sim$ 650, hence it cannot be an AGN."199 NDBIG63 was, XB163 was200profile. or it may be because the universalitv of the density profile (Navarroctal.1996:Mooreetal.1998) really fiuds its origin i accretion listorv (Conzález-Casadoetal.,"profile, or it may be because the universality of the density profile \citep{nfw,moore} really finds its origin in accretion history \citep{gsmh}."2012007).. Let us finally discuss what our newly fouud attractor may do for the mass-velocity anisotropy degeneracy., Let us finally discuss what our newly found attractor may do for the mass-velocity anisotropy degeneracy.202 When we observe the stellar ttics in a dwarf galaxy we can observe the stellar density aud the stellar dispersion., When we observe the stellar tics in a dwarf galaxy we can observe the stellar density and the stellar dispersion.203 Then. the jeans equation. eq. (3)).," Then, the jeans equation, eq. \ref{eq:jeans}) ),"204 tells us that for amy assumed. velocity anisotropy profile for the stars. oC). we can solve for the total eravitating mass.," tells us that for any assumed velocity anisotropy profile for the stars, $\beta(r)$, we can solve for the total gravitating mass."205 However. if woe had assumed a different (7) then we would rave found a differcut total mass profile (Strigari 2010).," However, if we had assumed a different $\beta(r)$ then we would have found a different total mass profile \citep{strigari}."206. If we for instance consider a Ieruquist density xofile with a > profile in agreement with uuucrical simulations aud observatious (ITansen&Piffaretti 2010).. hen the reconstructed nass is overestinated bv up to 1054. if the analvsis is made uncer the siuplifug assuniptiou j—0.," If we for instance consider a Hernquist density profile with a $\beta$ profile in agreement with numerical simulations and observations \citep{hansenpiff2007,host2009,wojtak2010}, then the reconstructed mass is overestimated by up to $40\%$, if the analysis is made under the simplifying assumption $\beta=0$."207 Also the derived inner density slope (from the tota lass) ds systematically found to be more shallow than the true slope is by up to 1056.," Also the derived inner density slope (from the total mass) is systematically found to be more shallow than the true slope is, by up to $10\%$."208 This means that if the true density slope is —1. then we will mcasure around 0.95. if we assumed ij= Oin the analysis.," This means that if the true density slope is $-1$, then we will measure around $-0.95$, if we assumed $\beta =0$ in the analysis."209 This max uo longer have to be the case., This may no longer have to be the case.210 If our attractor solutions also applies to stellar svstenis ii a dwarf ealaxy. or to the dynamics of the ealaxies in a galaxy cluster. then we have a unique connection between the 3 quantities. 5.& and ," If our attractor solutions also applies to stellar systems in a dwarf galaxy, or to the dynamics of the galaxies in a galaxy cluster, then we have a unique connection between the 3 quantities, $\gamma, \kappa$ and $\beta$."211Therefore. if we have measured (accurately) the stellar density aud dispersion profiles. then we do in principle kuow exactly what 2(r£0) looks like. aud we can then deduce the unique total mass profile.," Therefore, if we have measured (accurately) the stellar density and dispersion profiles, then we do in principle know exactly what $\beta(r)$ looks like, and we can then deduce the unique total mass profile."212 We have identified an attractor solution for dark matter structures., We have identified an attractor solution for dark matter structures.213 This implies that any dark matter structure which is repeatedly perturbed (e.g. through violent relaxtion durius mergi) and then allowed to relax (phase imix). will flow towards this l-cdimensional curve iu the 3-diuensional space spanned by the 2 radial derivatives of the deusity aud velocity dispersion. and the velocity anisotropy.," This implies that any dark matter structure which is repeatedly perturbed (e.g. through violent relaxtion during merging) and then allowed to relax (phase mix), will flow towards this 1-dimensional curve in the 3-dimensional space spanned by the 2 radial derivatives of the density and velocity dispersion, and the velocity anisotropy."214 This finding provides strong support for the idea that the universalities found in cosmological dark matter structures are a property of eravity. and not simply a result of similar accretion and morecr histories of different structures.," This finding provides strong support for the idea that the universalities found in cosmological dark matter structures are a property of gravity, and not simply a result of similar accretion and merger histories of different structures."215" This attractor solution effectively removes one deeree of freedom frou he Jeans equation. giving hope that we will eveutually be able to solve the Jeans equations analytically, and thereby truely understaud the origin of the universal profiles."," This attractor solution effectively removes one degree of freedom from the Jeans equation, giving hope that we will eventually be able to solve the Jeans equations analytically, and thereby truely understand the origin of the universal profiles."216 It is a pleasure to thaws Jens ITjorth for discussions., It is a pleasure to thank Jens Hjorth for discussions.217 The simulations were performed on the facilities provided by the Danish Center for Scientific Computing., The simulations were performed on the facilities provided by the Danish Center for Scientific Computing.218 The Dark Cosinology Centre is funded by the Danish National Research Foundation., The Dark Cosmology Centre is funded by the Danish National Research Foundation.219Moreover our results (e.g. temperatures) for C are consistent with those for C3 (?)..,Moreover our results (e.g. temperatures) for $C_2$ are consistent with those for $C_3$ \citep{Adamkovics2003}.220 The profile of the narrow DIB at apparently depends on the rotational temperature estimated from the dicarbon molecule. being broader for higher temperatures. which is characteristic of both C» and Cs.," The profile of the narrow DIB at apparently depends on the rotational temperature estimated from the dicarbon molecule, being broader for higher temperatures, which is characteristic of both $C_2$ and $C_3$ ."221 Cs was analysed by ? and ?.., $C_3$ was analysed by \citet{Adamkovics2003} and \citet{Oka2003}.222 They found that the column density and rotational temperature derived from Cs 15 well correlated with the same parameters of C»., They found that the column density and rotational temperature derived from $C_3$ is well correlated with the same parameters of $C_2$.223 It is interesting and important to analyse the simplest multicarbon chains (like, It is interesting and important to analyse the simplest multicarbon chains (like224(10)) of states into the linear combinations (11)) which are still vanishing.,\ref{eq:linearrel}) ) of states into the linear combinations \ref{eq:linearrel2}) ) which are still vanishing.225" Also. notice that ©, is self-adjoint. again because of the reality of 2? and it commutes with itself since The 2-netfunction W(s.s). defined in (2)). can be written now as More in general. we can celine so that Now. consider the free linear space A formed. by the (formal) linear combinations of spin networks. with complex coefficients There is a natural product defined on A bv ο5'=5Us’. and a natural star operation defined bv s*=s (Ilere we refer to spin networks labeled bv SU(2) representations ancl each representation of 90(2) is conjugate to itself."," Also, notice that $\hat\phi_{s}$ is self-adjoint, again because of the reality of $P$ and it commutes with itself since The 2-netfunction $W(s,s')$, defined in \ref{eq:W2}) ), can be written now as More in general, we can define so that Now, consider the free linear space $\cal A$ formed by the (formal) linear combinations of spin networks, with complex coefficients There is a natural product defined on $\cal A$ by $s\cdot s'=s\cup226s'$, and a natural star operation defined by $s^{*}=s$ (Here we refer to spin networks labeled by $SU(2)$ representations and each representation of $SU(2)$ is conjugate to itself."227 When spin networks are labeled by representations of groups which are not sell-conjugate the star operation should replaces representations will dual representations.), When spin networks are labeled by representations of groups which are not self-conjugate the star operation should replaces representations with dual representations.)228 We define the norm |L1]|= sup;|c;|., We define the norm $||A||=sup_{s}|c_{s}|$ .229 We obtain in this wav a C algebra structure on A., We obtain in this way a $C^{*}$ algebra structure on $\cal A$.230 The quantity W(s). defined in (17)). defines a linear functional on A.A straightforward. caleulation shows that the finelional is positive We can thus apply the Gelland-Naimark-Segal construction to the C algebra A and the positive linear functional WW. obtaining a Lilhert space H. a," The quantity $W(s)$, defined in \ref{eq:W}) ), defines a linear functional on $\cal A$.A straightforward calculation shows that the functional is positive We can thus apply the Gelfand-Naimark-Segal construction to the $C^*$ algebra $\cal A$ and the positive linear functional$W$ , obtaining a Hilbert space $\cal H$ , a"231The second column in Table 1 shows the small effect of this smoothing on the fit parameters.,The second column in Table \ref{tab:1t} shows the small effect of this smoothing on the fit parameters.232" The fit quality is slightly poorer after accounting for smoothing, suggesting that the emission lines may be emitted on a smaller scale than the broader X-ray emission(Section 3))."," The fit quality is slightly poorer after accounting for smoothing, suggesting that the emission lines may be emitted on a smaller scale than the broader X-ray emission(Section \ref{sect:profiles}) )."233 The goodness command gave 54.6 per cent of realisations of the best fitting model with a better fit than the data., The goodness command gave 54.6 per cent of realisations of the best fitting model with a better fit than the data.234" Since the lines are remarkably narrow, and the source extent is small, we can limit the velocity broadening."," Since the lines are remarkably narrow, and the source extent is small, we can limit the velocity broadening."235 We fitted the spectra in Section 2.1 with the model in assuming that the spectra were broadened by the thermal motion of the ions in the gas (version 12.5.0ah of fixed an error in the line widths) and line-of-sight velocity broadening added in quadrature., We fitted the spectra in Section \ref{sect:1t} with the model in assuming that the spectra were broadened by the thermal motion of the ions in the gas (version 12.5.0ah of fixed an error in the line widths) and line-of-sight velocity broadening added in quadrature.236 The spectral lines were used in the model., The spectral lines were used in the model.237 We show the change in fit statistic as a function of velocity in Fig. 3.., We show the change in fit statistic as a function of velocity in Fig. \ref{fig:broadening}.238 If we conservatively treat the cluster as a point source and do not use the model we obtain a 90 per cent upper limit of 274kms~!., If we conservatively treat the cluster as a point source and do not use the model we obtain a 90 per cent upper limit of $274\kmps$.239 This limit is improved to 182kms~! with the addition of the broadening model., This limit is improved to $182\kmps$ with the addition of the broadening model.240" Examining just the Fe-L spectral region between 9.6 and (rest), we obtain a limit of 214kms-!, and for the Ist order spectrum between 13.6 and22.4A, containing the strong O line, a limit of 3830kms~!."," Examining just the Fe-L spectral region between 9.6 and (rest), we obtain a limit of $214\kmps$, and for the 1st order spectrum between 13.6 and, containing the strong O line, a limit of $380\kmps$."241 We confirmed that such constraints are readily achievable with simulated spectra., We confirmed that such constraints are readily achievable with simulated spectra.242 We can limit how much gas can be cooling within temperature bins., We can limit how much gas can be cooling within temperature bins.243" We constructed a model using six cooling flow model components (constructed using models) with temperature ranges of 105.6>2.814070.35—0.0808 keV. We fixed the components to have the same metallicities, allowing individual elements to vary as in Section 2.1.."," We constructed a model using six cooling flow model components (constructed using models) with temperature ranges of $10 \rightarrow 5.6 \rightarrow 2.8 \rightarrow 1.4244\rightarrow 0.7 \rightarrow 0.35 \rightarrow 0.0808$ keV. We fixed the components to have the same metallicities, allowing individual elements to vary as in Section \ref{sect:1t}."245" The maximum cooling rates in Moyr~!, assuming isobaric cooling, were free parameters in the fit."," The maximum cooling rates in $\Msunpyr$, assuming isobaric cooling, were free parameters in the fit."246" The results are shown in Fig. 4,,"," The results are shown in Fig. \ref{fig:coolrate},"247 with and without the smoothing component., with and without the smoothing component.248" For comparison, we also fitted a model made up of an isothermal component plus a model cooling from its temperature to zero."," For comparison, we also fitted a model made up of an isothermal component plus a model cooling from its temperature to zero."249 The best fitting temperature of the isothermal component was 3.85 keV (close to the single temperature fits)., The best fitting temperature of the isothermal component was 3.85 keV (close to the single temperature fits).250 We obtain cooling rates from this temperature to zero of 70736Moyr!., We obtain cooling rates from this temperature to zero of $70^{+46}_{-56}\Msunpyr$.251 This rate is shown by the shaded bar in Fig. 4.., This rate is shown by the shaded bar in Fig. \ref{fig:coolrate}.252 The 90 per cent upper limit is 140Mayr! ., The 90 per cent upper limit is $140\Msunpyr$ .253 A Markov Chain Monte Carlo analysis produces, A Markov Chain Monte Carlo analysis produces254are systematically larger than the corresponding distances from the CLLA.,are systematically larger than the corresponding distances from the CLLA.255 This difference can be explained bv some unaccounted components. the presence of which may lead to an underestimation of photometric distances.," This difference can be explained by some unaccounted components, the presence of which may lead to an underestimation of photometric distances."256 Another reason lor this discrepancy is the photometric distance calibration adopted in the CLLA., Another reason for this discrepancy is the photometric distance calibration adopted in the CLLA.257 Most of the stars in our sample were examined for common proper molion components (Allenetal.2000:ZapateroOsorio&Martin2004).," Most of the stars in our sample were examined for common proper motion components \citep{allen,zapatero}."258". The speckle interferometric observations of 223 sample stars were carried out in 20062007 on the 6 m BTA telescope (Itastegaevetal.2007.2008) which dilfraction-limitec resolution is 0.023"" for A=550 nm and 0.033"" lor A=800 mnm."," The speckle interferometric observations of 223 sample stars were carried out in 2006–2007 on the 6 m BTA telescope \citep{rastegaev_2007,rastegaev_2008} which diffraction-limited resolution is $0.023 ''$ for $\lambda=550$ nm and $0.033 ''$ for $\lambda=800$ nm."259 Most of the observations were carried out using the svstem (Maksimovetal.2009) based on a 512x512 EAICCD (a CCD [eaturing on-chip multiplication gain) with high «quantum efficiency. ancl linearity.," Most of the observations were carried out using the system \citep{maksimov}260 based on a $\times$ 512 EMCCD (a CCD featuring on-chip multiplication gain) with high quantum efficiency and linearity."261" This svstem allowed us to detect objects with magnitude differences between the components ol up to Am=5"".", This system allowed us to detect objects with magnitude differences between the components of up to $\triangle m = 5^{m}$.262 Takine into account the limiting stellar magnitude of our sample (my« 127).detected secondary component can be as faint as 177.," Taking into account the limiting stellar magnitude of our sample $\mathrm{m_V}<12^{m}$ ),detected secondary component can be as faint as $17^{m}$."263" The 4.4” field of view of our system allows detection of secondary components at a separation of3"" from the primary star.", The $4.4''$ field of view of our system allows detection of secondary components at a separation of$3''$ from the primary star.264 The speckle interferograms were recorded using five filters: 545/30. 550/20. /600/40. 800/100/ and 800/110 nm (the first number indicates the central wavelength of the filler. the second the hall-width of the filters bandwidth) with the exposures of 5 to 20 ms.," The speckle interferograms were recorded using five filters: $545/30$, $550/20$, $600/40$, $800/100$ and $800/110$ nm (the first number indicates the central wavelength of the filter, the second — the half-width of the filter's bandwidth) with the exposures of 5 to 20 ms."265 For each object. we accumulated Irom 500 to 2000 short exposure images depending on weather conditions.," For each object, we accumulated from 500 to $2\ 000$ short exposure images depending on weather conditions."266" The observations were made with an average seeing of 1.5"".", The observations were made with an average seeing of $1.5''$.267" The accuracy of our speckle interferogram processing method (Dalegaοἱal.2002) may be as οσους as 0.02"". 0.001"". and 0.1* for the component maenitude dilference. angular separation and position angle respectively."," The accuracy of our speckle interferogram processing method \citep{balega_2002} may be as good as $0.02^m$ $0.001''$ , and $0.1^{\circ}$ for the component magnitude difference, angular separation and position angle respectively."268baselines.,baselines.269 Alternatively it may indicate that these largest scales are simply not present in the galaxies under consideration., Alternatively it may indicate that these largest scales are simply not present in the galaxies under consideration.270 It is clear though that at large scales there is good agreement., It is clear though that at large scales there is good agreement.271" As the spatial scale probed approaches the size of the natural-weighted beam however, the power in the lintegrated moment maps begins to decrease compared to the mmaps."," As the spatial scale probed approaches the size of the natural-weighted beam however, the power in the integrated moment maps begins to decrease compared to the maps."272 At about ~1.5 times the natural-weighted beam size we start to see a significant deviation., At about $\sim1.5$ times the natural-weighted beam size we start to see a significant deviation.273 Also shown in Figure 19 isa power law with slope —3 which is a reasonable description for the power spectrum at intermediate scales., Also shown in Figure \ref{fig:powerspec} is a power law with slope $-3$ which is a reasonable description for the power spectrum at intermediate scales.274 It is consistent with values found in other galaxies(?)., It is consistent with values found in other galaxies.275". We will not here to relate the power law to tthe turbulence or the attemptenergy input of the ISM, but we slopedraw attention to the fact that as the small-scale power in the mmaps starts to fall away, the ppower spectrum continues to follow this power-law behavior."," We will not attempt here to relate the power law slope to the turbulence or the energy input of the ISM, but we draw attention to the fact that as the small-scale power in the maps starts to fall away, the power spectrum continues to follow this power-law behavior."276 This indicates that the probe real small-scale structure more efficiently than the classical mmapsmmaps., This indicates that the maps probe real small-scale structure more efficiently than the classical maps.277 Note that the power in the mmaps only starts to drop away at scale sizes of half a natural-weighted beam., Note that the power in the maps only starts to drop away at scale sizes of half a natural-weighted beam.278 Scales probed in classical, Scales probed in classical279Since the late 1950s. it has been realized that neutrou-star iuterior may consist of a πο of quantum fiuids (see Shapiro Toukolsky 1983 for a review).,"Since the late 1950's, it has been realized that neutron-star interior may consist of a number of quantum fluids (see Shapiro Teukolsky 1983 for a review)."280 Curreuthe it is thought that both neutron superfiuid and proton superconductor are likely to coexist in the neutron-star cores (see. 6.8.. Link 2007 Or a discussion).," Currently, it is thought that both neutron superfluid and proton superconductor are likely to coexist in the neutron-star cores (see, e.g., Link 2007 for a discussion)."281 Several researchers have arec that if the proton supercouductivityv were| of the type IL then t1c superconductor fluxtuvos would couple strongv to the notron superfluid vortices.," Several researchers have argued that if the proton superconductivity were of the type II, then the superconductor fluxtubes would couple strongly to the neutron superfluid vortices."282 This line of rTOasonding is based ou the fact that nuclear forces contain velocity-depeuceut terms. which results i the οιαπλο! of protons in he jeutron supercurrent (Alpar. Langer. Sauls. 1958).," This line of reasoning is based on the fact that nuclear forces contain velocity-dependent terms, which results in the entrainment of protons in the neutron supercurrent (Alpar, Langer, Sauls, 1984)."283 Therefore. the vortices are sheathe« by charged curreits cutrained in the suοΠα flow. aud are strongv maenetized.," Therefore, the vortices are sheathed by charged currents entrained in the superfluid flow, and are strongly magnetized."284 Magnetic fiuxtubes interact stroely with t1C Wagjetized vortices. snilar to the wav in wuch ειο fluxtubes interact between each other (Rileruniui. Zhu. Chen 1998. and references therein).," Magnetic fluxtubes interact strongly with the magnetized vortices, similar to the way in which the fluxtubes interact between each other (Ruderman, Zhu, Chen 1998, and references therein)."285 As a result f this couphue. the vortices ect strongly piune« to +ie proton-clectron plasma in the core.," As a result of this coupling, the vortices get strongly pinned to the proton-electron plasma in the core."286 Such pininue would have nu»ortant iuplications for he neutron-star plienomenology., Such pinning would have important implications for the neutron-star phenomenology.287 Ruclerman. Zhu. Chen (1998) have argued. that t1ο vortex-pinning iu the core lay be responsible for he observed elitejes in the pulsar rotation rates.," Ruderman, Zhu, Chen (1998) have argued that the vortex-pinning in the core may be responsible for the observed glitches in the pulsar rotation rates."288 Link (2003) has considerec the effect of the vortex-fluxtube interaction ou the cdyvuamics of the precessineo pulsar PSR 1828-11 (observed by Stairs. Lyuc. Shelu 2000).," Link (2003) has considered the effect of the vortex-fluxtube interaction on the dynamics of the precessing pulsar PSR 1828-11 (observed by Stairs, Lyne, Shemar 2000)."289 Building ou the theoretica WNOYs by Shaham (1977) and Sedrakian. Wasscermial.. Cordes (1999). le has concluded that the interaction. if oesenut. would ultimaelv lead to the fast precession.," Building on the theoretical work by Shaham (1977) and Sedrakian, Wasserman, Cordes (1999), he has concluded that the interaction, if present, would ultimately lead to the fast precession."290 Since PSR JBN28-11 is precessing slowly and persisteuthy. Links (2003) las argued that the core vortex pinning is excluded by tιο observations and hence that either the proton sitperconductor migit be of tvpe L or that both proton aud neun condensates do uot coexist iuside that pulsar.," Since PSR 1828-11 is precessing slowly and persistently, Link (2003) has argued that the core vortex pinning is excluded by the observations and hence that either the proton superconductor might be of type I, or that both proton and neutron condensates do not coexist inside that pulsar."291 While Links argmeut is suggestive. we believe if is prenature ο rule out strong vortex piunius iu the cores of a] neutron stars.," While Link's argument is suggestive, we believe it is premature to rule out strong vortex pinning in the cores of all neutron stars."292 Iu this paper we consier hydromaguctic waves in the case when the neutron vortices are stronglv piined to the protou-clectrou plasma iu he core., In this paper we consider hydromagnetic waves in the case when the neutron vortices are strongly pinned to the proton-electron plasma in the core.293 We have 2 nun astroplivsical motivations or studying this problem., We have 2 main astrophysical motivations for studying this problem.294 The first one is due to the fairly receut observations, The first one is due to the fairly recent observations295"Even galaxies with little star formation activity continue to evolve, as evidenced by the substantial increase of their cosmic stellar mass density over the past 7 billion years","Even galaxies with little star formation activity continue to evolve, as evidenced by the substantial increase of their cosmic stellar mass density over the past 7 billion years"296D aud V stellar maeuitudes.,$B$ and $V$ stellar magnitudes.297" Using these parameters, we calculate the effective temperature and Imuuinositv via the Zaustra method (I&aloey 1983). ("," Using these parameters, we calculate the effective temperature and luminosity via the Zanstra method (Kaler 1983). ("2982) Dx locating the central star iu the logTig logL plane. we derive its mass (Mog) from comparison with a set of evolutionary tracks (Stanghellini Renzini 1993). (,"2) By locating the central star in the $\log T_{\rm eff}$ $\log L$ plane, we derive its mass $M_{\rm CS}$ ) from comparison with a set of evolutionary tracks (Stanghellini Renzini 1993). ("2993) Usine the initial massfinal mass relation. we colpute the progenitor mass. that is the stellar mass on the main sequence (A/S).,"3) Using the initial mass–final mass relation, we compute the progenitor mass, that is the stellar mass on the main sequence $M_{\rm MS}$ )."300 The stellar properties adopted for the PNe detected iu PCO and the derived. values of the progenitor mass are eiven in Table 3., The stellar properties adopted for the PNe detected in $^{13}$ CO and the derived values of the progenitor mass are given in Table 3.301 Details ou the individual objects are eiven iu the Appoeudix., Details on the individual objects are given in the Appendix.302 Let us exanune the uncertainty involved iu the final mass calculations., Let us examine the uncertainty involved in the final mass calculations.303 Estimates of the stellar temperature and Iuninositv eiven iu Table 3 are affected by errors in magnitudes. fixes. diameters aud extinctions.," Estimates of the stellar temperature and luminosity given in Table 3 are affected by errors in magnitudes, fluxes, diameters and extinctions."304 However. these quautitics are usually determined with οσους accuraev (~520%). so that the wnecrtainty in the erived mass of the central stars does not exceed ~15 or (0.02AZ.," However, these quantities are usually determined with good accuracy $\sim 5-20\%$ ), so that the uncertainty in the derived mass of the central stars does not exceed $\sim 15$, or $\sim 0.02$."305.. The values given iu the table do ixt oeiclude errors on the distances to the PN. which can e oeitrinsically hiehl (up to 505€)) but ave difficult to estimate on an individual basis.," The values given in the table do not include errors on the distances to the PN, which can be intrinsically high (up to ) but are difficult to estimate on an individual basis."306 To inter the main sequence masses. we have used the oeitial massfinal mass relation given by Uervig (1996).," To infer the main sequence masses, we have used the initial mass–final mass relation given by Hervig (1996)."307 This relation differs from that of Weideiiaun (LOST) vdopted in GSTP., This relation differs from that of Weidemann (1987) adopted in GSTP.308 We preferred Heorvigs prescription since it is based ou reliable observations of cluster white cdawarts. although the formal errors on the final mass are still substantial. and can amount to about 01 ALL.," We preferred Hervig's prescription since it is based on reliable observations of cluster white dwarfs, although the formal errors on the final mass are still substantial, and can amount to about 0.1 $M_\odot$."309 Since we derive initial masses from final masses. the errors ou the Ormer quantity can be even larger.," Since we derive initial masses from final masses, the errors on the former quantity can be even larger."310 Quantitatively. we assign a formal error to the main sequence mass of AAAsτσ1.5 stor low values of the initial mass GM<2 AL.) anda simaller error (AAAs&0.75 i) for higher masses.," Quantitatively, we assign a formal error to the main sequence mass of $\Delta M_{\rm MS}\simeq 1.5$ for low values of the initial mass $M_{\rm MS}<2$ ), and a smaller error $\Delta M_{\rm MS}\simeq3110.75$ ) for higher masses."312 This difference is due to the change of the slope of the initial mass - final mass relation at about 2M: sinaller masses are more sensitive to the adopted relation. and the uncertainty is correspondingly larger.," This difference is due to the change of the slope of the initial mass - final mass relation at about 2: smaller masses are more sensitive to the adopted relation, and the uncertainty is correspondingly larger."313 Tow do we interpret our results on the irafios in the framework of stellar uucleosvuthesis?, How do we interpret our results on the ratios in the framework of stellar nucleosynthesis?314 To help answering this question. we combine the information provided by the observed lisotopic ratios with the mass estimates of the progenitors of the PNe. and with the predictions of some representative stellar uucleosvuthesis models.," To help answering this question, we combine the information provided by the observed isotopic ratios with the mass estimates of the progenitors of the PNe, and with the predictions of some representative stellar nucleosynthesis models."315 Since the formation of a PN takes place at the eund of the AGB phase. the significant comparison is between the observed abundancees and those predicted for the stellar ejecta at the AGB tip.," Since the formation of a PN takes place at the end of the AGB phase, the significant comparison is between the observed abundances and those predicted for the stellar ejecta at the AGB tip."316 Uufortunately. no," Unfortunately, no"317p and d-sates of Bain.,$p$ and $d$ -sates of Ba.318 Then. we obtain the transition probabilities in the tensorial irreducible basis.," Then, we obtain the transition probabilities in the tensorial irreducible basis."319 Afterwards. these propabilities are integrated over impact parameters and Maxwellian distribution of relative velocities to obtain the depolarization and the transfer rates.," Afterwards, these propabilities are integrated over impact parameters and Maxwellian distribution of relative velocities to obtain the depolarization and the transfer rates."320 We perform calculations varying the temperature to obtain the best analytical fit to the collistonal rates., We perform calculations varying the temperature to obtain the best analytical fit to the collisional rates.321The spectra of our new program stars have been obtained in (he same fashion as (hose studied by C2003. using the Center for Astrophysics Digital Speedometers (Latham 1935. 1992). primarily with the 1.5-m Wveth reflector at (he Oak Ridge Observatory in Harvarel.Massachusetts’. as well as the 1.5-m Tillinghast reflector and the MMT. instruments atop Alt. Hopkins in Arizona.,"The spectra of our new program stars have been obtained in the same fashion as those studied by C2003, using the Center for Astrophysics Digital Speedometers (Latham 1985, 1992), primarily with the 1.5-m Wyeth reflector at the Oak Ridge Observatory in Harvard, as well as the 1.5-m Tillinghast reflector and the MMT instruments atop Mt. Hopkins in Arizona."322 The Tillinghast reflector was especially important for the stars south of —20° aand north of +62° iin declination., The Tillinghast reflector was especially important for the stars south of $-20$ and north of $+62$ in declination.323 Also as before. the wavelength coverage is 45A.. centered near 5187À.. with a resolution of 8.5+.," Also as before, the wavelength coverage is 45, centered near 5187, with a resolution of 8.5."324. The signal-to-noise ratio varied [rom 10 to 50 per resolution element. will a (vpical value of about 15.," The signal-to-noise ratio varied from 10 to 50 per resolution element, with a typical value of about 15."325 C2003 described in detail the measurement. of the radial velocities., C2003 described in detail the measurement of the radial velocities.326 A grid of model almospheres. defined by Zar. log g. and. |Fe/II] values. was computed.," A grid of model atmospheres, defined by $T_{\rm eff}$, log $g$, and [Fe/H] values, was computed."327 The grid spacing in temperature was 250 Ix. 0.5 dex in log g. and 0.5 dex in metallicity for [Fe/H] <—1.0.," The grid spacing in temperature was 250 K, 0.5 dex in log $g$, and 0.5 dex in metallicity for [Fe/H] $\leq\ -1.0$."328 functions and elemental abundances in whieh all the a elements (O. Ne. Meg. SiS. Ca. and Ti) were enhanced by 0.4 dex relative to the solar abundances.," functions and elemental abundances in which all the $\alpha$ "" elements (O, Ne, Mg, Si, S, Ca, and Ti) were enhanced by 0.4 dex relative to the solar abundances."329 More details may be found in Nórrdstrom et ((1994) and C2003., More details may be found in Nörrdstrom et (1994) and C2003.330 The program SYNTIIE was used to compute svnthetic spectra in Che waveleng(h range 5146-5229A.., The program SYNTHE was used to compute synthetic spectra in the wavelength range 5146-5229.331 In the case of the Sun. we obtained an excellent line-by-Iine match between the ]vurucz solar model andsynfhetie spectrum compared with theobserved solar [πι spectrum (Ixurucz et 11984).," In the case of the Sun, we obtained an excellent line-by-line match between the Kurucz solar model and spectrum compared with the solar flux spectrum (Kurucz et 1984)."332 SYNTIIE computes specilic intensity al 17 different emergent angles across the stellar disk. and integration over the disk. including the effects of stellar rotation. vields (he svntlietic [Inx spectrum.," SYNTHE computes specific intensity at 17 different emergent angles across the stellar disk, and integration over the disk, including the effects of stellar rotation, yields the synthetic flux spectrum."333 5YNTIIE enables us to include the effects of instrumental resolution. which we chose to be a Gaussian with a FWILM o£ 8.5f... which is appropriate io the instrumentation we enploved.," SYNTHE enables us to include the effects of instrumental resolution, which we chose to be a Gaussian with a FWHM of 8.5, which is appropriate to the instrumentation we employed."334 The adopted microturbulent. velocities were 2HL. and macroturbulent velocities were 3|.," The adopted microturbulent velocities were 2, and macroturbulent velocities were 3."335".. At each combination of temperature. eravity. and metallicity in our grid we computed svnthetic spectra employing a wide range of rotational broadening profiles. with Veo, = O0. 1. 2. 4. 6. 5. 10. 12. 16. 20. 25. 30. 35. 40. 50. 60. το SO. 90. LOO. 110. 120. and 1401."," At each combination of temperature, gravity, and metallicity in our grid we computed synthetic spectra employing a wide range of rotational broadening profiles, with $V_{\rm rot}$ = 0, 1, 2, 4, 6, 8, 10, 12, 16, 20, 25, 30, 35, 40, 50, 60, 70, 80, 90, 100, 110, 120, and 140."336 Once the stellar parameters. Zr. log g. and. [Fe/II] had been estimated. we relied on ihe model atmosphere grid point closest in (hese variables. paving special attention to the primary variable. temperature.," Once the stellar parameters, $T_{\rm eff}$, log $g$, and [Fe/H] had been estimated, we relied on the model atmosphere grid point closest in these variables, paying special attention to the primary variable, temperature."337 In the case of more than one close match. we emploved the template that gave the hiehest value for the peak correlation. averaged over all the observed," In the case of more than one close match, we employed the template that gave the highest value for the peak correlation, averaged over all the observed"338matter distribution.,matter distribution.339 The results of ? apply only to isolated galaxies. while the cross-correlation approach applies to any parent galaxy.," The results of \citet{chen_etal06} apply only to isolated galaxies, while the cross-correlation approach applies to any parent galaxy."340 This suggests that the environment of parent galaxies is relatively unimportant., This suggests that the environment of parent galaxies is relatively unimportant.341 My result is also generally consistent with cluster-sized simulations which include cooling and star formation., My result is also generally consistent with cluster-sized simulations which include cooling and star formation.342 In such simulations. while the distribution of dark matter subhalos is less concentrated than the dark matter. the stellar components of satellite galaxies more closely follow the dark matter distribution (??)..," In such simulations, while the distribution of dark matter subhalos is less concentrated than the dark matter, the stellar components of satellite galaxies more closely follow the dark matter distribution \citep{nagai_kravtsov05,maccio_etal06}."343 However. the possible discrepancy in values for f£ and fa. suggest that the spatial distribution of a class of satellite objects may depend on the size of the host halo: galaxy-sized halos. groups and clusters.," However, the possible discrepancy in values for $f$ and $f_{\rm cross}$ suggest that the spatial distribution of a class of satellite objects may depend on the size of the host halo: galaxy-sized halos, groups and clusters."344 ? suggest that the projected radial distribution of satellite galaxies depends upon satellite color. such that redder satellites have a significantly steeper density profile than bluer satellites.," \citet{chen08} suggest that the projected radial distribution of satellite galaxies depends upon satellite color, such that redder satellites have a significantly steeper density profile than bluer satellites."345 A cross-correlation analysis with a different selection function for satellites then should show such a color dependence., A cross-correlation analysis with a different selection function for satellites then should show such a color dependence.346 In order to achieve better constraints on the spatial distribution of satellite galaxies and its environmental and color dependency. larger data sets are required.," In order to achieve better constraints on the spatial distribution of satellite galaxies and its environmental and color dependency, larger data sets are required."347 One possibility for increasing the volume of data is to apply a similar analysis to the photometric redshift catalog. instead of limiting the sample to objects with spectroscopic redshifts.," One possibility for increasing the volume of data is to apply a similar analysis to the photometric redshift catalog, instead of limiting the sample to objects with spectroscopic redshifts."348 The SDSS photometric redshift catalog goes far deeper than the spectroscopic catalog and without the spectroscopic catalog’s fiber collision problem., The SDSS photometric redshift catalog goes far deeper than the spectroscopic catalog and without the spectroscopic catalog's fiber collision problem.349 However. photometric redshifts are less accurate and may complicate the statistical error analysis.," However, photometric redshifts are less accurate and may complicate the statistical error analysis."350 Given that future large surveys such as the Dark Energy Survey and the Panoramic Survey Telescope Rapid Response System (Pan-STARRS) which will produce significantly more photometric redshifts than currently available. it would be interesting to perform a analysis with photometric redshifts.," Given that future large surveys such as the Dark Energy Survey and the Panoramic Survey Telescope Rapid Response System (Pan-STARRS) which will produce significantly more photometric redshifts than currently available, it would be interesting to perform a cross-correlation analysis with photometric redshifts."351Internal dust might allect seriously the observed. properties of distant objects. ancl observations have indeed. found evidence of dust in high-z objects.,"Internal dust might affect seriously the observed properties of distant objects, and observations have indeed found evidence of dust in $z$ objects."352 Several active galaxies with z 2 have been detected at submillimeter wavelengths. suggesting the presence of large amounts of cust (Mausx1077 Mo: see Hughes 1996 for a recent review).," Several active galaxies with $z>$ 2 have been detected at submillimeter wavelengths, suggesting the presence of large amounts of dust $_{dust} \sim 10^{8-9}$ $_{\odot}$; see Hughes 1996 for a recent review)."353 The reddening of the background quamus and their metal abundances suggest the presence of dust in damped Lye absorption systems (Pettini et al., The reddening of the background quasars and their metal abundances suggest the presence of dust in damped $\alpha$ absorption systems (Pettini et al.354 1994: Per Fall 1996)., 1994; Pei Fall 1996).355 Also. the UV polarization of propertieshigh-z radio galaxies can be explained in term of dust scattering. suggesting a significant amount of dust in their ISM (Cimatti 1996 and references therein).," Also, the UV polarization properties of $z$ radio galaxies can be explained in term of dust scattering, suggesting a significant amount of dust in their ISM (Cimatti 1996 and references therein)."356 A substantial amount of cust. is also expected in evolutionary models of spheroidal galaxies at high-z (Franceschini et al., A substantial amount of dust is also expected in evolutionary models of spheroidal galaxies at $z$ (Franceschini et al.357 1994: Mazzei De Zotti 1996 and references therein)., 1994; Mazzei De Zotti 1996 and references therein).358 Although dust. is likely to be present in most high-z systems. no information is available neither on its spatial distribution. nor about howrealise dust. extinction. can aleet our view of distant galaxies.," Although dust is likely to be present in most $z$ systems, no information is available neither on its spatial distribution, nor about how dust extinction can affect our view of distant galaxies."359 In Lact. extinction is usually treated in a simplistic way. neelecting the contribution of scattering. and assuming naive spatial distributions (e.g. uniform foreground. screens or infinite slabs).," In fact, extinction is usually treated in a simplistic way, neglecting the contribution of scattering, and assuming naive spatial distributions (e.g. uniform foreground screens or infinite slabs)."360 Only recentlv. more realistic models have. been developed (IxvlIafis Baheall 1987: Bruzualet al.," Only recently, more realistic models have been developed (Kylafis Bahcall 1987; Bruzual et al."361 1988: Witt et al., 1988; Witt et al.362 1992: Byun et al., 1992; Byun et al.363 1994: Wise Silva 1996: Bianchi. Ferrara Giovanardi 1996. hereafter BEG).," 1994; Wise Silva 1996; Bianchi, Ferrara Giovanardi 1996, hereafter BFG)."364 Although these models have shown that cust scattering. plavs an important role by reducing the ellects. of dust. absorption. no systematic studies of the ellects of the extinction on the observed. colours of high-z. galaxies have been pertormed.," Although these models have shown that dust scattering plays an important role by reducing the effects of dust absorption, no systematic studies of the effects of the extinction on the observed colours of $z$ galaxies have been performed."365 Understanding these elfects is crucial in cosmology and galaxy evolution studies., Understanding these effects is crucial in cosmology and galaxy evolution studies.366 Vor instance. a relevant problem arises in the age estimates of high-z galaxies.," For instance, a relevant problem arises in the age estimates of $z$ galaxies."367 When deep spectroscopy. is not available. no information on the stellar continuum and absorption features are obtainable. and the age estipiates are based. solely on the fitting of the broad-band photomoetric Spectral Energy. Distributions (SEDs) with svnthetic stellar population spectra.," When deep spectroscopy is not available, no information on the stellar continuum and absorption features are obtainable, and the age estimates are based solely on the fitting of the broad-band photometric Spectral Energy Distributions (SEDs) with synthetic stellar population spectra."368 In a simplistic scenario where the extinction is entirely. due to absorption (the screen model). itis well known that the colours ofa reddened voung galaxy," In a simplistic scenario where the extinction is entirely due to absorption (the screen model), it is well known that the colours of a reddened young galaxy"369"As a measure of bar strength. we use Q,=max(hy;/«Pi>). the maximum of tangential force amplitude relative o the mean axisvmmetric radial force. evaluated at the region of the bar.","As a measure of bar strength, we use $Q_g = \max({F_T/<F_R>})$, the maximum of tangential force amplitude relative to the mean axisymmetric radial force, evaluated at the region of the bar."370 The mass-to-light ratio (AL/£) is assumed to be constant. and the vertical scale height of the disk (ancl bar) is estimated from the exponential scale length. using a Hubble tvpe dependent mean ratio.," The mass-to-light ratio $M/L$ ) is assumed to be constant, and the vertical scale height of the disk (and bar) is estimated from the exponential scale length, using a Hubble type dependent mean ratio."371 Also. the dillerent 3D density cüstribution of the bulge is corrected. based on bulge models obtained roni decompositions.," Also, the different 3D density distribution of the bulge is corrected, based on bulge models obtained from decompositions."372" The elfect of including dark halo force field was also investigated. but its influence on (Q, appeared to »e Insignificant at the bar region (63]. 17])."," The effect of including dark halo force field was also investigated, but its influence on $Q_g$ appeared to be insignificant at the bar region ([63], [17])."373" These calculations also give a proxy for the bar length. rQy. which is the radius where the maximum tangential force (Q, occurs."," These calculations also give a proxy for the bar length, $rQ_g$, which is the radius where the maximum tangential force $Q_g$ occurs."374" Dar lengths were estimated also from the phases of the ο Fourier amplitude » assuming that it is maintained nearly constant in the bar region (a correlation between this bar length and. rQ,y was shown by 48]).", Bar lengths were estimated also from the phases of the $A_2$ Fourier amplitude by assuming that it is maintained nearly constant in the bar region (a correlation between this bar length and $rQ_g$ was shown by [48]).375 The maximum of m-—2 Fourier amplitude. clo. was used as an estimate of the relative brightness of the bar.," The maximum of $m$ =2 Fourier amplitude, $A_2$, was used as an estimate of the relative brightness of the bar."376 These properties are shown in Figure 5 in Laurikainen et al. , These properties are shown in Figure 5 in Laurikainen et al. [377"51] as à function of Hubble: bars grow in length and in relative xightness (le) towards the earlv-tvpe galaxies. but for €, the trend is opposite.","51] as a function of Hubble: bars grow in length and in relative brightness $A_2$ ) towards the early-type galaxies, but for $Q_g$ the trend is opposite."378 Notice that although the bar ellipticity (shown in the same figure) correlates with Quy (91]). it has no systematic correlation with the Hubble tvpe.," Notice that although the bar ellipticity (shown in the same figure) correlates with $Q_g$ ([91]), it has no systematic correlation with the Hubble type."379 In Laurikainen et el. , In Laurikainen et el. [380"48 ] the tendency. of weakening bar strengths ((,) towards the early-type galaxies was explained by a dilution elfect due to the more massive bulges in the early-tvpe galaxies (c.g. the average Q, parameter may decrease even if the average se amplitude increases. since the bulge contribution to radial force becomes more important toward earlier tvpes).","48 ] the tendency of weakening bar strengths $Q_g$ ) towards the early-type galaxies was explained by a dilution effect due to the more massive bulges in the early-type galaxies (e.g. the average $Q_g$ parameter may decrease even if the average $A_2$ amplitude increases, since the bulge contribution to radial force becomes more important toward earlier types)."381" There exist à correlation also between Q, and slo. but for the above reason the correlations are dilferent for the early and late-type galaxies (sce Fig."," There exist a correlation also between $Q_g$ and $A_2$, but for the above reason the correlations are different for the early and late-type galaxies (see Fig."382 S in. 48])., 8 in [48]).383 The obtained tendeney for bar lengths was originally shown by Elmegreen & Elmoegreen *Ji, The obtained tendency for bar lengths was originally shown by Elmegreen $\&$ Elmegreen [28].384 For à sub-sample of 26 barred. galaxies in NIRSOS. the radial sto profiles were fitted by single (SC) ancl double (DC) Gaussian functions by Buta et al. ," For a sub-sample of 26 barred galaxies in NIRS0S, the radial $A_2$ profiles were fitted by single (SG) and double (DG) Gaussian functions by Buta et al. ["38518].,18].386 Ht appeared that 65% of the bars in SO0-S0/a galaxies have single Gaussian profiles. whereas 35% are best fitted by two Gaussian functions.," It appeared that $\%$ of the bars in S0-S0/a galaxies have single Gaussian profiles, whereas $\%$ are best fitted by two Gaussian functions."387 Evpical examples of such profiles are shown in Figure 4., Typical examples of such profiles are shown in Figure 4.388 The galaxies with DG bars tvpically have also significant higher Fourier modes in the bar region (m4. 6. 8). in addition to m=2 (51]).," The galaxies with DG bars typically have also significant higher Fourier modes in the bar region $m$ =4, 6, 8), in addition to $m$ =2 ([51])."389 Lt was discussed by Buta et al., It was discussed by Buta et al.390 that the DG-profiles are similar to those predicted by the simulation models (e.g... 5]) in which the bar transfers a large amount of angular momentum to the halo.," that the DG-profiles are similar to those predicted by the simulation models (e.g., [5]) in which the bar transfers a large amount of angular momentum to the halo."391 An attempt to associate the DC-profiles to specific morphological structures was made by Laurikainen et al. , An attempt to associate the DG-profiles to specific morphological structures was made by Laurikainen et al. [39251] who suggested that the fat or double-peaked Gaussian amplitude profiles ave due to two bar components. a long and narrow bar. and a shorter component in the inner parts of the bar (or an inner oval).,"51] who suggested that the fat or double-peaked Gaussian amplitude profiles are due to two bar components, a long and narrow bar, and a shorter component in the inner parts of the bar (or an inner oval)."393 DG bars were found to be more prominent. not only in terms of Qu. but also in ele and bar length 51].," DG bars were found to be more prominent, not only in terms of $Q_g$, but also in $A_2$ and bar length [51]."394 In Laurikainen et al. , In Laurikainen et al. [39554] these inner bar components were associated mainly with barlenses (though some of them can be ovals). which are found to appear in 30% of barred. SO-SO/a ealaxies in NIIRSOS. A good example is NGC 4314 (Fig.,"54] these inner bar components were associated mainly with barlenses (though some of them can be ovals), which are found to appear in $\%$ of barred S0-S0/a galaxies in NIRS0S. A good example is NGC 4314 (Fig."396 1). in which the barlens is the fat elongated structure inside the bar.," 1), in which the barlens is the fat elongated structure inside the bar."397 Erwin et al. , Erwin et al. [39830] have discussed two SOs. in which a superposition of a classical and a pseudo-bulge was suggested.,"30] have discussed two S0s, in which a superposition of a classical and a pseudo-bulge was suggested."399 These galaxies are NGC 2787 and NGC 3945. which form. part of NIBSOS. In both galaxies the component interpreted as a pseudo-bulge by Erwin et al.," These galaxies are NGC 2787 and NGC 3945, which form part of NIRS0S. In both galaxies the component interpreted as a pseudo-bulge by Erwin et al.,"400 is called as a fat inner bar component in 51]. and more recently delined as a barlens by us 54].," is called as a fat inner bar component in [51], and more recently defined as a barlens by us [54]."401 The above question was recently made by Buta et al. , The above question was recently made by Buta et al. [402S6] with the main emphasis to test the hypothesis by Bournaud. & Combes 14]. in which multiple bar episodes are expected in the Hubble time.,"86] with the main emphasis to test the hypothesis by Bournaud $\&$ Combes [14], in which multiple bar episodes are expected in the Hubble time."403 In this scenario bars form and evolve in galaxies when they have gas. and the evolution stops when the gas in used in star formation.," In this scenario bars form and evolve in galaxies when they have gas, and the evolution stops when the gas in used in star formation."404 These stars are then transferred. into he bulge. for example by bars or spiral arms in the central regions of the galaxies.," These stars are then transferred into the bulge, for example by bars or spiral arms in the central regions of the galaxies."405 When the central mass concentration ormed by star formation becomes very high. the bar will be destroved.," When the central mass concentration formed by star formation becomes very high, the bar will be destroyed."406" Therefore. if bar strength varies over time the relative requency of galaxies in each (2, bin tells us the relative amount of time a galaxy spends in a certain bar state (strong. weak. non-barred)."," Therefore, if bar strength varies over time the relative frequency of galaxies in each $Q_g$ bin tells us the relative amount of time a galaxy spends in a certain bar state (strong, weak, non-barred)."407 For spirals this was first tested by Block et al. , For spirals this was first tested by Block et al. [408SO].,80].409 They suggested that galaxies might have doubled their mass in 102 vears (see Fig., They suggested that galaxies might have doubled their mass in $^{10}$ years (see Fig.410 5. Left panel). evidenced by the extended: tail towards strong bars. and the lack of weak bars. which features are predicted in the strong gas accretion models by Bournaud & Combes 14].," 5, left panel), evidenced by the extended tail towards strong bars, and the lack of weak bars, which features are predicted in the strong gas accretion models by Bournaud $\&$ Combes [14]."411" This test was later repeated by Buta. Laurikainen and Salo 87] for the same galaxy sample. but using the refined. bar orque method described in the previous section (we also discussed why the obtained (, distribution was different [rom that w Block et al. "," This test was later repeated by Buta, Laurikainen and Salo [87] for the same galaxy sample, but using the refined bar torque method described in the previous section (we also discussed why the obtained $Q_g$ distribution was different from that by Block et al. ["412SO.,80]).413 In Buta et al. , In Buta et al. [41417] the bar and spiral llüxes were additionally separated from each other.,17] the bar and spiral fluxes were additionally separated from each other.415" Although the correction allected €, in à few individual cases having very strong spiral armis. it barely alfected the Q, distribution."," Although the correction affected $Q_g$ in a few individual cases having very strong spiral arms, it barely affected the $Q_g$ distribution."416" The refined Q, distribution (SY]. 17]) has a large number of weak bars lacking from that obtained by Block et al."," The refined $Q_g$ distribution ([87], [17]) has a large number of weak bars lacking from that obtained by Block et al."417 Likewise. it yas a slightly smaller number of very strong bars.," Likewise, it has a slightly smaller number of very strong bars."418" In fact. the obtained Q, distribution (sce Fig."," In fact, the obtained $Q_g$ distribution (see Fig."419 Sa in. 17]) largely resembles he non-accretion model by Bournaud & Combes shown in Figure 5 (left panel). thus supporting the view that bars in spirals are fairly robust.," 8a in [17]) largely resembles the non-accretion model by Bournaud $\&$ Combes shown in Figure 5 (left panel), thus supporting the view that bars in spirals are fairly robust."420 In Buta et al. , In Buta et al. [421"86] the (, distribution for NIIGSOS was calculated.",86] the $Q_g$ distribution for NIRS0S was calculated.422 Most importantly. a clear dillerence was ound between SOs and early-tvpe spirals (see Fig.," Most importantly, a clear difference was found between S0s and early-type spirals (see Fig."423 5. right panel).," 5, right panel)."424 This was suggested to support the view according to which SOs have not acercted gas for a long time. evidenced by the lack of the extended tail. and the existence of a large number of weak bars.," This was suggested to support the view according to which S0s have not accreted gas for a long time, evidenced by the lack of the extended tail, and the existence of a large number of weak bars."425 As discussed above (see Section :3.2) the smaller number of strong bars among the SOs can be due to a dilution elect caused by the more massive bulges and thicker disks in SOs., As discussed above (see Section 3.2) the smaller number of strong bars among the S0s can be due to a dilution effect caused by the more massive bulges and thicker disks in S0s.426 However. it was also discussed by Buta et al. ," However, it was also discussed by Buta et al. ["427"SG] that this cannot produce all of the dillerence in Q,y between the SOs and earlv-type spirals: although spirals have a larger number of",86] that this cannot produce all of the difference in $Q_g$ between the S0s and early-type spirals: although spirals have a larger number of428A key question in the study of the formation and evolution of galaxies concerns the relationship between their observational properties and the large-scale cosmological environment.,A key question in the study of the formation and evolution of galaxies concerns the relationship between their observational properties and the large–scale cosmological environment.429 In recent years. a flourishing of observational campaign has provided a detailed description of the evolution of the galaxy population in clusters.," In recent years, a flourishing of observational campaign has provided a detailed description of the evolution of the galaxy population in clusters."430 Indeed. galaxy clusters play a key role in the characterization of the galaxy evolution.," Indeed, galaxy clusters play a key role in the characterization of the galaxy evolution."431 Each cluster provides a large sample of galaxies. all placed at the same redshift.," Each cluster provides a large sample of galaxies, all placed at the same redshift."432 Furthermore. clusters offer the possibility of sampling a variety of environments. from their dense core regions. to the outskirts where the properties of the cluster galaxy population tends to approach that of the field.," Furthermore, clusters offer the possibility of sampling a variety of environments, from their dense core regions, to the outskirts where the properties of the cluster galaxy population tends to approach that of the field."433 A diversity of the galaxy population in nearby clusters. with respect to that in the field. was noticed already by ? and by ?..," A diversity of the galaxy population in nearby clusters, with respect to that in the field, was noticed already by \cite{1974ApJ...194....1O} and by \cite{1980ApJ...236..351D}."434 Rich clusters were shown to contain a higher fraction of bulge—dominated (early type and SO) galaxies. and a correspondingly lower fraction of star forming galaxies. than poor systems.," Rich clusters were shown to contain a higher fraction of bulge--dominated (early type and S0) galaxies, and a correspondingly lower fraction of star forming galaxies, than poor systems."435 ? noticed that moderately distant clusters ἐς~ 0.3) have a galaxy," \cite{1978ApJ...226..559B}436 noticed that moderately distant clusters $z\sim 0.3$ ) have a galaxy"437c3x1015erg.,$\simeq 3 \times 10^{48} \erg$.438" With no mass estimate to calculate the kinetic energy, we can only speculate that the total energy is in the order of 103:50erg, larger than that of other ILOTs with the same time scale."," With no mass estimate to calculate the kinetic energy, we can only speculate that the total energy is in the order of $\sim 10^{49-50} \erg$, larger than that of other ILOTs with the same time scale."439" As SN 2002bu is far out of the OTS, we conclude that it is not an ILOT, and the most likely explanation is that it is a peculiar class of SN."," As SN 2002bu is far out of the OTS, we conclude that it is not an ILOT, and the most likely explanation is that it is a peculiar class of SN."440 This demonstrates that the ETD can also be used to identify transients with different physical properties than the ones considered as accretion powered ILOTs., This demonstrates that the ETD can also be used to identify transients with different physical properties than the ones considered as accretion powered ILOTs.441" The transient M85 OT2006 was discovered in the lenticular galaxy M85 in Jan 7, 2006 (Kulkarni et al."," The transient M85 OT2006 was discovered in the lenticular galaxy M85 in Jan 7, 2006 (Kulkarni et al."442 2007a)., 2007a).443 It had a peak luminosity of ~2x10*°ergs! and total radiated energy of Eyaa(M85 OT2006) ~6x1045erg over a duration of ~180days (Kulkarni et al., It had a peak luminosity of $\sim 2 \times 10^{40} \erg \s^{-1}$ and total radiated energy of $E_{\rm{rad}}$ (M85 OT2006) $\simeq 6 \times 10^{46} \erg$ over a duration of $\sim 180 \days$ (Kulkarni et al.444 2007a)., 2007a).445 The progenitor’s mass was estimated to be <7Me (Ofek et al., The progenitor's mass was estimated to be $< 7~\rm{M_{\odot}}$ (Ofek et al.446 2008) and it was suggested that the origin of M85 OT2006 is a stellar merger (Kulkarni et al., 2008) and it was suggested that the origin of M85 OT2006 is a stellar merger (Kulkarni et al.447 2007a)., 2007a).448 Rau et al. (, Rau et al. (4492007) estimated that the effective temperature and the stellar radius at the time of peak luminosity were Teec4600K and R~3600Ro respectively.,2007) estimated that the effective temperature and the stellar radius at the time of peak luminosity were $T_{\rm{eff}} \simeq 4600\K$ and $R \simeq 3600~\rm{R_{\odot}}$ respectively.450 At a later time in the eruption the star cooled down and expanded to have Tig~950K and R~20000Ro (see their table 2).," At a later time in the eruption the star cooled down and expanded to have $T_{\rm{eff}} \simeq 950\K$ and $R \simeq 20\,000~\rm{R_{\odot}}$ (see their table 2)."451 The work of Shara et al. (, The work of Shara et al. (4522010a) extends the previous nova models of Yaron et al. (,2010a) extends the previous nova models of Yaron et al. (4532005).,2005).454 The data of six of their models is given in their tables 2-4., The data of six of their models is given in their tables 2–4.455 In the new extreme nova models of Shara et al. (, In the new extreme nova models of Shara et al. (4562010a) the mass of the white dwarf is 0.4 — 0.65Μο and the accretion rate is as low as 10? — 107?Moyr!.,2010a) the mass of the white dwarf is $0.4$ – $0.65~\rm{M_{\odot}}$ and the accretion rate is as low as $10^{-12}$ – $10^{-10}~\rm{M_{\odot} \yr^{-1}}$.457 The ejected mass in the novae ranges between 5.3x10* — 2.2x1073Mo at velocities of 150 — 480kms!., The ejected mass in the novae ranges between $5.3 \times 10^{-4}$ – $ 2.2 \times 10^{-3}~\rm{M_{\odot}}$ at velocities of $150$ – $480 \km \s^{-1}$.458 Shara et al. (, Shara et al. (459"2010a) claim that the eruption of M85 OT2006 was a nova, and suggest that it can be explained in the frame of the new extreme nova models.","2010a) claim that the eruption of M85 OT2006 was a nova, and suggest that it can be explained in the frame of the new extreme nova models."460 Their conclusion is mainly based on the following. (, Their conclusion is mainly based on the following. (4611) Their new results showing that novae can reach comparable luminosities of ILOTs such as M85 OT2006 and M31 RV (few x10’ Lo). (,1) Their new results showing that novae can reach comparable luminosities of ILOTs such as M85 OT2006 and M31 RV (few $\times 10^7~\rm{L_{\odot}}$ ). (462"2) The models produce red eruptions, as the spectra of ILOTs.","2) The models produce red eruptions, as the spectra of ILOTs."463 We hereby show that the mass ejected in the eruption of M85 OT2006 is much larger than the nova models can produce., We hereby show that the mass ejected in the eruption of M85 OT2006 is much larger than the nova models can produce.464 Let us consider the M85 OT2006 transient according to data given by Rau et al. (, Let us consider the M85 OT2006 transient according to data given by Rau et al. (465"2007), discussed above.","2007), discussed above."466" The column density required above the photosphere is given by where p is the average density above the photosphere, Ar is the thickness of the shell above the photosphere, and K is its average opacity."," The column density required above the photosphere is given by where $\rho$ is the average density above the photosphere, $\Delta r$ is the thickness of the shell above the photosphere, and $\kappa$ is its average opacity."467 We parameterize the thickness of the shell with Ar=GR with 8~ 0.1., We parameterize the thickness of the shell with $\Delta r = \beta R$ with $\beta \sim 0.1$ .468 The final mass we obtain for the shell above the photosphere changes by a factor <2 for 0.01«60.1., The final mass we obtain for the shell above the photosphere changes by a factor $<2$ for $0.01 \leqslant \beta \leqslant 0.1$.469 The reason is because the value of 6 determines the density that weakly influences the opacity., The reason is because the value of $\beta$ determines the density that weakly influences the opacity.470 We use opacities from Ferguson et al. (, We use opacities from Ferguson et al. (471"2005) (slightly extrapolated), using compositions from Asplund et al. (","2005) (slightly extrapolated), using compositions from Asplund et al. ("4722004) with hydrogen abundance X=0.7 and metallicity Z=0.1; other composition from Lodders (2003) gives very close opacity values.,2004) with hydrogen abundance $\rm{X}=0.7$ and metallicity $\rm{Z}=0.1$; other composition from Lodders (2003) gives very close opacity values.473 The total mass above the photosphere is Using the temperature at peak luminosity we find that the opacity at the peak luminosity is &~1.3x 10-?., The total mass above the photosphere is Using the temperature at peak luminosity we find that the opacity at the peak luminosity is $\kappa \simeq 1.3 \times 10^{-3}$ .474 Substituting theopacity and the radius at peak luminosity in equation (7)) we find that the mass above the photosphere at the time of peak luminosity is Mpn0.2 Mo., Substituting theopacity and the radius at peak luminosity in equation \ref{eq:M_ph}) ) we find that the mass above the photosphere at the time of peak luminosity is $M_{\rm{ph}} \sim 0.2~\rm{M_{\odot}}$ .475 Our results are summarized in Table 2.., Our results are summarized in Table \ref{tab:M_ph}.476" The real amount of mass is even larger, as some mass is well above the photosphere and we expect a large amount of mass to be below the photosphere as well."," The real amount of mass is even larger, as some mass is well above the photosphere and we expect a large amount of mass to be below the photosphere as well."477 The largest value of ejected mass in the nova models of Shara et al. (, The largest value of ejected mass in the nova models of Shara et al. (47820102) is 2.2x10?Mo or two orders of magnitude below the expected ejected mass we calculate.,2010a) is $ 2.2 \times 10^{-3}~\rm{M_{\odot}}$ or two orders of magnitude below the expected ejected mass we calculate.479 We therefore conclude that M85 OT2006 is not a nova., We therefore conclude that M85 OT2006 is not a nova.480" Previously in KFS10, we used the assumption of Ofek et al. ("," Previously in KFS10, we used the assumption of Ofek et al. ("481"2008) that the total ejected mass of M85 OT2006 is 0.1 Mo, and consequently obtained a total energy of ~1.4x1047 erg.","2008) that the total ejected mass of M85 OT2006 is $0.1~\rm{M_{\odot}}$ , and consequently obtained a total energy of $\sim 1.4 \times 10^{47} \erg$ ."482 Our new estimate of ejected mass is much higher and we update our estimate of the total energy in M85 OT2006., Our new estimate of ejected mass is much higher and we update our estimate of the total energy in M85 OT2006.483" Our new estimate, taking ejected mass velocity of ~870kms! (Rau et al."," Our new estimate, taking ejected mass velocity of $\sim 870 \km \s^{-1}$ (Rau et al."484" 2007) is 1.6x1048 — 4.6x107?erg, corresponding to 0.2Mo and 0.6Mo, respectively."," 2007) is $1.6 \times 10^{48}$ -- $4.6 \times 10^{48}\erg$, corresponding to $0.2~\rm{M_{\odot}}$ and $0.6~\rm{M_{\odot}}$, respectively."485 These values are summarized in Table 1 and the updated location of M85 OT2006 is shown in the ETD (Figure 1))., These values are summarized in Table \ref{tab:data} and the updated location of M85 OT2006 is shown in the ETD (Figure \ref{fig:totEvst}) ).486 We calculate the total (radiated and kinetic) energy of the new nova modelsof Shara el al. (, We calculate the total (radiated and kinetic) energy of the new nova modelsof Shara el al. (4872010a) andplot them in Figure 1..,2010a) andplot them in Figure\ref{fig:totEvst}. .488 We include all models from Yaron et al. (, We include all models from Yaron et al. (4892005),2005)490revisit the same image twice for a total of 3 snapshots per field of view.,revisit the same image twice for a total of 3 snapshots per field of view.491" Observing in 4 bands, for a total of 28 days integration time over 2 years, such a survey expects to see ~10 red supergiant PISNe and ~300 CCSNe."," Observing in 4 bands, for a total of 28 days integration time over 2 years, such a survey expects to see $\sim$ 10 red supergiant PISNe and $\sim$ 300 CCSNe."492" To the extent that PISN spectra can be represented as a distribution of blackbodies at different temperatures, since the temperature and redshift would be degenerate, it will be impossible to acquire photometric redshifts without further information about the SN epoch."," To the extent that PISN spectra can be represented as a distribution of blackbodies at different temperatures, since the temperature and redshift would be degenerate, it will be impossible to acquire photometric redshifts without further information about the SN epoch."493" However, our simulated spectra show significant deviations from a blackbody in the UV (I< 3500A)) due to metal-line blanketing in the SN photosphere, providing spectral and photometric signatures that could be used as redshift indicators, depending on their strength."," However, our simulated spectra show significant deviations from a blackbody in the UV $l < 3500$ ) due to metal-line blanketing in the SN photosphere, providing spectral and photometric signatures that could be used as redshift indicators, depending on their strength."494" Although the UV flux of PISNe is relatively short lived, the more massive PISNe stay bright in its rest frame visible band for over a year."," Although the UV flux of PISNe is relatively short lived, the more massive PISNe stay bright in its rest frame visible band for over a year."495" Given this brightness and long intrinsic duration, coupled with the (1+z) time dilation at high redshifts, it is conceivable that PISNe could contribute to the luminosity function of all objects at high redshifts whengalaxies were dim."," Given this brightness and long intrinsic duration, coupled with the $(1+z)$ time dilation at high redshifts, it is conceivable that PISNe could contribute to the luminosity function of all objects at high redshifts whengalaxies were dim."496" Figure 9 illustrates the luminosity function of PISNe at ~4000A,, calculated using the helium core progenitor models for PISN luminosity, and the Pop III Flat or Pop III Salpeter models for the star formation rate."," Figure \ref{luminosity_function} illustrates the luminosity function of PISNe at $\sim$, calculated using the helium core progenitor models for PISN luminosity, and the Pop III Flat or Pop III Salpeter models for the star formation rate."497" Shown for comparison are the projected galaxy luminosity functions at high redshifts, using the ? best fit Schechter parameterization for the UV luminosity function, and shifting to the visible band using U—Vz0.4, 0.3 for z—7, 8 respectively, measured using the Spitzer Infrared Array Camera (??).."," Shown for comparison are the projected galaxy luminosity functions at high redshifts, using the \citet{Bouwens2011b} best fit Schechter parameterization for the UV luminosity function, and shifting to the visible band using $U-V \approx 0.4$, $0.3$ for $z=7,$ $8$ respectively, measured using the Spitzer Infrared Array Camera \citep{Labb'e2010, Labb'e2010a}."498" Applying this U-V shift is a crude approximation, as luminous and faint galaxies have different rest frame UV-to-optical color; however, we are most interested in the bright end of the luminosity function, where this current U-V measurement is applicable."," Applying this U-V shift is a crude approximation, as luminous and faint galaxies have different rest frame UV-to-optical color; however, we are most interested in the bright end of the luminosity function, where this current U-V measurement is applicable."499 The luminosity function for PISNe implied by our Pop III IMF models overlaps with the galaxy luminosity function at the brightest magnitudes., The luminosity function for PISNe implied by our Pop III IMF models overlaps with the galaxy luminosity function at the brightest magnitudes.500" If a top-heavy Pop III IMF was solely responsible for reionization, PISNe will contaminate the brightest end of the galaxy luminosity function, unless great care is taken to remove these supernovae."," If a top-heavy Pop III IMF was solely responsible for reionization, PISNe will contaminate the brightest end of the galaxy luminosity function, unless great care is taken to remove these supernovae."501" Since the volumetric count of the brightest galaxies and PISNe is very low, it will take a wide infrared survey to observe this effect."," Since the volumetric count of the brightest galaxies and PISNe is very low, it will take a wide infrared survey to observe this effect."502" In our discussion we ignored complicating factors such as metallicity and rotation, and calculated the PISN and CCSN event rate using only the SN progenitor mass range along with the star formation rate."," In our discussion we ignored complicating factors such as metallicity and rotation, and calculated the PISN and CCSN event rate using only the SN progenitor mass range along with the star formation rate."503" However, at low redshifts z«1, the measured CCSN rate is a factor of ~2 smaller than that predicted by the analogous calculation using the measured cosmic star formation rate."," However, at low redshifts $z<1$, the measured CCSN rate is a factor of $\sim 2$ smaller than that predicted by the analogous calculation using the measured cosmic star formation rate."504 The discrepancy is likely due to many intrinsically low-luminosity or obscured SNe being missed in surveys (?).., The discrepancy is likely due to many intrinsically low-luminosity or obscured SNe being missed in surveys \citep{Horiuchi2011}.505" As this discrepancy is lower than the uncertainty in our SFR. model parameters, and we already account for lower intrinsic luminosities for the lower progenitor mass PISNe, we do not take obscuration into account for our predictions of the SN rate as seen by JWST."," As this discrepancy is lower than the uncertainty in our SFR model parameters, and we already account for lower intrinsic luminosities for the lower progenitor mass PISNe, we do not take obscuration into account for our predictions of the SN rate as seen by JWST."506 The IMF of early stellar populations responsible for reionization should also leave an imprint on the metal enrichment pattern via their SN products., The IMF of early stellar populations responsible for reionization should also leave an imprint on the metal enrichment pattern via their SN products.507" So far, the abundance patterns observed to date in extremely metal-deficient stars in the Galactic halo (?) are more consistent with an IMF that produced much more CCSNe instead of PISNe (?).."," So far, the abundance patterns observed to date in extremely metal-deficient stars in the Galactic halo \citep{Beers2005} are more consistent with an IMF that produced much more CCSNe instead of PISNe \citep{Joggerst2010}."508" However, in previous surveys, subtle selection effects might have disfavored finding PISN-enriched stars; the metal yields of PISNe are so high that the metal abundances of stars formed out of PISN ejecta (7) are already higher than the metallicity range targeted by metal-deficient star surveys (?).."," However, in previous surveys, subtle selection effects might have disfavored finding PISN-enriched stars; the metal yields of PISNe are so high that the metal abundances of stars formed out of PISN ejecta \citep{Greif2008} are already higher than the metallicity range targeted by metal-deficient star surveys \citep{Karlsson2008}."509" Large carbon enhancements observed in metal-poor stars, when interpreted as the outcome of pollution by winds from binary companions that have gone through the AGB phase, suggest the existence of a large number of intermediate-mass stars (~1— 8Mc) at high redshifts (??).."," Large carbon enhancements observed in metal-poor stars, when interpreted as the outcome of pollution by winds from binary companions that have gone through the AGB phase, suggest the existence of a large number of intermediate-mass stars $\sim 1-8 M_{\odot}$ ) at high redshifts \citep{Tumlinson2007,Tumlinson2007a}."510" Alternatively, nucleosynthesis in faint CCSNe from higher mass stars could also explain the observed carbon enhancement in metal-poor stars (?).."," Alternatively, nucleosynthesis in faint CCSNe from higher mass stars could also explain the observed carbon enhancement in metal-poor stars \citep{Iwamoto2005}."511" Observing the Type Ia SN rate during the epoch of reionization will be an complementary way to test these models, and constrain the number of intermediate-mass stars at high redshifts."," Observing the Type Ia SN rate during the epoch of reionization will be an complementary way to test these models, and constrain the number of intermediate-mass stars at high redshifts."512 The predicted initial mass range of ~140 to 260 Mo for PISN progenitors assumed the stars to be non-rotating (?).., The predicted initial mass range of $\sim$ 140 to 260 $M_{\odot}$ for PISN progenitors assumed the stars to be non-rotating \citep{Heger2002}. .513" However, observations find that at very low metallicities, stars rotate faster (?).."," However, observations find that at very low metallicities, stars rotate faster \citep{Martayan2007}."514" The fast rotation of the first stars is supported by the latest hydrodynamic simulations of their formation (?),, and also by observations of anomalously high abundances of Ba and La with respect to Fe in ancient stars (?),, which could originate in metal-poor fast- massive stars."," The fast rotation of the first stars is supported by the latest hydrodynamic simulations of their formation \citep{Stacy2011}, and also by observations of anomalously high abundances of Ba and La with respect to Fe in ancient low-mass stars \citep{Chiappini2011}, , which could originate in metal-poor fast-rotating massive stars."515" Generally, rotationshould increase the required PISN progenitor mass by increasing mass loss."," Generally, rotationshould increase the required PISN progenitor mass by increasing mass loss."516The general nucleosvnthesis associated with rapid acliabatic expansion and freeze-oul from Nuclear Statistical Equilibrinm (NSE) at high entropy-per-baryon. s. was first considered in the landmark paper bv Wagoner.Fowler.&Iovle(1967) (hereafter WEIL).,"The general nucleosynthesis associated with rapid adiabatic expansion and freeze-out from Nuclear Statistical Equilibrium (NSE) at high entropy-per-baryon, $s$, was first considered in the landmark paper by \cite{wag} (hereafter WFH)."517" Those authors concentrated on the environments associated wilh exploding supermassive objects. where s/h,~10. and on Dig Dang Nucleosvnthesis (BBN). where s/h,zzLOM."," Those authors concentrated on the environments associated with exploding supermassive objects, where $s/k_b\sim 10^3$, and on Big Bang Nucleosynthesis (BBN), where $s/k_b\approx10^{10}$."518 In both supermassive objects aud BBN the conditions are expected to be proton-rich (preponderance of protons over neutrons) and the characteristic expansion timescale large (e.g.. 7444~1008 for BBN).," In both supermassive objects and BBN the conditions are expected to be proton-rich (preponderance of protons over neutrons) and the characteristic expansion timescale large (e.g., $\tau_{\rm dyn}\sim 100{\rm s}$ for BBN)."519 The physics of relativistic outflows aud potential nucleosvinthesis in (hese sites has also been the subject of some recent. attention (Pruet.FullerBurrows.&Mever2001:Otsukiοἱal.," The physics of relativistic outflows and potential nucleosynthesis in these sites has also been the subject of some recent attention \citep{ultra,thomps,otsuki}."520 2000).. llere we extend the WFII study., Here we extend the WFH study.521 We consider [reeze-out. from NSE over a wide range of entropv-per-baryon spanning that in WEIL ancl a range of neutron-to-proton ratios. all the way [rom proton-rich. (o neutron-rich.," We consider freeze-out from NSE over a wide range of entropy-per-baryon spanning that in WFH, and a range of neutron-to-proton ratios, all the way from proton-rich to neutron-rich."522" NSE freezeout in these scenarios is calculated for expansion limescales ranging [rom those appropriate for relativistic flows from compact objects to those associated with BBN (10""s<zq4,1005).", NSE freezeout in these scenarios is calculated for expansion timescales ranging from those appropriate for relativistic flows from compact objects to those associated with BBN $10^{-6}{\rm s}<\tau_{\rm dyn}<100{\rm s})$.523 We find that the general nuclear physics features of BBN are a recurring theme throughout all of these parameter ranges., We find that the general nuclear physics features of BBN are a recurring theme throughout all of these parameter ranges.524 In particular. 71 svnthesis can be significant for relativistic flows.," In particular, ${^2}{\rm H}$ synthesis can be significant for relativistic flows."525 The expected possibly neutron-rich conditions of a GRB fireball are. in a sense. the isospin mirror of those for BBN.," The expected possibly neutron-rich conditions of a GRB fireball are, in a sense, the isospin mirror of those for BBN."526 However. we also show that it is not sufficient to simply consider only (he NSE freeze-out nucleosvnthesis in some parameter regimes appropriate for. e.g.. GRB fireballs.," However, we also show that it is not sufficient to simply consider only the NSE freeze-out nucleosynthesis in some parameter regimes appropriate for, e.g., GRB fireballs."527 In fact. the dvnaimics of these fireballs ean ciffer dramatically on a microscopic scale from (he conditions treated by WEIL.," In fact, the dynamics of these fireballs can differ dramatically on a microscopic scale from the conditions treated by WFH."528 In particular. for initially neutron-rich material (he few protons accelerate with the relativistic photon and. e= pair fluid to a very high Lorentz [actor (5~10?— 107).," In particular, for initially neutron-rich material, the few protons accelerate with the relativistic photon and $e^{\pm}$ pair fluid to a very high Lorentz factor $\gamma\sim10^2-10^3$ )."529 In this scenario. there will be high energy (~ GeV) collisions of protons on left-behind neutron stragglers (Fuller.Pruet.&Abazajian2000).," In this scenario, there will be high energy $\sim {\rm GeV}$ ) collisions of protons on “left-behind” neutron stragglers \citep{fpa}."530. It has been shown that high energy collisions of (his kind can result in significant destruction and/or production of light nuclei 1983).., It has been shown that high energy collisions of this kind can result in significant destruction and/or production of light nuclei \citep{dim}.531 We show below that these non-thermal nuclear reactions can modilv significantly the simple BBN-like NSE freeze-ont abundances of 7H and other species., We show below that these non-thermal nuclear reactions can modify significantly the simple BBN-like NSE freeze-out abundances of ${^2}{\rm H}$ and other species.532 In the right conditions. IL number fraction vields can approach YyL~1056.," In the right conditions, ${^2}{\rm H}$ number fraction yields can approach $Y_{\rm D}\sim 10\%$."533 This is a stageerinelySSSt. high5S vield given the known. and small. primordial ?II abundance.," This is a staggeringly high yield given the known, and small, primordial ${^2}\rm H$ abundance."534 The deuteron is a particularly interesting nucleus because it has a binding energv οἱ only 222.2 MeV. and is notoriously Iragile and difficult to synthesize., The deuteron is a particularly interesting nucleus because it has a binding energy of only $\approx$ 2.2 MeV and is notoriously fragile and difficult to synthesize.535 At present it is thought that essentially all of the deuterium in the universe is primordial., At present it is thought that essentially all of the deuterium in the universe is primordial.536 Estimates of the primordial “11/11 come from measurements of the 82 km/s isotope shift in the Lyman a line in a handful ol high redshift’ Lyman limit clouds., Estimates of the primordial $\hh$ /H come from measurements of the 82 km/s isotope shift in the Lyman $\alpha$ line in a handful of high redshift Lyman limit clouds.537 These are clouds that interseet our lines of sight to, These are clouds that intersect our lines of sight to538aud the stellaz component is caleulated using an N-body echuique.,and the stellar component is calculated using an N-body technique.539 The initial conditiois wore built according to Spriugel et al. (, The initial conditions were built according to Springel et al. (5402005). based ou the analytic work of Mo et al. (,"2005), based on the analytic work of Mo et al. ("5411998).,1998).542 The two model galaxies were chosen sucht hat trev represent a Milkv Way type and a simall spiral ooOalaxy. with the mass ratio of the two ealaxics Όσιο 8:1.," The two model galaxies were chosen such that they represent a Milky Way type and a small spiral galaxy, with the mass ratio of the two galaxies being 8:1."543 There‘Ore. the total mass of the model galaxies A and [2L) Is « 1072U ph1 M. aud L671Iss 4107 53 M;. respectivev.," Therefore, the total mass of the model galaxies A and B is $\times$ $^{12}$ $h^{-1}$ $M_{\sun}$ and $\times$ $^{11}$ $h^{-1}$ $M_{\sun}$, respectively."544 The combined N-bodv/SPII simulation then calculates 5 Cir of evoluion., The combined N-body/SPH simulation then calculates 5 Gyr of evolution.545 For every time step. we know he veocity of cach particle aud can hence extract realistic 2D veocitv fields of the eus.," For every time step, we know the velocity of each particle and can hence extract realistic 2D velocity fields of the gas."546ons Concerning the osoval. aliennneut aud impact )araucters. we follow he notation iutroduced by Duc et al. (," Concerning the spatial alignment and impact parameters, we follow the notation introduced by Duc et al. ("5472000) o describe the interaction geometry.,2000) to describe the interaction geometry.548 The )araiueer corresponds to the minim separation of he galaxies rajectories. as if they were ponit masses Ou Weplerian orbits.," The parameter corresponds to the minimum separation of the galaxies' trajectories, as if they were point masses on Keplerian orbits."549 Additionally. two angles. O aud ᾧ dchine the spalal orientation of the disc.," Additionally, two angles, $\Theta$ and $\Phi$ define the spatial orientation of the disc."550 For he complete sample. we selected. the aliguimeuts iu such a way as to cover as many ecoletrics as possible. including nuxx and major mergers aud fiv-bys (achieved lyVv increasing the minium separation).," For the complete sample, we selected the alignments in such a way as to cover as many geometries as possible, including minor and major mergers and fly-bys (achieved by increasing the minimum separation)."551 For this analvsis NXο just use three differeut mteraction eeonmetzies. which NXο list iu Table 1..," For this analysis we just use three different interaction geometries, which we list in Table \ref{sims}."552 These three simulations allow us to study clifferent classes of ποιαισα]. cistortiois. detailed in Sect. 3..," These three simulations allow us to study different classes of kinematical distortions, detailed in Sect. \ref{extract}."553" Additional important quantities of the simulations. as e.g. the particle αμαος, are smnunarized iu Table 2.."," Additional important quantities of the simulations, as e.g. the particle numbers, are summarized in Table \ref{galaxyproperties_resolution}."554 Throughout the paper we adopt the standard ACDAM cosinoloey wi hOQy=.7. 0enti=0.3. aud h=0.7.," Throughout the paper we adopt the standard $\Lambda$ CDM cosmology with $\Omega_\Lambda=0.7$, $\Omega_m=0.3$, and h=0.7."555 Il OCer to construct realistic 2D velocity fields. we projec all gas particles of the N-bodyv/SPII simulations [9]uto a Cartesian. equally spaced exid.," In order to construct realistic 2D velocity fields, we project all gas particles of the N-body/SPH simulations onto a Cartesian, equally spaced grid."556 The spacing is Crosea such. that the spatial resolution at the assumed redshit of the model galaxy correspouds to the angular resolujon of current state-ofthe-art observations.," The spacing is chosen such, that the spatial resolution at the assumed redshift of the model galaxy corresponds to the angular resolution of current state-of-the-art observations."557 For t lewrole investigated redshift range from z- to z=1.0 iMο adopt an aneular resolution typical for IFU or EPI [9]servations. namely 0.0 (ee. Chemin et al.," For the whole investigated redshift range from z=0 to z=1.0 we adopt an angular resolution typical for IFU or FPI observations, namely 0.4"" (e.g. Chemin et al."558 2006)., 2006).559 The angular resolution o|: SAURON. for example. would be IY or 0.97.," The angular resolution of SAURON, for example, would be 0.3"" or 0.9""."560 For iuterinediate redshüfts we additiona hose 0.257. which is the pixel scale along the FORS2 s sed bv Ziegler et al. (," For intermediate redshifts we additionally chose 0.25"", which is the pixel scale along the FORS2 slit as used by Ziegler et al. ("5612006) and 0.527. which is f et size of FLAMES/GIRAFFE at VET (c.g. Flores al.,"2006) and 0.52"", which is the lenslet size of FLAMES/GIRAFFE at VLT (e.g. Flores et al."562 2006)., 2006).563 We calculate for each redshift he physic 10360)utiion according o the eiven aieular resolution usi 16 concordance cosmological iiocel (see Table 3)., We calculate for each redshift the physical resolution according to the given angular resolution using the concordance cosmological model (see Table \ref{redshift}) ).564 T vclocity field of the galaxy is binned using this spatial 10360)mlon., The velocity field of the galaxy is binned using this spatial resolution.565 Frei the kuowledeoe of the full 3D velociv fields ancl the iucraction historv of the galaxies in our siuulations iMο define three kinematical classes. in principle analogue tc) classificatious in Flores et al. (," From the knowledge of the full 3D velocity fields and the interaction history of the galaxies in our simulations we define three kinematical classes, in principle analogue to classifications in Flores et al. ("5662006) or Krajnovic et al. (,2006) or Krajnović et al. (5672006): The main question tlat we investigate in this paper is how such classifications depend ou the redshift of re observed salaxy. ie. the actual spatial resolution of ie galaxy.,"2006): The main question that we investigate in this paper is how such classifications depend on the redshift of the observed galaxy, i.e. the actual spatial resolution of the galaxy."568 Iu this contest the κος plays a crucial 1¢de. as it typically exceeds the augular resolution of the oeistriuueit.," In this context the seeing plays a crucial role, as it typically exceeds the angular resolution of the instrument."569 To simulate secius effects on our velocity-field 1ieasureunieits. a c0wolution with a Gaussian poiut nixead function was appied.," To simulate seeing effects on our velocity-field measurements, a convolution with a Gaussian point spread function was applied."570 We adopted a value of 087 or the FWA of the Gaussian seeiug. which is a typical Vidue for eround. based observatious (see c.g. Jagecr et al.," We adopted a value of 0.8"" for the FWHM of the Gaussian seeing, which is a typical value for ground based observations (see e.g. Jägger et al."571 2)D., 2004).572 Tje appearance of the velocity fields is therefore clenunmated hy the secing., The appearance of the velocity fields is therefore dominated by the seeing.573 Note that we do not calculate a1 evolution of properics of the galaxy with redshift but Edy how a eiven kinematical state of a galaxy is observed a different recshifts., Note that we do not calculate an evolution of properties of the galaxy with redshift but study how a given kinematical state of a galaxy is observed at different redshifts.574 All studies imnoenutioned above are inainly ποσο cleuinated., All studies mentioned above are mainly seeing dominated.575 Wo addijonally study the possibility to XCentity mergers at high redshift (22) with an adaptive Qfies Iustrument such as SINFONT at the VET., We additionally study the possibility to identify mergers at high redshift $\sim$ 2) with an adaptive optics instrument such as SINFONI at the VLT.576" For this investigation we adop an angular resolution of 0.15"" as aclüeved by Cieuzel et al. ("," For this investigation we adopt an angular resolution of 0.15"" as achieved by Genzel et al. ("5772006).,2006).578 Iu the subsequent sub-sectious we will svstemiaticallv investigate velocity fields from cach of the kiueimatical classes defined im Sect. 3.., In the subsequent sub-sections we will systematically investigate velocity fields from each of the kinematical classes defined in Sect. \ref{extract}.579 The focus lies ou the visibility of kinematical distortions as a function of redshift., The focus lies on the visibility of kinematical distortions as a function of redshift.580 Iu order to «nantifv the distortions and to interpret the partly complex structures iu the velocity fields we use the kincimnetry package o ‘Wrajnovie et al. (, In order to quantify the distortions and to interpret the partly complex structures in the velocity fields we use the kinemetry package of Krajnović et al. (5812006).,2006).582 This analysis Is based on the assumption. that the mean velocity along best fitting cllipses can be reproduced by a cosine law. 1.6. Ww.here « is the lenet1 of senianajor axis of the cllipse. V Is for discs the azimuthal anele measured from the major axis iu the plane o| the Ooealaxv.," This analysis is based on the assumption, that the mean velocity along best fitting ellipses can be reproduced by a cosine law, i.e. where $a$ is the length of semi-major axis of the ellipse, $\Psi$ is for discs the azimuthal angle measured from the major axis in the plane of the galaxy."583 Note that in Oogeneral a siuple siue correction for the inclination of the dise is a»plied iu observations too., Note that in general a simple sine correction for the inclination of the disc is applied in observations too.584 The velocity fields presented, The velocity fields presented585The external galactic fiekl due to the bulge is represented by the potential of the Schwarzschild model (Sehwarzschild1979).,The external galactic field due to the bulge is represented by the potential of the Schwarzschild model \citep{scw}.586. The Schwarzschild model is a non-rotating. consistent triaxial ellipsoid with axis ratios 2: 1.25: 1. typical of many galaxies ," The Schwarzschild model is a non-rotating, self-consistent triaxial ellipsoid with axis ratios 2: 1.25: 1, typical of many galaxies \citep{ber91}."587"Defining adimensional units as rj being the bulge core radius. the potential (x.i.z) is expressed as the sum of a spherically simmetric term. ὃν (07= r/rj). which corresponds to the potential given by a density distribution following the modified IIubble's law pr)=p[1+(η. (po=My rj). plus two spherical armonics. 94(z.7) and Pole.ii): where and Ad, is the bulge"," Defining adimensional units as $r_b$ being the bulge core radius, the potential $\Phi(x^{\prime},y^{\prime},z^{\prime})$ is expressed as the sum of a spherically simmetric term, $\Phi_{r^{\prime}}(r^{\prime})$, $r^{\prime}=r/r_b$ ), which corresponds to the potential given by a density distribution following the modified Hubble's law $\rho(r^{\prime})=\rho_0 \left[1+\left(r^{\prime}\right)^2\right]^{-3/2}$, $\rho_0\equiv M_b/r_b^3 $ ), plus two spherical armonics, $\Phi_1(z^{\prime},r^{\prime})$ and $\Phi_2(x^{\prime},y^{\prime},r^{\prime})$ where and $M_b$ is the bulge"588There remain many pieces needed to assemble the circiumnuclear starburst ring puzzle.,There remain many pieces needed to assemble the circumnuclear starburst ring puzzle.589 A large-scale bar drives dust ancl gas towards the center of the galaxy., A large-scale bar drives dust and gas towards the center of the galaxy.590 At the boundary of the bulge and the disk. a starlormine ring develops.," At the boundary of the bulge and the disk, a starforming ring develops."591 Within (his ring. there is a fairly coherent loosely wound dust spiral.," Within this ring, there is a fairly coherent loosely wound dust spiral."592" [tis unclear. however. whether this nuclear dust spiral is composed of dust aud gas originating Irom the large-scale bar and/or the ring itself. or if it is native dust that is ""stiimred-up by the starburst ring."," It is unclear, however, whether this nuclear dust spiral is composed of dust and gas originating from the large-scale bar and/or the ring itself, or if it is native dust that is “stirred-up” by the starburst ring."593 Simulations by Regan&Teuben and Maciejewski(2004) imply that bars in low velocitv-dispersion svstemis are only capable ol driving gas down to the radius at which the ring formsbut not furthersuggesting that native dust comprises the nuclear spiral., Simulations by \citet{regan04} and \citet{maciejewski04b} imply that bars in low velocity-dispersion systems are only capable of driving gas down to the radius at which the ring forms—but not further—suggesting that native dust comprises the nuclear spiral.594 A galaxy. with a high central velocitv-dispersion. however. could potentially have mass transported into the ring (Maciejewski2004)..," A galaxy with a high central velocity-dispersion, however, could potentially have mass transported into the ring \citep{maciejewski04b}. ."595Following similar. transformations with the remaining Eqs. (,Following similar transformations with the remaining Eqs. (5966)-(9). we arrived at the linearized equations needed for stability analysis.,"6)-(9), we arrived at the linearized equations needed for stability analysis."597" We consider perturbations with the space- dependence «export—fk-r). where k=(k,.O.Kk.) Is the wavevector."," We consider perturbations with the space-time dependence $\propto \exp ( \sigma t - i \vec{k} \cdot \vec{r})$, where $\vec{k}= (k_{s}, 0, k_{z})$ is the wavevector."598" Then. the linearized MHD-equations read with accuracy in the lowest order in /l/s: where w,=gk."," Then, the linearized MHD-equations read with accuracy in the lowest order in $\lambda/s$: where $\omega_{\eta}= \eta k^2$."599 The dispersion equation. corresponding to Eqs. (," The dispersion equation, corresponding to Eqs. ("60018)-(22) is rather complex in the general case.,18)-(22) is rather complex in the general case.601 Therefore. we consider only a particular case when the wavevector is perpendicular to the magnetic field. k-B=0.," Therefore, we consider only a particular case when the wavevector is perpendicular to the magnetic field, $\vec{k} \cdot \vec{B} = 6020$."603 After some algebra. Eqs. (," After some algebra, Eqs. ("60418)-(22) can be combined into a sixth-order dispersion relation. The coefficients of this equation are expressed in terms of characteristic frequencies. where jt=- . The characteristic frequences are given by modes d,"18)-(22) can be combined into a sixth-order dispersion relation, The coefficients of this equation are expressed in terms of characteristic frequencies, where $\mu = k^{2}_{z}/k^{2}$ The characteristic frequences are given by where $c_{s} = \sqrt{\gamma p/ \rho}$ is the sound speed."605, Eq. (606escribes that can generally exist in a compressible: rotating. magnetized gas.,"23) describes six modes that can generally exist in a compressible, rotating, magnetized gas."607" In the non-dissipative. incompressible limit when the sound speed c, is very large. Eq. ("," In the non-dissipative, incompressible, limit when the sound speed $c_{s}$ is very large, Eq. ("6082) reduces to the dispersion relation for the inertial waves. This equation allows for unstable solutions if the Rayleigh criterion is fulfilled. «7<0.,"23) reduces to the dispersion relation for the inertial waves, This equation allows for unstable solutions if the Rayleigh criterion is fulfilled, $\kappa^{2} < 0$."609 Note that Eq. (, Note that Eq. (61023) does not describethe magnetorotational instability since it does not occur for perturbations with K-B=0.,23) does not describethe magnetorotational instability since it does not occur for perturbations with $\vec{k} \cdot \vec{B} = 0$.611 In non-dissipative limit. 7-0. Eq. (," In non-dissipative limit, $\eta \rightarrow 0$, Eq. ("61223) recovers the dispersion relation by Bonanno Urpin (2006) for MD Wwith Mk-B=0. The authors argued that Eq. (,"23) recovers the dispersion relation derived by Bonanno Urpin (2006) for perturbations with $\vec{k} \cdot \vec{B} = 0$, The authors argued that Eq. ("61325) has unstable solutions even when the criteria of the magnetorotational and Rayleigh instability are not satisfied.,25) has unstable solutions even when the criteria of the magnetorotational and Rayleigh instability are not satisfied.614" The condition of instability ts wpo*O and. hence. the shear-driven instability arises if B,#0."," The condition of instability is $\omega_{B \Omega} \neq 0$ and, hence, the shear-driven instability arises if $B_{s} \neq 0$."615 The paper by Bonanno Urpin (2007) considers the dispersion relation for any wavevectors. but does so in a conductivity limit 77—0.," The paper by Bonanno Urpin (2007) considers the dispersion relation for any wavevectors, but does so in a high-conductivity limit $\eta \rightarrow 0$ ."616 Eq. (, Eq. (617"23) can be simplified in many cases of interest if we take into account that. likely. the s-component of the magnetic field in disks 1s greater than the z-component. B,>>B-.","23) can be simplified in many cases of interest if we take into account that, likely, the $s$ -component of the magnetic field in disks is greater than the $z$ -component, $B_s \gg B_z$."618" For perturbations with k-B=0. we have k,=-k-B-/B,."," For perturbations with $\vec{k} \cdot 619\vec{B} = 0$, we have $k_s = - k_z B_z/ B_s$."620 Then. Substituting HR| into Eq. a(," Then, Substituting $\mu \approx 1$ into Eq. ("6212ONG can transform it into Two roots of this equation deseribe the inertial waves which can be unstable only if the Rayleigh criterion is satisfied.,"23), we can transform it into Two roots of this equation describe the inertial waves which can be unstable only if the Rayleigh criterion is satisfied."622 Other four modes. are deseribed by the dispersion relation This equation deseribes fast and slow magnetoacoustic waves. and we consider the stability of these modes.," Other four modes are described by the dispersion relation This equation describes fast and slow magnetoacoustic waves, and we consider the stability of these modes."623 Note that simplification (26) is made for mathematical convenience rather than for physical relevance., Note that simplification (26) is made for mathematical convenience rather than for physical relevance.624 In. fact. the dispersion relation (28) for fast and slow magnetosonie waves will change litle if wo#1.," In fact, the dispersion relation (28) for fast and slow magnetosonic waves will change little if $\mu \neq 1$."625 This particularly concerns the modes with jo]>Q because the dispersion equation for them will differ from Eq. (, This particularly concerns the modes with $|\sigma| > \Omega$ because the dispersion equation for them will differ from Eq. (62628) only by terms of the order of (ji—Lafor in coefficients.,28) only by terms of the order of $(\mu -1) \kappa^2/ \sigma^2$ in coefficients.627 Therefore. our results can be applied with a sufficient accuracy also for the perturbations with 4i#I.," Therefore, our results can be applied with a sufficient accuracy also for the perturbations with $\mu \neq 1$."628 The Hurwitz theorem states that an equation of the fourth order. has at least one root with a positive real part (unstable mode) if one of the following inequalities is fulfilled (see Aleksandrov et al.," The Hurwitz theorem states that an equation of the fourth order, has at least one root with a positive real part (unstable mode) if one of the following inequalities is fulfilled (see Aleksandrov et al."629 1985)., 1985).630 Conditions (30) are never satisfied 1n. protostellar disks., Conditions (30) are never satisfied in protostellar disks.631 Therefore. the instability arises rf one of the inequalities (31) or (32) are satisfied.," Therefore, the instability arises if one of the inequalities (31) or (32) are satisfied."632 These inequalities can be rewritten for Eq. (, These inequalities can be rewritten for Eq. (633"28) as Generally. the HOOexpressionsOj on the2, Lh.s.","28) as Generally, the expressions on the l.h.s."634 of both these inequalities can have a positive or negative sign. depending on the value of «go.," of both these inequalities can have a positive or negative sign, depending on the value of $\omega_{B \Omega}$."635 Therefore. Eq. (," Therefore, Eq. ("63633) and (34) impose restrictions on the rate of differential rotation which should be greater than some critical value.,33) and (34) impose restrictions on the rate of differential rotation which should be greater than some critical value.637 Note that. if the azimuthal field is generated winding up the radial field. then διο>0 and. hence. o;byBC;> Oas well.," Note that, if the azimuthal field is generated by winding up the radial field, then $B_s B_{\varphi} \Omega' >0$ and, hence, $\omega_{B \Omega}^3 > 0$ as well."638 Therefore. we consider only this case.," Therefore, we consider only this case."639Denoting 6=Casca/ 03. We have from Eq. (,"Denoting $\delta = c_{As} 640c_{A \varphi}/c_{A}^2$ , we have from Eq. ("64133 Thiscondition can ie satisfied only if the differential rotation is sufficiently strong.,33) Thiscondition can be satisfied only if the differential rotation is sufficiently strong.642 Since Q’<0 in astrophysical disks. we," Since $\Omega' < 0$ in astrophysical disks, we"643"For finite quark mass mg. F(r.T) remains finite lor r—x. since the ""string. between the two color charges ‘breaks’ when the corresponding potential enerey becomes equal to the mass AM, of the lowest hadron: bevond this point. it becomes energetically more favourable to produce an additional hadron.","For finite quark mass $m_q$ , $F(r,T)$ remains finite for $r \to \infty$, since the `string' between the two color charges `breaks' when the corresponding potential energy becomes equal to the mass $M_h$ of the lowest hadron; beyond this point, it becomes energetically more favourable to produce an additional hadron."644" Hence now £ no longer vanishes in the confined phase. but only becomes exponentially small there. L(T)——Al,/Th:: here M, is a tvpical hadron mass. of the order of 0.5 to 1.0 GeV. so that al 7.c170 MeV. L~107. rather than zero."," Hence now $L$ no longer vanishes in the confined phase, but only becomes exponentially small there, L(T); here $M_h$ is a typical hadron mass, of the order of 0.5 to 1.0 GeV, so that at $T_c \simeq 170$ MeV, $L \sim 10^{-2}$, rather than zero."645 Deconfinement is thus indeed much like the transition. for which the order parameter. the conductivity (7). also does not really vanish for T'>0. but with o(7)~exp{-AL/T} is only exponentially small. since thermal ionisation (with ionisation energv AZ) produces a small number of unbound electrons even in (he insulator phase.," Deconfinement is thus indeed much like the insulator-conductor transition, for which the order parameter, the conductivity $\sigma(T)$, also does not really vanish for $T>0$, but with $\sigma(T) \sim \exp\{-\Delta E/T\}$ is only exponentially small, since thermal ionisation (with ionisation energy $\Delta E$ ) produces a small number of unbound electrons even in the insulator phase."646 rel?ja illustrates (heschematically (he behavior of LCD) and of the corresponding susceptibility VLUD)e(L2)=(L7. as obtained in finite temperature lattice studies [2022].. for the case of two flavors of lisht quarks.," \\ref{2_4}$ $~\!$ a illustrates theschematically the behavior of $L(T)$ and of the corresponding susceptibility $\x_L(T) \sim \langle L^2 \rangle - \langle L \rangle^2$, as obtained in finite temperature lattice studies \cite{K&L,cheng75,647cheng77}, for the case of two flavors of light quarks."648 We note that L(7) undergoes the expected sudden increase from a small confinement to a much larger deconfinement value., We note that $L(T)$ undergoes the expected sudden increase from a small confinement to a much larger deconfinement value.649 The sharp peak of νε} defines «quite well a transition temperature 7). which we shall shortly specify in physical units.," The sharp peak of $\chi_L(T)$ defines quite well a transition temperature $T_L$, which we shall shortly specify in physical units."650 The next quantity to consider is (he effective quark mass: it is measured by (he expectation value of the corresponding term in the Lagrangian. (c0)(T).," The next quantity to consider is the effective quark mass; it is measured by the expectation value of the corresponding term in the Lagrangian, $\langle {\bar \psi} \psi \rangle(T)$."651 In the limit of vanishing current quark mass. (he Lagraneian becomes chirally νοτας and (σὺ(1) the corresponding order parameter.," In the limit of vanishing current quark mass, the Lagrangian becomes chirally symmetric and $\langle {\bar \psi} \psi \rangle(T)$ the corresponding order parameter."652" In the confined phase. with effective constituent. quark masses AL,c0.3 GeV. this chiral svuumetry is spontaneously broken. while in the deconfinecl phase. athieh enough temperature. we expect its restoration."," In the confined phase, with effective constituent quark masses $M_q \simeq 0.3$ GeV, this chiral symmetry is spontaneously broken, while in the deconfined phase, athigh enough temperature, we expect its restoration."653" Hence now (cc)(T) constitutes a genuine order parameter. [mite for 7<7), ancl vanishing for 7> T),. as shown in relchi.."," Hence now $\langle {\bar \psi} \psi \rangle(T)$ constitutes a genuine order parameter, finite for $T< T_m$ and vanishing for $T\geq T_m$ , as shown in \\ref{chi}. ."654Millisecond (recycled) radio pulsars are distinguished from ordinary pulsars by their very short and stable periods. P< ms. P—102!107? s sl,"Millisecond (recycled) radio pulsars are distinguished from ordinary pulsars by their very short and stable periods, $P\la 10$ ms, $\dot{P}\sim 10^{-21}-10^{-19}$ s $^{-1}$."655 [t is generally accepted that they are very old objects. with spin-down ages P/2P~10°—10? yr and low surface magnetic fields B(PP)!v40—10'? G (eg. Taylor. Manchester. Lyne 1993).," It is generally accepted that they are very old objects, with spin-down ages $\tau=P/2\dot{P}\sim 10^9-10^{10}$ yr and low surface magnetic fields $B\propto (P\dot{P})^{1/2}\sim 10^8-10^{10}$ G (e.g., Taylor, Manchester, Lyne 1993)."656 Similar to ordinary pulsars. a millisecond pulsar can emit nonthermal X-rays from its magnetosphere. with a hard power-law spectrum and sharp pulsations.," Similar to ordinary pulsars, a millisecond pulsar can emit nonthermal X-rays from its magnetosphere, with a hard power-law spectrum and sharp pulsations."657 In addition to this nonthermal radiation. thermal X-rays can be emitted from the neutron star (NS) surface. provided the surface is hot enough.," In addition to this nonthermal radiation, thermal X-rays can be emitted from the neutron star (NS) surface, provided the surface is hot enough."658 According to the models of NS thermal evolution (albeit rather uncertain at these old ages). recycled pulsars are too cold (surface temperature Τ<0.1 MK — see. e. g.. Tsuruta 1998) to be detectable in X-rays.," According to the models of NS thermal evolution (albeit rather uncertain at these old ages), recycled pulsars are too cold (surface temperature $T\la 0.1$ MK — see, e. g., Tsuruta 1998) to be detectable in X-rays."659 However. their polar caps can be heated up to X-ray temperatures by relativistic particles impinging onto the magnetic poles from the acceleration zones in the magnetosphere.," However, their polar caps can be heated up to X-ray temperatures by relativistic particles impinging onto the magnetic poles from the acceleration zones in the magnetosphere."660 The radio pulsar models (e.g.. Cheng Ruderman 1980: Arons 1981: Michel 1991; Beskin. Gurevich. Istomin 1993) predict polar cap radii Ay.~QzRPol? (where R~10 km ts the NS radius). i.e... Ry;— 1-5 km for millisecond pulsars. although different models predict quite different polar cap temperatures. in the range of 1-10 MK.," The radio pulsar models (e.g., Cheng Ruderman 1980; Arons 1981; Michel 1991; Beskin, Gurevich, Istomin 1993) predict polar cap radii $\rpc\sim (2\pi R^3/Pc)^{1/2}$ (where $R\approx 10$ km is the NS radius), i.e., $\rpc\sim 1$ –5 km for millisecond pulsars, although different models predict quite different polar cap temperatures, in the range of 1–10 MK."661 Detection of the polar cap thermal radiation would allow one to discriminate between various models of radio pulsars. study the properties of NS surface layers. and constrain the NS mass-to-radius ratio (Pavlov Zavlin 1997; Zavlin Pavlov 1998 [ZP98]).," Detection of the polar cap thermal radiation would allow one to discriminate between various models of radio pulsars, study the properties of NS surface layers, and constrain the NS mass-to-radius ratio (Pavlov Zavlin 1997; Zavlin Pavlov 1998 [ZP98])."662 However. just as in the case of ordinary pulsars. this radiation is detectable only if it is not buried under stronger nonthermal radiation.," However, just as in the case of ordinary pulsars, this radiation is detectable only if it is not buried under stronger nonthermal radiation."663 The current theoretical models are not elaborate enough to predict in which (if any) of millisecond pulsars the thermal component can be brighter than the nonthermal one (in particular. both the thermal and nonthermal luminosities are expected to increase with energy loss E. perhaps with different rates).," The current theoretical models are not elaborate enough to predict in which (if any) of millisecond pulsars the thermal component can be brighter than the nonthermal one (in particular, both the thermal and nonthermal luminosities are expected to increase with energy loss $\dot{E}$, perhaps with different rates)."664 Therefore. we have to rely upon the analysis of X-ray observations to distinguish the thermal and nonthermal components.," Therefore, we have to rely upon the analysis of X-ray observations to distinguish the thermal and nonthermal components."665 The X-ray observatoriesROSAT..ASCA.. and have detected 11 millisecond pulsars (nearly 1/3 of all X-ray-detected rotation-powered pulsars — see Becker Pavlov 2001 for a recent review).," The X-ray observatories, and have detected 11 millisecond pulsars (nearly 1/3 of all X-ray-detected rotation-powered pulsars — see Becker Pavlov 2001 for a recent review)."666 Five of these pulsars are identified in X-rays only by positional comeidence with the radio pulsars and. due to the low number of recorded counts. provide only crude flux estimates.," Five of these pulsars are identified in X-rays only by positional coincidence with the radio pulsars and, due to the low number of recorded counts, provide only crude flux estimates."667 The radiation from 3 pulsars — B1821-24 (Saito et al., The radiation from 3 pulsars — B1821--24 (Saito et al.668 1997). B1937+21 (Takahashi et al.," 1997), B1937+21 (Takahashi et al."669 2001). and JO218+4232 (Mineo et al.," 2001), and J0218+4232 (Mineo et al."670 2000) — is clearly nonthermal: their power-law spectra. detected with and up to energies of 5-10 keV. are very hard. with photon indices ~ |. and their pulse profiles show sharp. peaks.," 2000) — is clearly nonthermal: their power-law spectra, detected with and up to energies of 5–10 keV, are very hard, with photon indices $\gamma\sim 1$ , and their pulse profiles show sharp peaks."671 Interestingly. these 3 pulsars are characterized by particularly large E values. E=(2-20)«10? erg s!. and their magnetic fields at the light cylinder. By=ΒΚΔΙΣ~10° G. are close to that of the Crab pulsar.," Interestingly, these 3 pulsars are characterized by particularly large $\dot{E}$ values, $\dot{E}=(2-20)\times 10^{35}$ erg $^{-1}$, and their magnetic fields at the light cylinder, $B_{\rm lc}=B(R/R_{\rm lc})^3\sim 10^6$ G, are close to that of the Crab pulsar."672 The case for the other 3 pulsars — J0437-4715 (Becker Trümmper 1993. 1999 [BT93. BT99]|: ZP98). J2124—3358 (BT99). and JO030+0451 (Becker et al.," The case for the other 3 pulsars — J0437–4715 (Becker Trümmper 1993, 1999 [BT93, BT99]; ZP98), J2124–3358 (BT99), and J0030+0451 (Becker et al."673 2000) — is less certain., 2000) — is less certain.674 These pulsars show broad peaks of X-ray pulsations. but it does not necessarily mean that their radiation is thermal because broad peaks can be produced by nonthermal emission at some viewing angles.," These pulsars show broad peaks of X-ray pulsations, but it does not necessarily mean that their radiation is thermal because broad peaks can be produced by nonthermal emission at some viewing angles."675 High-quality spectra have been recorded for the brightest of these pulsars. J0437—4715. but their interpretation has been controversial — e.g.. ZP98 suggest that the radiation detected with and can be interpreted as thermal radiation from hot polar caps. whereas BT99 argue that the radiation is nonthermal (see $22).," High-quality spectra have been recorded for the brightest of these pulsars, J0437–4715, but their interpretation has been controversial — e.g., ZP98 suggest that the radiation detected with and can be interpreted as thermal radiation from hot polar caps, whereas BT99 argue that the radiation is nonthermal (see 2)."676 To resolve this controversy. the pulsar needed to be observed at energies abovethe soft and bands (EmS>2 keV). and with high spatial resolution to avoid," To resolve this controversy, the pulsar needed to be observed at energies abovethe soft and bands $E\ga 2$ keV), and with high spatial resolution to avoid"677a phase transition (22222)..,"a phase transition \citep{Baym:1995fk,Martin:1995su,Hindmarsh:1997tj,Boyanovsky:2002wa,Kahniashvili:2009qi}."678 Fields can also be produced by the production of non-linear vorticity from linear density perturbations(???????).. but the impact these have on the CMB ts complicated by their evolving. non-trivial nature.," Fields can also be produced by the production of non-linear vorticity from linear density \citep{Gopal:2004ut,Matarrese:2004kq,Takahashi:2005nd,Ichiki:2006cd,Siegel:2006px,Kobayashi:2007wd,Maeda:2008dv}, but the impact these have on the CMB is complicated by their evolving, non-trivial nature."679 The impact primordial magnetic fields have on the CMB and its anisotropies have been well-studied. with (22222222222922222?) being some instructive examples.," The impact primordial magnetic fields have on the CMB and its anisotropies have been well-studied, with \citep{Barrow:1997mj,Subramanian:1998fn,Durrer:1998ya,Koh:2000qw,Kahniashvili:2000vm,Mack:2001gc,Clarkson:2002dd,Lewis:2004ef,Kahniashvili:2006hy,Kahniashvili:2008hx,Yamazaki:2008gr,Finelli:2008xh,Paoletti:2008ck,Bonvin:2010nr,Giovannini:2009fu,Yamazaki:2010nf,Paoletti:2010rx,Kahniashvili:2010wm} being some instructive examples."680 While older literature tended to assume a homogeneous background component with an inhomogeneous perturbation. more recent work has typically focused on tangled configurations without a background component and | assume this throughout.," While older literature tended to assume a homogeneous background component with an inhomogeneous perturbation, more recent work has typically focused on tangled configurations without a background component and I assume this throughout."681 These studies fairly consistently suggest that the field is constrained to be of at most nano-Gauss in magnitude., These studies fairly consistently suggest that the field is constrained to be of at most nano-Gauss in magnitude.682 The spectral index is restricted to be approximately scale-invariant (??).. with limits growing extremely tight for a primordial magnetic field with index far from scale-invarianee (2)..," The spectral index is restricted to be approximately scale-invariant \citep{Yamazaki:2010nf,Paoletti:2010rx}, with limits growing extremely tight for a primordial magnetic field with index far from scale-invariance \citep{Caprini:2001nb}."683 A large-scale homogeneous field also introduces characteristic correlations between multipole moments with A/€[-2.0.2] and Am€(0.x1.£2] which vanish in the standard scenario (?)..," A large-scale homogeneous field also introduces characteristic correlations between multipole moments with $\Delta l\in\{-2,0,2\}$ and $\Delta m\in\{0,\pm 1,\pm2\}$ which vanish in the standard scenario \citep{Kahniashvili:2008sh}."684 However. the magnetic 2-point signal is overwhelmed on large-scales by the standard perturbations. with the B-mode polarisation being perhaps the most realistic option if we are to detect it directly.," However, the magnetic 2-point signal is overwhelmed on large-scales by the standard perturbations, with the $B$ -mode polarisation being perhaps the most realistic option if we are to detect it directly."685 The increasing accuracy of measurements of the CMB non-Gaussianity provides an alternative., The increasing accuracy of measurements of the CMB non-Gaussianity provides an alternative.686 The stress tensor of a magnetic field is implying that the statistics induced on matter perturbations are intrinsically non-Gaussian. regardless of the nature of the underlying magnetic field.," The stress tensor of a magnetic field is non-linear, implying that the statistics induced on matter perturbations are intrinsically non-Gaussian, regardless of the nature of the underlying magnetic field."687 Since the standard scenario contains relatively few sources of primordial it is possible that a magnetic signal is dominant.," Since the standard scenario contains relatively few sources of primordial non-Gaussianity, it is possible that a magnetic signal is dominant."688 Viewed another way. predicted signals from a magnetic field are likely to be of a characteristic nature. and must be found in and cleaned from the CMB data before any conclusions on early-universe physics can be made.," Viewed another way, predicted signals from a magnetic field are likely to be of a characteristic nature, and must be found in and cleaned from the CMB data before any conclusions on early-universe physics can be made."689 Aspects of the three-point moments have been studied in a series of papers in the last few years (22222222) ," Aspects of the three-point moments have been studied in a series of papers in the last few years \citep{Brown:2005kr,Brown:2006wv,Seshadri:2009sy,Caprini:2009vk,Trivedi:2010gi,Cai:2010uw,Shiraishi:2010yk,Kahniashvili:2010us}. ."690A bispectrum is set by three wavevectors. which we denote with k. p and q.," A bispectrum is set by three wavevectors, which we denote with $\mathbf{k}$, $\mathbf{p}$ and $\mathbf{q}$."691 Since to retain statistical isotropy these must form a closed triangle. this geometry can equivalently be expressed with the scalars &.7.o. where p=rk and ὁ is the angle between p and q.," Since to retain statistical isotropy these must form a closed triangle, this geometry can equivalently be expressed with the scalars $k,r,\phi$, where $p=rk$ and $\phi$ is the angle between $\mathbf{p}$ and $\mathbf{q}$ ."692 Employing these variables the bispectrum geometry can be written as a foliation of planes of constant rand for each constant angle @ we then have a one-dimensional line through the bispectrum which in broad terms is expected to act ina similar manner to the power spectra., Employing these variables the bispectrum geometry can be written as a foliation of planes of constant $r$ and for each constant angle $\phi$ we then have a one-dimensional line through the bispectrum which in broad terms is expected to act in a similar manner to the power spectra.693" The magnetic bispectra studied thus far have typically been along only three such lines. all in the 7=1 plane — the ""colinear case where k=p4/2 and so 6=0 (???.hereafterBCOS.BO6andCFPRO9).. the “equilateral” case wherek=p qand so 6=27/3 (???.hereafterSS09.CFPRO9andTSS10).. and the ""local"" or degenerate case where kοὐ=p. q=0 and so @=x (SS09. CFPRO9. TSS10)."," The magnetic bispectra studied thus far have typically been along only three such lines, all in the $r=1$ plane – the “colinear” case where $k=p=q/2$ and so $\phi=0$ \citep[hereafter BC05, B06 and CFPR09]{Brown:2005kr,Brown:2006wv,Caprini:2009vk}, the “equilateral” case where $k=p=q$ and so $\phi=2\pi/3$ \citep[hereafter SS09, CFPR09 and TSS10]{Seshadri:2009sy,Caprini:2009vk,Trivedi:2010gi}, and the “local” or degenerate case where $k\approx p$, $q\approx 0$ and so $\phi\approx\pi$ (SS09, CFPR09, TSS10)."694 TSS10 also considered configurations where @=0 but kzp., TSS10 also considered configurations where $\phi=0$ but $k\neq p$.695 The recent studies have expanded the previous results considerably., The recent studies have expanded the previous results considerably.696 SS09 considered the equilateral and degenerate lines of the bispectrum of the magnetic energy density for nearly scale-invariant magnetic fields. concluding that the degenerate line provides the greatest contribution to the integral and employing an approximation to this dominant term to estimate the CMB signal.," SS09 considered the equilateral and degenerate lines of the bispectrum of the magnetic energy density for nearly scale-invariant magnetic fields, concluding that the degenerate line provides the greatest contribution to the integral and employing an approximation to this dominant term to estimate the CMB signal."697 Likewise. CFPRO9 considered the bispectrum of the energy density and considered the colinear. equilateral and degenerate lines.," Likewise, CFPR09 considered the bispectrum of the energy density and considered the colinear, equilateral and degenerate lines."698 The authors generally relied on approximations that neglect angular terms in the integrations. or apply only on large scales.," The authors generally relied on approximations that neglect angular terms in the integrations, or apply only on large scales."699 Doing so recovers the scaling behaviour of the bispectrum at the expense of an accurate calculation of the relative amplitudes between lines., Doing so recovers the scaling behaviour of the bispectrum at the expense of an accurate calculation of the relative amplitudes between lines.700" Since the degenerate line was found to diverge as g-""~ as g—O this term ts likely to dominate.", Since the degenerate line was found to diverge as $q^{2n+3}$ as $q\rightarrow 0$ this term is likely to dominate.701 CFPRO9 also present exact solutions for the colinear case for both a causal field and a field relatively close to scale-invariance. which enable them to test their approximations.," CFPR09 also present exact solutions for the colinear case for both a causal field and a field relatively close to scale-invariance, which enable them to test their approximations."702 The approximations are certainly reasonable. but not ideal.," The approximations are certainly reasonable, but not ideal."703 In particular. since the bispectra are not positive-definite it is unclear whether there are strong cancellations to the degenerate line arising from other parts of the bispectrum.," In particular, since the bispectra are not positive-definite it is unclear whether there are strong cancellations to the degenerate line arising from other parts of the bispectrum."704 ? employed the approximations of SS09 and CFPRO9 and extended the treatment to full transfer functions., \citet{Cai:2010uw} employed the approximations of SS09 and CFPR09 and extended the treatment to full transfer functions.705 Morerecently. TSS10 considered the bispectrum of theanisotropic pressure of a near scale-invariant magnetic field.," Morerecently, TSS10 considered the bispectrum of theanisotropic pressure of a near scale-invariant magnetic field."706 Unlike the previous papers they evaluated the bispectrum along the degenerate line in full. without neglecting any angular," Unlike the previous papers they evaluated the bispectrum along the degenerate line in full, without neglecting any angular"707Finally we look at the RV data themselves.,"Finally, we look at the RV data themselves."708 The calculation of the formal uncertainties assumes that the RV errors scale linearly with the inverse of the spectrum SNR (Bouchyetal.2001)., The calculation of the formal uncertainties assumes that the RV errors scale linearly with the inverse of the spectrum SNR \citep{bou01}.709". On the other hand, we suspect that the dependency of the errors on the SNR is stronger than linear."," On the other hand, we suspect that the dependency of the errors on the SNR is stronger than linear."710" To test this, we divided the RV residuals shown in Fig."," To test this, we divided the RV residuals shown in Fig."711 3 by the SNR of the corresponding spectra., \ref{fig3} by the SNR of the corresponding spectra.712" The scatter of the lower-SNR data points is significantly larger than that of the high-SNR points, by a much larger amount than indicated by the formal error bars."," The scatter of the lower-SNR data points is significantly larger than that of the high-SNR points, by a much larger amount than indicated by the formal error bars."713 The same holds for the residuals relative to the mean., The same holds for the residuals relative to the mean.714 This observation clearly indicates that the dependency of the uncertainties on SNR is underestimated., This observation clearly indicates that the dependency of the uncertainties on SNR is underestimated.715" It can be modelled by assuming the presence of additional instrumental uncertainties with a steep dependence on SNR, at the ~5 mmss! level near the median SNR, increasing to about ~10 ! at the low end of the SNR range."," It can be modelled by assuming the presence of additional instrumental uncertainties with a steep dependence on SNR, at the $\sim 5$ $^{-1}$ level near the median SNR, increasing to about $\sim 10$ $^{-1}$ at the low end of the SNR range."716 The standard way of searching for the signature of planetary orbits in RV data is to use a Lomb-Scargle or generalized periodogram (Horne&Baliunas1986;PressRybicki1989;Zechmeister&Kürster 2009).," The standard way of searching for the signature of planetary orbits in RV data is to use a Lomb-Scargle or generalized periodogram \citep{Hor86,Pre86,Zec09}."717. The periodogram of the HARPS RV data for CoRoT-7 is highly complex., The periodogram of the HARPS RV data for CoRoT-7 is highly complex.718" In the course of their pre-whitening analysis, Q09 identify no less than eleven peaks, all of them highly ‘significant’ in the sense that they correspond to low formal false alarm probabilities."," In the course of their pre-whitening analysis, Q09 identify no less than eleven peaks, all of them highly `significant' in the sense that they correspond to low formal false alarm probabilities."719" However, on should bear in mind that the false alarm probability expresses the probability that a given peak is due to Gaussian white noise, but neither activity nor instrumental noise are expected to be white or Gaussian."," However, on should bear in mind that the false alarm probability expresses the probability that a given peak is due to Gaussian white noise, but neither activity nor instrumental noise are expected to be white or Gaussian."720" The very irregular sampling of the data also implies that one or more of the peaks may arise from signal at a single, apparently unrelated frequency."," The very irregular sampling of the data also implies that one or more of the peaks may arise from signal at a single, apparently unrelated frequency."721" Nonetheless, the main peaks in the RV periodogram are clearly related to the signal from activity, being near the rotation period, its harmonics, and their one-day aliases."," Nonetheless, the main peaks in the RV periodogram are clearly related to the signal from activity, being near the rotation period, its harmonics, and their one-day aliases."722" There is also a peak corresponding to the period of the transit signal detected in the CoRoT photometry, P;=0.854 dd. Fig."," There is also a peak corresponding to the period of the transit signal detected in the CoRoT photometry, $P_b=0.854$ d. Fig."723" 4 shows the semi-amplitude of the best-fit sinusoid at the period and phase of the photometric transit, as a function of SNR threshold, as measurements derived from low-SNR spectra are progressively discarded, starting with the lowest SNR."," \ref{fig6} shows the semi-amplitude of the best-fit sinusoid at the period and phase of the photometric transit, as a function of SNR threshold, as measurements derived from low-SNR spectra are progressively discarded, starting with the lowest SNR."724" For the most stringent threshold (SNR>1.15SNRmea), about one third of the measurements remain."," For the most stringent threshold ${\rm SNR}>1.15\,{\rm SNR}_{\rm725med}$ ), about one third of the measurements remain."726" We performed this calculation with both the raw RV, and the residuals from our activity models."," We performed this calculation with both the raw RV, and the residuals from our activity models."727" The stellar rotation and planetary orbital frequencies are widely separated, but the latter is close to the one-day alias of the third harmonic of the former."," The stellar rotation and planetary orbital frequencies are widely separated, but the latter is close to the one-day alias of the third harmonic of the former."728" As a result, it is not clear a priori whether correcting for the activity signal improves the semi-amplitude measurements, or on the contrary adds noise to them."," As a result, it is not clear a priori whether correcting for the activity signal improves the semi-amplitude measurements, or on the contrary adds noise to them."729" Fortunately, the two methods give very similar results."," Fortunately, the two methods give very similar results."730" In both cases, the measured orbital semi-amplitude depends strongly on the SNR threshold: including lower-SNR measurements favours a higher value."," In both cases, the measured orbital semi-amplitude depends strongly on the SNR threshold: including lower-SNR measurements favours a higher value."731" Low-SNR measurements are more likely to be outliers (as their formal uncertainties are underestimated), and would favour a higher amplitude for all fitted features: the higher number of measurements is offset by their poorer quality."," Low-SNR measurements are more likely to be outliers (as their formal uncertainties are underestimated), and would favour a higher amplitude for all fitted features: the higher number of measurements is offset by their poorer quality."732" It is therefore not clear whether the most reliable value of Ky is the one derived from all the RV measurements, or from only the best third or half."," It is therefore not clear whether the most reliable value of $K_b$ is the one derived from all the RV measurements, or from only the best third or half."733 The monotonic trend in Ky versus SNR threshold suggests that SNR-dependent effects play a large role in the detected amplitude., The monotonic trend in $K_b$ versus SNR threshold suggests that SNR-dependent effects play a large role in the detected amplitude.734" Figure 5 shows the radial-velocity data corrected by one of our variability models and phased to the transit signal, together with the best-fit Keplerian orbit with and without an SNR cut."," Figure \ref{phasedvr}735 shows the radial-velocity data corrected by one of our variability models and phased to the transit signal, together with the best-fit Keplerian orbit with and without an SNR cut."736" 'To evaluate the effect of the steeper rise in the total uncertainty at the lower end of the SNR range, we re-calculate the constraints on K in the following way: we add a new term to the radial-velocity uncertainties, with a quadratic rather than linear dependence on the inverse of the signal-to-noise ratio."," To evaluate the effect of the steeper rise in the total uncertainty at the lower end of the SNR range, we re-calculate the constraints on $K$ in the following way: we add a new term to the radial-velocity uncertainties, with a quadratic rather than linear dependence on the inverse of the signal-to-noise ratio."737 We set the magnitude of this term so that the reduced x? of the residuals of the, We set the magnitude of this term so that the reduced $\chi^2$ of the residuals of the738among different observables (MHDO3: McHardy et al.,among different observables (MHD03; McHardy et al.739 2006). supporting the notion that AGN are indeed scaled-up galactic black holes.," 2006), supporting the notion that AGN are indeed scaled-up galactic black holes."740 However. physical models for the disc-jet coupling in BHXRB based on the observed correlations between radio and X-ray luminosity (Fender.Gallo&Jonker2003:Kórding.Fender&Migliari2006) all face a large uncertainty due to the lack of reliable measurements of the jet kinetic power.," However, physical models for the disc-jet coupling in BHXRB based on the observed correlations between radio and X-ray luminosity \cite{fender:03,koerding:06b} all face a large uncertainty due to the lack of reliable measurements of the jet kinetic power."741 In this context. it is interesting to include in Figure 2. the only GBH for which a measurement of the kinetic output has been made. Cyg X-I (Galloetal.2005)... which turns out to be consistent with the relationship derived from the AGN sample.," In this context, it is interesting to include in Figure \ref{fig:lkinledd} the only GBH for which a measurement of the kinetic output has been made, Cyg X-1 \cite{gallo:05}, which turns out to be consistent with the relationship derived from the AGN sample."742 Clearly. systematic efforts to estimate kinetic power of BHXRB in the low/hard state are needed in order to assess their similarity with radio mode AGN.," Clearly, systematic efforts to estimate kinetic power of BHXRB in the low/hard state are needed in order to assess their similarity with radio mode AGN."743 If the above correlation (3)) directly reveals fundamental physical properties of jet-producing AGN of lowpower. it still shows a non- intrinsic scatter.," If the above correlation \ref{eq:lklambda}) ) directly reveals fundamental physical properties of jet-producing AGN of lowpower, it still shows a non-negligible intrinsic scatter."744 On the other hand. one should expect a more direet relationship between the nuclear radio core emission and the larger scale Kinetic power. as both originate from the jet.," On the other hand, one should expect a more direct relationship between the nuclear radio core emission and the larger scale kinetic power, as both originate from the jet."745 All theoretical models for AGN flat-spectrum compact jet cores (Blandford&Kónigl1979:FalekeBiermann1996:HeinzSunyaev2003) predict a dependence of the radio luminosity on the jet power in the form Ly-x ," All theoretical models for AGN flat-spectrum compact jet cores \cite{blandford:79,falcke:96,heinz:03} a dependence of the radio luminosity on the jet power in the form $L_{\rm R} \propto746L_{\rm kin}^{17/12}$."747The current sample provides by far the best opportunity to test Lus.these predictions., The current sample provides by far the best opportunity to test these predictions.748 A Kendalls tau correlation test reveals that the kinetic power is correlated with theobserved radio core luminosity Lion... with Pun=9210 (see the empty circles in Figure +)).," A Kendall's tau correlation test reveals that the kinetic power is correlated with the radio core luminosity $L_{\rm749R,obs}$, with $P_{\rm null}= 9.2\times 10^{-5}$ (see the empty circles in Figure \ref{fig:pboth}) )."750 We have fitted the data with a linear relationship:Ly. once again making use of a symmetric regression algorithm that takes into account errors on both variables.," We have fitted the data with a linear relationship:, once again making use of a symmetric regression algorithm that takes into account errors on both variables."751 We obtain clan.=(22.1$3.5). Ba=(0.54+0.09). with a large intrinsic scatter of y=0.41.," We obtain $A_{\rm obs}= (22.1 \pm 3.5)$, $B_{\rm obs}=(0.54 \pm 0.09)$, with a large intrinsic scatter of $\sigma=0.47$."752 Such a correlation. however. must be at some level biased by relativistic Doppler boosting of the radio emission in the relativistic jets.," Such a correlation, however, must be at some level biased by relativistic Doppler boosting of the radio emission in the relativistic jets."753" An alternative way to proceed would be to use estimators of the nuclear radio core luminosity which are less affected by relativistic beaming (Heinz&Grimm2005).. as. for example. the multivariate relation between BH mass. radio core and hard X-ray luminosity. the so-called ""fundamental plane’ (FP) of active black holes (MHDO3)."," An alternative way to proceed would be to use estimators of the nuclear radio core luminosity which are less affected by relativistic beaming \cite{heinz:05}, as, for example, the multivariate relation between BH mass, radio core and hard X-ray luminosity, the so-called `fundamental plane' (FP) of active black holes (MHD03)."754 Recent analysis of this correlation (Heinz Merloni 2004: Kórrding. Faleke Corbel 2006 [KFCO6]: lerloni et al.," Recent analysis of this correlation (Heinz Merloni 2004; Körrding, Falcke Corbel 2006 [KFC06]; Merloni et al."755 2006) have shown that both Doppler boosting and sample selection play a crucial role in the exact determination of he intrinsic correlation coefficients of the FP. which also need to be accounted for.," 2006) have shown that both Doppler boosting and sample selection play a crucial role in the exact determination of the intrinsic correlation coefficients of the FP, which also need to be accounted for."756 In the Appendix. we diseuss in detail a possible way © overcome such a bias with the aid of a Monte Carlo simulation of the samples used to derive the FP relation.," In the Appendix, we discuss in detail a possible way to overcome such a bias with the aid of a Monte Carlo simulation of the samples used to derive the FP relation."757" That study allows us to estimate statistically the intrinsic (un-boosted) radio core uminosity of the AGN jets as a function of their (mean) Lorentz actor. Li, in a way that ean be approximated by the following expression: where Lip ds the intrinsic (un-boosted) radio core luminosity of the jet at 5 GHz. Lx the nuclear 2-10 keV intrinsic (un-absorbed) luminosity and Ly, the mean Lorentz factor of the jets."," That study allows us to estimate statistically the intrinsic (un-boosted) radio core luminosity of the AGN jets as a function of their (mean) Lorentz factor, $\Gamma_{\rm m}$ in a way that can be approximated by the following expression: where $ L_{\rm R,FP}$ is the intrinsic (un-boosted) radio core luminosity of the jet at 5 GHz, $L_{\rm X}$ the nuclear 2-10 keV intrinsic (un-absorbed) luminosity and $\Gamma_{\rm m}$ the mean Lorentz factor of the jets."758 In what follows. we will adopt the specitic version of the FP relation derived from a sample of low luminosity AGN only. i.e. free from the bias introduced by the inclusion of bright. radiatively efficient AGN or QSOs (see discussion in KFCO6).," In what follows, we will adopt the specific version of the FP relation derived from a sample of low luminosity AGN only, i.e. free from the bias introduced by the inclusion of bright, radiatively efficient AGN or QSOs (see discussion in KFC06)."759 For that. the correlation coefficients are ax=0.71. £n;=0.62. slightly different (but only at the 1-7 level) from those found in MHDO3.," For that, the correlation coefficients are $\xi_{\rm RX}=0.71$, $\xi_{\rm RM}=0.62$, slightly different (but only at the $\sigma$ level) from those found in MHD03."760 Given Eq. (4)).," Given Eq. \ref{eq:fp_corr}) ),"761" assuming a distribution of Lorentz factors for the AGN jets tor its mean. provided that the distribution is not too broad). we determine the ""true"" relationship between the Kinetic luminosity and radio core luminosityby fitting the 15 data points in our sample with the linear relationship The fitted values for the intrinsic slope. Di. as a function of Puy. are shown as a dot-dashed line in the bottom panel of Figure 3.."," assuming a distribution of Lorentz factors for the AGN jets (or its mean, provided that the distribution is not too broad), we determine the “true” relationship between the Kinetic luminosity and radio core luminosityby fitting the 15 data points in our sample with the linear relationship The fitted values for the intrinsic slope, $B_{\rm int}$, as a function of $\Gamma_{\rm m}$, are shown as a dot-dashed line in the bottom panel of Figure \ref{fig:bobs}."762 From this we can see that the higher the mean Lorentz ‘actor of the jets. the steeper must the intrinsic correlation between Kinetic power and jet core luminosity be. and the larger the discrepancy with the measured slope. 4...=0.54+0.09 (solid ines in Figure 39). obtained using simply the observed radio core uminosity. without any attempt to correct for relativistic beaming.," From this we can see that the higher the mean Lorentz factor of the jets, the steeper must the intrinsic correlation between kinetic power and jet core luminosity be, and the larger the discrepancy with the measured slope, $B_{\rm763obs}=0.54 \pm 0.09$ (solid lines in Figure \ref{fig:bobs}) ), obtained using simply the observed radio core luminosity, without any attempt to correct for relativistic beaming."764 In fact. such a discrepancy between the intrinsic and the observed slopes of the Ly - Li; relation is indeed expected if the 15 sources of our sample harbor relativistic jet randomly oriented with respect o the line of sight?.," In fact, such a discrepancy between the intrinsic and the observed slopes of the $L_{\rm R}$ - $L_{\rm kin}$ relation is indeed expected if the 15 sources of our sample harbor relativistic jet randomly oriented with respect to the line of ."765. In order to show this quantitatively. we have simulated (107 imes) the observed sample. assuming an underlying relationship," In order to show this quantitatively, we have simulated $10^4$ times) the sample, assuming an underlying relationship"766"CH, mixing ratio with altitude with a scale height of ~20 km. due to photolysis.","$_4$ mixing ratio with altitude with a scale height of $\sim$ 20 km, due to photolysis."767 We obtained the same column density as above. indicating a partial pressure of methane of 9.843.7 nbar. 1.8. a surface density of (1.940.7) x 107 οι”.," We obtained the same column density as above, indicating a partial pressure of methane of $\pm$ 3.7 nbar, i.e. a surface density of $\pm$ 0.7) x $^{12}$ $^{-3}$."768" This appears to be 43. times larger than inferred from Voyager in 1989. adopting the CH, number densities of Herbert and Sandel (1991) and Strobel and Summers (1995) (4.7x10!! em"". within a factor 1.7. averaging ingress and egress)."," This appears to be $^{+5}_{-2.5}$ times larger than inferred from Voyager in 1989, adopting the $_4$ number densities of Herbert and Sandel (1991) and Strobel and Summers (1995) $\times$ $^{11}$ $^{-2}$, within a factor 1.7, averaging ingress and egress)."769" An even larger enhancement factor (57°) is indicated 1f the Krasnopolsky and Cruikshank (1995) reanalysis of the Voyager UV data. giving CH; = 3.1z0.8x 10"" em? at the surface. is used."," An even larger enhancement factor $^{+6}_{-2}$ ) is indicated if the Krasnopolsky and Cruikshank (1995) reanalysis of the Voyager UV data, giving $_4$ = $\pm$ $\times$ $^{11}$ $^{-3}$ at the surface, is used."770 Results are independent on the surface pressure. as collisional broadening is negligible.," Results are independent on the surface pressure, as collisional broadening is negligible."771" They clearly demonstrate that the CH, partial pressure has increased in the last 20 years.", They clearly demonstrate that the $_4$ partial pressure has increased in the last 20 years.772 The 2335-2365 nm part of the Triton spectrum (see excerpts in Fig., The 2335-2365 nm part of the Triton spectrum (see excerpts in Fig.773 3) shows the detection of 8 lines due to the CO(2-0) band (R2-R5. P2. P3. P5 and P5). providing the first detection of CO in its atmosphere.," 3) shows the detection of 8 lines due to the CO(2-0) band (R2-R5, P2, P3, P5 and P8), providing the first detection of CO in its atmosphere."774 An accurate determination of the CO abundance is particularly difficult. as at infinite spectral resolution. these features are very narrow. saturated Doppler-shaped lines.," An accurate determination of the CO abundance is particularly difficult, as at infinite spectral resolution, these features are very narrow, saturated Doppler-shaped lines."775 Nonetheless. assuming a vertically uniform CO distribution. and utilizing the whole set of CO lines (see Fig.," Nonetheless, assuming a vertically uniform CO distribution, and utilizing the whole set of CO lines (see Fig."776 + on-line). we determine a CO column of 0.30 em-am. re. a CO partial pressure of 24 nbar. within a factor of 3.," 4 on-line), we determine a CO column of 0.30 cm-am, i.e. a CO partial pressure of 24 nbar, within a factor of 3."777" The column density CO/CH, ratio is nominally ~3.75 (surface partial pressure ratio CO/CH, -2.5). with a factor of 4 uncertainty."," The column density $_4$ ratio is nominally $\sim$ 3.75 (surface partial pressure ratio $_4$ $\sim$ 2.5), with a factor of 4 uncertainty."778 Deriving the CO/N> and CHj/N» mixing ratio is complicated by the fact that the surface pressure in 2009 is unknown., Deriving the $_2$ and $_4$ $_2$ mixing ratio is complicated by the fact that the surface pressure in 2009 is unknown.779 Stellar occultation results (Olkin et al., Stellar occultation results (Olkin et al.780" 1997, Sicardy et al."," 1997, Sicardy et al."781 1998. Elliotet al.," 1998, Elliot et al."782 1998. 20002) indicate that the pressure has been doubling in ~ 10 years from the 14 jibar value determined by Voyager in 1989 (Gurrola 1995).," 1998, 2000a) indicate that the pressure has been doubling in $\sim$ 10 years from the 14 $\mu$ bar value determined by Voyager in 1989 (Gurrola 1995)."783 A reasonable assumption for 2009 is 40 jibar. providing CO/Ns ~ 6x 1077 and CH4/N> -2.4x 1077 at the surface. within factors of3 and 1.4 respectively.," A reasonable assumption for 2009 is 40 $\mu$ bar, providing $_2$ $\sim$ $\times$ $^{-4}$ and $_4$ $_2$ $\sim$ $\times$ $^{-4}$ at the surface, within factors of 3 and 1.4 respectively."784 The CO abundance we determine is many times less than previous upper limits (Broadfoot et al., The CO abundance we determine is many times less than previous upper limits (Broadfoot et al.785 1989. Young et al.," 1989, Young et al."786 2001)., 2001).787" Near-infrared observations indicate that CO and CH, are present on Triton’s surface with mixing ratios of 0.05 and 0.1 relative to N». and at least for CH. mostly in solid solution in N? (Cruikshank et al."," Near-infrared observations indicate that CO and $_4$ are present on Triton's surface with mixing ratios of 0.05 and 0.1 relative to $_2$, and at least for $_4$, mostly in solid solution in $_2$ (Cruikshank et al."788 1993. Quirico et al.," 1993, Quirico et al."789" 1999, Grundy et al."," 1999, Grundy et al."790 2010)., 2010).791 In this situation. the expected partial pressure of each species is the product of its solid mole fraction and its pure vapor pressure (Raoult's law for an ideal mixture).," In this situation, the expected partial pressure of each species is the product of its solid mole fraction and its pure vapor pressure (Raoult's law for an ideal mixture)."792" This scenario leads to atmospheric abundances of CO and CH, that are about | and 3 orders of magnitude lower than observed. respectively (Fig."," This scenario leads to atmospheric abundances of CO and $_4$ that are about 1 and 3 orders of magnitude lower than observed, respectively (Fig."793 5)., 5).794" Although Henry's law may be more applicable than Raoult's in the case of the N:-CH, system."," Although Henry's law may be more applicable than Raoult's in the case of the $_2$ $_4$ system,"795noise and resolution properties and the presence of damped systems for the flux power.,noise and resolution properties and the presence of damped systems for the flux power.796 We summarize here the constraints found., We summarize here the constraints found.797" For the pdf: B=0.08£0.05 and 6<0.19 (2e C.L.) by using the flux pdf alone in the range F=[0.1— 0.8], with a reduced x/v=1.09 (35 d.o.£);"," For the pdf: $\beta=0.08\pm7980.05$ and $\beta<0.19$ $\sigma$ C.L.) by using the flux pdf alone in the range $F=[0.1-0.8]$ , with a reduced $\chi^2/\nu=1.09$ (35 d.o.f.);"799" 8=0.04+ and 8«0.1 (2c C.L.) by using the flux pdf alone in the whole range F=[0—1], with a reduced y?/v=1.21 (53 d.o.f.)."," $\beta=0.04\pm 0.04$ and $\beta<0.1$ $\sigma$ C.L.) by using the flux pdf alone in the whole range $F=[0-1]$, with a reduced $\chi^2/\nu=1.21$ (53 d.o.f.)."800" The ranges at low and high transmissivity are those that are most difficult to model due to the presence of strong systems and continuum fitting errors, respectively."," The ranges at low and high transmissivity are those that are most difficult to model due to the presence of strong systems and continuum fitting errors, respectively."801" Thus, we regard the first result presented as more conservative even though we do model continuum fitting errors and correct for numerical resolution (?).."," Thus, we regard the first result presented as more conservative even though we do model continuum fitting errors and correct for numerical resolution \citep{vbh09}."802" For the flux power only we obtain: 6=0.07+0.04 and 6<0.14 (2e C.L.) using all the 132 data points (x?/v=1.16, for 120 d.o.f.)."," For the flux power only we obtain: $\beta=0.07\pm 0.04$ and $\beta<0.14$ $\sigma$ C.L.) using all the 132 data points $\chi^2/\nu=1.16$, for 120 d.o.f.)."803 All these numbers are reasonable and demonstrate that the regions of high transmissivity have a constraining power which is stronger than the power spectrum alone., All these numbers are reasonable and demonstrate that the regions of high transmissivity have a constraining power which is stronger than the power spectrum alone.804" If we combine the two measurements we find the same trends as in ?:: there is not a very good fit to the data (x?=200 for 164 d.o.f.),"," If we combine the two measurements we find the same trends as in \cite{vbh09}: there is not a very good fit to the data $\chi^2=200$ for 164 d.o.f.),"805 and a reasonable x? is obtained only when neglecting the three highest redshift bins of the SDSS flux power., and a reasonable $\chi^2$ is obtained only when neglecting the three highest redshift bins of the SDSS flux power.806" In this case, we obtain 8=0.05-£0.03 and 8«0.1 (2e C.L.) with a reduced x?/v=1.09 (146 d.o.£.)."," In this case, we obtain $\beta=0.05\pm 0.03$ and $\beta<0.1$ $\sigma$ C.L.) with a reduced $\chi^2/\nu=1.09$ (146 d.o.f.)."807 All the other parameters are not affected significantly by the new parameter introduced and there are not strong degeneracies for B., All the other parameters are not affected significantly by the new parameter introduced and there are not strong degeneracies for $\beta$.808 From the analysis performed we can conclude that robust 2c upper limits on the coupling constant are in the range 6«0.1—0.2 (depending on the subset of data chosen)., From the analysis performed we can conclude that robust $2\sigma$ upper limits on the coupling constant are in the range $\beta<0.1-0.2$ (depending on the subset of data chosen).809 These bounds are exclusively derived by the analysis of the observed properties of the IGM and represent a completely new and independent test of cDE cosmologies w.r.t., These bounds are exclusively derived by the analysis of the observed properties of the IGM and represent a completely new and independent test of cDE cosmologies w.r.t.810 previous constraints (ase.g.???)..," previous constraints \citep[as \eg811][]{Bean_etal_2008,LaVacca_etal_2009,xia09}."812 We regard a 2c limit of 8<0.15 as a conservative overall bound once the statistical limitations of the different samples are taken into account., We regard a $2\sigma $ limit of $\beta \lesssim 0.15$ as a conservative overall bound once the statistical limitations of the different samples are taken into account.813 In this work we have explored the possibility of constraining the coupling 9 between CDM and DE through the statistical properties of the transmitted flux in forest QSO spectra at z—24.2., In this work we have explored the possibility of constraining the coupling $\beta $ between CDM and DE through the statistical properties of the transmitted flux in forest QSO spectra at $z=2-4.2$.814" For this purpose, we have performed the first high-resolution hydrodynamical simulations with gas cooling and star formation in the context of cDE models and quantitatively exploited the capabilities of flux 1-pt and 2-pt functions to constrain the strength of the coupling 6 between DE and CDM."," For this purpose, we have performed the first high-resolution hydrodynamical simulations with gas cooling and star formation in the context of cDE models and quantitatively exploited the capabilities of flux 1-pt and 2-pt functions to constrain the strength of the coupling $\beta$ between DE and CDM."815" 'The main results can be summarized as 'This work quantitatively shows that the range of scales and redshifts, where the growth of structures can be radically different from that measured from a naive extrapolation of either local or very high redshift probes, is promising for constraining coupled dark energy cosmologies."," The main results can be summarized as This work quantitatively shows that the range of scales and redshifts, where the growth of structures can be radically different from that measured from a naive extrapolation of either local or very high redshift probes, is promising for constraining coupled dark energy cosmologies."816" The increasing number of QSO spectra that are being collected (e.g.BOSS?]. offers the exciting prospect of further improvingX-Shooteif]) the numbers and of understanding in a more refined way, by performing simulations and by addressing systematic errors, the impact that coupled dark energy cosmologies can have on the diffuse gas at high redshift."," The increasing number of QSO spectra that are being collected (e.g., ) offers the exciting prospect of further improving the numbers and of understanding in a more refined way, by performing simulations and by addressing systematic errors, the impact that coupled dark energy cosmologies can have on the diffuse gas at high redshift."817" MB is supported by the DFG Cluster of Excellence ""Origin and Structure of the Universe"" and partly supported by the TRR Transregio Collaborative Research Network on the “Dark Universe”."," MB is supported by the DFG Cluster of Excellence “Origin and Structure of the Universe"" and partly supported by the TRR Transregio Collaborative Research Network on the “Dark Universe""."818" MV is partly supported by ASI/AAE, INFN-PD51 and PRIN/MIUR."," MV is partly supported by ASI/AAE, INFN-PD51 and PRIN/MIUR."819 Numerical simulations have been performed atRZG Computing Centre in Garching., Numerical simulations have been performed atRZG Computing Centre in Garching.820" Post processing and data analysis have been carried out at COSMOS and HPCS (Cambridge), and CINECA thanks to a CINECA/INAF grant."," Post processing and data analysis have been carried out at COSMOS and HPCS (Cambridge), and CINECA thanks to a CINECA/INAF grant."821Fieure G compares the joint distribution of colors and absolute magnitudes for ALLOL. AJSI. and the MIOL clusters found by Chandaretal.(2004).,"Figure \ref{cmd} compares the joint distribution of colors and absolute magnitudes for M101, M81, and the M101 clusters found by \citet{cwl04}."822.. Chandaretal.(2004) were specifically attempting to select old. elobular clusters in M1OL. while were searching onlv [or ‘compact’ star clusters in M81. with no selection on age.," \citet{cwl04} were specifically attempting to select old, globular clusters in M101, while \citet{cft1} were searching only for `compact' star clusters in M81 with no selection on age."823 The ellect of our {ραπ magnitude cut is particularly clear in the right-hand panel of this figure: our cluster sample is missing faint. blue clusters.," The effect of our $I$ -band magnitude cut is particularly clear in the right-hand panel of this figure: our cluster sample is missing faint, blue clusters."824 It is clear that the excess of red clusters in M81 compared to ATLOL is mostly at brighter cluster Iuminositües: brighter than Af;—3.4 (the expected peak of the globular cluster luminosity function). the M81 clusters are mostly red. while the M1OI clusters are mostly blue.," It is clear that the excess of red clusters in M81 compared to M101 is mostly at brighter cluster luminosities: brighter than $M_V=-7.4$ (the expected peak of the globular cluster luminosity function), the M81 clusters are mostly red, while the M101 clusters are mostly blue."825 We also confirm the detection by of a number of faint. red clusters in MIOI.," We also confirm the detection by \citet{cwl04} of a number of faint, red clusters in M101."826 Some of these could be faint. backeround ealaxies. especially ellipticals or the bulges of spirals whose disks are (oo [aint to observe.," Some of these could be faint background galaxies, especially ellipticals or the bulges of spirals whose disks are too faint to observe."827 Such contamination is unlikelv (o account for all of the faint red. clusters. whose nature is discussed further in Section 3.2..," Such contamination is unlikely to account for all of the faint red clusters, whose nature is discussed further in Section \ref{sec:spatdist}. ."828 The cluster candidate luminosity distribution is shown in Figure 7.. for the full sample and (he ved and blue subsamples.," The cluster candidate luminosity distribution is shown in Figure \ref{lum-dist}, for the full sample and the red and blue subsamples."829 The strong fall-olff at V.2»23 (Mj.=—6.1) is due to our imposed magnitude limit (see above)., The strong fall-off at $V>23$ $M_V>-6.1$ ) is due to our imposed magnitude limit (see above).830 Brighter (han this limit. we find that the blue ]usters are about. 0.25 mae brighter in the median than the red clusters.," Brighter than this limit, we find that the blue clusters are about 0.25 mag brighter in the median than the red clusters."831 From population svnthesis models. such a difference is consistent with the blue clusters being vounger. by a few Gyr if both populations are Z8 Gyr old. or less if the clusters are younger.," From population synthesis models, such a difference is consistent with the blue clusters being younger, by a few Gyr if both populations are $\gtrsim 8$ Gyr old, or less if the clusters are younger."832 Inferred ages are of course stronglv dependent on additional factors such as metallicity. recldenine. and initial mass function.," Inferred ages are of course strongly dependent on additional factors such as metallicity, reddening, and initial mass function."833" As ciseussed below. there are many more τος clusters than the expected number of globular clusters for a galaxy of MIOIs huninositw: (he smooth curve in the left. panel of the figure shows a ‘standard’ elobular cluster luminosity function (a Gaussian with mean A4,=—T.4 and standard deviation ¢= 1.3) scaled to the number of clusters with Ay«—1.4 (also see Section 3.2))."," As discussed below, there are many more red clusters than the expected number of globular clusters for a galaxy of M101's luminosity: the smooth curve in the left panel of the figure shows a `standard' globular cluster luminosity function (a Gaussian with mean $M_V=-7.4$ and standard deviation $\sigma=1.3$ ) scaled to the number of clusters with $M_V<-7.4$ (also see Section \ref{sec:spatdist}) )."834" The bright end of the LF for the full sample is consistent with the power-law distribution of luminosity d:N(L)/dLxL? CN(L)xL 4), similar to the distributions seen for voung clusters in mergers (Whitmoreetal.1999:οἱal.1997). and Z/57-lentiliel (not necessarily voung) clusters in other spirals 2002).. as well as the bright end of the GCLF for the Milkv Way ancl M31. 1994)."," The bright end of the LF for the full sample is consistent with the power-law distribution of luminosity $dN(L)/dL \propto L^{-2}$ $N(L) \propto L^{-1}$ ), similar to the distributions seen for young clusters in mergers \citep{whi99,mwsf97}835 and -identified (not necessarily young) clusters in other spirals \citep{lar02b}, as well as the bright end of the GCLF for the Milky Way and M31 \citep{hp94,mcl94}."836. Do the blue and red subsamples of M1OI correspond to voung! aud ‘lobular cluster eroups?, Do the blue and red subsamples of M101 correspond to `young' and `globular' cluster groups?837 One way to find out is to compare (he number of clusters per unit galaxy. Iuminosity in MIOLto values for other galaxies., One way to find out is to compare the number of clusters per unit galaxy luminosity in M101to values for other galaxies.838 The total magnitudes of MIOLI as given by, The total magnitudes of M101 as given by839 (Skrutskieetal., \citep{Skrutskie90}.8401990).. ~30 µια. (Brownal., $\sim$ $\mu$ \citep{Brown07}.8412007).. ~10? (Alexanderetal.2006).. (Liu&Papaloizou1979).. (e.8Jensen&Mathieu1997:Dutrev2005:White&Wil," $\sim10^5$ \citep{Alexander06}, \citep{Lin79}. \citep[e.g][]{Jensen97,Beust05,White05}."842lebrand2005).. (Espaillatetal.2007) (Careutheretal., \citep{Espaillat07} \citep{Guenther07}.8432007).. (e.g.Furlanetal.2007) in the Taurus star-forming region (l~2 MM. —115 ppc).," \cite[e.g.][]{Furlan07} in the Taurus star-forming region $\sim$ Myr, $\sim$ pc)."844 οσα απο was discovered to have a large ~20- JS0;0u excess and no excess at wavelengths <& san throughTelescope Tutrared Spectrograph observatious (Forrestetal.2001)., CoKu Tau/4 was discovered to have a large $\sim$ $\mu$ m excess and no excess at wavelengths $<8$ $\mu$ m through Infrared Spectrograph observations \citep{Forrest04}.845. The disk has been modeled as having an inner hole of radius —10 AAU (D'Alessioetal.2005).. with suggestions that this hole is due to a giant planet (Quillenctal.2001).," The disk has been modeled as having an inner hole of radius $\sim$ AU \citep{DAlessio05}, with suggestions that this hole is due to a giant planet \citep{Quillen04}."846. Tn this paper. we describe nem-iufrared aperture-nasking interferometry and niaegnmeg observations that demonstrate that Colu Tau/l1 is a nuear-equal mass ünuarv star system.," In this paper, we describe near-infrared aperture-masking interferometry and imaging observations that demonstrate that CoKu Tau/4 is a near-equal mass binary star system."847 We show that the predicted mner role size from dynamical models is comparable to. mt larecr than. that predicted from radiative transfer nodels.," We show that the predicted inner hole size from dynamical models is comparable to, but larger than, that predicted from radiative transfer models."848 Finally. we discuss whether other so-called “transitional” disks could be circiuunubiuarv disks.," Finally, we discuss whether other so-called “transitional” disks could be circumbinary disks."849 We conclude that for candidates in Taurus. much of the lass ratio-separation space for stellar companions cau by ruled out bv existiug observations. but that defuitively ruling out the possibility of binaritv is generally difficult or iudividual so-called “transitional” disks.," We conclude that for candidates in Taurus, much of the mass ratio-separation space for stellar companions can by ruled out by existing observations, but that definitively ruling out the possibility of binarity is generally difficult for individual so-called “transitional” disks."850 Colku αι was observed with the NIRC2 camera behind Adaptive Optics (AO) at the heck IE telescope ou 2007 Nov 23 as part of an ongoing aperture-masking survey of nearby vouug star-forming associations., CoKu Tau/4 was observed with the NIRC2 camera behind Adaptive Optics (AO) at the Keck II telescope on 2007 Nov 23 as part of an ongoing aperture-masking survey of nearby young star-forming associations.851 Apertureauaskiug interferometry (e.e.Tuthilletal.2000)x is a well established technique for achieving the full diffraction limit of a single telescope. receutly applied to observations belind adaptive optics svstenis (6.8.Lloydal.2006:Krauset 2008)..," Aperture-masking interferometry \citep[e.g.][]{Tuthill00} is a well established technique for achieving the full diffraction limit of a single telescope, recently applied to observations behind adaptive optics systems \citep[e.g.][]{Lloyd06,Kraus08a}."852. A 9-hole mask was placed in a filter wheel near a pupil plane in the NIRC2 camera. euablius interfercuce fringes ou 236 baseline to be simultaneously recorded ou the cameras Waging array.," A 9-hole mask was placed in a filter wheel near a pupil plane in the NIRC2 camera, enabling interference fringes on 36 baselines to be simultaneously recorded on the camera's imaging array."853 The observations of Colxu Tau/1 consisted of two iniage setstakenthroughaIC filter. each with eight 20 secondexposures.calibratedbytwointerleaved image sets of CN Tan.," The observations of CoKu Tau/4 consisted of two image setstakenthroughaK' filter, each with eight 20 secondexposures,calibratedbytwointerleaved image sets of CX Tau."854 The airmass of Col&u Tau/1 observations varied, The airmass of CoKu Tau/4 observations varied855demonstrates that the periodicity is appareut. but is nof detected with sufficicut significance to allow us to obtain a useful dprovement in the precision of the orbital period.,"demonstrates that the periodicity is apparent, but is not detected with sufficient significance to allow us to obtain a useful improvement in the precision of the orbital period."856 Seven separate peaks in the inteusitv of this source are apparent in the latter part of the ASM light curve., Seven separate peaks in the intensity of this source are apparent in the latter part of the ASM light curve.857 Their average separation is close to 111 davs and is probably not consistent with the period determined hrough pulse timing., Their average separation is close to 114 days and is probably not consistent with the period determined through pulse timing.858 The outbursts of some other Be starfneutron star binaries do not occur at precisely the orbital periods determined by pulse timing. so there is 10 reason in this case to doubt the accuracy of the pulse uius value.," The outbursts of some other Be star/neutron star binaries do not occur at precisely the orbital periods determined by pulse timing, so there is no reason in this case to doubt the accuracy of the pulse timing value."859 is à huumnous IIMXND in the Large Magellanic Cloud which highly likely comprises astcllar-nass black hole and its normal OB-type companion., is a luminous HMXB in the Large Magellanic Cloud which highly likely comprises astellar-mass black hole and its normal OB-type companion.860 The properties of the svstem are described in detail by Oroszetal.(2009)., The properties of the system are described in detail by \citet{orosz09}.861. Oroszetal.(2009). also present the ASM results and how they relate to other observations iucludiug other period deterünations., \citet{orosz09} also present the ASM results and how they relate to other observations including other period determinations.862 Oroszotal.(2009) adopt the value of 3.909175E0.00005 days as their best estimate of the period based on optical photoimoetry and spectroscopy.," \citet{orosz09} adopt the value of $3.90917863\pm 0.00005$ days as their best estimate of the period based on optical photometry and spectroscopy."864 An updated ASAD power spectrma is shown in Fieure 13) aud revised estimates. based solely ou the ASM power spectrum. of the orbital frequency and period are eiven in Table 3..," An updated ASM power spectrum is shown in Figure \ref{fig:pdslmcx1} and revised estimates, based solely on the ASM power spectrum, of the orbital frequency and period are given in Table \ref{tbl:detect}."865 We note that the ASAT detection provides the onlv reported evidence to date of the orbital period i N-aravs aud that our N-arav-based period estimate of P=3.90898+0.00021[0.00153] days is fully consistent with the opticallv-based value of Oroszetal. (2009).," We note that the ASM detection provides the only reported evidence to date of the orbital period in X-rays and that our X-ray-based period estimate of $P = 3.90898 \pm8660.00021 [\pm 0.00153]$ days is fully consistent with the optically-based value of \citet{orosz09}."867 is a trausicut X-ray pulsar in a Be/X-vay binary discovered with the BATSE experimceut ontheO, is a transient X-ray pulsar in a Be/X-ray binary discovered with the BATSE experiment on the.868bservatory2007).. Shraderct used ~2 vears of ASM data to obtain au carly estimate of the outburst period of P~135 d. Levine&Corbet(2006) reported a relatively precise determination of the outburst period. {κε=218.940.5 d. that was based on the time intervals between widely spaced outbursts.," \citet{shrad99} used $\sim2$ years of ASM data to obtain an early estimate of the outburst period of $P \sim 135$ d. \citet{lcatel06}869 reported a relatively precise determination of the outburst period, $P_{outburst} = 248.9 \pm 0.5$ d, that was based on the time intervals between widely spaced outbursts."870 We believe this is still the best current estimate of the outburst cycle time., We believe this is still the best current estimate of the outburst cycle time.871 The ASM power spectrum is shown in the top panel of Figure 1L., The ASM power spectrum is shown in the top panel of Figure \ref{fig:pdsj1008}.872 Though the peaks due to the outburst periodicity are clear. they are not sufficiently well-defined to vield a superior period estimate.," Though the peaks due to the outburst periodicity are clear, they are not sufficiently well-defined to yield a superior period estimate."873 Coeetal.(2007) have estimated the orbital period through timing of the— 93-s pulsations secu in the BATSE data., \citet{coe07} have estimated the orbital period through timing of the 93-s pulsations seen in the BATSE data.874 They obtained the result P—217.5XU. d and noted that is it is iu good agreement with the result of Levine&Corbet(2006).," They obtained the result $P =875247.8 \pm 0.4$ d and noted that is it is in good agreement with the result of \citet{lcatel06}."876. is also a transicut pulsar in a BesX-ray binary system., is also a transient pulsar in a Be/X-ray binary system.877 A brief description of the early history and of a pulse-timine analysis usine BATSE data may be found in Fingeretal.(1996)., A brief description of the early history and of a pulse-timing analysis using BATSE data may be found in \citet{fwc96}.878. The BATSE timine analysis vielded an estimate of the orbital period of P=12.12£0.03 d as well as the projected seminajor axis. eccentricity. epoch of periastron passage. aud other orbital elements.," The BATSE timing analysis yielded an estimate of the orbital period of $P =87942.12 \pm 0.03$ d as well as the projected semimajor axis, eccentricity, epoch of periastron passage, and other orbital elements."880 A slight revision to the orbital period and time of periastrou passage have been eiven by Iunii L).., A slight revision to the orbital period and time of periastron passage have been given by \citet{inam04}.881 Evidence of outbursts recurring at intervals that are close iu duration to the orbital period (or nmltiples thereof) is secu in the ASAD power spectrum (see Fie. 15))., Evidence of outbursts recurring at intervals that are close in duration to the orbital period (or multiples thereof) is seen in the ASM power spectrum (see Fig. \ref{fig:pds1417}) ).882 Frequency aud period estimates obtaimed from the spectrum shown in Fie., Frequency and period estimates obtained from the spectrum shown in Fig.883 15 are listed iu Table but are not as precise as the pulse-timing values., \ref{fig:pds1417} are listed in Table \ref{tbl:detect} but are not as precise as the pulse-timing values.884 is au X-ray pulsar with the rather long pulse period of ~1300 s (Lutovinovetal. 2005).., is an X-ray pulsar with the rather long pulse period of $\sim 1300$ s \citep{lutov05}. .885 (e.gDuquennoy&Aavor&Marcy1992) (c.getal.2002).. (Goodwin," \citep[e.g][]{Duquennoy1991,Fischer1992} \citep[e.g][]{Patience2002}, \citep{Goodwin2009}."886" AQ Ado. Af,>Af. AA, Mo Ξἀλλ. 4 ou the primary mass", $M_{1}$ $M_{2}$ $M_{1}>M_{2}$ $M_{1}$ $M_{2}$ $q=M_{2}/M_{1}$ $q$ on the primary mass.887 —CTraditionallv these classes of nodels have been divided iuto capture aud fragmnenutatiou scenarios., Traditionally these classes of models have been divided into capture and fragmentation scenarios.888 Capture refers to the tidal capture of two nabound objects on a timescale that is long compares o the collapse time of cach component (c.g.MeDon-ald&Clarke 1993)., Capture refers to the tidal capture of two unbound objects on a timescale that is long compared to the collapse time of each component \citep[e.g.][]{McDonald1993}.889. For cach primary star the mass of he secondary is chosen randomly from the single star nass function aud the secondary-lmass distribution woul reflect the IMIF., For each primary star the mass of the secondary is chosen randomly from the single star mass function and the secondary-mass distribution would reflect the IMF.890 While tidal capture appears to be too ineticicut in reproducing high binary fractious. it has con noticed that. particularly in small eroups of stars. stir-disk encounters may form binaries (MeDonald&Clarke 1995).," While tidal capture appears to be too inefficient in reproducing high binary fractions, it has been noticed that, particularly in small groups of stars, star-disk encounters may form binaries \citep{McDonald1995}."891. Tn iu case even this disk assisted capture. whereby a star passing through the disk of another which dissipates enough kinetic cucrey to form a bound svsten. is unlikely to be the most relevant binary formation nechamsin (Doffinctal.1998).," In any case even this disk assisted capture, whereby a star passing through the disk of another which dissipates enough kinetic energy to form a bound system, is unlikely to be the most relevant binary formation mechanism \citep{Boffin1998}."892. Fragiucntation scenarios are the preferred mechanism or the formation of multiple systems., Fragmentation scenarios are the preferred mechanism for the formation of multiple systems.893 The so-called rasnientation models are usually classified as pronuipt ragiueutation (60.8.Boss1986:Bounell&Basticu1992) and disk fraeioeutation (e.g.Bonnell1991:Stamatellos&Wlitworth 2009).," The so-called fragmentation models are usually classified as prompt fragmentation \citep[e.g.][]{Boss1986,Bonnell1992} and disk fragmentation \citep[e.g.][]{Bonnell1994,Stamatellos2009}."894. In the prowpt fragmentation scenario voth primary aud secondarystarsform byfragmentationofthesame collapsing molecular cloud core., In the prompt fragmentation scenario both primary and secondarystarsform byfragmentationofthesame collapsing molecular cloud core.895 Disk ragiueutation takes place in a uewly formed star-disk, Disk fragmentation takes place in a newly formed star-disk896times. the high-mass mass spectrum converges to a slope between = Bands=2.5 consistent with the accretion during the stellar dominated phase. and to the Salpeter slope of=235.,"times, the high-mass mass spectrum converges to a slope between $\gamma=-2$ and $\gamma=-2.5$ consistent with the accretion during the stellar dominated phase, and to the Salpeter slope of $\gamma=-2.35$."897 1n order ascertain whether we are correct in our interpretation of a two-power law LAL resulting from accretion in. gas-clominatecl ancl stellardominated: regimes. we evaluated. how much. of the eventual stellar. mass. was added: in each regime.," In order ascertain whether we are correct in our interpretation of a two-power law IMF resulting from accretion in gas-dominated and stellar-dominated regimes, we evaluated how much of the eventual stellar mass was added in each regime."898 Figure 4 plots the amount. of mass accumulated by cach star during the stellar-dominate yhase against the final stellar mass.," Figure \ref{accdyn}899 plots the amount of mass accumulated by each star during the stellar-dominated phase against the final stellar mass."900 Llieh-mass stars accumulate the majority of their mass during this phase whereas most of the low-mass stars only acerete a smal raction of their mass during this phase., High-mass stars accumulate the majority of their mass during this phase whereas most of the low-mass stars only accrete a small fraction of their mass during this phase.901 The break between he two regimes occurs at approximately the same fina mass (2 LM.) where the mass function displavs a break tween the two slopes., The break between the two regimes occurs at approximately the same final mass $\approx 1 \solm$ ) where the mass function displays a break between the two slopes.902 This supports the assertion that the wo different power-laws in the LAL derive. from. dilferen ohvsical regimes which allect how the stars accrete., This supports the assertion that the two different power-laws in the IMF derive from different physical regimes which affect how the stars accrete.903 The ow-niass stars derive thei mass from tical-lobe accretion during the gas-domünated regime whereas the high-mass stars derive the majority of their mass from a Boncdi-LHovle ype accretion that occurs in the inner parts of the cluster where the potential is dominated by the stars themselves., The low-mass stars derive their mass from tidal-lobe accretion during the gas-dominated regime whereas the high-mass stars derive the majority of their mass from a Bondi-Hoyle type accretion that occurs in the inner parts of the cluster where the potential is dominated by the stars themselves.904 In addition to producing the two power-law mass spectrum. competitive accretion naturally results in a certain degree of mass segregation.," In addition to producing the two power-law mass spectrum, competitive accretion naturally results in a certain degree of mass segregation."905 This arises due to the accretion in the gas-dominated phase where there is a strong correlation between aceretion rate. and thus the final mass. and position in the cluster. (sce equations(16)) ancl (20)).," This arises due to the accretion in the gas-dominated phase where there is a strong correlation between accretion rate, and thus the final mass, and position in the cluster (see \ref{tidmsvsrad}) ) and \ref{massvsrad}) )."906 This. direct correlation between the final mass and position in the eluster neelects variations in the initial masses and the relative movements of the stars due to their interactions., This direct correlation between the final mass and position in the cluster neglects variations in the initial masses and the relative movements of the stars due to their interactions.907" If the cluster is mass segregated entering the stellar. dominated ohase. then the Ag,xAZ? implies that the mass segregation =μαill persist."," If the cluster is mass segregated entering the stellar dominated phase, then the $\macc \propto \ms^2$ implies that the mass segregation will persist."908 Simulations of accretion in clusters show that 10 Mass segregation does result. but that there is not a one to one correlation between mass and radius (Bonnell 2000)., Simulations of accretion in clusters show that the mass segregation does result but that there is not a one to one correlation between mass and radius (Bonnell 2000).909 In fact. low-mass stars are located throughout 1e cluster. including in the core. but the high-mass stars are edominantly located in the central regions as is found in voung stellar clusters such as the ONC (Llillenbranc 1997).," In fact, low-mass stars are located throughout the cluster, including in the core, but the high-mass stars are predominantly located in the central regions as is found in young stellar clusters such as the ONC (Hillenbrand 1997)."910 lt ds also worth noting that although the moclels oesented. here are meant to. consider accretion onto voung stars. they are equally appropriate for the erowth of clumps in a molecular cloud.," It is also worth noting that although the models presented here are meant to consider accretion onto young stars, they are equally appropriate for the growth of clumps in a molecular cloud."911 As the clumps evolve owards eravitational instability. they will accrete from the surrounding σας and this accretion will be governed by the ohvsies described here.," As the clumps evolve towards gravitational instability, they will accrete from the surrounding gas and this accretion will be governed by the physics described here."912 Thus. for example. the clump mass-tunction found by Motte. André Neri (1998) for the p Oph molecular cloud. could be due to the accretion by the pre-stellar clumps as the whole system collapses down to form a cluster.," Thus, for example, the clump mass-function found by Motte, André Neri (1998) for the $\rho$ Oph molecular cloud could be due to the accretion by the pre-stellar clumps as the whole system collapses down to form a cluster."913 The observed y=3/2 slope would imply that the whole svstem is in a px27 density configuration and is subvirial (dominated by the diffuse gas not in the elumps)., The observed $\gamma=-3/2$ slope would imply that the whole system is in a $\rho \propto R^{-2}$ density configuration and is subvirial (dominated by the diffuse gas not in the clumps).914 The steeper slope found by. Motte (1998) at the hieh-mass end of the mass spectrum can be interpreted as arising from a regionὃν of near-uniform oσας density., The steeper slope found by Motte (1998) at the high-mass end of the mass spectrum can be interpreted as arising from a region of near-uniform gas density.915 A test of such a possibility is to estimate the degree of mass segregation of the clumps in this pre-stellar cluster system., A test of such a possibility is to estimate the degree of mass segregation of the clumps in this pre-stellar cluster system.916" Finally. it is possible that the mass spectrum for massive stars. Al,<LOAL.. is significantly different than that for lower mass stars if they do not form in a similar fashion."," Finally, it is possible that the mass spectrum for massive stars, $\ms917\simgreat 10 \solm$, is significantly different than that for lower mass stars if they do not form in a similar fashion."918 Lf massive stars cannot acercte above 1041. due to the cllect of radiation pressure on the infalling dust (Yorke Ixrüsgecl 1977: Yorke 1993) but form through a merger process in a dense core (Bonnell 1998) then the expected mass spectrum could be significantly dilferent from that presented here., If massive stars cannot accrete above $10 \solm$ due to the effect of radiation pressure on the infalling dust (Yorke Krüggel 1977; Yorke 1993) but form through a merger process in a dense core (Bonnell 1998) then the expected mass spectrum could be significantly different from that presented here.919 Competitive accretion in voung stellar clusters results in a two power-law mass spectrum with a slope of 5=3Mle for low-mass stars and a steeper slope of 252.5 for high-mass stars., Competitive accretion in young stellar clusters results in a two power-law mass spectrum with a slope of $\gamma\approx - 3/2$ for low-mass stars and a steeper slope of $-2 \ge \gamma \simgreat -2.5$ for high-mass stars.920 The dillerent slopes are due to whether the eas or the stars dominate the cluster. potential., The different slopes are due to whether the gas or the stars dominate the cluster potential.921 When eas dominates the cluster potential. the accretion rates are given by a tidal-DIobe accretion radius and results in a," When gas dominates the cluster potential, the accretion rates are given by a tidal-lobe accretion radius and results in a"922The detailed variability study on the AGB. stars in NGC 1846 undertaken in Lebzelter Wood (2007)) allows us to investigate the observed star-to-star differences in the C/O ratio in light of the pulsational properties of these stars.,The detailed variability study on the AGB stars in NGC 1846 undertaken in Lebzelter Wood \cite{LW07}) ) allows us to investigate the observed star-to-star differences in the C/O ratio in light of the pulsational properties of these stars.923 In refpl we show the location of our sample stars in the P-L-diagram of the cluster., In \\ref{pl} we show the location of our sample stars in the P-L-diagram of the cluster.924 Data are taken from Lebzelter Wood., Data are taken from Lebzelter Wood.925 LE9 could not be plotted due to the lack of any detected periodicity in its small-amplitude light. variability., LE9 could not be plotted due to the lack of any detected periodicity in its small-amplitude light variability.926 The other stars can be easily attributed to a pulsation mode: LEI3 and LE16 are first overtone pulsators. while LES and H39 pulsate in the second overtone.," The other stars can be easily attributed to a pulsation mode: LE13 and LE16 are first overtone pulsators, while LE8 and H39 pulsate in the second overtone."927 Long secondary periods beyond the fundamental mode period have been found in the light curves of several of these stars. but no shorter secondary periods.," Long secondary periods beyond the fundamental mode period have been found in the light curves of several of these stars, but no shorter secondary periods."928 As already pointed out in Lebzelter Wood (2007)). LE17 cannot be attributed unequivocally to any pulsation sequence.," As already pointed out in Lebzelter Wood \cite{LW07}) ), LE17 cannot be attributed unequivocally to any pulsation sequence."929 Stars located on the first overtone sequence show a higher C/O ratio than objects on the second overtone sequence refpl)., Stars located on the first overtone sequence show a higher C/O ratio than objects on the second overtone sequence \\ref{pl}) ).930 As can be seen this cannot be attributed to a simple difference in luminosity as the two groups clearly overlap., As can be seen this cannot be attributed to a simple difference in luminosity as the two groups clearly overlap.931 This is further illustrated in refemd2 where we marked the sample stars in the color- diagram according to their pulsation mode., This is further illustrated in \\ref{cmd2} where we marked the sample stars in the color-magnitude diagram according to their pulsation mode.932 For LE9 (C/O=0.2). no periodic change was found by Lebzelter Wood (2007)).," For LE9 (C/O=0.2), no periodic change was found by Lebzelter Wood \cite{LW07}) )."933 Assuming that it would have been nore likely to find a first overtone pulsation than a shorter period higher overtone pulsation. we attributed this star to the group of higher overtone pulsators.," Assuming that it would have been more likely to find a first overtone pulsation than a shorter period higher overtone pulsation, we attributed this star to the group of higher overtone pulsators."934 These findings suggest that in addition to a relation between the C/O ratio and the luminosity. seems to exist some dependency of this ratio on the pulsation mode.," These findings suggest that in addition to a relation between the C/O ratio and the luminosity, seems to exist some dependency of this ratio on the pulsation mode."935 At this point it ts not clear how to explain the existence of such a relation., At this point it is not clear how to explain the existence of such a relation.936 As mentioned above. all sample stars are on the AGB. we can exclude an RGB contamination.," As mentioned above, all sample stars are on the AGB, we can exclude an RGB contamination."937 From the difference in the C/O ratio. we can assume that the sample stars are in at least three different thermal pulse cycles.," From the difference in the C/O ratio, we can assume that the sample stars are in at least three different thermal pulse cycles."938 Both from our models and from the literature (e.g. Vassiliadis Wood 1993)). we would expect that the stars become cooler from cycle to cycle.," Both from our models and from the literature (e.g. Vassiliadis Wood \cite{VW93}) ), we would expect that the stars become cooler from cycle to cycle."939 As the two C enhanced stars LEI6 and LEI3 would form a sequence at higher temperatures than the other stars. it can be excluded that a difference in the thermal pulse cyele is responsible for the observed differences in pulsation mode and temperature.," As the two C enhanced stars LE16 and LE13 would form a sequence at higher temperatures than the other stars, it can be excluded that a difference in the thermal pulse cycle is responsible for the observed differences in pulsation mode and temperature."940 A difference in pulsation mode between two stars of the same luminosity requires a difference in some other global characteristic of the star., A difference in pulsation mode between two stars of the same luminosity requires a difference in some other global characteristic of the star.941 This could be a difference in mass., This could be a difference in mass.942 Mackey Broby Nielsen (2007)) report indications for two stellar populations in NGC 1846 separated in age by about 300 Myr., Mackey Broby Nielsen \cite{MB07}) ) report indications for two stellar populations in NGC 1846 separated in age by about 300 Myr.943 At a given time we may then find AGB stars of different masses at the same time on the AGB., At a given time we may then find AGB stars of different masses at the same time on the AGB.944 Indeed such an age or mass difference could easily account for the observed color difference refemd2))., Indeed such an age or mass difference could easily account for the observed color difference \\ref{cmd2}) ).945 This would also be compatible with the observed C/O ratios as the more massive stars may show a more efficient dredge up., This would also be compatible with the observed C/O ratios as the more massive stars may show a more efficient dredge up.946 Furthermore. the starting abundances of the two populations may be different. which could enhance the observed difference in C/O and  C/UC. It has to be stressed that. even in the case of a single stellar population. a scatter in the abundances of C and O on the main sequence and after the first dredge up may be present.," Furthermore, the starting abundances of the two populations may be different, which could enhance the observed difference in C/O and $^{12}$ $^{13}$ C. It has to be stressed that, even in the case of a single stellar population, a scatter in the abundances of C and O on the main sequence and after the first dredge up may be present."947 Such an abundance scatter has been found in various clusters., Such an abundance scatter has been found in various clusters.948 An interesting case is LEI3., An interesting case is LE13.949 This star has been classified as spectral type 83/3 by Lloyd Evans (1983))., This star has been classified as spectral type S3/3 by Lloyd Evans \cite{LE83}) ).950 Indeed the ZrO band head at is clearly visible in the spectrum presented by him (see his 1H)., Indeed the ZrO band head at is clearly visible in the spectrum presented by him (see his 1).951 According to its spectral classification. TiO-bands and ZrO bands should be of similar strength (Keenan Boeshaar 1980)).," According to its spectral classification, TiO-bands and ZrO bands should be of similar strength (Keenan Boeshaar \cite{KB80}) )."952 YO should occur. but is not covered by the spectra shown by Lloyd Evans.," YO should occur, but is not covered by the spectra shown by Lloyd Evans."953 As expected, As expected954"factor Wer is the fraction of the eveuts produced above a given threshold £i, which are able to reach a cistance equal to the distance from the source to us. d. with an cucrey still above that sae threshold.","factor $W_{GZK}$ is the fraction of the events produced above a given threshold $E_{th}$ which are able to reach a distance equal to the distance from the source to us, $d$, with an energy still above that same threshold."955" Assuming a power law spectrum at the sources dN ανE* this factor is just where ζειο) is the iitial cucrey that a CR misthave in order to survive with £=Ey, after traveling a distance d."," Assuming a power law spectrum at the sources $N$ $E\propto956E^{-s}$ this factor is just where $E_i(E_{th},d)$ is the initial energy that a CR musthave in order to survive with $E=E_{th}$ after traveling a distance $d$."957" The resulting factors for Ly,=60 aud SO EeV. (keep in imiud that there are still significant systematic macertaintics in the energy recoustructed in CR experiments) are shown in fie. ὃν,"," The resulting factors for $E_{th}=60$ and 80 EeV (keep in mind that there are still significant systematic uncertainties in the energy reconstructed in CR experiments) are shown in fig. \ref{fracp.fig},"958 computed following reference |0].., computed following reference \cite{horizon}.959 We asstmed a source spectral index of 2.2. but the results are not much scnusitive to the particular value adopted. aud considered a proton composition.," We assumed a source spectral index of 2.2, but the results are not much sensitive to the particular value adopted, and considered a proton composition."960 The suppressious are qualitatively simular for Fe uuclei but are stronger for intermediate mass nuclei:, The suppressions are qualitatively similar for Fe nuclei but are stronger for intermediate mass nuclei.961 Iucludius these factors iu the previous comparisons changes the results with respect to those obtained using the simple cutoff at LOO Mpe. but not iu a drastic way. although with larger unuuber of eveuts the proper iuclusion of the GZIN factors could become more nuportaut.," Including these factors in the previous comparisons changes the results with respect to those obtained using the simple cutoff at 100 Mpc, but not in a drastic way, although with larger number of events the proper inclusion of the GZK factors could become more important."962 Another aspect that iav affect the source models is the fact that the catalogs considered here ouly contain sources brighter than a certain lint. aud although they may be quite complete above those liuiting briehtuesses. the fraction of sources which are faint increases significantly with increasing distance. and those wuebservecl sources may actually contribute to the CR fluxes. »ossdblv eiviug a more diffuse background.," Another aspect that may affect the source models is the fact that the catalogs considered here only contain sources brighter than a certain limit, and although they may be quite complete above those limiting brightnesses, the fraction of sources which are faint increases significantly with increasing distance, and those unobserved sources may actually contribute to the CR fluxes, possibly giving a more diffuse background."963 In models where the sources are weighted by their fluxes he expected contribution from the unobserved sources is however reduced., In models where the sources are weighted by their fluxes the expected contribution from the unobserved sources is however reduced.964 In the past some works 11.0] using the IRAS galaxy catalog. for which the selection effects as a function of distance are shown. corrected for the incompleteness of the catalog by dividiug the observed deusity of ealaxics at a eiven distance by the corresponding selection. fiction in order to obtain a more complete represcutation of the ealaxy distribution.," In the past some works \cite{wax,cuoco}965 using the IRAS galaxy catalog, for which the selection effects as a function of distance are known, corrected for the incompleteness of the catalog by dividing the observed density of galaxies at a given distance by the corresponding selection function in order to obtain a more complete representation of the galaxy distribution."966 A possible drawback of this approach is that one assigus he unobserved galaxies to the same locations where bright ealaxies are observed. aud this way not ο Very precise when the galaxies are sparscly siuupled. as happens at large distances.," A possible drawback of this approach is that one assigns the unobserved galaxies to the same locations where bright galaxies are observed, and this may not be very precise when the galaxies are sparsely sampled, as happens at large distances."967 For sinaller catalogs. such as the SWIFT one. even if a selection function were kuown it would not be very realistic to asstune that the unobserved faint AGN are in the same locations as the bright ones.," For smaller catalogs, such as the SWIFT one, even if a selection function were known it would not be very realistic to assume that the unobserved faint AGN are in the same locations as the bright ones."968 Auvway. if one restricts the sources to the ucarby ones (ising a distance cutoff or weighting them by the GZIs attenuation factors). iu the models where the sources are weighted bv their fluxes the contribution from the faint unobserved sources is not large. so that the model expectations obtained should still be reasonably accurate.," Anyway, if one restricts the sources to the nearby ones (using a distance cutoff or weighting them by the GZK attenuation factors), in the models where the sources are weighted by their fluxes the contribution from the faint unobserved sources is not large, so that the model expectations obtained should still be reasonably accurate."969 For instance. if onemodels the Ipuuinositv distribution of the sources according to a Schechter fuuctiou.. «N/AxLSexptEf E). where a typical value," For instance, if onemodels the luminosity distribution of the sources according to a Schechter , $N/$ $L\propto L^a\exp(-L/L_*)$ , where a typical value"970"where ff.) is the Dubble parameter at redshift τν e is fF(0.0) is the as a thefunction speedof ofpositionlight, ou the sky. aud O is samplingthe solid fractionangle.","where $H(z)$ is the Hubble parameter at redshift $z$, $c$ is the speed of light, $f(\theta,\phi)$ is the sampling fraction as a function of position on the sky, and $\Omega$ is the solid angle."971 We take the function. f(0.0) to be a coustant over the eutire 8032 square deerce areal coverage of the SDSS DR? Leeacy survey.," We take the function $f(\theta, \phi)$ to be a constant over the entire 8032 square degree areal coverage of the SDSS DR7 Legacy survey."972 The augular part of Equation 7 thus becomes (1/17)(8032/112253)(13) = 0.19£7., The angular part of Equation \ref{vmax-integral} thus becomes $(1/4\pi) (8032/41253) (4\pi)$ = 0.1947.973" V, values are given in the data table presented in Section ??..", $V_{max}$ values are given in the data table presented in Section \ref{bigtables}.974 We based the quality assessment of our measured structural paraicters on three science-notivated metres., We based the quality assessment of our measured structural parameters on three science-motivated metrics.975 Figure 7? shows the size-Iununositv of disks iu the g-baud., Figure \ref{disk-lumsize} shows the size-luminosity of disks in the $g$ -band.976 Disks are expected to follow a well-defined, Disks are expected to follow a well-defined977dvnamically in the merger remnant.,dynamically in the merger remnant.978 Dynamical friction might be the cause of the slower velocity. of rotation of the GC's of GI compared to the gaseous component., Dynamical friction might be the cause of the slower velocity of rotation of the GCs of G1 compared to the gaseous component.979 The peak at |Fe/II] — - 1 in all three groups may be an indication of a significant contribution of 10? to LOMM. haloes during the hierarchical galaxy formation. assuming that these haloes produced GCs with roughly the same metallicity ancl (hat thev follow the mass-metallicitv relation al z  1 (Lamareille et al.," The peak at [Fe/H] = - 1 in all three groups may be an indication of a significant contribution of $10^{9}$ to $10^{10} M_{\odot}$ haloes during the hierarchical galaxy formation, assuming that these haloes produced GCs with roughly the same metallicity and that they follow the mass-metallicity relation at z $\simeq$ 1 (Lamareille et al."980 2009)., 2009).981 For comparison. we mention (hat the metallicity distribution of M87. another giant. elliptical galaxy with an active nucleus. also peaks around - 1 (Cohen et al.1993). but. unlike NGC 5128. does not have a very. low metallicity tail.," For comparison, we mention that the metallicity distribution of M87, another giant elliptical galaxy with an active nucleus, also peaks around - 1 (Cohen et al.1998), but, unlike NGC 5128, does not have a very low metallicity tail."982 There are in [act other analogies between the metallicity distributions of the three groups and those of GC's in nearby galaxies: for example between the GCs of AIL 31 (Darmby et al., There are in fact other analogies between the metallicity distributions of the three groups and those of GCs in nearby galaxies: for example between the GCs of M 31 (Barmby et al.983 2000) and GI. or the GCs of the LAICG (Beasley et al.," 2000) and G1, or the GCs of the LMC (Beasley et al."984 2002) and the metal-rich subgroup of G3., 2002) and the metal-rich subgroup of G3.985 However. (hese comparisons can only give order of magnitude estimates. since the LMC may have been metal-enrichecl by interactions with our Galaxy. and M 31 might have a very different evolutionary. history [rom the progenitor galaxy of Gl. which merged at least 200 Myrs ago with NGC 5123.," However, these comparisons can only give order of magnitude estimates, since the LMC may have been metal-enriched by interactions with our Galaxy, and M 31 might have a very different evolutionary history from the progenitor galaxy of G1, which merged at least 200 Myrs ago with NGC 5128."986 In summary. the above scenario rests mainlv on the assuniption of accretion events which shaped the evolution of NGC 5128 and its GCs.," In summary, the above scenario rests mainly on the assumption of accretion events which shaped the evolution of NGC 5128 and its GCs."987 The latter have several possible origins: the GCs of G2 were produced in a major merger. while the GC's of the (wo other eroups were pre-existing in smaller galaxies Chat were subsequently accreted and disrupted.," The latter have several possible origins: the GCs of G2 were produced in a major merger, while the GCs of the two other groups were pre-existing in smaller galaxies that were subsequently accreted and disrupted."988 This favors the categories ii) and v) listed in the introduction lor the lormation of the galaxy itself. namely a major merger and several accretions and in-situ merging.," This favors the categories ii) and v) listed in the introduction for the formation of the galaxy itself, namely a major merger and several accretions and in-situ merging."989 The proposed scenario remains highlv speculative. in the absence of spectroscopically determined ages and metallicities for most GCs in the galaxy.," The proposed scenario remains highly speculative, in the absence of spectroscopically determined ages and metallicities for most GCs in the galaxy."990"a modified Planck function, which has some given dust temperature and opacity.","a modified Planck function, which has some given dust temperature and opacity."991" We apply an opacity law with the standard form &οςA~°, where a value of 8=0 would correspond to a perfect blackbody."," We apply an opacity law with the standard form $\kappa \propto \lambda^{-\beta}$, where a value of $\beta=0$ would correspond to a perfect blackbody."992" In our model we set B—1.5 (Dunne Eales, 2001)."," In our model we set $\beta=1.5$ (Dunne Eales, 2001)."993" Emission from polycyclic aromatic hydrocarbons (PAHs) is an important spectral feature between 3 and 13 microns (Leger Puget, 1984; Allamandola et al.,"," Emission from polycyclic aromatic hydrocarbons (PAHs) is an important spectral feature between 3 and 13 microns (Leger Puget, 1984; Allamandola et al.,"994 1985)., 1985).995" In our spiral galaxy model we include a fairly schematic representation of the PAH features, designed to broadly resemble the spectra of galaxies observed in the Spitzer Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al.,"," In our spiral galaxy model we include a fairly schematic representation of the PAH features, designed to broadly resemble the spectra of galaxies observed in the Spitzer Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al.,"996" 2003), presented by Draine et al. ("," 2003), presented by Draine et al. ("9972007).,2007).998 We show the IR spectra of our early- and late-types in reffig:modelseds.., We show the IR spectra of our early- and late-types in \\ref{fig:modelseds}.999" While the modelling of the PAH features in the late-types is somewhat rudimentary, we find it is sufficient to reproduce the observed colours and sources counts of galaxies which sample this part of the spectrum, so we consider it sufficient for the purposes of this paper."," While the modelling of the PAH features in the late-types is somewhat rudimentary, we find it is sufficient to reproduce the observed colours and sources counts of galaxies which sample this part of the spectrum, so we consider it sufficient for the purposes of this paper."1000" reffig:modelseds also includes a third SED, which represents a starburst component modelled on Markarian 33 (Mrk33)."," \\ref{fig:modelseds} also includes a third SED, which represents a starburst component modelled on Markarian 33 (Mrk33)."1001" This is an additional spectral type which we introduce when dealing with observations at A>24um ((see refssec:24,, where the motivation for including this component is discussed)."," This is an additional spectral type which we introduce when dealing with observations at $\lambda \ge 24$ (see \\ref{ssec:24}, where the motivation for including this component is discussed)."1002" We assume a dust temperature of 30K, which is found to be a typical temperature of interstellar dust in star-forming galaxies (Farrah et al.,"," We assume a dust temperature of 30K, which is found to be a typical temperature of interstellar dust in star-forming galaxies (Farrah et al.,"1003" 2003; Pope et al.,"," 2003; Pope et al.,"1004" 2006; Coppin et al.,"," 2006; Coppin et al.,"1005" 2008; Elbaz et al.,"," 2008; Elbaz et al.,"1006 2010)., 2010).1007" We take Ap=0.3 magnitudes for the normalisation of the 1/A dust absorption law; this is the value determined by Metcalfe et ((2001) when using this model at optical wavelengths, and is a fairly conservative amount of extinction."," We take $A_B=0.3$ magnitudes for the normalisation of the $1/\lambda$ dust absorption law; this is the value determined by Metcalfe et (2001) when using this model at optical wavelengths, and is a fairly conservative amount of extinction."1008" The total integrated flux of the dust emission, in the form of both the PAH features and the blackbody, is normalised to be equal to the total absorbed flux."," The total integrated flux of the dust emission, in the form of both the PAH features and the blackbody, is normalised to be equal to the total absorbed flux."1009" For each galaxy type, we use a luminosity function (LF) defined locally in the B band and correct it into the appropriate IR band using a rest-frame z=0 colour."," For each galaxy type, we use a luminosity function (LF) defined locally in the B band and correct it into the appropriate IR band using a rest-frame $z=0$ colour."1010 The resulting LF parameters for the K band are given in Table 1.., The resulting LF parameters for the $K$ band are given in Table \ref{t-lf}.1011" Then using k+e corrections determined by the BC03 evolution code, the LF is evolved back from the present day."," Then using $k+e$ corrections determined by the BC03 evolution code, the LF is evolved back from the present day."1012" The model is a simplification of the model of Metcalfe et ((2001, 2006) since rather than 5 independent colours, one for each type, we use only two, one for all of the early- (E/SO and Sab) and one for all of the late-types (Sbc, Scd and Sdm)."," The model is a simplification of the model of Metcalfe et (2001, 2006) since rather than 5 independent colours, one for each type, we use only two, one for all of the early-types (E/S0 and Sab) and one for all of the late-types (Sbc, Scd and Sdm)."1013 This is justified on the basis that the band-band colour differences in the NIR are smaller than in the optical., This is justified on the basis that the band-band colour differences in the NIR are smaller than in the optical.1014 The z=0 colours are taken from the BC03 model predictions., The $z=0$ colours are taken from the BC03 model predictions.1015 Table 2 shows the values of M* , Table \ref{t-cols} shows the values of $M^*$ 1016P. the epoch of transit centre Το. the RV. semi-amplituc[4 K. ecosw esinc (e being the eccentricity and w the angle of the periastron). and VsinZcosp Vsindsin. with VsinI being the projection of the stellar equatorial rotation. and 6 the projection of the angle between the stellar spin axis and the planetary orbit axis.,"$P$, the epoch of transit centre $T_0$, the RV semi-amplitude $K$, $e\,cos\,\omega$ $e\,sin\,\omega$ $e$ being the eccentricity and $\omega$ the angle of the periastron), and $V\,sin\,I\,cos \,\beta$ $V\,sin\,I\,sin\,\beta$, with $V\,sin\,I$ being the projection of the stellar equatorial rotation, and $\beta$ the projection of the angle between the stellar spin axis and the planetary orbit axis."1017" Di addition. we employed free normalization factors. for each lightcurve (WASP and Euler) and each set of radial velocity (y,, for and y, for ». which enablec variations to be made in instrumental zero points."," In addition, we employed free normalization factors for each lightcurve (WASP and Euler) and each set of radial velocity $\gamma_\textit{\tiny H}$ for and $\gamma_\textit{\tiny C}$ for ), which enabled variations to be made in instrumental zero points."1018 From these parameters. parameters were derived to characterise the planetary system.," From these parameters, parameters were derived to characterise the planetary system."1019 The best-fit set of parameters that minimize the y? (reduced y is 0.86) are listed in Table 2. as well as their related computed physical parameters., The best-fit set of parameters that minimize the $\chi^2_r$ (reduced $\chi^2$ is 0.86) are listed in Table \ref{tab:params} as well as their related computed physical parameters.1020 With this best-fit solution one computes for the data y=204 with 48 measurements. and for HARPS data y=188 with 82 measurements which implies that additional jittering is present that is not accounted for by the fitted model.," With this best-fit solution one computes for the data $\chi^2=204$ with 48 measurements, and for HARPS data $\chi^2=188$ with 82 measurements which implies that additional jittering is present that is not accounted for by the fitted model."1021 Since the main deviation is related to the data. the uncertainties in the orbital solutions are most likely underestimated.," Since the main deviation is related to the data, the uncertainties in the orbital solutions are most likely underestimated."1022 However. the error bars in the Rossiter parameters are driven mostly by the HARPS on-transit data. and one can assume that they are almost correct.," However, the error bars in the Rossiter parameters are driven mostly by the HARPS on-transit data, and one can assume that they are almost correct."1023 Our best-fit solution corresponds to a giant planet with an eccentric (e= 0.3)) 8.16-day orbit and an additional long-term radial-velocity drift of!.., Our best-fit solution corresponds to a giant planet with an eccentric $e=0.3$ ) 8.16-day orbit and an additional long-term radial-velocity drift of.1024" The planet is dense with 2.25M; and a radius of 1.04 in contrast to the substantial fraction of ""inflated"" hot Jupiters."," The planet is dense with $2.25\,M_j$ and a radius of $1.04\,R_j$, in contrast to the substantial fraction of `inflated"" hot Jupiters."1025"Rj. Surprisingly. the projected angle between the orbital and stellar spin axes is found to be 6=123.3"". indicative of à retrograde orbit."," Surprisingly, the projected angle between the orbital and stellar spin axes is found to be $\beta = 123.3^\circ$, indicative of a retrograde orbit."1026" We note that Vsind21.59kms""! iis in accordance with the line rotation broadening (in Table 1)) derived by the spectral analysts."," We note that $V\,sin\,I=1.59$ is in accordance with the line rotation broadening (in Table \ref{wasp8-params}) ) derived by the spectral analysis."1027 We checked whether the partial defocusing of HARPS during the transit spectroscopic sequence had any effect on our result., We checked whether the partial defocusing of HARPS during the transit spectroscopic sequence had any effect on our result.1028 We divided the series into two subsets and considered for each of them an independent offset (y)., We divided the series into two subsets and considered for each of them an independent offset $\gamma$ ).1029 We obtain a solution with a marginal improvement in the y., We obtain a solution with a marginal improvement in the $\chi^2$.1030 By comparing the solution obtained from these two sets with that for the complete set. the angle6 was changed by 1.5c.," By comparing the solution obtained from these two sets with that for the complete set, the angle $\beta$ was changed by $1.5\,\sigma$."1031 The defocusing problem does not affect the results of this paper., The defocusing problem does not affect the results of this paper.1032 The detection of a hot Jupiter on an eccentric orbit that is misaligned with the. stellar rotation axis and moving m a retrograde directior raises many questions about the origins of this system., The detection of a hot Jupiter on an eccentric orbit that is misaligned with the stellar rotation axis and moving in a retrograde direction raises many questions about the origins of this system.1033 Although the answer is beyond the scope of this paper. the visual faint companion and the drifting y velocity of the system are key components of the puzzle.," Although the answer is beyond the scope of this paper, the visual faint companion and the drifting $\gamma$ velocity of the system are key components of the puzzle."1034 From the observed separation between the A and B components. one can derive a most likely orbital semi-major axis (a=1.35x600 AAU) (Duquennoy&Mayor1991)..," From the observed separation between the A and B components, one can derive a most likely orbital semi-major axis $a=1.35\rho\approx600$ AU) \citep{1991A&A...248..485D}."1035 The observed radial-velocity drift is therefore unlikely to be related to the B component of the binary (Y«GMa*<Ems yr). suggesting that these is an additional closer companion of both unknown mass and period.," The observed radial-velocity drift is therefore unlikely to be related to the B component of the binary $\dot{\gamma}<GMa^{-2}<1$ ), suggesting that these is an additional closer companion of both unknown mass and period."1036 The lack of curvature indicates that the companion is more massive than the transiting planet., The lack of curvature indicates that the companion is more massive than the transiting planet.1037 This intermediate body is very likely to play a significant dynamical role in the system., This intermediate body is very likely to play a significant dynamical role in the system.1038"lightcone, although typically at different redshifts and so with different properties and at offset positions (due to large-scale motions).","lightcone, although typically at different redshifts and so with different properties and at offset positions (due to large-scale motions)."1039" ? suggested applying a series of transformations (rotations, translations and inversions) when tiling space with periodic replications."," \citet{Blaizot2005} suggested applying a series of transformations (rotations, translations and inversions) when tiling space with periodic replications."1040" This does not, of course, prevent multiple appearances of a given object within the lightcone, but these duplicates are then viewed from different directions and no longer fall on a (nearly) regular lattice."," This does not, of course, prevent multiple appearances of a given object within the lightcone, but these duplicates are then viewed from different directions and no longer fall on a (nearly) regular lattice."1041" Unfortunately, this technique also introduces discontinuities in large-scale structure at the boundaries between replications, affecting clustering statistics in a way which is at least as difficult to model as that of the original periodicity."," Unfortunately, this technique also introduces discontinuities in large-scale structure at the boundaries between replications, affecting clustering statistics in a way which is at least as difficult to model as that of the original periodicity."1042" ? showed that for lightcones of relatively small solid angle, the central line-of-sight can be chosen to pass through the lattice of periodic replications in such a direction that multiple images of the same object are minimised or eliminated altogether."," \citet{Kitzbichler2007} showed that for lightcones of relatively small solid angle, the central line-of-sight can be chosen to pass through the lattice of periodic replications in such a direction that multiple images of the same object are minimised or eliminated altogether."1043" The latter is not possible if the comoving volume of the lightcone exceeds that of the simulation, but this technique can still be used to ensure that multiple appearances occur as far apart as possible both on the sky and in redshift."," The latter is not possible if the comoving volume of the lightcone exceeds that of the simulation, but this technique can still be used to ensure that multiple appearances occur as far apart as possible both on the sky and in redshift."1044 We therefore use the method of ? in this paper., We therefore use the method of \citet{Kitzbichler2007} in this paper.1045" Space is filled with periodic replications of the simulation, a position is chosen for the observer, and the central line-of-sight of the survey field is given a previously chosen orientation."," Space is filled with periodic replications of the simulation, a position is chosen for the observer, and the central line-of-sight of the survey field is given a previously chosen orientation."1046 Galaxies whose positions intercept the lightcone are selected and their comoving distance is converted into a redshift., Galaxies whose positions intercept the lightcone are selected and their comoving distance is converted into a redshift.1047" As explained in ?,, the time between stored snapshots for the Millennium Simulation varies between 100 and 380 Myr."," As explained in \citet{Kitzbichler2007}, the time between stored snapshots for the Millennium Simulation varies between 100 and 380 Myr."1048 This means that the intrinsic properties of galaxies are not generally available at the time corresponding to their comoving distance., This means that the intrinsic properties of galaxies are not generally available at the time corresponding to their comoving distance.1049 Rather they must be taken from the stored snapshot which is closest to their light-cone position., Rather they must be taken from the stored snapshot which is closest to their light-cone position.1050" Hence, galaxies with redshift (z+zi-1)/2«z«(%+2i41)/2 are assigned the physical properties stored at z;."," Hence, galaxies with redshift $(z_i+z_{i-1})/2<z<(z_i+z_{i+1})/2$ are assigned the physical properties stored at $z_i$."1051" The resulting discontinuity in galaxy population properties, at the boundaries between snapshots, could be reduced by interpolating, but this works poorly for positions and velocities since the output separation is comparable to orbital times within groups and clusters."," The resulting discontinuity in galaxy population properties, at the boundaries between snapshots, could be reduced by interpolating, but this works poorly for positions and velocities since the output separation is comparable to orbital times within groups and clusters."1052" Moreover, it is not straightforward for other galaxy properties either since these change discontinuously on timescales shorter than the output spacing, for example through mergers and starbursts."," Moreover, it is not straightforward for other galaxy properties either since these change discontinuously on timescales shorter than the output spacing, for example through mergers and starbursts."1053 We thus follow ? and do not attempt any interpolation., We thus follow \citet{Kitzbichler2007} and do not attempt any interpolation.1054 The semi-analytic calculations are perfomed on these intermediate time-steps that vary between 5 and 15 Myr., The semi-analytic calculations are perfomed on these intermediate time-steps that vary between 5 and 15 Myr.1055" This means, for example, that a burst of star formation will have this duration and can happen anywhere between (or at) output snapshots, with the corresponding increase in flux being reflected in galaxy properties at the snapshot."," This means, for example, that a burst of star formation will have this duration and can happen anywhere between (or at) output snapshots, with the corresponding increase in flux being reflected in galaxy properties at the snapshot."1056 Theapparent luminosities and colours of galaxies depend strongly on their redshifts through the conversion between rest- and observed-frame photometric bands and through the inverse square dependence of apparent luminosity on distance., The luminosities and colours of galaxies depend strongly on their redshifts through the conversion between rest- and observed-frame photometric bands and through the inverse square dependence of apparent luminosity on distance.1057 The final redshift of the galaxy in the lightcone is not available at the time observed-frame luminosities are computed in the semi-analytic model., The final redshift of the galaxy in the lightcone is not available at the time observed-frame luminosities are computed in the semi-analytic model.1058" However, there will be two extreme redshifts that bracket it."," However, there will be two extreme redshifts that bracket it."1059" We compute apparent observed-frame luminosities (for fixed intrinsic properties) using these upper and lower limits, and once the galaxy is placed in the lightcone, we interpolate to obtain final observed-frame quantities."," We compute apparent observed-frame luminosities (for fixed intrinsic properties) using these upper and lower limits, and once the galaxy is placed in the lightcone, we interpolate to obtain final observed-frame quantities."1060 For this paper we construct lightcones for square areas of 1.4x1.4 deg? out to high redshift with no faint magnitude cut.," For this paper we construct lightcones for square areas of $1.4\times10611.4$ $^2$ out to high redshift with no faint magnitude cut."1062 They are however limited by the mass resolution of the dark matter simulation (1.7x10'°h7'Me in halo mass) corresponding to stellar masses of ~109? at z=0., They are however limited by the mass resolution of the dark matter simulation $1.7\times10^{10}h^{-1}\rm{M}_{\sun}$ in halo mass) corresponding to stellar masses of $\sim10^{9.5}\rm{M}_{\sun}$ at z=0.1063" While this does not matter for the questions we studyM in this paper, it should be borne in mind if the lightcones are used for other Semi-analytic models predict intrinsic properties of galaxies, such as stellar mass, star formation history, gas and dust content and metallicity."," While this does not matter for the questions we study in this paper, it should be borne in mind if the lightcones are used for other Semi-analytic models predict intrinsic properties of galaxies, such as stellar mass, star formation history, gas and dust content and metallicity."1064" In order to convert these into observed spectral energy distributions (SED) or broad-band photometry, evolutionary population synthesis and dust models are required."," In order to convert these into observed spectral energy distributions (SED) or broad-band photometry, evolutionary population synthesis and dust models are required."1065" The former predict the evolution of the light associated with a single short burst of star formation of given metallicity and with an assumed Initial Mass Function (IMF), a so-called Simple Stellar Population (SSP)."," The former predict the evolution of the light associated with a single short burst of star formation of given metallicity and with an assumed Initial Mass Function (IMF), a so-called Simple Stellar Population (SSP)."1066 The intrinsic stellar emission from a model galaxy is then represented as a superposition of SSPs weighted, The intrinsic stellar emission from a model galaxy is then represented as a superposition of SSPs weighted1067with masses below 1.3 will on average be unevolved.,with masses below 1.3 will on average be unevolved.1068 The ratio of double degenerate svstems to those with MS companions has been addressed in detailed population synthesis studies (Iben et al., The ratio of double degenerate systems to those with MS companions has been addressed in detailed population synthesis studies (Iben et al.1069 1997)., 1997).1070 Iben et al., Iben et al.1071 assumed (hat only WDs with companions less than 0.3 wwere detectable. and such low mass svstems would be of the detectable total.," assumed that only WDs with companions less than 0.3 were detectable, and such low mass systems would be of the detectable total."1072 For a flat relative mass function this implies that in the underlving population of all svstems would have MS companions with masses less (han 1.3., For a flat relative mass function this implies that in the underlying population of all systems would have MS companions with masses less than 1.3.1073Af... However. for hotter WDs one would detect à larger fraction of systems.," However, for hotter WDs one would detect a larger fraction of systems."1074 Out of all systems would be undetected (CM> 0.65. )). detected with an MS companion GV< )). and. would be double degenerates.," Out of all systems would be undetected $M_{\rm secondary} >$ 0.65 ), detected with an MS companion $M_{\rm secondary} <$ ), and would be double degenerates."1075 With this assumption an observed relative formation rate of [or binary WDs becomes an intrinsic population of5., With this assumption an observed relative formation rate of for binary WDs becomes an intrinsic population of.10769%.. This is entirely consistent with the predicted binary fraction of from Duquennoy Mawor (1991) period distribution., This is entirely consistent with the predicted binary fraction of from Duquennoy Mayor (1991) period distribution.1077 We can use the 241ASS data to check on the fraction of systems with low mass companions., We can use the 2MASS data to check on the fraction of systems with low mass companions.1078 With a field binary fraction of we would expect half of the “single” WDs to actually have distant. and unresolved non-interacting binary companions: (hese companions will tvpically be low mass stars., With a field binary fraction of we would expect half of the “single” WDs to actually have distant and unresolved non-interacting binary companions; these companions will typically be low mass stars.1079 We expect of the low mass WDs produced from binary interactions to have close MIS companions and the other to be double degenerate svstems., We expect of the low mass WDs produced from binary interactions to have close MS companions and the other to be double degenerate systems.1080 In the Liebert οἱ al., In the Liebert et al.1081 sample there are 14 low mass WDs with τω>25.000 IX that were used to estimate the formation rate without applving an incompleteness correction.," sample there are 14 low mass WDs with $T_{\rm eff} > 25,000$ K that were used to estimate the formation rate without applying an incompleteness correction."1082 We would expect 2 double degenerate svstems. 5 close binary companions. ancl 3.5 distant binary. companions.," We would expect 2 double degenerate systems, 5 close binary companions, and 3.5 distant binary companions."1083 Three systems are known {ο have degenerate companions. aud we find significant J and Ix band excesses [or 4 of the svstems.," Three systems are known to have degenerate companions, and we find significant J and K band excesses for 4 of the systems."1084 Dased upon a comparison with isochrones. we infer masses ol 0.3-0.35 {for 3 companions and ~0.5 {for one.," Based upon a comparison with isochrones, we infer masses of 0.3-0.35 for 3 companions and $\sim$ 0.5 for one."1085 This indicates a deficit of companions relative to expectations. and in particular we note that there are no hieher mass companions seen even (though a significant number would be naturally produced by binary evolution.," This indicates a deficit of companions relative to expectations, and in particular we note that there are no higher mass companions seen even though a significant number would be naturally produced by binary evolution."1086 If all svstems were [rom interacting binaries. aud the sample was complete. we would expect LO MS companions.," If all systems were from interacting binaries, and the sample was complete, we would expect 10 MS companions."1087 We therefore conclude that even (he hot sample of He WDs is likely to be incomplete due to the presence of nearby companions., We therefore conclude that even the hot sample of He WDs is likely to be incomplete due to the presence of nearby companions.1088 Until a radial velocity survey. is completed. il is also not clear whether the MS companions detected are close binaries (attributed to the binary lormation channel) or distant ones (attributed) to the single Formation channel).," Until a radial velocity survey is completed, it is also not clear whether the MS companions detected are close binaries (attributed to the binary formation channel) or distant ones (attributed to the single formation channel)."1089 Follow-up work of this (wpe will be essential to empirically set the intrinsic space densities., Follow-up work of this type will be essential to empirically set the intrinsic space densities.1090considering that the photon index changes by 0.5 below and above the cooling break frequency. the maximum photon index allowed in the X-ray band is ay«a40.5=2. irrespective of by.,"considering that the photon index changes by 0.5 below and above the cooling break frequency, the maximum photon index allowed in the X-ray band is $\alpha_X < \alpha_X' + 0.51091= 2$, irrespective of $b_X$."1092 Now we compare these constraints on ay with the range required to 10nize the jet., Now we compare these constraints on $\alpha_X$ with the range required to ionize the jet.1093 The observed X-ray flux at the day 5 should not be much different at the day ~10-13. and we extrapolate the X-ray luminosity down to the UV band and compare to Loree.," The observed X-ray flux at the day 5 should not be much different at the day $\sim$ 10–13, and we extrapolate the X-ray luminosity down to the UV band and compare to $L_{\rm ph, rec}$."1094 Note that only a fraction of the X-ray luminosity is directed to the jet material. and this fraction is given by ~b/by from a geometrical consideration.," Note that only a fraction of the X-ray luminosity is directed to the jet material, and this fraction is given by $\sim b/b_X$ from a geometrical consideration."1095 We found that the spectral index must be extremely soft as ayz 5 or 4 for by= 1 or by=b0.1. respectively. in order that the extrapolated flux down to vp is equal to ἔρμιος.," We found that the spectral index must be extremely soft as $\alpha_X \gtrsim$ 5 or 4 for $b_X = $ 1 or $b_X = b = 0.1$, respectively, in order that the extrapolated flux down to $\nu_T$ is equal to $L_{\rm ph, rec}$."1096 Therefore we can safely exclude the possibility that the nonthermal radiation producing the observed X-rays is tonizing the jet., Therefore we can safely exclude the possibility that the nonthermal radiation producing the observed X-rays is ionizing the jet.1097 Secondly we consider a possibility that a hot. UV-radiating star close to the SN 2002ap may ionize the jet.," Secondly we consider a possibility that a hot, UV-radiating star close to the SN 2002ap may ionize the jet."1098 It is expected that SN 2002ap occurred in a massive star forming region where young massive stars are clustering., It is expected that SN 2002ap occurred in a massive star forming region where young massive stars are clustering.1099 A close binary system ts a candidate for the type Ic supernova progenitors (Nomoto. Filippenko. Shigeyama 1990). and it may provide even stronger ionization source.," A close binary system is a candidate for the type Ic supernova progenitors (Nomoto, Filippenko, Shigeyama 1990), and it may provide even stronger ionization source."1100 However. the total ionization lummosity given in eq. (16))," However, the total ionization luminosity given in eq. \ref{eq:L_rec}) )"1101" is even larger by a factor of several than the ionization flux. ~107""?s7!. above the frequency ry~6«10"" Hz for the most luminous and hottest stars (Schaere de Koter 1997)."," is even larger by a factor of several than the ionization flux, $\sim 10^{49.5} \ \rm s^{-1}$, above the frequency $\nu_T \sim 6 \times 10^{15}$ Hz for the most luminous and hottest stars (Schaere de Koter 1997)."1102 It should be noted that this luminosity is for all direction. but only the radiation within the solid angle of the jet viewed from the tonizing star is available for the Jet ionization. which is expected to be a small fraction.," It should be noted that this luminosity is for all direction, but only the radiation within the solid angle of the jet viewed from the ionizing star is available for the jet ionization, which is expected to be a small fraction."1103 If the region around SN 2002ap is filled up by radiation field with ~Fig) to a radius of rj. the region should have a luminosity of at least Ars corresponding. bolometric luminosity of 2.1«10b16irlΟΕ assuming the spectral energy distribution of the most luminous O stars.," If the region around SN 2002ap is filled up by radiation field with $\sim F_{\rm ion}$ to a radius of $r_{\rm jet}$, the region should have a luminosity of at least $4 \pi F_{\rm ion} r_{\rm1104jet}^2$, corresponding bolometric luminosity of $2.1 \times 10^8 b_{-1}^{-2}1105\zeta_{-1}^{-1} \mu_{14}^{-1} t_{10}^{-3} (f_{\rm el}/0.3) L_\odot$, assuming the spectral energy distribution of the most luminous O stars."1106 Such a huge luminosity is apparently ruled out by the prediscovery image of the SN 2002ap field reported by Smartt et al. (, Such a huge luminosity is apparently ruled out by the prediscovery image of the SN 2002ap field reported by Smartt et al. (11072002).,2002).1108 To conclude. ionization by nearby young stars Is Impossible.," To conclude, ionization by nearby young stars is impossible."1109 Since photoionization of the jet seems difficult. the only way to ionize the jet is enhanced collisional ionization by external heating.," Since photoionization of the jet seems difficult, the only way to ionize the jet is enhanced collisional ionization by external heating."1110 If the jet is generated at the central compact object. it might include a significant amount of radioactive nuclei such as ??Ni.," If the jet is generated at the central compact object, it might include a significant amount of radioactive nuclei such as $^{56}$ Ni."1111 Asymmetric explosion induced by the jet should also affect nucleosynthesis. and ??Ni production along the jet direction is enhanced (Nagataki 2000: Maeda et al.," Asymmetric explosion induced by the jet should also affect nucleosynthesis, and $^{56}$ Ni production along the jet direction is enhanced (Nagataki 2000; Maeda et al."1112 2002)., 2002).1113 °°Ni decays by electron capture and gamma-ray emission to ??Co. with an exponential decay time scale of fy;=fj2/1n(2)8.5 d and decay energy is ew;=2.1 MeV. When the material is optically thick. the radioactive heat is quickly thermalized into optical radiation field. as generally seen for supernovae.," $^{56}$ Ni decays by electron capture and gamma-ray emission to $^{56}$ Co, with an exponential decay time scale of $t_{\rm Ni} = t_{1/2} /1114\ln (2) = 8.5$ d and decay energy is $\epsilon_{\rm Ni} = 2.1$ MeV. When the material is optically thick, the radioactive heat is quickly thermalized into optical radiation field, as generally seen for supernovae."1115 On the other hand. if the material is mildly optically thin. the gamma-rays emitted by decaying ?*Ni scatter electrons with a probability ~ and since the gamma-ray energy is comparable with the electronπα. rest mass. the scattered electrons acquire mildly relativistic speed and energy.," On the other hand, if the material is mildly optically thin, the gamma-rays emitted by decaying $^{56}$ Ni scatter electrons with a probability $\sim \tau_{\rm jet}$, and since the gamma-ray energy is comparable with the electron rest mass, the scattered electrons acquire mildly relativistic speed and energy."1116" Such high energy electrons would lose their energy by ionization loss in the jet plasma. with a time scale of where c, 1s the initial velocity of high energy electrons."," Such high energy electrons would lose their energy by ionization loss in the jet plasma, with a time scale of where $\upsilon_e$ is the initial velocity of high energy electrons."1117 Here we used the ionization loss formulae of Longair (1992) and the logarithmic factor is set to be 15., Here we used the ionization loss formulae of Longair (1992) and the logarithmic factor is set to be 15.1118 Therefore the energy deposited by radioactive gamma-rays Is used to ionize the jet material within the time scale of interest. giving an efficient ionization process.," Therefore the energy deposited by radioactive gamma-rays is used to ionize the jet material within the time scale of interest, giving an efficient ionization process."1119" When optical depth is very low. this process would be dominated by positrons emitted from decay of οσο, which has a longer exponential lifetime of fe,= 111.26 days and energy fraction given to positrons is of the total decay energy (Arnett 1979; Woosley. Pinto. Hartmann 1989)."," When optical depth is very low, this process would be dominated by positrons emitted from decay of $^{56}$ Co, which has a longer exponential lifetime of $t_{\rm Co} = $ 111.26 days and energy fraction given to positrons is of the total decay energy (Arnett 1979; Woosley, Pinto, Hartmann 1989)."1120 The tonizing balance is determined by the energy balance between radioactive heating and recombination cooling (see also Graham 1988) as: where fw; is the “CNi mass fraction in the jet. and τν the ionization. potential.," The ionizing balance is determined by the energy balance between radioactive heating and recombination cooling (see also Graham 1988) as: where $f_{\rm Ni}$ is the $^{56}$ Ni mass fraction in the jet, and $w$ the ionization potential."1121 The recombination rate. coefficient depends on the electron gas temperature. which is determined by balance between the radioactive heating and cooling processes.," The recombination rate coefficient depends on the electron gas temperature, which is determined by balance between the radioactive heating and cooling processes."1122" We can estimate the minimum amount of ??Ni by taking the minimum value of o as where the adopted value of aj, is minimum value of doubly ionized oxygen or carbon at temperature of ~10""K (Nahar Pradhan 1997; Nahar 1999) and too=w/(20 eV).", We can estimate the minimum amount of $^{56}$ Ni by taking the minimum value of $\alpha_{\rm rec}$ as where the adopted value of $\alpha_{\rm rec}$ is minimum value of doubly ionized oxygen or carbon at temperature of $\sim 10^4$ K (Nahar Pradhan 1997; Nahar 1999) and $w_{20} = w$ /(20 eV).1123 The jet may include significant amount of heavier nuclei that are difficult to ionize for a fixed value of fy. but on the other hand. it may also include considerable helium that is easier to ionize.," The jet may include significant amount of heavier nuclei that are difficult to ionize for a fixed value of $f_{\rm el}$, but on the other hand, it may also include considerable helium that is easier to ionize."1124 The helium could be mixed from remaining helium layer of the progenitor. or it may be newly synthesized.," The helium could be mixed from remaining helium layer of the progenitor, or it may be newly synthesized."1125 Production of helium is also enhanced along the jet direction in energetic jet-like nucleosynthesis (Maeda et al., Production of helium is also enhanced along the jet direction in energetic jet-like nucleosynthesis (Maeda et al.1126 2002)., 2002).1127 We also note that highly ionized heavy nuclei. such as “Ni. should produce observable line emission in X-ray bands. and hence the observed weak X-ray flux gives a constraint on the species of ionized elements. (," We also note that highly ionized heavy nuclei, such as $^{56}$ Ni, should produce observable line emission in X-ray bands, and hence the observed weak X-ray flux gives a constraint on the species of ionized elements. ("1128See $5.2. for possible connection to X-ray line features often observed in GRB afterglows.),See \ref{section:X-ray-line} for possible connection to X-ray line features often observed in GRB afterglows.)1129 Whatever the jet composition ts. the above result indicates that. 1f the jet ts kept tonized by radioactive heating. it must have a considerable amount of ??Ni (mass fraction of order unity).," Whatever the jet composition is, the above result indicates that, if the jet is kept ionized by radioactive heating, it must have a considerable amount of $^{56}$ Ni (mass fraction of order unity)."1130 This estimate is. however. very uncertain especially about the composition of the jet. o and electron temperature.," This estimate is, however, very uncertain especially about the composition of the jet, $\alpha_{\rm rec}$ and electron temperature."1131 More sophisticated treatment is necessary to determine the ionization status. but it is beyond the scope of the paper.," More sophisticated treatment is necessary to determine the ionization status, but it is beyond the scope of the paper."1132 Therefore. it 1s difficult to conclude that the jet should be ionized. but it seems the best candidate of ionization process among others.," Therefore, it is difficult to conclude that the jet should be ionized, but it seems the best candidate of ionization process among others."1133 The jet may be ionized by gamma-rays of ?Ni decay leaking. from the photosphere of SN 2002ap. even if the jet does not have radioactive nuclei.," The jet may be ionized by gamma-rays of $^{56}$ Ni decay leaking from the photosphere of SN 2002ap, even if the jet does not have radioactive nuclei."1134 In fact. ionization. of helium envelope above the photosphere. which is required to explain the observed He lines in SN 1987A and type Ib supernovae. is ascribed to the leaking gamma-rays from photosphere (Graham 1988; Lucy 1991).," In fact, ionization of helium envelope above the photosphere, which is required to explain the observed He lines in SN 1987A and type Ib supernovae, is ascribed to the leaking gamma-rays from photosphere (Graham 1988; Lucy 1991)."1135 The mass of *°Ni produced by SN 2002ap Is estimated to be 0.0743:0.02M|. from the light curve modeling by Mazzali et al. (, The mass of $^{56}$ Ni produced by SN 2002ap is estimated to be $0.07 \pm 0.02 M_\odot$ from the light curve modeling by Mazzali et al. (11362002). which is larger than the jet mass.,"2002), which is larger than the jet mass."1137 However. the efficiency for gamma-rays to hit the jet is reduced by the beaming factor. 5b. and it is further reduced by escaping fraction from photosphere.," However, the efficiency for gamma-rays to hit the jet is reduced by the beaming factor, $b$, and it is further reduced by escaping fraction from photosphere."1138 Although some supernovae. including SN 1998bw. showed evidence that a significant amount of gamma-rays are leaking in late phase (= 30 days) (Nakamura et al.," Although some supernovae, including SN 1998bw, showed evidence that a significant amount of gamma-rays are leaking in late phase $\gtrsim$ 30 days) (Nakamura et al."1139 2001: Patat et al., 2001; Patat et al.1140 2001). the leaking fraction should not be large in early phase of ~ 10 days. when the optical luminosity ts still glowing up by," 2001), the leaking fraction should not be large in early phase of $\sim$ 10 days, when the optical luminosity is still glowing up by"1141percentage of events for which the OT would not be observed even without dust extinction. considering the light curves of our simulated GRBs and fixing BR hnuuitiug inaenitudeo of linον=20.5.,"percentage of events for which the OT would not be observed even without dust extinction, considering the light curves of our simulated GRBs and fixing R limiting magnitude of $R_{lim}=20.5$."1142 Solid bold line in Fi, Solid bold line in Fig.1143e.o 8 shows that for the majority of data, \ref{senzaz1} shows that for the majority of data1144Uulike the previous systems. Fomalhaut bas not vet been subject to detailed uumerical simulations.,"Unlike the previous systems, Fomalhaut has not yet been subject to detailed numerical simulations."1145 The Foimallhaut system also differs markedly in that 850 SSCUBA observatious suggest the system is seen nearly edge-on aud consists of a dusty. torus. rather than a thin disk (Hollandetal.1998).," The Fomalhaut system also differs markedly in that 850 SCUBA observations suggest the system is seen nearly edge-on and consists of a dusty torus, rather than a thin disk \citep{holland98}."1146. The unproved spatial resolution offered by the 150 SSCUBA images of Hollandetal.(2003) confirm the torus structure. while revealing a previously uuseen arc of emission near or withliu the torus.," The improved spatial resolution offered by the 450 SCUBA images of \citet{holland03} confirm the torus structure, while revealing a previously unseen arc of emission near or within the torus."1147 This departure from a uuilorm structure is strougly suggestive of a planetary. presence., This departure from a uniform structure is strongly suggestive of a planetary presence.1148 The generation of a single arc of emission can be achieved by a 1:1 resonance. as suggested by Hollandetal.(2003).," The generation of a single arc of emission can be achieved by a 1:1 resonance, as suggested by \citet{holland03}."1149. Inspection of results for resonance occupaucy obtained iu the construction of our synthetic catalogue. however. indicate that the 1:1 resonance is difficult to populate under normal circumstances.," Inspection of results for resonance occupancy obtained in the construction of our synthetic catalogue, however, indicate that the 1:1 resonance is difficult to populate under normal circumstances."1150 An example is shown in Figure 13.. where a Jupiter mass planet populates the 1:1 resonance when parent bodies exist interior to the planet. but does not populate the 1:1 resonance when the pareut bodies are all external to the planet.," An example is shown in Figure \ref{fig:11resonance}, where a Jupiter mass planet populates the 1:1 resonance when parent bodies exist interior to the planet, but does not populate the 1:1 resonance when the parent bodies are all external to the planet."1151 We consider it uulikely that parent bodies with sslightly less than the planets wwould survive for the required length. of time., We consider it unlikely that parent bodies with slightly less than the planet's would survive for the required length of time.1152" As an alternative explanation of the single are feature. we find that systems iucludingo a massive CÀZ,(abt2M{1 j,,). rnoderately eccentric 0.5) planet can iuduce a sugle emission arc over a backgrouud ring by trapping many particles in the n: resonances. where n71."," As an alternative explanation of the single arc feature, we find that systems including a massive $M_{pl} \geq M_{Jup}$ ), moderately eccentric $0.3 \leq e_{pl} \leq 0.5$ ) planet can induce a single emission arc over a background ring by trapping many particles in the $n:1$ resonances, where $n>1$."1153" Alter examining our synthetic catalogue. we chose a model with the followiug parameters: AL,—2.34AL. (appropriateforanΑΝstarsuchasFomalbaut.BarradoyNavascuesetal.1997).. Al=2M yup. ερ=OL. and ay~39 AU."," After examining our synthetic catalogue, we chose a model with the following parameters: $M_{\star} = 2.3 M_{\odot}$ \citep[appropriate for an A3V star such as Fomalhaut -- ][]{bar97}, $M_{pl} = 2 M_{Jup}$ , $e_{pl}=0.4$, and $a_{pl}1154\sim 59$ AU."1155" Parent bodies were distributed. with initial orbital parameters in the following ranges: 0<ey,«0.3. 100<ap,180 AU. 0<7,<25.."," Parent bodies were distributed with initial orbital parameters in the following ranges: $0<e_{pb}<0.3$, $100<a_{pb}<180$ AU, $0<i_{pb}<25$."1156. The higher parent body inclinatious naturally. generates a torus structure. which is believed to exist in the Fomalhaut system.," The higher parent body inclinations naturally generates a torus structure, which is believed to exist in the Fomalhaut system."1157 The simulation used 1000 test particles with 9=0.05. since tle mass of dust erains in the Fomalhaut system with diameters > ((corresponding to values of 8> 0.1) is believed to be negligible (Dentetal.2000).," The simulation used 1000 test particles with $\beta = 0.05$, since the mass of dust grains in the Fomalhaut system with diameters $>$ (corresponding to values of $\beta > 0.1$ ) is believed to be negligible \citep{dent00}."1158. The results of our simulations. which ran until no particles remained after 212 million years. are shown iu Figure 1L.," The results of our simulations, which ran until no particles remained after 212 million years, are shown in Figure \ref{fig:fomalhaut}."1159 The simulated observatious bear a close resemblance to the craw’ image of Fomalhaut presented in Hollandetal.(2003)., The simulated observations bear a close resemblance to the `raw' image of Fomalhaut presented in \citet{holland03}.1160. This planetary configurationOm does egenerates a dust distribution which does uot rotate with the planet. but appears fixed [rom a viewpoint external to the system over au orbital period.," This planetary configuration does generates a dust distribution which does not rotate with the planet, but appears fixed from a viewpoint external to the system over an orbital period."1161 Thus. observations of Fomalhaut over time would slow no eliauge iu emissiou if this model is correct.," Thus, observations of Fomalhaut over time would show no change in emission if this model is correct."1162 The effects of planetary phase are shown in Figure 15.., The effects of planetary phase are shown in Figure \ref{fig:fomalphase}.1163 Whilst the planetary configuration we have preseuted here displays a close similarity to the observations of Fomalhaut to date. as noted above results from our syuthetic catalogue suggest tliat," Whilst the planetary configuration we have presented here displays a close similarity to the observations of Fomalhaut to date, as noted above results from our synthetic catalogue suggest that"1164primary at birth.,primary at birth.1165" The initial conditions described in section (2) use q = wsEi as a definition of the mass-ratio, but since the gainer has Mabecome the most massive component of the Algol-system, we use q = “4 as the definition of mass-ratio of an Algol-system."," The initial conditions described in section \ref{sec_Initial}) ) use q = ${M_{g}\over M_{d}}$ as a definition of the mass-ratio, but since the gainer has become the most massive component of the Algol-system, we use q = ${M_{d}\over M_{g}}$ as the definition of mass-ratio of an Algol-system."1166 The observedMg Algols combine a large fraction of systems where Algol characteristics are produced, The observed Algols combine a large fraction of systems where Algol characteristics are produced1167"orientations probed, as shown in the third row, left panel of Figure 7..","orientations probed, as shown in the third row, left panel of Figure \ref{fig:dist2}."1168" Here, a represents the angle between the angular momentum vectors of 92 and the IBH at the start of the simulations."," Here, $\alpha$ represents the angle between the angular momentum vectors of S2 and the IBH at the start of the simulations."1169" This incidentally argues against the possibility that the observed deviations are due to Kozai oscillations, since the mechanism requires large relative inclinations to operate."," This incidentally argues against the possibility that the observed deviations are due to Kozai oscillations, since the mechanism requires large relative inclinations to operate."1170" However, fits with the same sets of parameters (q,e,a) but different a can have different values of x?."," However, fits with the same sets of parameters $(q,\,e,\,a)$ but different $\alpha$ can have different values of $\chi^2$."1171" This means that the knowledge of (q,e,a) is insufficient to predict the reduced x? but information on the sky position is necessary."," This means that the knowledge of $(q,\,e,\,a)$ is insufficient to predict the reduced $\chi^2$ but information on the sky position is necessary."1172" Finally, the right panel in the third row of Figure 7 shows that the minimum 3D distance between S2 and the IBH also correlates with the reduced x?."," Finally, the right panel in the third row of Figure \ref{fig:dist2} shows that the minimum 3D distance between S2 and the IBH also correlates with the reduced $\chi^2$."1173" In general it holds that the smaller the minimum distance, the worse the corresponding fit."," In general it holds that the smaller the minimum distance, the worse the corresponding fit."1174" Clearly, this parameter is not independent of the semi-major axis."," Clearly, this parameter is not independent of the semi-major axis."1175 The initial parameters adopted for the black hole binary are not sampled homogeneously., The initial parameters adopted for the black hole binary are not sampled homogeneously.1176" This is obvious for the first three panels in Figure 7 showing the reduced x? as a function of the binary parameters (q,e,a)."," This is obvious for the first three panels in Figure \ref{fig:dist2} showing the reduced $\chi^2$ as a function of the binary parameters $(q,\,e,\,a)$."1177 But it also holds (and is less obvious) for the plot investigating the minimum distance between 82 and the IBH., But it also holds (and is less obvious) for the plot investigating the minimum distance between S2 and the IBH.1178 A comparison of the goodness of fit for the runs starting with the IBH at periapsis and apoapsis is shown in Figure 9.., A comparison of the goodness of fit for the runs starting with the IBH at periapsis and apoapsis is shown in Figure \ref{fig:cfr}. .1179" In both cases, the semi-major axis of the black hole orbit is a=10mpc."," In both cases, the semi-major axis of the black hole orbit is $a=10\mpc$."1180 We find a modest worsening of the x? in the case of an IBH initially at the apoapsis of its orbit., We find a modest worsening of the $\chi^2$ in the case of an IBH initially at the apoapsis of its orbit.1181" This can be attributed to the fact that S2's apoapsis, where the star spends most of its time, is about 10mpc."," This can be attributed to the fact that S2's apoapsis, where the star spends most of its time, is about $10\mpc$."1182" Finally, we also investigated the minimum time required for the IBH to become detectable."," Finally, we also investigated the minimum time required for the IBH to become detectable."1183" For this purpose, we repeated the orbital fits for a few cases assuming that the observations span 10,15, 20,25, 30, 35, 40, 45 or 50 years (Figure 10))."," For this purpose, we repeated the orbital fits for a few cases assuming that the observations span 10,15, 20,25, 30, 35, 40, 45 or 50 years (Figure \ref{fig:time}) )."1184" Our initial conditions are such that the first periapse passage of S2 happens after 10 years, the second after 26 years and the third after 42 years."," Our initial conditions are such that the first periapse passage of S2 happens after 10 years, the second after 26 years and the third after 42 years."1185" Figure 10 shows that the reduced X? starts to increase beyond our threshold of 1.22 after the second periapsis passage for fits that show a large reduced x? after 50 simulated years (red/dashed, blue/long-dashed and green/solid curves)."," Figure \ref{fig:time} shows that the reduced $\chi^2$ starts to increase beyond our threshold of 1.22 after the second periapsis passage for fits that show a large reduced $\chi^2$ after 50 simulated years (red/dashed, blue/long-dashed and green/solid curves)."1186 Only for fits that after 50 simulated years have a reduced χ΄<3 is the threshold passed after the third periapse passage(black curve)., Only for fits that after 50 simulated years have a reduced $\chi^2 \simless 3$ is the threshold passed after the third periapse passage(black curve).1187 The discrete nature of periapse passages also is the reason why, The discrete nature of periapse passages also is the reason why1188"For direct comparison purposes we also show MDF data for the halo stars in three locations of the NGC 5128 halo, from Harris et al. (1999, 2000,, 2002)).","For direct comparison purposes we also show MDF data for the halo stars in three locations of the NGC 5128 halo, from Harris et al. \cite{har99}, \cite{har00}, \cite{har02}) )."1189" These were all taken with the HST WFPC2 camera in (V,7) and have rather similar limiting absolute magnitudes to our M87 ACS-based photometry, though they are less affected by crowding."," These were all taken with the HST WFPC2 camera in $(V,I)$ and have rather similar limiting absolute magnitudes to our M87 ACS-based photometry, though they are less affected by crowding."1190 The MDFs were derived in all cases with the same RGB grid of tracks and interpolation code., The MDFs were derived in all cases with the same RGB grid of tracks and interpolation code.1191" The mid-to-outer fields at projected galactocentric distances of 21 and 31 kpc have MDFs that are virtually identical to each other and are combined in the upper panel of Figure 12,, while the inner-halo 8 kpc field is shown by itself in the middle panel."," The mid-to-outer fields at projected galactocentric distances of 21 and 31 kpc have MDFs that are virtually identical to each other and are combined in the upper panel of Figure \ref{fehhisto}, while the inner-halo 8 kpc field is shown by itself in the middle panel."1192 Our corrected MDF for the M87 inner halo — at a mean projected distance of 10 kpc - clearly resembles the 8-kpc NGC 5128 field more closely than the outer fields., Our corrected MDF for the M87 inner halo – at a mean projected distance of 10 kpc – clearly resembles the 8-kpc NGC 5128 field more closely than the outer fields.1193" Over the range [m/H] «—0.3 where we can make the comparison, we conclude that the inner halos of both these giant ellipticals have basically similar MDFs that are broad, predominantly metal-rich, and with very small numbers of metal-poor stars."," Over the range [m/H] $< -0.3$ where we can make the comparison, we conclude that the inner halos of both these giant ellipticals have basically similar MDFs that are broad, predominantly metal-rich, and with very small numbers of metal-poor stars."1194" The MDF for M87 reach a peak near [m/H] ~—0.4, but for the present this metallicity should be viewed as a lower limit to the peak value."," The MDF for M87 reach a peak near [m/H] $\sim -0.4$, but for the present this metallicity should be viewed as a lower limit to the peak value."1195" We have used HST Archive data for an unusually deep set of F606W,F814W ACS images to probe the red-giant stellar population in the inner halo of M87."," We have used HST Archive data for an unusually deep set of $F606W, F814W$ ACS images to probe the red-giant stellar population in the inner halo of M87."1196" Although the crowding levels of the faint halo stars in this field are severe at best, the regions from 115""—155"" (9.3 to 12.5 kpc projected galactocentric distance) give a useful first look at the brightest 1.5 magnitudes of the red-giant branch."," Although the crowding levels of the faint halo stars in this field are severe at best, the regions from $115'' - 155''$ (9.3 to 12.5 kpc projected galactocentric distance) give a useful first look at the brightest 1.5 magnitudes of the red-giant branch."1197" This corresponds to distances of 1.4R,—1.9R,, where R, is the effective radius in the I band of 81"" (6.3 kpc) (Zeilinger, 1993))."," This corresponds to distances of $1.4 R_e -11981.9 R_e$, where $R_e$ is the effective radius in the I band of 81” (6.3 kpc) (Zeilinger, \cite{zei93}) )."1199 We have used this material to obtain a preliminary TRGB-calibrated distance to M87 for the first time., We have used this material to obtain a preliminary TRGB-calibrated distance to M87 for the first time.1200" We find a distance d=(16.7+0.9) Mpc, in good agreement with the few other relatively direct methods available, including the planetary nebula luminosity function, the Cepheid-calibrated SBF method, and the linear diameters of globular clusters."," We find a distance $d = (16.7 \pm 0.9)$ Mpc, in good agreement with the few other relatively direct methods available, including the planetary nebula luminosity function, the Cepheid-calibrated SBF method, and the linear diameters of globular clusters."

Showing the first 1,200 of 10416 lines. Download the file for the rest.