ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2" 1 will therefore scale it with o,=5.", I will therefore scale it with $\beta_s=5$ .3" Sealing the different variables of upper AGB stars gives where M, is the envelope mass.", Scaling the different variables of upper AGB stars gives where $M_{\rm env}$ is the envelope mass.4 For the relevant dynamical time I take (Ga)F?. where pay is the average density. of the entire star.," For the relevant dynamical time I take $(G \rho_{\rm av})^{-1/2}$, where $\rho_{\rm av}$ is the average density of the entire star."5" Scaling wilh tvpical number gives During the evolution along the AGB before the star starts to contract in radius. the luminosity increases aid the mass decreases. such that 74, increases and Τικony decreases."," Scaling with typical number gives During the evolution along the AGB before the star starts to contract in radius, the luminosity increases and the mass decreases, such that $\tau_{\rm dyn}$ increases and $\tau_{\rm KH-env}$ decreases."6 Therefore. their ratio QΞτιντικenv Increases.," Therefore, their ratio $Q \equiv \tau_{\rm dyn}/\tau_{\rm KH-env}$ increases."7 For evaluating the value of Q on the upper AGB. when the envelope mass is low. I take A.=M in equation (4)).," For evaluating the value of $Q$ on the upper AGB, when the envelope mass is low, I take $M_c=M$ in equation \ref{tkh2}) )."8 This gives for the upper AGB The thermal time Fy)egy d8 less (han an order of magnitude longer than the dvnamica lime during this late AGB phase., This gives for the upper AGB The thermal time $\tau_{\rm KH-env}$ is less than an order of magnitude longer than the dynamical time during this late AGB phase.9 Soker Πατρα (1999) noted that this relatively short thermal Gime must result in a strong irregular behavior of the envelope because dynamica molions. suchas pulsations aud convective motion. can cause large thermal perturbations," Soker Harpaz (1999) noted that this relatively short thermal time must result in a strong irregular behavior of the envelope because dynamical motions, suchas pulsations and convective motion, can cause large thermal perturbations"10continuum Lares and their origins in the core or in a nascent jet.,continuum flares and their origins in the core or in a nascent jet.11 We observed Alrk 348 on 2000 Alay 2 using MEHRLIN. which has a maximum baseline of 217 km. giving a beam size of 12 mas.," We observed Mrk 348 on 2000 May 2 using MERLIN, which has a maximum baseline of 217 km, giving a beam size of 12 mas."12 In order to obtain images of Alrk 348 as rapidly as possible we observed for a single 17 hr run in the maximum 16 MlIZ bandwidth. which corresponds to a total velocity width of ~200 km  and only covers the red-shifted half of the line seen by ? plus sullicient. line-free channels for continuum subtraction.," In order to obtain images of Mrk 348 as rapidly as possible we observed for a single 17 hr run in the maximum 16 MHz bandwidth, which corresponds to a total velocity width of $\sim200$ km $^{-1}$ and only covers the red-shifted half of the line seen by \scite{Falcke00} plus sufficient line-free channels for continuum subtraction."13" We observed Alrk 348 ab oa fixed frequeney of 5, = 21891.6 Mllz alternately with phase reference source J0057|3021 for 4 and 2 min respectively.", We observed Mrk 348 at a fixed frequency of $\nu_{\rm o}$ = 21891.6 MHz alternately with phase reference source J0057+3021 for 4 and 2 min respectively.14 36273. which had a [ux of 21.09Jy at that time Clerássranta. private communication). was observed. for 40 min and used to set the [ux scale for all sources.," 3C273, which had a flux of 21.09 Jy at that time (Terässranta, private communication), was observed for 40 min and used to set the flux scale for all sources."15 Further data processingὃνwas performed usingo (?).., Further data processingwas performed using \cite{Greisen94}.16 We applied the instrumental corrections from the calibrators and the phase reference source. solutions to Alrk 348., We applied the instrumental corrections from the calibrators and the phase reference source solutions to Mrk 348.17 We then adjusted the data to a constant velocity. putting τν = 4641.6 km “in channel 30 o£ 60 usable channels. with a separation of 3.37 knis +.," We then adjusted the data to a constant velocity, putting $V_{\rm LSR}$ = 4641.6 km $^{-1}$ in channel 30 of 60 usable channels, with a separation of 3.37 km $^{-1}$."18 We averaged all data for Mk 348 and made a map using natural weighting of the visibility data. which gave a beam size of 31 mas ο 21 mas.," We averaged all data for Mrk 348 and made a map using natural weighting of the visibility data, which gave a beam size of 31 mas $\times$ 21 mas."19" The three components (CO) above 398,5. from the initial map of Alrk 348 lay. in a region 6 by 18 mas elongated: north-south.", The three components ) above $3\sigma_{\rm rms}$ from the initial map of Mrk 348 lay in a region 6 by 18 mas elongated north-south.20 We used these as a model for phase self-calibration. applied. the solutions ancl examined. the spectrum of Mrk 348.," We used these as a model for phase self-calibration, applied the solutions and examined the spectrum of Mrk 348."21 The complex. visibilities were vector averaged in time channel by channel., The complex visibilities were vector averaged in time channel by channel.22 The flux density increased noticably with increasing frequency if we used the more southerly positions as the origin of phase., The flux density increased noticably with increasing frequency if we used the more southerly positions as the origin of phase.23 Fig., Fig.24 1 shows the spectrum on the baselines to the Cambridge antenna using the mean position of the maser emission (Table 12) as the phase origin., \ref{xan-fig1.eps} shows the spectrum on the baselines to the Cambridge antenna using the mean position of the maser emission (Table \ref{position}) ) as the phase origin.25 The low-frequeney end: of the band appeared to be continuum only., The low-frequency end of the band appeared to be continuum only.26 We used the first 4.25 MllIz of data for phase and. amplitude selt-calibration and applied the same corrections to the lince-only data às we did for the continuum cata., We used the first 4.25 MHz of data for phase and amplitude self-calibration and applied the same corrections to the line-only data as we did for the continuum data.27 The multi-channel \Irk 348 data were averaged: over every 0.75. MlIz ancl Fourier transformed. to make a 20-channel total emission dirty data cube., The multi-channel Mrk 348 data were averaged over every 0.75 MHz and Fourier transformed to make a 20-channel total emission dirty data cube.28 The average of the first 4.50 MllIz was subtracted pixel by pixel from the whole cube to leave a line-only dirty data cube which was €CLEANed to produce a line-onlv cube., The average of the first 4.50 MHz was subtracted pixel by pixel from the whole cube to leave a line-only dirty data cube which was ed to produce a line-only cube.29 The linc-only dirty data cube was subtracted from the total emission dirty data cube and the result was CLEANed to give a continuum-only. cube., The line-only dirty data cube was subtracted from the total emission dirty data cube and the result was ed to give a continuum-only cube.30 We also made a total emission cube., We also made a total emission cube.31" We fitted 2D Gaussian components to emission above 360,4. in each channel of every cube to determine the position and. peak [lux ορ.", We fitted 2D Gaussian components to emission above $3\sigma_{\rm rms}$ in each channel of every cube to determine the position and peak flux $S_{\rm p}$.32 We also measured the position and flux density of components fitted to the total [lux from Mrk 348 mapped using 15-MlIz bande-width prior to self-calibration and the 4.25-MlIz bandwidth self-calibrated continuum map., We also measured the position and flux density of components fitted to the total flux from Mrk 348 mapped using 15-MHz band-width prior to self-calibration and the 4.25-MHz bandwidth self-calibrated continuum map.33 Phe brightest continuum components was resolved after deconvolving the beam. so we could measure its FWIIM s and the integrated [ux δι.," The brightest continuum components was resolved after deconvolving the beam, so we could measure its FWHM $s$ and the integrated flux $S_{\rm i}$."34 Phe position anc component sizeuncertainties (A... 7) are proportional to the beame-wicdth divided by the signal-to-noise ratio allowing for the sparse baseline coverage (?:: 2)).," The position and component sizeuncertainties $\sigma_{\rm pos}$, $\sigma_{\rm s}$ ) are proportional to the beam-width divided by the signal-to-noise ratio allowing for the sparse baseline coverage \pcite{Condon98}; \pcite{Richards99}) )."35 All maps are presented in total intensity., All maps are presented in total intensity.36 Phe Dux scale should have «10/4 error but ALERLIN only has 5 antennas operating at 22 Gllz. and the sparse we coverage means it is possible for extended: continuum. emission. to appear brighter and more compact than it actually is. although the peak positions should be accurate.," The flux scale should have $<10\%$ error but MERLIN only has 5 antennas operating at 22 GHz, and the sparse $uv$ coverage means it is possible for extended continuum emission to appear brighter and more compact than it actually is, although the peak positions should be accurate."37 Moreover as only one side of the bandpass is line-free. errors. in baseline subtraction or bandpass calibration max. allect the maser [lux measurements.," Moreover as only one side of the bandpass is line-free, errors in baseline subtraction or bandpass calibration may affect the maser flux measurements."38 The absolute position accuracy is ~50 mas. mostly due to uncertainty in the position of J0057|3021 (7)..," The absolute position accuracy is $\sim50$ mas, mostly due to uncertainty in the position of J0057+3021 \cite{Wilkinson98}."39 A single patch of maser emission. was detected. in. 11 channel maps of the linc-onlv. data cube from. 4543.8. to 4655.0 km +. shown bv the contours in Fig. 2..," A single patch of maser emission was detected in 11 channel maps of the line-only data cube from 4543.8 to 4655.0 km $^{-1}$, shown by the contours in Fig. \ref{xan-fig2.eps}."40" These are in linear multiples of 30,4.", These are in linear multiples of $3\sigma_{\rm rms}$.41" Phe erev-scale shows the continuum-only emission above 305,4.", The grey-scale shows the continuum-only emission above $3\sigma_{\rm rms}$.42 Phe continuum peaks are marked ancl and the white crosses show their positions. which agree in cach channel to within 0.1 mas.," The continuum peaks are marked and and the white crosses show their positions, which agree in each channel to within 0.1 mas."43 The positions ancl Dux densities of the peaks are given in Table L.., The positions and flux densities of the peaks are given in Table \ref{position}. .44 Phe peak of the maser emission is Consistently olfset from the continuum peak: its mean position is 0.7 mas north ofS., The peak of the maser emission is consistently offset from the continuum peak; its mean position is $2.7\pm0.7$ mas north of.45 The total angular size ofthe maser region is X3 mas and the individual components are unresolved., The total angular size of the maser region is $\le3$ mas and the individual components are unresolved.46 Vhere is no significanto svstematic angularὃν separation- gradient. greater than .2mas in. 111 kmsn ο., There is no significant systematic angular separation-velocity gradient greater than 2mas in 111 km$^{-1}$ .47 appears to be resolved. ancl was fitted with à component," appears to be resolved, and was fitted with a component"48found the best fit spin to be αλ=0.46. due to ils harder spectrum.,"found the best fit spin to be $a/M=0.46$, due to its harder spectrum."49 Because limb carkenine has some spin-dependence. we also redid the fits with [ree normalization. but still got very similar best-fit spins.," Because limb darkening has some spin-dependence, we also redid the fits with free normalization, but still got very similar best-fit spins."50 The advent of thermodvuamically selbeonsistent simulations of MRI turbulence in verlically stratified shearing boxes (Turner2004:Hirose.ντο.&Stone2005) has finally opened the door to the construction of spectral models which are based on the physics of the turbulence itself.," The advent of thermodynamically self-consistent simulations of MRI turbulence in vertically stratified shearing boxes \citep{tur04,hir05} has finally opened the door to the construction of spectral models which are based on the physics of the turbulence itself."51 Surprisinglv. our investigation in (his paper indicates that the thermodwvnanmics itself. ie. the vertical profile of dissipation. has little effect on the emergent spectrum compared to models based on the the standard assumption of a constant dissipation rate per unit mass.," Surprisingly, our investigation in this paper indicates that the thermodynamics itself, i.e. the vertical profile of dissipation, has little effect on the emergent spectrum compared to models based on the the standard assumption of a constant dissipation rate per unit mass."52 This conclusion was also reached in (he earlier investigation of (2005)., This conclusion was also reached in the earlier investigation of \citet{dav05}.53. Unless an annulus in the disk is only moderately. effectively thick. most of the dissipation occurs deeper than the effective photosphere. and ils vertical profile therefore has little effect on the emergent spectrum.," Unless an annulus in the disk is only moderately effectively thick, most of the dissipation occurs deeper than the effective photosphere, and its vertical profile therefore has little effect on the emergent spectrum."54 The reader should bear in mind. however. (hat ihe numerical dissipation proliles. e.g. equation (1)). are based on simulations in which mechanical enerev is simply lost at the grid scale and replaced by internal energy.," The reader should bear in mind, however, that the numerical dissipation profiles, e.g. equation \ref{eqdfdmfit}) ), are based on simulations in which mechanical energy is simply lost at the grid scale and replaced by internal energy."55 More work needs (o be done to investigate whether these profiles are robust to changes in grid resolution., More work needs to be done to investigate whether these profiles are robust to changes in grid resolution.56 In addition. the simulations were [for gas pressure dominated annuli. aud large uncertainties still remain to be fully investigated in the radiation pressure dominated case. which is more relevant lor the innermost annuli of Iuninous accretion disks around black holes and neutron stars.," In addition, the simulations were for gas pressure dominated annuli, and large uncertainties still remain to be fully investigated in the radiation pressure dominated case, which is more relevant for the innermost annuli of luminous accretion disks around black holes and neutron stars."57 In contrast to the thermodsvnamies. the actualdynamics. i.e. Che contribution of magnetic yressure lo the vertical support of the disk againste exgravity. does produce significante changese in (he emergent spectrum when compared to models based on the standard ad hoc assumptions.," In contrast to the thermodynamics, the actual, i.e. the contribution of magnetic pressure to the vertical support of the disk against gravity, does produce significant changes in the emergent spectrum when compared to models based on the standard ad hoc assumptions."58 Magnetic pressure support at high altitude results in a much more extended density. scale heieht. because gas pressure is not required to support the atinosphere against (he increasing vertical gravitv at high altitude.," Magnetic pressure support at high altitude results in a much more extended density scale height, because gas pressure is not required to support the atmosphere against the increasing vertical gravity at high altitude."59 Because absorption opacily generally increases with density. a larger density scale height results in a smaller densitv at the effective photosphere in order io produce the same effective optical depth of munity.," Because absorption opacity generally increases with density, a larger density scale height results in a smaller density at the effective photosphere in order to produce the same effective optical depth of unity."60 As anticipated bv IHirose. (2005).. this enhances non-LTE effects in the spectrum because collisional processes in ihe plasma are (hen diminished compared to radiative processes.," As anticipated by \citet{hir05}, this enhances non-LTE effects in the spectrum because collisional processes in the plasma are then diminished compared to radiative processes."61 Even in LTE. lower density at the effective photosphere generally implies that the plasma becomes more ionized. reducing the bound-lree opacity.," Even in LTE, lower density at the effective photosphere generally implies that the plasma becomes more ionized, reducing the bound-free opacity."62 Electron scattering is enhanced compared to absorption. driving the emergent radiation spectrum closer to a modilied blackbocdsy.," Electron scattering is enhanced compared to absorption, driving the emergent radiation spectrum closer to a modified blackbody."63 All of these effects, All of these effects64considered. which may not be true either.,"considered, which may not be true either."65 Probing smaller subfields is uo solution for this problem as the number of data points becomes too sinall with respect to the uunuber of free parameters., Probing smaller subfields is no solution for this problem as the number of data points becomes too small with respect to the number of free parameters.66 A third aremment against dominant offsets in the data is the high. quality of the determination of RAL. Le. a linear o(A?)-relation with a low 47.," A third argument against dominant offsets in the data is the high quality of the determination of $RM$, i.e. a linear $\phi(\lambda^2)$ -relation with a low $\chi^2$."67 Of all pixels with high chough polarized intensity (CP)>20 uJv/beam). ~ (Gin Auriga) and ~ (in Horologimn) has a reduced v65ox2.," Of all pixels with high enough polarized intensity $\left<P\right> >6820$ mJy/beam), $\sim$ (in Auriga) and $\sim$ (in Horologium) has a reduced $\chi^2 < 2$."69 Π .offsets of. the same order of the data would exist. FAMs could not be so well-deteriined over such a lage part of the fields.," If offsets of the same order of the data would exist, $RM$ s could not be so well-determined over such a large part of the fields."70" For ideal data with constant 7. random offsets⋅ cause à 4727 ""nif the offsets are larecr than cs%.."," For ideal data with constant $P$, random offsets cause a $\chi^2 > 2$ if the offsets are larger than $\sim$."71 Finally. inodels of depolarization m a svuchrotrou-cluittineOo and Faraday-rotatingC» medi. which are presented im a companion paper (Iaverkorn et al.," Finally, models of depolarization in a synchrotron-emitting and Faraday-rotating medium, which are presented in a companion paper (Haverkorn et al."72 200La). do no show average Q or C values 210 mJy/beam (29).," 2004a), do not show average $Q$ or $U$ values $\ga7310$ mJy/beam $\sigma$ )."74 From the laree ouy.. the good quality of the RAL determinations. the depoluization models. aud from solving+ for+ offsets. that minimize∙∙H Dy. we conclude that the preseuce of considerable undoetected. large-scale structure due to the mussing short spacings is uulikelv.," From the large , the good quality of the $RM$ determinations, the depolarization models, and from solving for offsets that minimize $\chi^2$, we conclude that the presence of considerable undetected large-scale structure due to the missing short spacings is unlikely."75 This conclusion can also be checked. with absolutely calibrated polarized iuteusitv maps at los MIIZ. by Berkhuijscu and Droww (1963)., This conclusion can also be checked with absolutely calibrated polarized intensity maps at 408 MHz by Berkhuijsen and Brouw (1963).76 This frequency is close chough to 350 MIIz to allow comparison. although the volarized intensity at [08 ΛΠΙΣ is expected to be shelthy Heher because the polarization horizon is further away at his frequency.," This frequency is close enough to 350 MHz to allow comparison, although the polarized intensity at 408 MHz is expected to be slightly higher because the polarization horizon is further away at this frequency."77 We have smoothed our data to the FFWIIM of Berkhuijsen and DBroww. aud derived any nissing large-scale structure by comparing the two data sets.," We have smoothed our data to the FWHM of Berkhuijsen and Brouw, and derived any missing large-scale structure by comparing the two data sets."78 The polarized brightucss temperatures at Los MITz at the positions of the Auriga and IHorologi fields are 1.5 I aud 2.7 Ik. respectively.," The polarized brightness temperatures at 408 MHz at the positions of the Auriga and Horologium fields are 1.8 K and 2.7 K, respectively."79 Using a power law spectral index of 2.7. this corresponds to 2 I& aud 3 I at 350 MII.," Using a power law spectral index of 2.7, this corresponds to 2 K and 3 K at 350 MHz."80 The polarized Ivightucss temperatures derived from the sunoothed data are 0.07 I& and 0.12 I. respectively.," The polarized brightness temperatures derived from the smoothed data are 0.07 K and 0.12 K, respectively."81 Converting from Welvin to Jausky per beam (see Table 13) and taking into account that offsets in Q and C are on average a factor 2 sinaller than those in P. this meaus hat anv missing large-scale compoucuts in Stokes Q and U απο smaller than 10.6 mJy | for the Airiga region. and 13.7 ταν 1 for Torologimm.," Converting from Kelvin to Jansky per beam (see Table \ref{t3:data}) ) and taking into account that offsets in $Q$ and $U$ are on average a factor $\sqrt{2}$ smaller than those in $P$, this means that any missing large-scale components in Stokes $Q$ and $U$ are smaller than 10.6 mJy $^{-1}$ for the Auriga region, and 13.7 mJy $^{-1}$ for Horologium."82" For both ποια», this corresponds to about 2 to 3 signal-to-noise in Q and C. although it is not known what the influence of he difference in polarization horizon is."," For both fields, this corresponds to about 2 to 3 signal-to-noise in $Q$ and $U$, although it is not known what the influence of the difference in polarization horizon is."83 We conclude that hese data are not in disagreement with our conclusion hat ulssiug large-scale structure does uot play a major role in these observations., We conclude that these data are not in disagreement with our conclusion that missing large-scale structure does not play a major role in these observations.84 Therefore. the structure in polarized intensity niust v0 due wholly to depolarization mechauisuis.," Therefore, the structure in polarized intensity must be due wholly to depolarization mechanisms."85 For a pure Faraday screen the only kind of depolarization that is oossible is beam cepolarization. because the observed values of RAL nuplv that bandwidth depolarization is iof imuportaut. while depth depolarization requires that he rotating miedimun oenüts as well," For a pure Faraday screen the only kind of depolarization that is possible is beam depolarization, because the observed values of $RM$ imply that bandwidth depolarization is not important, while depth depolarization requires that the rotating medium emits as well."86 ILowewver. beam depolarization can ouly explain structure in P ou beau- scales.," However, beam depolarization can only explain structure in $P$ on beam-size scales."87 Therefore we are led to cousider the more realistic situation in which we observe a poluized, Therefore we are led to consider the more realistic situation in which we observe a polarized88must be invoked to bring them planets to their current locations.,must be invoked to bring them planets to their current locations.89 We uote that none of these svstenis are candidates for formation bv Nozai evcles with tidal friction due to perturbations *oa distant stellar companion (FabryvckyaudTremaine 2007).. since the presence of the other jauet would slut off Iozai effects.," We note that none of these systems are candidates for formation by Kozai cycles with tidal friction due to perturbations by a distant stellar companion \citep{fabr2007}, since the presence of the other planet would shut off Kozai effects."90 The two major classes of remaiming theories or qnoving these planets im are planet-planet scattering (RasioaudFord1996) ando disk uieration (ColdreichandTremaine1980:Liuetal. 1996).," The two major classes of remaining theories for moving these planets in are planet-planet scattering \citep{rasi1996} and disk migration \citep{gold1980,lin1996}."91". Planet scattering teuds to excite orbital eccentricity and inclination aud has difficulty wierating planets into short period orbits,", Planet scattering tends to excite orbital eccentricity and inclination and has difficulty migrating planets into short period orbits.92 Disk uieration is able to migrate multiple planets iuto short period orbits and tends to damp iuclination., Disk migration is able to migrate multiple planets into short period orbits and tends to damp inclination.93 The nearresonanut ratios of the observed systems avors the disk iieration hpothesis., The near-resonant ratios of the observed systems favors the disk migration hypothesis.94 Towever. continued nuieration after resonant trapping excites eccentricities and inclinations (LeeaudPeale2002:LeeandThomues 2009).," However, continued migration after resonant trapping excites eccentricities and inclinations \citep{lee2002,lee2009}."95. We consider the plivsical properties of planets (ο lnass and composition) that cau be derived frou the photometry.," We consider the physical properties of planets (e.g., mass and composition) that can be derived from the photometry."96 The primary motivation for this effort is to identify plausible masses for these planets in order to conduct a Moute Carlo study of TTY signals (Section 5))., The primary motivation for this effort is to identify plausible masses for these planets in order to conduct a Monte Carlo study of TTV signals (Section \ref{TTV}) ).97 For a eiven radius we estimate a range of lasses that depends heavily ou the possible bulk composi10115., For a given radius we estimate a range of masses that depends heavily on the possible bulk compositions.98 For planets with radi larecr than Saturn. the planet mass is larecly iundeteriuinate because of the transition iu the mass-racdius relation frou a Coulomb to an clectron degeueracy. dominated equation of state.," For planets with radii larger than Saturn, the planet mass is largely indeterminate because of the transition in the mass-radius relation from a Coulomb to an electron degeneracy dominated equation of state."99" The imass-radius curve turns over. so an object with a LR, radius could be auvthiug between a sub-Saturni mass planet (6.9.. ITAT-P-12h. 0931. Παπαefαἱ, (2009))) to a brown dwarf (e.g... CoRoT-3b. 21.7;. Deleuilctal, (2008)))."," The mass-radius curve turns over, so an object with a $1 R_J$ radius could be anything between a sub-Saturn mass planet (e.g., HAT-P-12b, $0.2 M_J$, \citet{hart2009}) ) to a brown dwarf (e.g., CoRoT-3b, $21.7 M_J$, \citet{dele2008}) )."100 For a planet radius up to that of Neptune. he planet mass is constraiue better between low-density objects rich idu oeas aud volatiles and the rocky. iron-rich aud ligh-density super-Earths.," For a planet radius up to that of Neptune, the planet mass is constrained better between low-density objects rich in gas and volatiles and the rocky, iron-rich and high-density super-Earths."101 Iuterpreting the bulk composition of such. planets is inore difficult due to the degeneracies thi uooarise with materials having different equations of state., Interpreting the bulk composition of such planets is more difficult due to the degeneracies that arise with materials having different equations of state.102 The measured dlanetaryradius aud the expected mass ranges of the candidate plauet sis shown in Table L., The measured planetaryradius and the expected mass ranges of the candidate planets is shown in Table \ref{planetproperties}.103 Under thesc| limitations. we estimate a dass range for each object using heoretical models. that are consisent with this level of observational uucertaintv (Valeuciaetal.2106.2007:Fortuev2009).," Under these limitations, we estimate a mass range for each object using theoretical models, that are consistent with this level of observational uncertainty \citep{vale2006,vale2007,fort2007,seag2007,gras2009}."104. These nodels of planeary interiors cover a wide range oD physical constitutionsΠο pure hydrogen to pure irou plaucts., These models of planetary interiors cover a wide range of physical constitutions—from pure hydrogen to pure iron planets.105 Obviously. there would be extremes that could iot arise qu nature.," Obviously, there would be extremes that could not arise in nature."106 One can consrain lasses based ou pure irou and pure water super-Earths. arening that plauct formation sceuarios would uot allow for such pure constitutions (Valenciactal.2007:Marcus2010α.0)...," One can constrain masses based on pure iron and pure water super-Earths, arguing that planet formation scenarios would not allow for such pure constitutions \citep{vale2007,marc2010a,marc2010b}."107" Iu rnv planet formation of any flavor. oejut imupacts and late water delivery are the ouly plausible way to ""purifv au initially iuixed-nuacrials formation iu a protoplanetary disk."," In primary planet formation of any flavor, giant impacts and late water delivery are the only plausible way to “purify” an initially mixed-materials formation in a protoplanetary disk."108 For exanuple. the iron-eulhauced bulk composition of Mercury is explained by an carly head-on iupact wit ioa similar bod.," For example, the iron-enhanced bulk composition of Mercury is explained by an early head-on impact with a similar body."109 Alareusctal.(20104) find that a uass-depeudeut limit ou final mean density {101100 radius) shouk exist for super-Earth planets WOKE Massive fiui l Mg. which 1s significantly less deuse thau pure iron.," \citet{marc2010a} find that a mass-dependent limit on final mean density (hence, radius) should exist for super-Earth planets more massive than 1 $M_E$, which is significantly less dense than pure iron."110 On the low-density bound (high radius). Marcusetal.(201No) show that more tiui about by nass chrichiment iu pure water is not possible. but here the upper envelope is not easily coustrained due to the possible addition of a II/Ie euvelope and/or extended: atuxsphere for a hot planet (Adamsetal.2008:RoecrsandSeager 2010)..," On the low-density bound (high radius), \citet{marc2010b} show that more than about by mass enrichment in pure water is not possible, but here the upper envelope is not easily constrained due to the possible addition of a H/He envelope and/or extended atmosphere for a hot planet \citep{adam2008,roge2010}. ."111" Stirting with the smalksize objects. KOT 191.2. with R,=2.0g. may well be a superEarth."," Starting with the small-size objects, KOI 191.02, with $R_p = 2.0 R_E$, may well be a super-Earth."112 It is half the size of he ice eiut Urauus (LRE) and stnaller than 11211 (2.7 Rr)., It is half the size of the ice giant Uranus $4 R_E$ ) and smaller than GJ1214b $2.7 R_E$ ).113 Towever. the 1nass range of5LAL spans the range between a water-vich world aud au won-rich remnant of a eiut impact collision (Marcusetal. 2010a).," However, the mass range of $5-18 M_E$ spans the range between a water-rich world and an iron-rich remnant of a giant impact collision \citep{marc2010a}."114". We use a mass of M,=LOAL¢ for KOT 191.02.", We use a mass of $M_p = 10 M_E$ for KOI 191.02.115" Next we have NOD 877.01 aud 2 have radii of 2.6Rp and 2.3/8gp. respectively,"," Next we have KOI 877.01 and 02 have radii of $2.6 R_E$ and $2.3 R_E$, respectively."116 Tjose planets are near the transition to the ice eiaits Uranus aud Neptune. but may be volatile-ricjsub-Neptunes or super-Eartls like 6121tb (Clhiarbonneaucfal. 2009)..," These planets are near the transition to the ice giants Uranus and Neptune, but may be volatile-richsub-Neptunes or super-Earths like GJ1214b \citep{char2009}. ."117 The estimated mass ranec in 610ME for ΟΙ 877.01 and 525Mpg for IKOI 877.02., The estimated mass range is $6-40 M_E$ for KOI 877.01 and $5-25 M_E$ for KOI 877.02.118 Since the high-imass. high-density hits are difficult to explain by existing planet formation scenarios we," Since the high-mass, high-density limits are difficult to explain by existing planet formation scenarios we"119data as well.,data as well.120 Nevertheless. the energy resolution of the PSPC is very limited (ALsi~ 50 per cent) at 1 keV. RS spectra have been fit to the data by Saracco Ciliegi (1995) and Pilelis et al. (," Nevertheless, the energy resolution of the PSPC is very limited $\Delta E/ E \sim$ 50 per cent) at 1 keV. RS spectra have been fit to the data by Saracco Ciliegi (1995) and Pildis et al. ("1211094).,1994).122 They derive a temperature of KI 1 keV. Instead. E fit a power-LIaw spectrum which is the standard model for AGN spectra at least in the small band.," They derive a temperature of $\sim$ 1 keV. Instead, I fit a power-law spectrum which is the standard model for AGN spectra at least in the small band."123 Lo obtain a spectral index of P=2.6301. much. steeper than the fit., I obtain a spectral index of $\Gamma=2.63^{+0.04}_{-0.16}$ much steeper than the fit.124" The column density is Vy=3.4""10"" em= higher than the Galactic column (C=82.11/82 dof).", The column density is $N_H=3.4^{+1.0}_{-0.4}\times 10^{20}$ $\rm cm^{-2}$ higher than the Galactic column $\chi^2=82.1/82$ dof).125 Finally. for the sake of completeness. E perform joint fits to the and the data.," Finally, for the sake of completeness, I perform joint fits to the and the data."126 However. bear in mind that. this joint analvsis has to be viewed with great. caution.," However, bear in mind that this joint analysis has to be viewed with great caution."127 IecentlIvy Iwasawa Nandra Fabian (1999) made spectral fits on and data of NOC5548., Recently Iwasawa Nandra Fabian (1999) made spectral fits on and data of NGC5548.128 Thev demonstrate that. the power-law fits may dilfer as much as ALzmO04 even in the common 0.5-2 keV band., They demonstrate that the power-law fits may differ as much as $\Delta\Gamma\approx 0.4$ even in the common 0.5-2 keV band.129 The reason for this large discrepancy may. be related with uncertainties in the calibration of both the and. detectors., The reason for this large discrepancy may be related with uncertainties in the calibration of both the and detectors.130 A single. power-law fit. to the combined. and data of HUXS00317-2142 vields P=2.00ο. Ny=19402.10° em? (=227.2/197 dol).," A single power-law fit to the combined and data of IRAS00317-2142 yields $\Gamma=2.00^{+0.07}_{-0.07}$, $N_H=1.9\pm 0.2\times 10^{20}$ $\rm cm^{-2}$ $\chi^2=227.2/197$ dof)."131 The power-law normalization is allowed to vary [ους between the and the observation epoch., The power-law normalization is allowed to vary freely between the and the observation epoch.132 Next. a RS component is added to the model.," Next, a RS component is added to the model."133 The temperature is constrained to have an upper limit of 1 keV. otherwise the resulting temperature becomes unrealistically high ( 15 keV).," The temperature is constrained to have an upper limit of 1 keV, otherwise the resulting temperature becomes unrealistically high $\sim$ 15 keV)."134 The best fit temperature is kT=0.07.111 keV. while the power-luy slope is P=1.9ρα., The best fit temperature is $\rm kT=0.07^{+0.02}_{-0.01}$ keV while the power-law slope is $\Gamma=1.9^{+0.12}_{-0.03}$.135 Despite the inclusion of the additional RS component the fit did. not improve (47=231.4/195 dot)., Despite the inclusion of the additional RS component the fit did not improve $\chi^2= 231.4/195$ dof).136 The detection of variability between the ancl the data clearly suggests an ACN origin for the X- emission., The detection of variability between the and the data clearly suggests an AGN origin for the X-ray emission.137" The data are well represented with a single power-law D—LS with no absorption above the Galactic Ny~1.5107"" em ", The data are well represented with a single power-law $\Gamma \sim 1.8$ with no absorption above the Galactic $N_H\sim 1.5 \times 10^{20}$ $\rm cm^{-2}$.138llence. the X-ray spectrum alone suggests a Sevfert-1 tvpe AGN.," Hence, the X-ray spectrum alone suggests a Seyfert-1 type AGN."139 This. is again compatible with the high X-ray. luminosity of this object. L.~1075Lo during the observation.," This is again compatible with the high X-ray luminosity of this object, $L_x\sim 10^{43}$, during the observation."140 Llowever. this interpretation comes in stark contrast with the optical spectrum which is indicative of a low luminosity or obscured ACN.," However, this interpretation comes in stark contrast with the optical spectrum which is indicative of a low luminosity or obscured AGN."141 Then a few possibilities arise for the nature of the AGN in LRASOO317-2142., Then a few possibilities arise for the nature of the AGN in IRAS00317-2142.142 We could view a Sevfert-1 nucleus overwhelmed in the optical by the emission of a powerful starforming galaxy., We could view a Seyfert-1 nucleus overwhelmed in the optical by the emission of a powerful starforming galaxy.143 In the X-ray band the emission. [roni the Sevfert-1: nucleus should. dominate over that arising from star-forming processes., In the X-ray band the emission from the Seyfert-1 nucleus should dominate over that arising from star-forming processes.144 Still. according to the spectral fits. the luminosity of the RS component alone is 5.107 erest: this classifies LLASOO317-2142 as one of the most. powerful X-ray star-forming galaxies known with a luminosity more than an order of magnitude above that of MS2 (Ptak et al.," Still, according to the spectral fits, the luminosity of the RS component alone is $5\times 10^{41}$ $\rm erg~s^{-1}$; this classifies IRAS00317-2142 as one of the most powerful X-ray star-forming galaxies known with a luminosity more than an order of magnitude above that of M82 (Ptak et al."145 1997)., 1997).146 The above scenario for the composites. which was originally proposed by Moran et al. (," The above scenario for the composites, which was originally proposed by Moran et al. ("147"1996) can be tested. by comparing the level of the Jf, in respect. with the hard X-ray luminosity.",1996) can be tested by comparing the level of the $H_{\alpha}$ in respect with the hard X-ray luminosity.148 Indeed. Ware et al. (," Indeed, Ward et al. ("1491988) found a strong correlation between the two quantities in a sample of LRAS selected Sevlert-L galaxics.,1988) found a strong correlation between the two quantities in a sample of IRAS selected Seyfert-1 galaxies.150" Then according to the scenario above. the composites should follow the same relation between £, ancl broad L(1£,)."," Then according to the scenario above, the composites should follow the same relation between $L_x$ and broad $L(H_\alpha)$ ."151" The luminosity of the broad ££, component in our object is about half of the total ZZ, luminosity (Moran et al.", The luminosity of the broad $H_{\alpha}$ component in our object is about half of the total $H_\alpha$ luminosity (Moran et al.152 1996)., 1996).153" Then the observed broad. Z4, luminosity is L(44,)~2104 L(Coziol et al.", Then the observed broad $H_\alpha$ luminosity is $L(H_\alpha) \sim 2\times 10^{41}$ (Coziol et al.154 1993)., 1993).155" This roughly translates to an X-ray luminosity of L,4.1077 according to Ward et al. (", This roughly translates to an X-ray luminosity of $L_x\sim 4\times 10^{42}$ according to Ward et al. (156"1988) not far olf the observed X-ray luminosity of L,~21075..",1988) not far off the observed X-ray luminosity of $L_x\sim 2\times 10^{42}$.157 Of course the long term X-ray. variability observed. introduces. some level of uncertainty in the above test., Of course the long term X-ray variability observed introduces some level of uncertainty in the above test.158 Note that Bassani et al. (, Note that Bassani et al. (1591999) reported the detection of a few Sevfert galaxics which possibly have a weak or absent Broad Line Region.,1999) reported the detection of a few Seyfert galaxies which possibly have a weak or absent Broad Line Region.160 Our object could. in principle belong to this category., Our object could in principle belong to this category.161 lHlowever. the L.£L(ILU) ratio rather argues against this hypothesis.," However, the $L_x/L(H_\alpha)$ ratio rather argues against this hypothesis."162 Again the possibility that the X-ray emission remains relatively unabsorbed while the optical sullers from additional obscuration cannot be ruled. out., Again the possibility that the X-ray emission remains relatively unabsorbed while the optical suffers from additional obscuration cannot be ruled out.163" The Balmer decrement in our object (using the Hux of the narrow LL, ancl Lf; from Moran et al.", The Balmer decrement in our object (using the flux of the narrow $H_\alpha$ and $H_\beta$ from Moran et al.164 1900) is about 5 corresponding toa column of ⋅⊽Ng~10721 em- (Bohlin et al., 1996) is about 5 corresponding to a column of $N_H\sim 10^{21}$ $\rm cm^{-2}$ (Bohlin et al.165 LOTS) ic an order of magnitude higher than the column derived above in the case of cold absorption., 1978) ie an order of magnitude higher than the column derived above in the case of cold absorption.166 The reason for this discrepancy is not obvious., The reason for this discrepancy is not obvious.167 One possibility is that the absorbing column is ionised: indeed in the case of a warm absorber the derived column is Ng~107 cm., One possibility is that the absorbing column is ionised: indeed in the case of a warm absorber the derived column is $N_H\sim 10^{22}$ $\rm cm^{-2}$.168" Then some fraction of the ""warm column should be located further away from the nucleus. where the narrow ff, ancl ff. lines originate."," Then some fraction of the 'warm' column should be located further away from the nucleus, where the narrow $H_\alpha$ and $H_\beta$ lines originate."169 Finally. LRASOOS17-2142 could be a heavily obscured (Compton thick) GN like eg Νέας10605.," Finally, IRAS00317-2142 could be a heavily obscured (Compton thick) AGN like eg NGC1068."170 Then a large fraction of the X-ray. emission could be due to scattered X-ravs. on a warm electron medium which should be situated well above the obseuring torus.," Then a large fraction of the X-ray emission could be due to scattered X-rays, on a warm electron medium which should be situated well above the obscuring torus."171" Phe broad £4, wing observed could. arise [rom scattered. radiation.", The broad $H_\alpha$ wing observed could arise from scattered radiation.172" However. the narrow ancl broad. 44, components have comparable Uuxes (Moran et al."," However, the narrow and broad $H_\alpha$ components have comparable fluxes (Moran et al."173 1996) arguing against this interpretation., 1996) arguing against this interpretation.174 Note that. the peculiar combination of an un-absorbed. X-ray spectrum with a narrow-line dominated. obscured optical spectrum. was also encountered in NOC3147 (Ptak et al.," Note that, the peculiar combination of an un-absorbed X-ray spectrum with a narrow-line dominated, obscured optical spectrum was also encountered in NGC3147 (Ptak et al."175 1996)., 1996).176 This object has a tvpe-2 twpe nucleus according to its optical spectrum which presents a relatively broac [NH] line (FWIIMAA00 kms. |)., This object has a type-2 type nucleus according to its optical spectrum which presents a relatively broad [NII] line $\sim$ 400 $\rm km~s^{-1}$ ).177 Phe X-ray spectrum of NCGCOS3147 is again very similar to our object as it shows no intrinsic absorption and a steep spectral index. EsL8., The X-ray spectrum of NGC3147 is again very similar to our object as it shows no intrinsic absorption and a steep spectral index $\Gamma \approx 1.8$.178" Hs observed X-ray luminosity is far below (L,~104 1)) that of our object: an Fe line at G4 keV (rest-frame) was clearly detected in NGCS3147 (Ptak et al.", Its observed X-ray luminosity is far below $L_x\sim 10^{41}$ ) that of our object; an Fe line at 6.4 keV (rest-frame) was clearly detected in NGC3147 (Ptak et al.179 1996) with an equivalent width greater than 130eV. Ptak ct al. (, 1996) with an equivalent width greater than 130eV. Ptak et al. (1801996) favoured a scenario where NGCSI47 harbours an obscured ΑΝ.,1996) favoured a scenario where NGC3147 harbours an obscured AGN.181 Lf indeed LRASOO317-2142 is a heavily obscured AGN the, If indeed IRAS00317-2142 is a heavily obscured AGN the182The results of the photoionization modeling of our nine svstenis are presented here and a summary can be found in Table 1..,The results of the photoionization modeling of our nine systems are presented here and a summary can be found in Table \ref{table:UandZ}.183 Voigt profile fit results with errors obtained using Mifit are presented in Table 2.., Voigt profile fit results with errors obtained using Minfit are presented in Table \ref{table:bandN}.184 The best-fit models are superimposed on kev constraint transitions in Figures 1- 9.., The best-fit models are superimposed on key constraint transitions in Figures \ref{fig:S1}- \ref{fig:S9}.185 This is a single cloud weak absorber. i.e. there is only one resolved component of absorption inMglt.," This is a single cloud weak absorber, i.e. there is only one resolved component of absorption in."186. A second phase is required to reproduce the observed aabsorption., A second phase is required to reproduce the observed absorption.187 The ionization parameteris constrained to be —3.0.xlogU<—2.5 for the pphase and log€>—1.8 lor the pphase., The ionization parameteris constrained to be $-3.0 \le \log{U} \le -2.5$ for the phase and $\log{U} \ge -1.8$ for the phase.188 Higher values for (he ionization parameter in the low pliase give rise to an overproduction of andAliu.. and lower values fail to produce enoughAlu.," Higher values for the ionization parameter in the low phase give rise to an overproduction of and, and lower values fail to produce enough."189 A lower ionization parameter in the hieh phase overproducesSilv., A lower ionization parameter in the high phase overproduces.190". Using a number density of photons of logn,=—4.83 αἱ this redshift (Haardt&Maclan 2001).. this corresponds to —2.3<lognj;X—1.8[em?] for ihe pphase and lognj€—3.1[em7] for the pphase."," Using a number density of photons of $\log{n_{\gamma}} = -4.83$ at this redshift \citep{Haardt01}, , this corresponds to $-2.3 \le \log{n_H} \le -1.8~[\cc]$ for the phase and $\log{n_H} \le -3.1~[\cc]$ for the phase."191 This constraint is based upon a solar ratio of Fe to Me., This constraint is based upon a solar ratio of Fe to Mg.192 However. higher densities would apply if there is a-enhancement (i.e. if the Fe to Meg ratio is lower than the solar value).," However, higher densities would apply if there is $\alpha$ -enhancement (i.e. if the Fe to Mg ratio is lower than the solar value)."193 iis not covered due to the low redshift of (his svstem. so there is no constraint on metallicity.," is not covered due to the low redshift of this system, so there is no constraint on metallicity."194 In the optically thin regime. our constraints on logÜ are insensitive (to the assumed metallicity.," In the optically thin regime, our constraints on $\log U$ are insensitive to the assumed metallicity."195 This is à single cloud absorber inMgr., This is a single cloud absorber in.196 A second phase is required to reproduce the observed aabsorption., A second phase is required to reproduce the observed absorption.197 The ionization parameter is constrained to be —3.5xlogÜ<—3.0 for the low ionizalion phase. and —1.5<logU1.3 for the high ionization phase.," The ionization parameter is constrained to be $-3.5 \le \log{U} \le -3.0$ for the low ionization phase, and $-1.8 \le \log{U} \le 1.3$ for the high ionization phase."198 Lower ionization parameters lor the low phase overproduceOr. and higher values overproduce Sirv..," Lower ionization parameters for the low phase overproduce, and higher values overproduce ."199 The constraint on thehigh phase is based upon the strength of the, The constraint on thehigh phase is based upon the strength of the200FN)mBisnes|(0=Tey).,$R_p(\lambda) = R_{planet}+z(\tau= \tau_{eq})$.201 Using a model atinosphere. and nuuericallv integrating over the whole trausluceut atmosphere. we can calculate the effective planet radius. aud then obtain the correspouding optical thickuess at the effective radius. 74.," Using a model atmosphere, and numerically integrating over the whole translucent atmosphere, we can calculate the effective planet radius, and then obtain the corresponding optical thickness at the effective radius, $\tau_{eq}$."202 For various atmospheric scale heights. we caleulate Tey by nunerical iutegration.," For various atmospheric scale heights, we calculate $\tau_{eq}$ by numerical integration."203" For a wide range of:mosphoeric scale height. provided that R,/II is between ~30 and +3000, the resulting τοι is roughly constant at a value ty%0.56 (Pig. D))."," For a wide range of atmospheric scale height, provided that $_p$ /H is between $\sim$ 30 and $\sim$ 3000, the resulting $\tau_{eq}$ is roughly constant at a value $\tau_{eq}\approx 0.56$ (Fig. \ref{Tau_eq}) )."204" Iu the case of 11897323b. R,/IT varies between 280 aud 560 heu the temperature varies from 1000 to 2000 Ik: therefore. the approximation of a coustaut τος at 0.56 tally applies."," In the case of 189733b, $_p$ /H varies between 280 and 560 when the temperature varies from 1000 to 2000 K; therefore, the approximation of a constant $\tau_{eq}$ at 0.56 fully applies."205 This demonstrates that. for a eiven atmospheric structure aud coluposition. estimating the altitude at which r=r40.56 is all that is needed to calculate the effective radius of the plauct at a given waveleneth.," This demonstrates that, for a given atmospheric structure and composition, estimating the altitude at which $\tau=\tau_{eq}=0.56$ is all that is needed to calculate the effective radius of the planet at a given wavelength."206" For a giveu atmospheric structure aud composition. the effective. altitude of the atimosphere at a waveleneth A is calculated by finding τετ=7.,.A). which solves the equation TA.2)=τοι."," For a given atmospheric structure and composition, the effective altitude of the atmosphere at a wavelength $\lambda$ is calculated by finding $z(\tau = \tau_{eq} ,\lambda )$, which solves the equation $\tau(\lambda,z)=\tau_{eq}$."207" Using the quantities defined above. the effective altitude + is given by where Gan, ancl £454, are the cross section aud abundance of the dominant absorbing species."," Using the quantities defined above, the effective altitude $z$ is given by where $\sigma_{abs}$ and $\xi_{abs}$ are the cross section and abundance of the dominant absorbing species."208 If the variation of the cross section as a function of wavelength is kuown. the observation of the altitude as a function of wavelength allows the derivation of 11 aud. therefore. of the temperature T. eiven by: Usine Eq. 1..," If the variation of the cross section as a function of wavelength is known, the observation of the altitude as a function of wavelength allows the derivation of $H$ and, therefore, of the temperature $T$ , given by: Using Eq. \ref{z_lambda},"209 the partial pressure of the main absorbent at the reference altitude is estimated by where A.y ds he wavelength at which the effective planetary radius corresponds to (or is used to define) 2=0., the partial pressure of the main absorbent at the reference altitude is estimated by where $\lambda_{z=0}$ is the wavelength at which the effective planetary radius corresponds to (or is used to define) $z=0$.210 We note that there is a deecneracy between the abuucdance am the total pressure in the atinosphere., We note that there is a degeneracy between the abundance and the total pressure in the atmosphere.211 From the measurement of the effective radius using transit spectroscopy. oue needs to assume an abundance of the absorbent to derive the pressure. or alternatively to assunie a pressure to derive the abundance.," From the measurement of the effective radius using transit spectroscopy, one needs to assume an abundance of the absorbent to derive the pressure, or alternatively to assume a pressure to derive the abundance."212 Iun sunuuaryv. the effective planetary radius is characteristic of the pressure and abundauce. aud the variation iu this radius as a function of waveleueth is characteristicof the temperature.," In summary, the effective planetary radius is characteristic of the pressure and abundance, and the variation in this radius as a function of wavelength is characteristicof the temperature."213 Iu 1159733b. the plot of altitude as a function of the wavelengthshows au increase in absorption toward shorter waveleugths (Fig. 2)).," In 189733b, the plot of altitude as a function of the wavelengthshows an increase in absorption toward shorter wavelengths (Fig. \ref{H_vs_Lambda}) )."214 Using Eq., Using Eq.215" 2. and assmuine a scaling law for the cross section iu the form 06=89(ÀA/Ag)"". the slope of the planet radius as a function of the wavoeleugth is given by dR,/dluA=difdluAaff. Therefore. we have For IIDI1159733b. the slope is measured το be + E106 ((svst.)"," \ref{T_dz} and assuming a scaling law for the cross section in the form $\sigma= \sigma_0 ( \lambda/ \lambda_0)^{\alpha}$ , the slope of the planet radius as a function of the wavelength is given by $216dR_p/d\ln\lambda = dz/d\ln\lambda = \alpha H.217$ Therefore, we have For 189733b, the slope is measured to be $\pm$ $\pm$ (syst.)"218 from 600. to l000nun (Pout ot 22008)., from 600 to nm (Pont et 2008).219" This corresponds to GR,fduiAz -920 kin", This corresponds to $dR_p/d\ln\lambda$ $\approx$ $-920$ km.220" Wethere foreobtaina T στ Sslo. Theteiperaturcin 1l189133bisdeterininedtobeinthera K(Deiningctal. 22006: νομος, 22007):there fore. is found to be close toe zz1. which is typical of Raleigh scatterme."," We therefore obtain $\alpha T$ $\approx$ $-5840$ K. The temperature in 189733b is determined to be in the range K (Deming et 2006; Knutson et 2007); therefore, $\alpha$ is found to be close to $\alpha\approx -4$, which is typical of Rayleigh scattering."221 Assuming a= as Expectoc for Ravleigh scattering alc sing Eq. 2.. ," Assuming $\alpha=-4$ as expected for Rayleigh scattering and using Eq. \ref{T_dz}, ,"222we derive a temperature where (stat.), we derive a temperature where (stat.)223 aud. (svst.), and (syst.)224 correspond to statistical aud systematics error bars on fü. as quoted in Pout ct ((2008).," correspond to statistical and systematics error bars on $dR_p$, as quoted in Pont et (2008)."225 For coudeusates. the scale height Z7. is often found o be significantly smaller than the gaseous scale height GL.~ HT/3) (Ackermmau Marley 2001: Fortueyv. 2005).," For condensates, the scale height $H_c$ is often found to be significantly smaller than the gaseous scale height $(H_c\sim H /3$ ) (Ackerman Marley 2001; Fortney 2005)."226 If lis were the case. because of the observed slope iu the spectrum. the temperature (or a) should be about three nes higher than found above. which seeiis unlikely.," If this were the case, because of the observed slope in the spectrum, the temperature (or $\alpha$ ) should be about three times higher than found above, which seems unlikely."227 We herefore couclude that. if produced by dust condeusates. he observed transit spectrum of bb shows that hese condensates aust be παπαπαν well-mixed vertically with the atmospheric eas with a similar scale height.," We therefore conclude that, if produced by dust condensates, the observed transit spectrum of b shows that these condensates must be unusually well-mixed vertically with the atmospheric gas with a similar scale height."228 Oue possible carricy ofthe Ravleigh scattering is the most abundant molecule. molecular hydrogeu.," One possible carrier of the Rayleigh scattering is the most abundant molecule, molecular hydrogen."229 The effective altitude at which the Ravieigh scattering of inolecular hydrogen dominates only depends on the mean density., The effective altitude at which the Rayleigh scattering of molecular hydrogen dominates only depends on the mean density.230" Therefore. measureiieuts of the altitude im the regime where Ravleigh scattering bw Πω dominates over other absorbeuts allows the determination of not onlv the temperature. but because the abundance of IT) is close to 1 (£j, 1) also the total density aud consequently the total pressure at the reference zero altitude."," Therefore, measurements of the altitude in the regime where Rayleigh scattering by $_2$ dominates over other absorbents allows the determination of not only the temperature, but because the abundance of $_2$ is close to 1 $\xi_{H_2}\sim 1$ ) also the total density and consequently the total pressure at the reference zero altitude."231 Following Eq. 3.. ," Following Eq.\ref{xiP}, ,"232the pressure Lyat the altitude corresponding to the radius at waveleneth Ay iswhere a is the Ravleigh scattering cross section at Ay., the pressure $P_0$at the altitude corresponding to the radius at wavelength $\lambda_0$ iswhere $\sigma_0$ is the Rayleigh scattering cross section at $\lambda_0$ .233 Using the refractive iudex of molecular lydrogen, Using the refractive index of molecular hydrogen234We start from au exact solution ou the uniform grkl haviug oulv one cell.,We start from an exact solution on the uniform grid having only one cell.235 The solution is used as an initial value to obtain the solution ou the other uniform grid., The solution is used as an initial value to obtain the solution on the other uniform grid.236 The obtained solution is used to obtain the solution on the nested grid of (NV.0) = (2°.Dαμ).," The obtained solution is used to obtain the solution on the nested grid of $ (N, \, h) $ = $ (2^2, \, 2 ^{n-2} h _0) $."237 This solution is obtained by the 5iccessive Over Relaxation (SOR.see.e.g;P‘essοἱal.1986.forthebasicofSOR)..," This solution is obtained by the Successive Over Relaxation \citep[SOR, see, e.g.][for the basic of SOR]{press86}."238" The computaion cost is very small since this nested eric coutaius only 616,444 cells."," The computation cost is very small since this nested grid contains only $ 64 \, \ell _{\rm max} $ cells."239 We obtained a solution «3 the coarsest erid alter Linas times Iteratiol in a typical model., We obtained a solution on the coarsest grid after $ 4 \ell _{\rm max} $ times iteration in a typical model.240 The iteration shotld reduces tie residual by à factor of 10.oE7. since £p:N4/2 1lines iteration reduces the 1residual by a factOr of 10 Pi1 SOR (Pressetal.1986). when the grkl has Ny cells in one dimeusion.," The iteration should reduces the residual by a factor of $ 10 ^{-3} $, since $ 4 p N _1 / 3 $ times iteration reduces the residual by a factor of $ 10 ^{-p} $ in SOR \citep{press86} when the grid has $ N _1 $ cells in one dimension."241" By interpolating ue solution we obtain the initial data for iteration oi the erid of CV.A) (25.9""hoy."," By interpolating the solution we obtain the initial data for iteration on the grid of $ (N, \, h) $ = $ (2 ^3, \, 2 ^{n-3} h _0) $."242 After a few times (typically two times) red-black Catuss-Seiclel iteratiou. we obtaiu an apxoximnate solution on the fine eric.," After a few times (typically two times) red-black Gauss-Seidel iteration, we obtain an approximate solution on the fine grid."243 We «jbtain au approxinate solution on a finer nested ος successively by combination of inerrolaion. the red-black Ciauss-Seidel iteration auc| restriction (averaging).," We obtain an approximate solution on a finer nested grid successively by combination of interpolation, the red-black Gauss-Seidel iteration and restriction (averaging)."244 We use the bilinear fo“Tdla 1ο luterpoate the solutiou on a coarse eric for that ou a fine grid., We use the bilinear formula to interpolate the solution on a coarse grid for that on a fine grid.245 Ou he other haud we use {e ftll weightiug (the simple volume average) to restrict the solution ou a fue grkl to that ou a coarse grkl., On the other hand we use the full weighting (the simple volume average) to restrict the solution on a fine grid to that on a coarse grid.246 We yerforim these operatious on the order of the FAIG., We perform these operations on the order of the FMG.247 The schedule of the ο»eratlolns are sciematically shown in Figure 2.., The schedule of the operations are schematically shown in Figure \ref{mg.eps}.248 The arrows cirectine right upward denote interpolation while those directing rigu downward do restriction., The arrows directing right upward denote interpolation while those directing right downward do restriction.249 The symbol. C. denotes the red-lack Catss-Seiclel Heration while the symbol. S. dose SOR.," The symbol, G, denotes the red-black Gauss-Seidel iteration while the symbol, S, dose SOR."250 Atthe points of the syiubol. E. au exac| solution is given.," At the points of the symbol, E, an exact solution is given."251 The red-]dack Gatss-Seidel iteraion is performed twiὁ both after interpolation aud interpolation ina typical model., The red-black Gauss-Seidel iteration is performed twice both after interpolation and interpolation in a typical model.252 Lh other words. the number ¢XE pre- αι post-iterations is two in our com»utation.," In other words, the number of pre- and post-iterations is two in our computation."253 Each restriction is performed by simply taking tlie average in the cells involved iu a coarse cell., Each restriction is performed by simply taking the average in the cells involved in a coarse cell.254 Restriction is done from fine (= (µας) to coarse (f= 1).," Restriction is done from fine $ \ell \, = \, \ell _{\rm max} $ ) to coarse $ \ell \, = \, 1 $ )."255 Each interpolation is performed in the following procedures., Each interpolation is performed in the following procedures.256 Each red-black iteration is performec in tle followiug procedures., Each red-black iteration is performed in the following procedures.257to the conclusions drawn from the asvuunetric Fe Ίνα line (Tanaka et al.,to the conclusions drawn from the asymmetric Fe $\rm K\alpha$ line (Tanaka et al.258 1995)., 1995).259 Persistent. strong aud rapid variability secs to be a rare phenomenon among radio-quiet ACN aud las previously been detected oulv in the uarrow-line Sevtert 1 ealaxy IRAS 132213809 (Boller et al.," Persistent, strong and rapid variability seems to be a rare phenomenon among radio-quiet AGN and has previously been detected only in the narrow-line Seyfert 1 galaxy IRAS 13224–3809 (Boller et al."260 1997)., 1997).261 The object shows amplitude X-ray variability of a factor ercater than 10. ou time scales less han about 1 dav.," The object shows amplitude X-ray variability of a factor greater than 10, on time scales less than about 1 day."262 The oivsical origin of elaut aud rapid N-rav flares ix most probably different roni the X-rav variability occuring on much louser time scales of vears. or longer. im active aud active ealaxies.," The physical origin of giant and rapid X-ray flares is most probably different from the X-ray variability occuring on much longer time scales of years, or longer, in active and non-active galaxies."263 A value of the Unbble coustaut of Ly=TOlans+Mpe| and a cosmological deceleration parameter of quy=4 have been adopted throughout., A value of the Hubble constant of $H_0 \rm = 70\ km\ s^{-1}\ Mpc^{-1}$ and a cosmological deceleration parameter of $q_0 \rm = \frac{1}{2}$ have been adopted throughout.264 The X-ray properties of LES 1927|651 discussed. below were obtained with the ROSAT aud Chaucdra satellites., The X-ray properties of 1ES 1927+654 discussed below were obtained with the ROSAT and Chandra satellites.265 The ROSAT data cover the survey observations carried out between July 1990 aud July 1991 as well as à ROSAT pointed observation obtained in December 1998., The ROSAT data cover the survey observations carried out between July 1990 and July 1991 as well as a ROSAT pointed observation obtained in December 1998.266 For the ROSAT AI-Skv. Survev observations the source counts were obtained using a circular source cell with a radius of 1 .larcnmdn., For the ROSAT All-Sky Survey observations the source counts were obtained using a circular source cell with a radius of 11.1 arcmin.267 The uuuber of source plu4. backeroux counts within this cell was 1255 + 65., The number of source plus background counts within this cell was 4255 $\pm$ 65.268" The backeroux was determined from a source-free cell with a radius of 11.5 arci. located in the sean direction through the centroid position of the source and ceutered at R.A.(2000) = 19'22™26.1°, Dec.(2000) = 6°LI17""."," The background was determined from a source-free cell with a radius of 11.5 arcmin, located in the scan direction through the centroid position of the source and centered at R.A.(2000) = $\rm 19^h22^m26.1^s$, Dec.(2000) = $\rm 65^{\circ}44^{\prime}17^{\prime\prime}$."269 The muuber of backgrouud counts normalizes to the source cell size Was 338 + Ls., The number of background counts normalized to the source cell size was 338 $\pm$ 18.270 The uct counts are therefore 3917 + 62. resulting in a mean count rate of 1.22 £0.02countss|.," The net counts are therefore 3917 $\pm$ 62, resulting in a mean count rate of 1.22 $\rm \pm 0.02\ counts\ s^{-1}$."271 For cach path of the source through the ROSAT PSPC detector the corresponding backeroune was subtracted., For each path of the source through the ROSAT PSPC detector the corresponding background was subtracted.272 LES 1927|651 was again observed during the final observation period of the ROSAT satellite iu. 1998 December 8 between 18:17:21. UT and 18:11:01. UT with the PSPC detector., 1ES 1927+654 was again observed during the final observation period of the ROSAT satellite in 1998 December 8 between 18:17:21 UT and 18:41:01 UT with the PSPC detector.273 The total exposure time was 1338 seconds., The total exposure time was 1338 seconds.274 This observation was affected by some anomalies (cf, This observation was affected by some anomalies (cf.275 Sect., Sect.276 2.1 of Boller et al., 2.1 of Boller et al.277 2000) aud we have performed a careful data quality check to ensure that our results were not affected., 2000) and we have performed a careful data quality check to ensure that our results were not affected.278 Tn addition. LES 19271651 was observed with Chandra during the guaranteed time programe.," In addition, 1ES 1927+654 was observed with Chandra during the guaranteed time programme."279 The Chandra observation was performed using the low cherey transmission erating spectrograph (LETCS. see Brinkman et al.," The Chandra observation was performed using the low energy transmission grating spectrograph (LETGS, see Brinkman et al."280 2000). which comprises the low energv transiuission erating (LETC) aud the high resolution caluera iuicrochannel plate detector for spectroscopy (IIRC-S).," 2000), which comprises the low energy transmission grating (LETG) and the high resolution camera microchannel plate detector for spectroscopy (HRC-S)."281 IESI9271651 was observed from March 20. 2001 at 13:20 UT until March 21. 2001 at 06:58 UT.," 1ES1927+654 was observed from March 20, 2001 at 13:20 UT until March 21, 2001 at 06:58 UT."282 The total exposure obtained from all eood time intervals. corrected for deadtime is 63110ssec.," The total exposure obtained from all good time intervals, corrected for deadtime is sec."283 During the whole observation uo times with high backerouud were detected., During the whole observation no times with high background were detected.284 The raw data was reprocessed and the spectrum extracted usns CTIÀO version 2À and CALDB release 2.15 following the correspoucding science threads available at the Chandra N-rayCounter., The raw data was reprocessed and the spectrum extracted using CIAO version 2.2.1 and CALDB release 2.15 following the corresponding science threads available at the Chandra X-ray.285 The overlapping higher spectral orders were taken iuto account bv using the LETCS erating response matrices (version July. 2002) which include 1st to 6th order.," The overlapping higher spectral orders were taken into account by using the LETGS grating response matrices (version July, 2002) which include 1st to 6th order."286 Separate files are provided for the negative aud positiveders?., Separate files are provided for the negative and positive.287. The fits with the same parameters were performed sinultaucouslv with the extracted negative and positive orders., The fits with the same parameters were performed simultaneously with the extracted negative and positive orders.288 Ότο inass fractious Va322.105.,$^3$ He mass fractions $X_3 \gtrsim 2\E{-3}$.289 These differeut nova rigecrs change the pre-outburst hpuuinositv of the nova system which we detail in refseciobs..," These different nova triggers change the pre-outburst luminosity of the nova system, which we detail in \\ref{sec:obs}."290 The ignition mass ds also affected by the rigecring iaechanisn. which results im a previously nuconsidered dependeuce of ealactic nova rates ou conrposition.," The ignition mass is also affected by the triggering mechanism, which results in a previously unconsidered dependence of galactic nova rates on composition."291 This effect is especially relevant eiven he upconmüug flood of data from optical transient survevs such as Pan-STARRS-1. Pan-STARRS-1. aud he Large Synoptic Survey Telescope. which will mcasure j0va rates in external galaxies with ereater accuracy han available iu current data.," This effect is especially relevant given the upcoming flood of data from optical transient surveys such as Pan-STARRS-1, Pan-STARRS-4, and the Large Synoptic Survey Telescope, which will measure nova rates in external galaxies with greater accuracy than available in current data."292 We speculate on the observational consequences audsunmnmniuize our work iu refseciconc.., We speculate on the observational consequences andsummarize our work in \\ref{sec:conc}.293 Corc-euvelope mixing models differ ou the mechliauisum by which core material is brought iuto the accreted laver., Core-envelope mixing models differ on the mechanism by which core material is brought into the accreted layer.294 Chemical diffusion (Prialnik&Iovetz1981:Ixovetz and shear musing caused by differeutial rotation of the acereted material (Iippeulhahu&Thomas19758:Livio2001). result in pre-convective penetration of a simall amount of hwdrosen into the uuderbiug material aud vice versa.," Chemical diffusion \citep{pk84,kp85,pk95,yaron05} and shear mixing caused by differential rotation of the accreted material \citep{kt78,lt87,ks89,alex04} result in pre-convective penetration of a small amount of hydrogen into the underlying material and vice versa."295" df the material below the laver is απο, p1 PC yeactions trigecr couvection. which homogenizes the envelope and the cutrained core material."," If the material below the layer is C-rich, $p+^{12}$ C reactions trigger convection, which homogenizes the envelope and the entrained core material."296 These pre-convective curichment models differ frou mechanisms ij which convection is triggered above the core iu accreted iiaterial whose composition is determined by that of the donor star and is thus relatively C-poor., These pre-convective enrichment models differ from mechanisms in which convection is triggered above the core in accreted material whose composition is determined by that of the donor star and is thus relatively C-poor.297 The introduction of core material iuto the envelope for these convective curichment models is caused by the couvective motion itself. either via couvective overshoot iuto the core (WoosleyL986) or shear mixing induced by the convective eddies (Cdasueretal.1997:Rosner 2001).," The introduction of core material into the envelope for these convective enrichment models is caused by the convective motion itself, either via convective overshoot into the core \citep{woos86} or shear mixing induced by the convective eddies \citep{glt97,ros01}."298. As of vet. no mixing mechanisin has definitively proven itself successful in explaining the curichiments of all uovae.," As of yet, no mixing mechanism has definitively proven itself successful in explaining the enrichments of all novae."299 For example. the multi-cvele diffusion studies of DPrialuik&Isovetz(1995) and Yaronetal.(2005). are initiated with matter accreted directly onto naked C/O cores. vet the accreting WD mav be O/Ne iu as may as 1/23 of all observed novae (πια&Livio1956:Rittereal.1991:Livio&TruranCül-Ponset 2003).," For example, the multi-cycle diffusion studies of \cite{pk95} and \cite{yaron05} are initiated with matter accreted directly onto naked C/O cores, yet the accreting WD may be O/Ne in as many as 1/3 of all observed novae \citep{tl86,rit91,lt94,gil03}."300. For these systems. diffusion would notlead to the imitiation of convection below the accreted laver aud subsequenu core dredee-up because the underlying material is not rich.," For these systems, diffusion would notlead to the initiation of convection below the accreted layer and subsequent core dredge-up because the underlying material is not C-rich."301 The studies of Livio&Truran (1987).. Fujimoto(19088. 1993).. and Piro&Bildsten(2001). rule ou any significant differential rotation between the accretec laver aud the core. which casts doubt on accretion-induced shear mixing mechanisius.," The studies of \cite{lt87}, , \cite{fuji88,fuji93}, and \cite{pb04} rule out any significant differential rotation between the accreted layer and the core, which casts doubt on accretion-induced shear mixing mechanisms."302 Werceketal.(1998.1999) find that couvective overshoot aud shearing do uot sufficiently: curich the envelope to produce a fast uova. although possible problems with their boundary conditions are pointed out by Clasnoretal.(2005).," \cite{kht98,kht99} find that convective overshoot and shearing do not sufficiently enrich the envelope to produce a fast nova, although possible problems with their boundary conditions are pointed out by \cite{glt05}."303. Moreover. sole recurrent nova. which are novae with recurrence times ~30 vr. do uot show over-albuudauces of mnetals in their ejecta (Williamsetal.1981:Williams1982:Warner1995:Tachisu&Kato 2001).. possibly due to a large Ποτά buffer above the core.," Moreover, some recurrent nova, which are novae with recurrence times $\sim 30 $ yr, do not show over-abundances of metals in their ejecta \citep{wil81,wil82,war95,hk01}, possibly due to a large helium buffer above the core."304 Thus. it is unclear how inuch the envelope will be enriched iu inetals prior to nova ignition for AM>109AL.vr|.," Thus, it is unclear how much the envelope will be enriched in metals prior to nova ignition for $\dot{M}>10^{-9} \ \smpy$ ."305 Some nova studies that include the accretion phase assmnue that the accreting euvelope is pre-enriched bv the core and consists of up to core material by mass οJosé&Ieruauz1998:Star-rheldetal.1998).," Some nova studies that include the accretion phase assume that the accreting envelope is pre-enriched by the core and consists of up to core material by mass \citep[e.g.,][]{jh98,starr98}."306. However. the uncertainties involved in the imixiug mechanisius coupled with the lack of observed metal enrichment iu recurrent novae lead us to exanune the consequences of assuming no C-euricluneut in the accreted envelope prior to the ouset of convection.," However, the uncertainties involved in the mixing mechanisms coupled with the lack of observed metal enrichment in recurrent novae lead us to examine the consequences of assuming no C-enrichment in the accreted envelope prior to the onset of convection."307" There are several previous studies that also follow this treatinent (Starrfieldetal.1985.1988.2000:Truran 2007).. but none sutficicutly samples the full parameter space in which we are interested: Starrficldetal.(1985.1988) ancl ναιetal.(1988). consider Z=0.02 accretion onto AL>1.35AL. WDs: the inodels of Starvfieldetal.(2000) have metallicity Z=10.2 or 0.02. as motivated by novae in the Large Magellauic Cloud. with lower accretion rates <0?Af.vy| than what we study: Piersautictal.(2000) exanune accretion with 3 metallicitics (Z=0.02. 10 andl )onto WDs with masses <0.68AL... lower than our parameter range: aud Joséetal.(2007).. in their study of novae in primordial binaries. have ouly solar aud very sub-solar metallicity 1.35AL. models with Z=0.02. 2.10© or 10.* with 322«101""3f,veto which is lower than our range."," There are several previous studies that also follow this treatment \citep{starr85,starr88,starr00,tru88,pier00,jose07}, but none sufficiently samples the full parameter space in which we are interested: \cite{starr85,starr88} and \cite{tru88} consider $Z=0.02$ accretion onto $ M \geq1.35 \ M_\odot$ WDs; the models of \cite{starr00} have metallicity $Z=10^{-3}$ or $0.02$, as motivated by novae in the Large Magellanic Cloud, with lower accretion rates $\lesssim 10^{-9} \ \smpy$ than what we study; \cite{pier00} examine accretion with 3 metallicities $Z=0.02$, $10^{-3}$, and $10^{-4}$ ) onto WDs with masses $<0.68 \ M_\odot$, lower than our parameter range; and \cite{jose07}, in their study of novae in primordial binaries, have only solar and very sub-solar metallicity $1.35 \ M_\odot$ models with $Z=0.02$, $2\E{-6}$, or $10^{-7}$, with $\dot{M}=2\E{-10} \ \smpy$, which is lower than our range."308" Moreover. none of these studies consider the effect of πο, which can play a dominant role in trigecring the nova (Shara1980:Townsley&200 1)."," Moreover, none of these studies consider the effect of $^3$ He, which can play a dominant role in triggering the nova \citep{shara80,tb04}."309. Iu our study. we assume that the effect of chemical diffision is neglieible.," In our study, we assume that the effect of chemical diffusion is negligible."310 This assumption. aud thus our results. are invalid if the material directly below the accreted laver is C-nieh.," This assumption, and thus our results, are invalid if the material directly below the accreted layer is C-rich."311 However. as we have described above. any nova svstenis exist in which the uuderlviug material is C-poor and diffusion is indeed negligible.," However, as we have described above, many nova systems exist in which the underlying material is C-poor and diffusion is indeed negligible."312 Our results ouly apply to these svstenis., Our results only apply to these systems.313 Tf the accretion rate ina CV is lower than the niuinuinuuvate for stability. the result will be a hydrogen shell flash.," If the accretion ratein a CV is lower than the minimumrate for stability, the result will be a hydrogen shell flash."314 Iu this section. we calculate the ignition conditions for these thermonuclear novae.," In this section, we calculate the ignition conditions for these thermonuclear novae."315" Throughout this study. we make the asstuuption that the accreted laver is thin. with a pressure scale height at the euvelope base. 7—P,pay. much less than the WD core radius. A."," Throughout this study, we make the assumption that the accreted layer is thin, with a pressure scale height at the envelope base, $h=P_b/\rho_b g$, much less than the WD core radius, $R$ ."316 The subscript b refers to the base of the accreted envelope. g=GAL/R? is the exavitational acceleration. assumed coustaut because h«Ro and M is the WD core mass.," The subscript $b$ refers to the base of the accreted envelope, $g=GM/R^2$ is the gravitational acceleration, assumed constant because $h \ll R$, and $M$ is the WD core mass."317" The ratio of the scale height to the WD radius for an ideal gas equation of state is the ratio of the thermal energy to eravitational energy. oq where T+ is the basetempcrature in units of 10* IN, AL, is the WD core mass in units of AY... R4=R/LO? cn. the proton niass is aad the atomic mass por particle is go=0.6 for solar ην."," The ratio of the scale height to the WD radius for an ideal gas equation of state is the ratio of the thermal energy to gravitational energy, = = , where $T_7$ is the basetemperature in units of $10^7$ K, $ M_1 $ is the WD core mass in units of $ M_\odot $ , $R_9=R/10^9$ cm, the proton mass is $m_p$ , and the atomic mass per particle is $\mu=0.6$ for solar composition."318 Fortypical ignition conditious. 7;2ον101 Is. so the composition.shell is very thin for the," Fortypical ignition conditions, $T_b \simeq 2\times 10^7$ K, so the shell is very thin for the"319 , 320"they should give identical answers, and minimizing the sensitivity of the match to the complexities of atmospheric detail.","they should give identical answers, and minimizing the sensitivity of the match to the complexities of atmospheric detail."321" Optical depth is sensitive to temperature (the opacity is κ.οςT? here), so that the visible surface is a complex structure (see Fig."," Optical depth is sensitive to temperature (the opacity is $\kappa \propto T^9$ here), so that the visible surface is a complex structure (see Fig."322 24 in Stein&Nordlund (1998)))., 24 in \cite{sn98}) ).323" For example, a fluctuation in temperature implies a change in 2.4 in the opacity."," For example, a fluctuation in temperature implies a change in 2.4 in the opacity."324" The optical depth of the photosphere occurs at different radii for different positions on the solar surface, so that fitting it with a single radius is difficult."," The optical depth of the photosphere occurs at different radii for different positions on the solar surface, so that fitting it with a single radius is difficult."325" At greater depths we expect the 3D atmospheres and the 1D models to agree, but near the surface it is not clear that the 3D and 1D definiinitions of optical depth are consistent."," At greater depths we expect the 3D atmospheres and the 1D models to agree, but near the surface it is not clear that the 3D and 1D definiinitions of optical depth are consistent."326 Pressure should be a better coordinate for matching 3D results to a 1D model., Pressure should be a better coordinate for matching 3D results to a 1D model.327" Unlike the optical depth, the pressure is a weaker function of angular position on the solar surface."," Unlike the optical depth, the pressure is a weaker function of angular position on the solar surface."328" Hydrodynamic flow tends to smooth pressure variations, making the definition of a mean pressure-radius relation more meaningful."," Hydrodynamic flow tends to smooth pressure variations, making the definition of a mean pressure-radius relation more meaningful."329" Figure 5 plots convective velocities versus log pressure for models A through E. We can see that the Asplund,etal.(2005) model smoothly joins onto model D. Let us construct a simple model of the motion in this region to see how hydrodynamic arguments might give modifications to the purely hydrostatic boundary conditions used in models A through E. We will assume that the velocity is dominated by flow at the largest scales of turbulence.", Figure \ref{figasp2} plots convective velocities versus log pressure for models A through E. We can see that the \cite{aspat} model smoothly joins onto model D. Let us construct a simple model of the motion in this region to see how hydrodynamic arguments might give modifications to the purely hydrostatic boundary conditions used in models A through E. We will assume that the velocity is dominated by flow at the largest scales of turbulence.330" These scales contain most of the kinetic energy, and are least non-laminar."," These scales contain most of the kinetic energy, and are least non-laminar."331 Convective motions are driven by the sinking of matter which is cooling due to transparency near the surface., Convective motions are driven by the sinking of matter which is cooling due to transparency near the surface.332 This generates gravity waves in the near-surface region., This generates gravity waves in the near-surface region.333" We will approximate the large scale average of this motion by g-modes (Landau&Lifshitz(1959),, see 12) whose amplitude falls off exponentially with pressure scale height."," We will approximate the large scale average of this motion by g-modes \cite{ll59}, see 12) whose amplitude falls off exponentially with pressure scale height."334 This implied a scaling with, This implied a scaling with335value of σι at three reference maguituce levels are also reported in each panel to provide au easier aud quantitative Comparison.,value of $\sigma_m$ at three reference magnitude levels are also reported in each panel to provide an easier and quantitative comparison.336 Each of the figuresOm presented above provide a complete and statistically robust characterization ol the photometric errors and of the effects of bleuding affecting our observations., Each of the figures presented above provide a complete and statistically robust characterization of the photometric errors and of the effects of blending affecting our observations.337 lu what follows we will refer to Fig., In what follows we will refer to Fig.338 19. as au exemple to describe the general treuds that are present. in all these plots., \ref{dm_v} as an exemple to describe the general trends that are present in all these plots.339 We will 1ake full use of the extensive set of artificial star experiments lor the production of proper svutlieic CMDs that will allow the recoustruction of the star formation history of the galaxy., We will make full use of the extensive set of artificial star experiments for the production of proper synthetic CMDs that will allow the reconstruction of the star formation history of the galaxy.340 This part ofthe analysis will be described in a companion paper (Aunibali et al., This part of the analysis will be described in a companion paper (Annibali et al.341 in preparation)., in preparation).342 Iu the followiip. we shortly comment ou the effects of blending aud incompleteness to assess the overall accuracy of our plotomery. with particular atteution to the impact ou the distauce estimate discussed in Section 5.," In the following, we shortly comment on the effects of blending and incompleteness to assess the overall accuracy of our photometry, with particular attention to the impact on the distance estimate discussed in Section 5."343 From the inspection of Fig. 19..," From the inspection of Fig. \ref{dm_v},"344 Fi, Fig.345eο 20 and Fig., \ref{dm_i} and Fig.346 21 Hineear that the very coucentratecd populatio1 of bright sars in Region 7 procduces a hieh degree of iucinpleteuess at faint magnitudes iu this region., \ref{dm_jh} it is clear that the very concentrated population of bright stars in Region 7 produces a high degree of incompleteness at faint magnitudes in this region.347 The high degree of crow«ing aud the high level of scy background. produced by the wings of lie uauv bdeht stars restltS in much larger photomet‘ic errors aud brighter limitiug magitude than in al the other Ree1ous (0-6)., The high degree of crowding and the high level of sky background produced by the wings of the many bright stars results in much larger photometric errors and brighter limiting magnitude than in all the other Regions (0-6).348 Ouside of this ve'w central part. the behavior of Anas sa Auction of 1iagnitude is etile sinilar and very well sampded everywhere.," Outside of this very central part, the behavior of $\Delta m $ as a function of magnitude is quite similar and very well sampled everywhere."349 I is particularly interesting to cleck at which maguuitucle he effects of blending become signifiCall. le. at wha level Xin signiicautly cleviates [rom zero.," It is particularly interesting to check at which magnitude the effects of blending become significant, i.e. at what level $\Delta m $ significantly deviates from zero."350" First of all we uote that the mean value ol Ai ds always lower than 6,5. thus raucdoim errors dominate over the systematics associated with lie lending at every magnitu«e."," First of all we note that the mean value of $\Delta m $ is always lower than $\sigma_m$, thus random errors dominate over the systematics associated with the blending at every magnitude."351 Second. in the Reegious [roi 10 to 6. the mean Ar is lower," Second, in the Regions from 0 to 6, the mean $\Delta m $ is lower"352Light curves of three intermediate polars TGR J15091-6619. TOR J16500-3307 aud ICR J17195-1100 Were obtained with the SAAO αι aud 0.7642 telescopes aud the UCT CCD (this is a frame-transter CCD. miplviug that there is no dead time between exposures).,"Light curves of three intermediate polars IGR J15094-6649, IGR J16500-3307 and IGR J17195-4100 were obtained with the SAAO 1-m and 0.76-m telescopes and the UCT CCD (this is a frame-transfer CCD, implying that there is no dead time between exposures)."353 Detailed description of these data is preseuted in Pretorius(2009)., Detailed description of these data is presented in \cite{pretorius09}.354 Ileh tiuc-resolution observatiows of SSer were taken at the prime focus of the newly available 10-14 class (SALT:Buckley.Swart&Aleiring2016:οὭωμουimeetal.66) curing its performance verification phase οji 2008 May. 13 aud June 12.," High time-resolution observations of Sgr were taken at the prime focus of the newly available 10-m class \citep[SALT;][]{Buck06,Dono06} during its performance verification phase on 2008 May 13 and June 12."355 Observations were done with the iuaeiugC»c» camera (SALTICAN:O'Doroghuectal.2005.x200) dn «ot mode (O'Donoghueetal.2006).," Observations were done with the imaging camera \citep[SALTICAM;][]{Dono03,Dono04} in `slot mode' \citep{Dono06}."356. Tιο SALTICAM is a CCD mosaic of two 20184 1102. pixels CCDs with full nuage size of about 6\9.6 arcu. but t1ο noluinal scleice field is 8 arcuunm.," The SALTICAM is a CCD mosaic of two $\times$ 4102 pixels CCDs with full image size of about $\times$ 9.6 arcmin, but the nominal science field is 8 arcmin."357 In “slot mode’. a mask with a «lot o the width of 20 arcsec (11Lo unbiuned rows) is projectiie on he CCDs.," In 'slot mode', a mask with a slot of the width of 20 arcsec (144 unbinned rows) is projecting on the CCDs."358 Oulv the ceitral 72 rows reccive all he ight., Only the central 72 rows receive all the light.359 The slot extends across the fu] 58 arcmin iuf16 horizonal direction., The slot extends across the full 8 arcmin in the horizontal direction.360 After cach exposure. onlv 111 LOWS are ransterred across the framie-transter boundary. which srortens the vertical clocking overheac to Ll τι».," After each exposure, only 144 rows are transferred across the frame-transfer boundary, which shortens the vertical clocking overhead to 14 ms."361 The iustrluneit thus provided a sequence οf images ὃ arcu ong.|mt only 20 arcsec wide.," The instrument thus provided a sequence of images 8 arcmin long, but only 20 arcsec wide."362 These 1nages were split iuto four distinct sections. correspoudine to the four readout amplifiers of the two CCDs.," These images were split into four distinct sections, corresponding to the four readout amplifiers of the two CCDs."363 Total racking time was ~27T00 sec on Aav 13 and 1300 sec ou June 12., Total tracking time was $\sim$ 2700 sec on May 13 and $\sim$ 1300 sec on June 12.364 The observatious were mace with the B filter., The observations were made with the $B$ filter.365 In order to facilitate the shortest possible exposures. «6 Mune. or 1.2 arcsec/pixcl was used.," In order to facilitate the shortest possible exposures, $\times$ 6 binning, or 1.2 arcsec/pixel was used."366 The exposure time for each nuage was 112 ms (the time to move each image behind the mask. during which iuage smncaring takes place. is Jtans).," The exposure time for each image was 112 ms (the time to move each image behind the mask, during which image smearing takes place, is 14 ms)."367 The observed data were all referenced to the nearN*ostar of similar magnitude. 63 arcsec to the cas and 7 arcsec south of the target star.," The observed data were all referenced to the nearby star of similar magnitude, 63 arcsec to the east and 7 arcsec south of the target star."368 The comparison star aud the target star were nuaged within the same amplificr., The comparison star and the target star were imaged within the same amplifier.369 No fiat-Beld calibration frames were obtained (the calibration svsteni was not available at the time of the observations) aud the frames were simply subtracted with overscaured pixels frou cach row., No flat-field calibration frames were obtained (the calibration system was not available at the time of the observations) and the frames were simply bias-subtracted with overscanned pixels from each row.370 The data reduction of “slot mode” was done with the the user package., The data reduction of “slot mode” was done with the the user package.371 This tool is primarily written iu python/PyRAF with some additional IRAFcode?., This tool is primarily written in python/PyRAF with some additional IRAF.372. Analysis of the liehteurves obtained with SALTICAAT in the clescribed mode revcaled problems iu power spectra of celestial sources at Fourier frequencies higher than lll. while at lower frequeucies the results remained reliable.," Analysis of the lightcurves obtained with SALTICAM in the described mode revealed problems in power spectra of celestial sources at Fourier frequencies higher than $\sim$ 1Hz, while at lower frequencies the results remained reliable."373 Therefore iun our subsequent analysis we have not considered power spectra of the considered source at Fourier frequency higher than 0.7 Iz., Therefore in our subsequent analysis we have not considered power spectra of the considered source at Fourier frequency higher than 0.7 Hz.374 Power spectra of V1223 Ser obtained iun different observations were virtually incdistineuishable from each other. therefore we averaged thon.," Power spectra of V1223 Sgr obtained in different observations were virtually indistinguishable from each other, therefore we averaged them."375the Drake&Testa(2005) Ne/O ratio.,the \citet{Drake.Testa:05} Ne/O ratio.376 Another important clue to the solar neon content is provided by past coronal Χαν Ne/O measurements., Another important clue to the solar neon content is provided by past coronal X-ray Ne/O measurements.377 Results reviewed by Drake&Testa(2005) showed considerable scatter. and an extensive study of active regions has demonstrated variations in (he coronal Ne/O ratio bv [actors of more than 2 (Mcelxenzie&Feklhnan1992).," Results reviewed by \citet{Drake.Testa:05} showed considerable scatter, and an extensive study of active regions has demonstrated variations in the coronal Ne/O ratio by factors of more than 2 \citep{McKenzie.Feldman:92}."378. Such variations indicate coronal fractionation of Ne relative to O bv mechanisms that are not vet firmly identified or understood. and cast serious doubt on the validity of conclusions based on a simple average ol observed values to infer the Ne content of the Sun.," Such variations indicate coronal fractionation of Ne relative to O by mechanisms that are not yet firmly identified or understood, and cast serious doubt on the validity of conclusions based on a simple average of observed values to infer the Ne content of the Sun."379 The variation in observed solar Ne/O ratios motivates re-examination of existing observations using modern atomic data lor futher insights into the behaviour of Ne in the solar outer abmosphere., The variation in observed solar Ne/O ratios motivates re-examination of existing observations using modern atomic data for further insights into the behaviour of Ne in the solar outer atmosphere.380 Unfortunately. some earlier studies of Ne in the solar X-ray spectrum cid nol list spectral line [Iuxes. rendering re-analysis somewhat diffieult.," Unfortunately, some earlier studies of Ne in the solar X-ray spectrum did not list spectral line fluxes, rendering re-analysis somewhat difficult."381 Observations made between 1979 March 23 anc November 29 bx the SOLEX spectrometer on the US Department of Delense PT3-1 satellite analysed by Melxenzie&Feldman(1992).. and the recent study by Schmelzetal.(2005) of (SAINI) Flat Crvstal Spectrometer (FCS) observations obtained between 1986 May. 20 and 1957 December 18 are (wo exceptions (hat form the basis of this work.," Observations made between 1979 March 23 and November 29 by the SOLEX spectrometer on the US Department of Defense P78-1 satellite analysed by \citet{McKenzie.Feldman:92}, and the recent study by \citet{Schmelz.etal:05b} of (SMM) Flat Crystal Spectrometer (FCS) observations obtained between 1986 May 20 and 1987 December 18 are two exceptions that form the basis of this work."382" The analvsis presented here is similar to (hat originally applied to Ne and O lines bv Actonetal.(1975).. ancl duplicated in suecessive papers including Melxenzie&Feld-man(1992) and Schmelzetal.(2005) (see also Drake&Testa 20091): Ne/O abundance ratios are inferred. from observed ratios of the intensities of the transitions Ne IX. AL3.47 132p!P,—1s?tS) and O VILE AT8.97 2p?Pus4;αυ and comparison with theoretical predictions."," The analysis presented here is similar to that originally applied to Ne and O lines by \citet{Acton.etal:75}, and duplicated in successive papers including \citet{McKenzie.Feldman:92} and \citet{Schmelz.etal:05b} (see also \citealt{Drake.Testa:05}) ): Ne/O abundance ratios are inferred from observed ratios of the intensities of the transitions Ne IX $\lambda 13.47$ $1s2p\, ^1P_1 \rightarrow3831s^2\, ^1S_0$ and O VIII $\lambda 18.97$ $2p\, ^2P_{3/2,1/2}384\rightarrow 1s\, ^2S_{1/2}$ and comparison with theoretical predictions."385 Line intensities were taken [rom Table 2 of Mcelxenzie&Feldman(1992) ancl Table 2 ol Schmelzetal.(2005).., Line intensities were taken from Table 2 of \citet{McKenzie.Feldman:92} and Table 2 of \citet{Schmelz.etal:05b}.386 One possible complication to the analvsis of (he Ne EX. A13.47 line is (he potential presence of Fe XIX blends (see.e.g.Nessetal.2003): both MelxenzieFelkhnan(1992) and Schinelzetal.(2005). (take pains to note this and assess that for their non-ILaring spectra such blends are negligible.," One possible complication to the analysis of the Ne IX $\lambda 13.47$ line is the potential presence of Fe XIX blends \citep[see, e.g.][]{Ness.etal:03}; both \citet{McKenzie.Feldman:92} and \citet{Schmelz.etal:05b} take pains to note this and assess that for their non-flaring spectra such blends are negligible."387 This assessment is verified below when specific account of the Fe XIX contribution to Ne IN. AT3.4T is taken., This assessment is verified below when specific account of the Fe XIX contribution to Ne IX $\lambda 13.47$ is taken.388 The theoretical Ne IX/O VIII line ratio has significant temperature dependence aud, The theoretical Ne IX/O VIII line ratio has significant temperature dependence and389r=22.0 should allow detection of targets with a S/N of 10. kr=19.5 will have S/N of 100.,"$r=22.0$ should allow detection of targets with a $S/N$ of 10, $r=19.5$ will have $S/N$ of 100."390 Table 2. shows the distance oul to which certain fiducial EB (vpes can be detected., Table \ref{EB_table_interesting_targets} shows the distance out to which certain fiducial EB types can be detected.391 For exaniple. a pair of eclipsing M2 dwarls will be detected out to 1 kpe with S/N of 100.," For example, a pair of eclipsing M2 dwarfs will be detected out to 1 kpc with $S/N$ of 100."392 We can estimate the munber of EBs that LSST will be able to fully characterize., We can estimate the number of EBs that LSST will be able to fully characterize.393 Experience modeling ED light curves shows that S/N~10 per data point (vpically sullices for the determination of physical and geometric parameters {ο a lew percent., Experience modeling EB light curves shows that $S/N \sim 10$ per data point typically suffices for the determination of physical and geometric parameters to a few percent.394 Gaia will observe 1 billion stars down to r20.5 over the whole sky (?7).., Gaia will observe $\sim$ 1 billion stars down to $r\sim20.5$ over the whole sky \citep{lindegren2008}.395 We can expect that LSsT will observe 70.5 billion stars to this same depth in the southern hemisphere: scaling this lor~22.0 vields ~2 billion stars., We can expect that LSST will observe $\sim$ 0.5 billion stars to this same depth in the southern hemisphere; scaling this to $r \sim 22.0$ yields $\sim$ 2 billion stars.396 Extrapolating the results [romAepler (2). Chat observed 18179 EBs in the sample of 156.000 stars (1.27)). the LSST sample will contain ~24 million EBs with S/N>10.," Extrapolating the results from \citep{prsa2011} that observed 1879 EBs in the sample of 156,000 stars ), the LSST sample will contain $\sim$ 24 million EBs with $S/N \geq 10$."397 The average detection rate (see section 3.3)) for EBs over all periods will be around2854.. bringing the total nunber to ~6.7 million EBs.," The average detection rate (see section \ref{sec:detrate}) ) for EBs over all periods will be around, bringing the total number to $\sim$ 6.7 million EBs."398 Roughly of those will have components of similar luminosities (double-limed svstemis). vielding 71.7 million EBs with S/N>10 for readx detailed modeling.," Roughly of those will have components of similar luminosities (double-lined systems), yielding $\sim$ 1.7 million EBs with $S/N \geq 10$ for ready detailed modeling."399 The success rate seems low. but the following needs to be taken into consideration:," The success rate seems low, but the following needs to be taken into consideration:"400profile midpoint) and SI which arrive late.,profile midpoint) and 81 which arrive late.401" H£ we take the rate between phases 0.18 and 0.78 in Fig 5. as ""background? we estimate that no more than 3 pulses are false positives in each of the early and late phase windows.", If we take the rate between phases 0.18 and 0.78 in Fig \ref{loc2} as `background' we estimate that no more than 3 pulses are false positives in each of the early and late phase windows.402 Phe carly arriving eiants have a mean flux density of 10.43 mv. higher than the 7.5 mJy for the mean flux density of the late arrivals.," The early arriving giants have a mean flux density of 10.4 mJy, higher than the 7.5 mJy for the mean flux density of the late arrivals."403 The three strongest eiants from the sample are part ofthe carly eroup., The three strongest giants from the sample are part of the early group.404 Fig ο shows the unbinnecd cumulative probability distribution of our giant pulses on a log-og plot., Fig \ref{logns} shows the unbinned cumulative probability distribution of our giant pulses on a log-log plot.405 A Ixolmogorov-Smirnoy comparison of the early ancl late giants eives a probability that these are drawn from the same distribution., A Kolmogorov-Smirnov comparison of the early and late giants gives a probability that these are drawn from the same distribution.406 Accordingly we have used a Bavesian approach (7) to fit a power law clistribution to the carly and. [ate giants separately. as well as to the combined set of giant »ulses.," Accordingly we have used a Bayesian approach \cite{whe04} to fit a power law distribution to the early and late giants separately, as well as to the combined set of giant pulses."407 These best-fit power laws are plotted on Figure 6.., These best-fit power laws are plotted on Figure \ref{logns}.408 'To facilitate comparison with other pulsars we express these giant pulse intensity clistributions in units of20.- the average subse [lux Say.=Sep(2068S)aca). since this is often aken as a threshold defining giant pulse emission.," To facilitate comparison with other pulsars we express these giant pulse intensity distributions in units of $20\times$ the average pulse flux $S_{20\times} = S_{GP}/(20\langle S_{{\rm 1.4GHz}} \rangle )$, since this is often taken as a threshold defining giant pulse emission."409 Lor POLScH)=500 pJv., For $20\langle S_{{\rm 1.4GHz}} \rangle= 500$ $\mu$ Jy.410 With these units. he best-fit cumulative probability. intensity distributions are The limited statistics ancl refractive variations imply a 20 per cent uncertainty in the normalizations.," With these units, the best-fit cumulative probability intensity distributions are The limited statistics and refractive variations imply a $\sim20$ per cent uncertainty in the normalizations."411 bor comparison. the Crab giant pulses have f.—610S477Du (at OS Cllz: Lundgren et al 1995) while PSR D1937|21 shows giant pulses with f.=1.10sah (at 1.4 Gllz. Ixinkhabwala Thorsett 2000)).," For comparison, the Crab giant pulses have $f_> = 6 \times 10^{-3} S_{20\times}^{-2.3}$ (at 0.8 GHz; Lundgren et al 1995) while PSR B1937+21 shows giant pulses with $f_> = 1 \times 10^{-5} S_{20\times}^{-1.8}$ (at 1.4 GHz, Kinkhabwala Thorsett \nocite{kt00}) )."412 Note that the giant pulse distributions in show some evidence for a break to a [latter index for S«7 pdx. but with a limited. dynamic range it is cilliculty to quantify the significance of such a break.," Note that the giant pulse distributions in show some evidence for a break to a flatter index for $S < 7~\mu$ Jy, but with a limited dynamic range it is difficulty to quantify the significance of such a break."413 There is no evidence of a cut-olf at. large. pulse. amplitudes., There is no evidence of a cut-off at large pulse amplitudes.414 As noted below there is some evidence for refractive interstellar scintillation modulation of both the integrated. profile [ux and the giant pulse rate., As noted below there is some evidence for refractive interstellar scintillation modulation of both the integrated profile flux and the giant pulse rate.415 We do not have sullicicnt signal-o-noise to calibrate this modulation., We do not have sufficient signal-to-noise to calibrate this modulation.416 Lt does not strongly allect the fitted pulse distributions. although we note that areest (wo (carly) giant pulses were seen at periods of low integrated Dux.," It does not strongly affect the fitted pulse distributions, although we note that largest two (early) giant pulses were seen at periods of low integrated flux."417 Correction for this elfect would. tend. to urther Dlatten the early giant pulse distribution., Correction for this effect would tend to further flatten the early giant pulse distribution.418 The largest pulse (from. the carly group) has an equivalent continuum Ilux of 64 midy. some the mean lux ancl the integrated [ux density for this particular observing epoch.," The largest pulse (from the early group) has an equivalent continuum flux of 64 mJy, some $\times$ the mean flux and $\times$ the integrated flux density for this particular observing epoch."419 Overall. we obtain a el lO mJy pulse every wo hours. about equally from the two components.," Overall, we obtain a $\gapp$ 10 mJy pulse every two hours, about equally from the two components."420 We expect one ο 100 mv. pulse in ~9 d of observing: this will ikelv occur from the carly component., We expect one $\gapp$ 100 mJy pulse in $\sim$ 9 d of observing; this will likely occur from the early component.421 Armed with a timinge solution we can now sum the ata together. to. determine. the integrated: profile of, Armed with a timing solution we can now sum the data together to determine the integrated profile of.42269.. Fig 7 shows the profile after summation of 1e data taken in 2003 August., Fig \ref{integprof} shows the profile after summation of the data taken in 2003 August.423 Phe pulsar is clearly detected with an integrated. Hux density of 24 pdx., The pulsar is clearly detected with an integrated flux density of 24 $\mu$ Jy.424 We also formed 1e integrated prolile for each of the other observing sessions from 2003 and from 2001., We also formed the integrated profile for each of the other observing sessions from 2003 and from 2001.425 With less integration time. the signal to noise ratio is not as good as in Fig 7. but the flux densities we obtain are 16 peJv (2001). 19 μὴν (Alay/June 2003) and 20 pdx (Sept. 2003).," With less integration time, the signal to noise ratio is not as good as in Fig \ref{integprof} but the flux densities we obtain are 16 $\mu$ Jy (2001), 19 $\mu$ Jy (May/June 2003) and 20 $\mu$ Jy (Sept 2003)."426 We note that the variability in these values is consistent. with the eiant pulse rate which also varies slightly from session to session., We note that the variability in these values is consistent with the giant pulse rate which also varies slightly from session to session.427 There is therefore some evidence for variability in the [lux density with a, There is therefore some evidence for variability in the flux density with a428well.,well.429 Figure 2 shows the evolution of the radiative fluxes arising from different temperature components in the beam and no-beam cases., Figure 2 shows the evolution of the radiative fluxes arising from different temperature components in the beam and no-beam cases.430 In general. all of the radiative fluxes fluctuate with time.," In general, all of the radiative fluxes fluctuate with time."431" Nevertheless for the case of D,=1G and πο>=1.3 km/s. the electron beam ol Fi,=107 erg ? !1 gives rise to only tiny effects on the radiations compared to the no-beam case. as expected from the snapshot structures shown in Figure 1."," Nevertheless for the case of $B_*=1$ G and $<v_\perp>=1.8$ km/s, the electron beam of $F_{bi}=10^2$ erg $^{-2}$ $^{-1}$ gives rise to only tiny effects on the radiations compared to the no-beam case, as expected from the snapshot structures shown in Figure 1."432" In contrast. the case of F5;=10! erg ? ! corresponding to 2B,=5G exhibits noticeable differences from (he no-beam case. which can also be expected [rom (he snapshot structures in Figure 1."," In contrast, the case of $F_{bi}=10^4$ erg $^{-2}$ $^{-1}$ corresponding to $B_*=5$ G exhibits noticeable differences from the no-beam case, which can also be expected from the snapshot structures in Figure 1."433" The raciative fluxes from the hot chromosphere (7000-200001x) and the transition region 5x ΠΙΑ) are ereatlv enhanced by the warm region in the vicinity of r=1.001—1.01R,.", The radiative fluxes from the hot chromosphere (7000-20000K) and the transition region $5\times 10^5$ K) are greatly enhanced by the warm region in the vicinity of $r=1.001-1.01R_*$.434" since (he warm region is on (he average denser (han the usual chromosphere and transition region in rZ1.003/2, (see Figure 1). the radiative flux corresponding to the chromospheric temperature is on average increased by a [actor of ~10 and (he radiative Εαν associated with the (ransiGion-region temperature is intensilied by a factor of ~100."," Since the warm region is on the average denser than the usual chromosphere and transition region in $r\gtrsim 1.003 R_*$ (see Figure 1), the radiative flux corresponding to the chromospheric temperature is on average increased by a factor of $\sim10$ and the radiative flux associated with the transition-region temperature is intensified by a factor of $\sim100$."435" :The bottom panel of. Figure. i2 shows that the beam of. £j;=10! erg >7s + associated. with D,= 5G is able to enhance the radiation [rom the hot eas of T>5x10? occasionally by a factor of =100.", The bottom panel of Figure 2 shows that the beam of $F_{bi}=10^4$ erg $^{-2}$ $^{-1}$ associated with $B_*=5$ G is able to enhance the radiation from the hot gas of $T>5\times 10^5$ K occasionally by a factor of $\gtrsim 100$.436 This radiation is mainlv from the warm region that intermittently develops around r=1.001—1.00322., This radiation is mainly from the warm region that intermittently develops around $r=1.001-1.003R_*$.437 With this beam heating. the temperature of the warm region can sometimes go up and down between LO? and LO“. This temperature variability is due (ο the thermal instability. (Landini&Monsignori-Fossi1990:Suzuki2007)...," With this beam heating, the temperature of the warm region can sometimes go up and down between $10^5$ K and $10^6$ K. This temperature variability is due to the thermal instability \citep{LM90,Suzuki07}."438 The electron beam is energetic enough (o continuously heat up the warm region. while the wave dissipation heats it up in a stochastic manner.," The electron beam is energetic enough to continuously heat up the warm region, while the wave dissipation heats it up in a stochastic manner."439 A small change of the wave heating rate (riggers violent [Inctuations of temperature in the thermally unstable regime of 10. As is shown in Figure 2.. the thermal response of the stellar atmosphere to the onset of the beam-heating to evolve to the hot state is nearly instantaneously.," A small change of the wave heating rate triggers violent fluctuations of temperature in the thermally unstable regime of $10^5 < T < 10^6$ K. As is shown in Figure \ref{fig2}, the thermal response of the stellar atmosphere to the onset of the beam-heating to evolve to the hot state is nearly instantaneously."440 In 82.1... we estimated the duration of the electron beam as 100 min. which is rather short in comparison with (he simulation (ime presented in Figure 2..," In \ref{sec:ebi}, we estimated the duration of the electron beam as $\sim 100$ min, which is rather short in comparison with the simulation time presented in Figure \ref{fig2}. ."441 To estimate the area of the hot spot. it is equally important to examine how long the hot atmosphere can be sustained after the electron beam passes bv.," To estimate the area of the hot spot, it is equally important to examine how long the hot atmosphere can be sustained after the electron beam passes by."442" Figures 2. 4 show the results for the D,= 5G case but the beam is switched olf at /=109 min.", Figures \ref{fig4} \ref{fig5} show the results for the $B_*=5$ G case but the beam is switched off at $t=109$ min.443 Figure 3.. which presents the snapshot structure at 219 min. namely 110 min after the beam stops. illustrates (hat while (he warm reeion in the chromosphere is still present as shown by the temperature prolile. its density has dropped to the original lower-densitv level.," Figure \ref{fig4}, which presents the snapshot structure at $219$ min, namely 110 min after the beam stops, illustrates that while the warm region in the chromosphere is still present as shown by the temperature profile, its density has dropped to the original lower-density level."444 Figure 4 shows that (he radiative fluxes from the hot chromosphere (lop) ancl (ransiGion regions (middle) have declined lor 22 50-60 minutes since the beam is switched off. which indicates (hat the atmosphere takes only a fraction of the heating time lo revert to the normal state in the absence of the electron beam.," Figure \ref{fig5}445 shows that the radiative fluxes from the hot chromosphere (top) and transition regions (middle) have declined for $\approx$ 50-60 minutes since the beam is switched off, which indicates that the atmosphere takes only a fraction of the heating time to revert to the normal state in the absence of the electron beam."446 Therefore. the size of the hot spot estimated in terms of (he projected area of theplanets magnetosphere along a flux," Therefore, the size of the hot spot estimated in terms of the projected area of theplanet's magnetosphere along a flux"4471997).,.448. IniGally. planetesinials erow slowly.," Initially, planetesimals grow slowly."449 As thev grow. dvnanmical friction ceireularizes the orbits of the largest objects: viscous stirring excites the orbits of the smallest objects.," As they grow, dynamical friction circularizes the orbits of the largest objects; viscous stirring excites the orbits of the smallest objects."450 Slow. orderly growth ends when the gravitational cross-sections of the largest objects exceed their geometric cross-sections.," Slow, orderly growth ends when the gravitational cross-sections of the largest objects exceed their geometric cross-sections."451 Because dvnamical Drietion is faster than accretion. the largest objects stay on circular orbits and grow faster ancl faster relative to the smallest planetesimals.," Because dynamical friction is faster than accretion, the largest objects stay on circular orbits and grow faster and faster relative to the smallest planetesimals."452 This runaway growth rapidly concentrates solid material into a few protoplanets(e.g..Greenberg1993:Kokubo&Ida./|. 2000).," This runaway growth rapidly concentrates solid material into a few protoplanets\citep[e.g.,][]{gre78,gre84,wet89,ws93,wei97,kl98,kok00}."453. During runaway growth. protoplanets stir the orbits of the leftover planetesimals.," During runaway growth, protoplanets stir the orbits of the leftover planetesimals."454 Stirring reduces gravitational focusing [actors and slows the growth of the protoplanets (Wetherill&1993:Ida&MakinolxenvonBromley2002;Ralikoy 2003c).," Stirring reduces gravitational focusing factors and slows the growth of the protoplanets \citep{ws93,ida93,kb02,raf03c}."455. Although protoplanets grow slowly. (hev grow faster than the leftover planetesimals aud. intermediate mass objects do.," Although protoplanets grow slowly, they grow faster than the leftover planetesimals and intermediate mass objects do."456 These large ‘oligarchs’ slowly clear their orbits of smaller objects and reach a Maxim ‘isolation mass (hat depends on (the initial surface density of solid material (LissauerLOST:Lissauer&Stewart1993:IXokuboIda1998:Ralikov2003€: ).," These large `oligarchs' slowly clear their orbits of smaller objects and reach a maximum `isolation mass' that depends on the initial surface density of solid material \citep{lis87,lis93,kok98,raf03c,gol04b}."457. Throughout oligarchic growth. protoplanets repel other oligarchs in the disk (IxokuboIda1996.1993.2000:IXominami&2002:Thommes.Duncan.Levison|2003 ).," Throughout oligarchic growth, protoplanets repel other oligarchs in the disk \citep{kok96,kok98,kok00,kom02,tho03}."458. Eventually. however. dvinamical interactions produce collisions and gravitational scattering among the oligarchs (Chambers|.2001:Kokubo&Ida|2002:Nominami—hommes.Duncan.&Levison2002:Ralikov2003€:Goldreich.LithwickSari 2004b).," Eventually, however, dynamical interactions produce collisions and gravitational scattering among the oligarchs \citep{cha01,kok02,kom02,tho02,raf03c,gol04b}."459. During (his phase of ‘chaotic growth. the largest oligarchs merge wilh aud clear away other oligarchs and smaller objects to become Full-f[ledeed planets (e.g..Goldreich. 2004a).," During this phase of `chaotic growth', the largest oligarchs merge with and clear away other oligarchs and smaller objects to become full-fledged planets \citep[e.g.,][]{gol04a}."460. Within this framework. the transition from oligarchic growth to chaotic growth is poorly understood.," Within this framework, the transition from oligarchic growth to chaotic growth is poorly understood."461 Although analytic estimates provide a guide to the late evolutionary stages (e.g..Goldreich.Lithwiek&Sari 2004a).. numerical calculations are necessary to test the basic theory ancl to derive the end states and Gmescales as a function of initial conditions (e.g.Nominami&Ida2002).," Although analytic estimates provide a guide to the late evolutionary stages \citep[e.g.,][]{raf03c,gol04a}, numerical calculations are necessary to test the basic theory and to derive the end states and timescales as a function of initial conditions \citep[e.g.][]{kom02}."462. Published ealeulations do not test this regime of the theory in much detail., Published calculations do not test this regime of the theory in much detail.463 Pure coagulation and simple hvbricl caleulations cannot follow the transition accurately (Wetherill&Stewart1993:Weidenschillingetal.1997).," Pure coagulation and simple hybrid calculations cannot follow the transition accurately \citep{ws93,wei97}."464. Direct n-body calculations with planetesimals are computationally expensive and tend to locus on evolution during oligarchic growth. when orbits of individual objects are easier to track (Alexander&Aenor1905:IXokuboIda2002.andreferences therein)..," Direct $n$ -body calculations with planetesimals are computationally expensive and tend to focus on evolution during oligarchic growth, when orbits of individual objects are easier to track \citep[][and references therein]{ale98,kok02}. ."465 Direct. n-body caleulations without planetesinals follow dynamical interactions after (he (transition and, Direct $n$ -body calculations without planetesimals follow dynamical interactions after the transition and466"andoy: X-ray bright QSOs (with small 6,44)) tend to have flat (lind)a.",and; X-ray bright QSOs (with small ) tend to have flat (hard).467.. At these and sliehtlv lower huninosities (Sevtert ls. e.g.. Boller et al.," At these and slightly lower luminosities (Seyfert 1s, e.g., Boller et al."468 1996). ouly objects with flat Chard) have large FWHALID)).," 1996), only objects with flat (hard) have large )."469 Tf is representative. the combination of these two trends would predict that NB QSOs should on average have woader Cluission lines. counter to the overall treuds sccu in our samples.," If is representative, the combination of these two trends would predict that XB QSOs should on average have broader emission lines, counter to the overall trends seen in our samples."470" Of the emission lines to which we cau fit wo componcuts(παν,CIV...MgIL. aud 1). it is the line for which the percentage of total line flux in the jurow line coniponeut may beleryer in the NF composite."," Of the emission lines to which we can fit two components, and ), it is the line for which the percentage of total line flux in the narrow line component may be in the XF composite."471 Since appears to be unique among the larger variety of emission ines studied here. it is clear that auax not be the best. or only represcutative of BLR line widths.," Since appears to be unique among the larger variety of emission lines studied here, it is clear that may not be the best, or only representative of BLR line widths."472 Fitting the line iu our optical composites with a Cassia component leads to similar conchisious to these previous studies (nie. that EWIIMEITL2)) is larecr in the AF composite).," Fitting the line in our optical composites with a Gaussian component leads to similar conclusions to these previous studies (i.e., that ) is larger in the XF composite)."473 ILowewver. fitting narrow and broad line colmpoucuts separately reveals that has a stronger broad componcut in the NB composite. and a larger broad/uarrow line ratio. while the actual FWIINM is similar to the ΝΕ composite.," However, fitting narrow and broad line components separately reveals that has a stronger broad component in the XB composite, and a larger broad/narrow line ratio, while the actual FWHM is similar to the XF composite."474 This reveals that single Gaussian fits to compound diues must be interpreted with caution., This reveals that single Gaussian fits to compound lines must be interpreted with caution.475 Furthermore. we also caution that the huuinositv ranges for these (PC and LBQS) salples are digjoiut.," Furthermore, we also caution that the luminosity ranges for these (PG and LBQS) samples are disjoint."476 Towever. since and not seclus to have the more fundamental correlation with both and iu Laor et al. (," However, since and not seems to have the more fundamental correlation with both and in Laor et al. ("4771997). we highlight the need for correlation of with a wider range of eunission line micasurements.,"1997), we highlight the need for correlation of with a wider range of emission line measurements."478" The parameters enipliasized here. and narrow line cussion. also appear to be linked to the following quautitics: X-ray spectral slopea... streneth. luminosity. radio loudness (e.gο, Laor ct al."," The parameters emphasized here, and narrow line emission, also appear to be linked to the following quantities: X-ray spectral slope, strength, luminosity, radio loudness (e.g., Laor et al."479 1997. Lawrence et al.," 1997, Lawrence et al."480 1997. Cieen ct al.," 1997, Green et al."481 1995. (021.," 1995, BG92)."482 We sueeest here that the as vet iivsterious plivsical link )etween these diverse properties is intimately related to Heh velocity. outflowing winds near the nucleus that. bv absorbing the intrinsic nuclear coutimuun. strongly affect he radiation observed at larger distances.," We suggest here that the as yet mysterious physical link between these diverse properties is intimately related to high velocity, outflowing winds near the nucleus that, by absorbing the intrinsic nuclear continuum, strongly affect the radiation observed at larger distances."483 The continua nupineine on the NLR is closest to that received by distant observers like ourselves. but is quite different from hat arriving at the BLR.," The continuum impinging on the NLR is closest to that received by distant observers like ourselves, but is quite different from that arriving at the BLR."484 Material in the DLR itself nay reprocess the intrinsic contiuuun issuing from the uushrouded QSO nucleus., Material in the BLR itself may reprocess the intrinsic continuum issuing from the unshrouded QSO nucleus.485 Ouly some of the correlations between measured line and coutinuun parameters are intriusic. and others siuplv add dispersion to more primary correlations.," Only some of the correlations between measured line and continuum parameters are intrinsic, and others simply add dispersion to more primary correlations."486 The principal physical processes niust be extracted from the principle observational eigeuvectors iu a multivariate. multiwavelength approach. with careful attention to contimmiun slopes and detailed cussion line fits.," The principal physical processes must be extracted from the principle observational eigenvectors in a multivariate, multiwavelength approach, with careful attention to continuum slopes and detailed emission line fits."487 We are accunmmlatiug a lieh-quality homogeneous database inchiding all this information for a large sample of QSOs. primarily from the ROSAT and UST archives.," We are accumulating a high-quality homogeneous database including all this information for a large sample of QSOs, primarily from the ROSAT and HST archives."488 We believe that multiwaveleugth studies such as this show promise for sienificant advances in our understanding., We believe that multiwavelength studies such as this show promise for significant advances in our understanding.489 Thauks to Marianne Vestergaard for performing the iron subtraction. to Craig Foltz for the LBOS spectra. Ken Lauzetta for the IVE QSO atlas. aud Todd Boroson for the optical template.," Thanks to Marianne Vestergaard for performing the iron subtraction, to Craig Foltz for the LBQS spectra, Ken Lanzetta for the IUE QSO atlas, and Todd Boroson for the optical template."490" The author eratefully acknowledges support provided by NASA through Grant NACG5-1253.πει and Coutract NASS-39073 (ASC). as well as 201-92A awarded by the Space Telescope Scicuce Institute. which is operated by the Association of Universities for RE in Astronomy. Πιο, under NASA contract NAS5-3555."," The author gratefully acknowledges support provided by NASA through Grant NAG5-1253, and Contract NAS8-39073 (ASC), as well as HF-1032.01-92A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555."491 οThis research has made use of the NASA/TPAC Database (NED) which i9 operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Achuinistration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."492intensity of the X-ray [Iux.,intensity of the X-ray flux.493 Oue stream crossing occurs just before periastron (orbital phase e-0.93)., One stream crossing occurs just before periastron (orbital phase $\sim$ 0.93).494 GN301-2 is nearing ils peak speed to overtake the stream. but (he stream is near its highest clensity and lowest racial velocity resulting in a large increase in luminosity.," GX301-2 is nearing its peak speed to overtake the stream, but the stream is near its highest density and lowest radial velocity resulting in a large increase in luminosity."495 The second stream erossing is at orbital phase ~0.55., The second stream crossing is at orbital phase $\sim$ 0.55.496 As GN301-2 approaches apastron it slows {ο ils most leisurely pace. so (he stream is able {ο overtake the neutron star.," As GX301-2 approaches apastron it slows to its most leisurely pace, so the stream is able to overtake the neutron star."497 Since (he racial wind speed is highest aud stream density lowest al apastron. a significantly lower peak in luminosity occurs.," Since the radial wind speed is highest and stream density lowest at apastron, a significantly lower peak in luminosity occurs."498 Physically one expects the stream: width to increase with radial distance from the companion star. due to expansion of the higher density. overpressured stream in the lower density surrounding stellar wind.," Physically one expects the stream width to increase with radial distance from the companion star, due to expansion of the higher density, overpressured stream in the lower density surrounding stellar wind."499 Here we use a Gaussian densitv profile with variable width., Here we use a Gaussian density profile with variable width.500 The angular width (7) is taken as a power law function of distance from the center of the companionstar: o(r)=σι)Γρ). with ση the width normalized αἱ periastron distance ry.," The angular width $\sigma$ ) is taken as a power law function of distance from the center of the companionstar: $\sigma (r)= \sigma_0 (r/r_{per})^{-\kappa}$, with $\sigma_0$ the width normalized at periastron distance $r_{per}$."501 &«0 corresponds to a stream with increasing physical width as the stream propagates oulwarel., $\kappa<0$ corresponds to a stream with increasing physical width as the stream propagates outward.502 We ensure mass conservation bx emploving the continuity. equation. so the density of the stream varies with r depending on the value of &.," We ensure mass conservation by employing the continuity equation, so the density of the stream varies with $r$ depending on the value of $\kappa$."503" The angular width of the stream (viewed [rom the companion star) depends on (r/ry,)loso H7—] corresponds lo a stream wilh decreasing angular width as the stream propagates outward."," The angular width of the stream (viewed from the companion star) depends on $(r/r_{per})^{-\kappa-1}$, so $\kappa>-1$ corresponds to a stream with decreasing angular width as the stream propagates outward."504 GA301-2 has a significant. absorption bv its stellar wind in soft. A-ravs with column densities several times 107ni 27., GX301-2 has a significant absorption by its stellar wind in soft X-rays with column densities several times $10^{23}$ $^{-2}$.505 This shows up well in Fig., This shows up well in Fig.506 1 above: Band 1 and Band 2 data are affected siegnilicantlv by absorption but. Band 3 (5-12.1 keV) is mostly [ree of absorption elfects since the photoelectric cross-section is very small above 3-4 keV. This is also confirmed by the consisteney in shape of the Band 3 light eurve with the DATSE light curve (Ixohetal.1997).. although the RNTE/ ASAI light eurve here is of significantly better statistical eualitv.," 1 above: Band 1 and Band 2 data are affected significantly by absorption but Band 3 (5-12.1 keV) is mostly free of absorption effects since the photoelectric cross-section is very small above 3-4 keV. This is also confirmed by the consistency in shape of the Band 3 light curve with the BATSE light curve \citep{b13}, although the RXTE/ ASM light curve here is of significantly better statistical quality."507 Thus the 5-12.1 keV band flux is taken here to be proportional to the X-ray luminosity of the pulsar., Thus the 5-12.1 keV band flux is taken here to be proportional to the X-ray luminosity of the pulsar.508 The comparison of our wind plus stream models to the RNTE/ ASAI orbital light curve is made by 47 minimization using the non-linear conjugate gradient method., The comparison of our wind plus stream models to the RXTE/ ASM orbital light curve is made by $\chi^2$ minimization using the non-linear conjugate gradient method.509 For each minimizalion (i.e. fitting). some parameters were (taken as fixed parameters (stellar radius and system inclination) and the remaining parameters were taken as [ree parameters.," For each minimization (i.e. fitting), some parameters were taken as fixed parameters (stellar radius and system inclination) and the remaining parameters were taken as free parameters."510 The mass-radius constraints on WRAY 977 were discussed in detail in LO2., The mass-radius constraints on WRAY 977 were discussed in detail in L02.511 IIere we replot the radius contraints in Fig., Here we replot the radius contraints in Fig.512 5., 5.513 We use inclination rather (han mass as (he independent variable. since radius and inclination are eritical inputs to our model caleulations.," We use inclination rather than mass as the independent variable, since radius and inclination are critical inputs to our model calculations."514 Also we have extended the upper limit of T;jy to 22500 Ix. as suggested by the data of (1995).. and show the most relevant region of the radius vs. inclination diagram.," Also we have extended the upper limit of $T_{eff}$ to 22500 K as suggested by the data of \cite{b9}, , and show the most relevant region of the radius vs. inclination diagram."515 Allowable, Allowable516Gamma Ray Bursts (GRB) are catastrophic explosions of cosmological origin that illuminate the sky once or twice a day with relatively short. intense. flashes of y-rays on the order of the MeV and duration that ranges from 107* to about 10° seconds.,"Gamma Ray Bursts (GRB) are catastrophic explosions of cosmological origin that illuminate the sky once or twice a day with relatively short, intense, flashes of $\gamma$ -rays on the order of the MeV and duration that ranges from $10^{-3}$ to about $10^3$ seconds."517 Until the launch of Swift. the most widely accepted taxonomy of GRBs is the division between short-hard and long-soft bursts. after Dezalay (1992) and Kouveliotou (1993). who found in BATSE data a bimodal distribution in the burst duration defined as the time it takes the 90% of the flux to arrive (Too). with respect to the local time of the detector.," Until the launch of Swift, the most widely accepted taxonomy of GRBs is the division between short-hard and long-soft bursts, after Dezalay (1992) and Kouveliotou (1993), who found in BATSE data a bimodal distribution in the burst duration defined as the time it takes the $90\%$ of the flux to arrive $T_{90}$ ), with respect to the local time of the detector."518" These two distinct GRB distributions are separated by a minimum located at To,=2 s. This sharp division is supposed to be also due to a selection effect of the instrumentation onboard the satellite.", These two distinct GRB distributions are separated by a minimum located at $T_{90} \simeq 2$ s. This sharp division is supposed to be also due to a selection effect of the instrumentation onboard the satellite.519 In fact. new Swift observations permit to include in the classification scheme. as short bursts. also some distant events having apparently Tyo>2 s with respect to Swift’s proper time.," In fact, new Swift observations permit to include in the classification scheme, as short bursts, also some distant events having apparently $T_{90}>2$ s with respect to Swift's proper time."520 The determination of the redshift permits in fact to correct the Το by the relativistic effects (Cheetal..1997.Ruffini 2009)).," The determination of the redshift permits in fact to correct the $T_{90}$ by the relativistic effects \cite{che97,ruf09}) )."521 This suggests the need of a more solid classification of GRBs on the basis on a broader set of criteria. beyond the mere burst duration and the verification of the present taxonomy of long/short GRBs (Donaghyetal..2006.Zhang2007.Bloom2008.Belezynskietal.. 2008)).," This suggests the need of a more solid classification of GRBs on the basis on a broader set of criteria, beyond the mere burst duration and the verification of the present taxonomy of long/short GRBs \cite{don06,zha07,blo08,bel08}) )."522 Long and short GRBs are thought to have been generated by different. progenitors., Long and short GRBs are thought to have been generated by different progenitors.523 Anyway. different progenitors are thought to. produce ᾱ black hole accreting material from a dise. or from a torus with the emission of gamma rays (Ghirlandaetal.. 2009)).," Anyway, different progenitors are thought to produce a black hole accreting material from a disc or from a torus with the emission of gamma rays \cite{ghi09}) )."524 The basic model for long-duration GRBs. related to the catastrophic. release. of energy from the collapse of massive stars (Woolsey.2001.Fryer&Kalogera. 2001)). has received a strong support. from. the observations of their X (Gehrelsetal..2008.Nysewander 2005)). optical and radio counterparts. and the association with supernova detections. of which the GRB980425/SN1998bw association was the first clear example (Galamaetal..1998.Kulkarni.1998.vanParadis. 1999)).," The basic model for long-duration GRBs, related to the catastrophic release of energy from the collapse of massive stars \cite{woo01,fry01}) ), has received a strong support from the observations of their X \cite{geh08,nys08}) ), optical and radio counterparts, and the association with supernova detections, of which the GRB980425/SN1998bw association was the first clear example \cite{gal98,kul98,vpa99}) )."525 Thus. X-ray flashes. associated to long-duration GRBs. are probably produced by the highly-relativistic jets ejected in core-collapse supernova explosions.," Thus, X-ray flashes, associated to long-duration GRBs, are probably produced by the highly-relativistic jets ejected in core-collapse supernova explosions."526 The relativistic fireball model (Goodman.1986.Paezynski.1986.1990.Castro-Tiradoetal.. 2001)) accounts for a reasonable description of the observed afterglow spectrum. powered by the synchrotron emission of electrons accelerated in a relativistic shock with an estimated total energy budget roughly the same order of magnitude as that of supernovae Ib/c. The discovery of the slowly fading X-ray. from optical and radio afterglows of GRBs and the identification of host galaxies αί cosmological distances. gave further support to describe their progenitors as the product of the evolution of short-lived massive stars in different cosmological epochs. whose detection 1s limited by the BATSE threshold. that. sets the observational limit at distances slightly larger. than zo4 (Witersetab.1997.)Blain&Natarajan. 2000)).," The relativistic fireball model \cite{goo86,pac86,pac90,cas01}) ) accounts for a reasonable description of the observed afterglow spectrum, powered by the synchrotron emission of electrons accelerated in a relativistic shock with an estimated total energy budget roughly the same order of magnitude as that of supernovae Ib/c. The discovery of the slowly fading X-ray, from optical and radio afterglows of GRBs and the identification of host galaxies at cosmological distances, gave further support to describe their progenitors as the product of the evolution of short-lived massive stars in different cosmological epochs, whose detection is limited by the BATSE threshold that sets the observational limit at distances slightly larger than $z \sim 4$ \cite{wij97,bla00}) )."527 Now the new data from Swift passed the barrier of z=8 (Salvaterraetal.. 2009))., Now the new data from Swift passed the barrier of $z=8$ \cite{sal09}) ).528 Short GRBs are thought to be produced by highly - relativistic jets ejected in different. processes. as during neutron star - neutron star (NS-NS) or black hole - neutron star (BH-NS) binary mergers (Narayanetal.. 1992)). whose averaged redshift distribution. in the Swift-era. seems to be (2~1.0 (see e.g. Maghiocchettt. Ghirlanda Celotti. 2003. Tanvir et al.," Short GRBs are thought to be produced by highly - relativistic jets ejected in different processes, as during neutron star - neutron star (NS-NS) or black hole - neutron star (BH-NS) binary mergers \cite{nar92}) ), whose averaged redshift distribution, in the Swift-era, seems to be $\langle z \rangle \sim1.0$ (see e.g. Magliocchetti, Ghirlanda Celotti, 2003, Tanvir et al."529 2005. and Ghirlanda et al.," 2005, and Ghirlanda et al."530 2006)., 2006).531 Anyway. recent Swift and HETE-2 observations put in evidence the presence of clear similarities with the afterglows detected in the correspondence of long GRBs. thanks to the detection of X-ray and optical afterglows.," Anyway, recent Swift and HETE-2 observations put in evidence the presence of clear similarities with the afterglows detected in the correspondence of long GRBs, thanks to the detection of X-ray and optical afterglows."532 Other similarities found between the initial stage of the spectra of short and long GRBs together with the presence. in some cases. of X-ray flares. seem to indicate the presence of a common mechanism during the first few seconds (Barthelmyetetal..2005.Coward. 2007)).," Other similarities found between the initial stage of the spectra of short and long GRBs together with the presence, in some cases, of X-ray flares, seem to indicate the presence of a common mechanism during the first few seconds \cite{bar05,cow07}) )."533 Short GRBs have been observed mostly in elliptical galaxies. but even less frequently also 1n nearby irregular and in star-forming galaxies. confirming as progenitor the binary merging scenario (Meszaros.2006.Narayanetal..2001.Belezynskietal.. 2008)) also. through the mechanism of (Grindlay. 2006)).," Short GRBs have been observed mostly in elliptical galaxies, but even less frequently also in nearby irregular and in star-forming galaxies, confirming as progenitor the binary merging scenario \cite{mes06, nar01,bel08}) ) also through the mechanism of \cite{gri06}) )."534 Α similar. behavior Is expected for NS-BH binaries (Paezynski. 1991))., A similar behavior is expected for NS-BH binaries \cite{pac91}) ).535 The angular and spatial distributions of BATSE GRBs look quite isotropic (Briggsetal.. 1996)). with only very," The angular and spatial distributions of BATSE GRBs look quite isotropic \cite{bri96}) ), with only very"536In order to determine the characteristics of planets that reveal their existence in microlensing events. to model stellar. lens binaries. and to distinguish between these contigurations. efficient algorithms for ealeulating light curves of extended source stars due to microlensing of a binary lens system are required.,"In order to determine the characteristics of planets that reveal their existence in microlensing events, to model stellar lens binaries, and to distinguish between these configurations, efficient algorithms for calculating light curves of extended source stars due to microlensing of a binary lens system are required."537 For a single point-like lens. the magnification of a point source is an analytic function (2).. the magnification of a uniformly bright circular source can be expressed by means of elliptic integrals (?). and is therefore a one-dimensional integral over a semi-analytic function in general. which can be approximated by the product of the point-source magnification with a characteristic function of a single variable. depending on the source brightness profile (2)..," For a single point-like lens, the magnification of a point source is an analytic function \citep{Ein36}, the magnification of a uniformly bright circular source can be expressed by means of elliptic integrals \citep{WM94}, and is therefore a one-dimensional integral over a semi-analytic function in general, which can be approximated by the product of the point-source magnification with a characteristic function of a single variable, depending on the source brightness profile \citep{Gouldapprox}."538" For a binary lens. however. the complexity increases enormously,"," For a binary lens, however, the complexity increases enormously."539 Whereas a point source affected by a single point-mass lens always has two images. with the only exception of lens and source being perfectly aligned. binary lenses vield either three or five images depending on the source position relative to the lens. where three-image and five-image regions are separated by extended caustics.," Whereas a point source affected by a single point-mass lens always has two images, with the only exception of lens and source being perfectly aligned, binary lenses yield either three or five images depending on the source position relative to the lens, where three-image and five-image regions are separated by extended caustics."540 This implies that one cannot as easily integrate over the source position as in the single-lens case., This implies that one cannot as easily integrate over the source position as in the single-lens case.541 Moreover. the point-source magnification cannot be expressed as analytic or semi-analytic function. but needs to be determined by means of solving a fifth-order complex polynomial (2)..," Moreover, the point-source magnification cannot be expressed as analytic or semi-analytic function, but needs to be determined by means of solving a fifth-order complex polynomial \citep{WM95:fifth}."542 While a two-dimensional integration over the source is possible in principle. one carefully needs to keep track of the caustic and needs to deal with a divergent integrand as it is approached.," While a two-dimensional integration over the source is possible in principle, one carefully needs to keep track of the caustic and needs to deal with a divergent integrand as it is approached."543 While ray-shooting algorithms (22). overcome the need for determining the images by creating a magnification map based on mapping the image plane to the corresponding true source positions. summing the hits within a source still is a discretization of a two-dimensional integration.," While ray-shooting algorithms \citep{KRS:rayshooting,SchneiWei:AGN} overcome the need for determining the images by creating a magnification map based on mapping the image plane to the corresponding true source positions, summing the hits within a source still is a discretization of a two-dimensional integration."544 On the other hand. ? have found that contour plot routines offer a simple way for plotting the images of an extended source. which avoids the inversion of the lens equation.," On the other hand, \citet{SK87} have found that contour plot routines offer a simple way for plotting the images of an extended source, which avoids the inversion of the lens equation."545 ?. later showed that in addition to providing plots. the same technique ean be used to efficiently derive numerical values. and in particular to determine the magnification of extended sources as well as of the incurred astrometric shift in the centroid of their unresolved imges by applying Green's theorem (?)..," \citet{Do95:Num} later showed that in addition to providing plots, the same technique can be used to efficiently derive numerical values, and in particular to determine the magnification of extended sources as well as of the incurred astrometric shift in the centroid of their unresolved imges by applying Green's theorem \citep{Do98:NumSrc}."546 In fact. determining the image area from a contour plot is similar to the ray-shooting magnitication map approach in mapping an image plane grid to the respective source position. but rather than the magnitication being determined over the source area. an integral in the image plane is discretized.," In fact, determining the image area from a contour plot is similar to the ray-shooting magnification map approach in mapping an image plane grid to the respective source position, but rather than the magnification being determined over the source area, an integral in the image plane is discretized."547 In this paper. an algorithm is presented that determines the magnification from the image contour using an adaptive grid rather han a static one.," In this paper, an algorithm is presented that determines the magnification from the image contour using an adaptive grid rather than a static one."548 By recognizing the positions of the images of he source centre. which can be calculated by solving a order complex polynomial (23.. the fact that holes in the images must include the positions of the point-like deflectors. and by tinding all images stretching over critical curves using a parametric representation based on the classification of the topology of the critical curves and causties by ?.. it is ensured that all images for a binary lens are appropriately considered.," By recognizing the positions of the images of the source centre, which can be calculated by solving a fifth-order complex polynomial \citep{WM95:fifth}, the fact that holes in the images must include the positions of the point-like deflectors, and by finding all images stretching over critical curves using a parametric representation based on the classification of the topology of the critical curves and caustics by \citet{Erdl}, it is ensured that all images for a binary lens are appropriately considered."549 While Sect., While Sect.550 2. reviews the basic properties of gravitational, \ref{sec:gravlens} reviews the basic properties of gravitational551increases rapidly for / above 45°.,increases rapidly for $i$ above $^\circ$.552" Even with this consideration, however, the orbit parameter and mass errors are likely to be underesumated because of the non-physical assumpton of a circular orbit required in this very preliminary modeling."," Even with this consideration, however, the orbit parameter and mass errors are likely to be underestimated because of the non-physical assumption of a circular orbit required in this very preliminary modeling."553" From near infrared spectroscopy, Preibisch (1999) estimated a spectral type K2 + 2 for EC 95, and an extinction-corrected bolometric luminosity L = 60!M L.."," From near infrared spectroscopy, Preibisch (1999) estimated a spectral type K2 $\pm$ 2 for EC 95, and an extinction-corrected bolometric luminosity $L$ = $^{+30}_{-20}$ $L_\odot$."554 This was obtained assuming a distance to Serpens of310 pe; using the value of the distance obtained here increases the luminosity to an SL., This was obtained assuming a distance to Serpens of 310 pc; using the value of the distance obtained here increases the luminosity to $^{+54}_{-36}$ $L_\odot$.555".More recently, Doppmann et ((2005) estimated the effective temperature of EC 95 to be 4400!iD K. in agreement with the spectral type proposed by Preibisch (1999), and its luminosity (calculated assuming a distance of 259 pc) to be 23 L.."," More recently, Doppmann et (2005) estimated the effective temperature of EC 95 to be $^{+115}_{-57}$ K, in agreement with the spectral type proposed by Preibisch (1999), and its luminosity (calculated assuming a distance of 259 pc) to be 23 $L_\odot$."556" Scaled to the distance determined here, this would imply a luminosity of 59 L.., also consistent with the determination of Preibisch (1999),"," Scaled to the distance determined here, this would imply a luminosity of 59 $L_\odot$, also consistent with the determination of Preibisch (1999)."557 The extinction determined by Preibisch (1999) 1s Ay = 36 + 2: an alternative determination by Pontoppidan et (2004) yields Ay ~ 37., The extinction determined by Preibisch (1999) is $A_V$ = 36 $\pm$ 2; an alternative determination by Pontoppidan et (2004) yields $A_V$ $\sim$ 37.558 The mass and the age of EC 95 can be estimated by placing it on an HR diagram., The mass and the age of EC 95 can be estimated by placing it on an HR diagram.559" Using that method, Preibisch (1999) estimated a mass of about 4 M. and an age of 0.2!n Myr."," Using that method, Preibisch (1999) estimated a mass of about 4 $M_\odot$ and an age of $^{+0.2}_{-0.1}$ Myr."560 Pontoppidan et (2004) proposed a similar mass (3.5 M.). and a marginally larger age (0.4 Myr) because the shorter distance they used implies a correspondingly smallerluminosity.," Pontoppidan et (2004) proposed a similar mass (3.5 $M_\odot$ ), and a marginally larger age (0.4 Myr) because the shorter distance they used implies a correspondingly smallerluminosity."561" In any case, all these characteristics suggest that EC 95 is a very young precursor of à ~~ 4 M. star."," In any case, all these characteristics suggest that EC 95 is a very young precursor of a $\sim$ 4 $M_\odot$ star."562" Using our new distance determination, the position of EC 95 on the HR diagram moves up somewhat, and its position suggests a mass of about 5 M. , and an"," Using our new distance determination, the position of EC 95 on the HR diagram moves up somewhat, and its position suggests a mass of about 5 $M_\odot$ , and an"563one.,one.564 The combined PO+RS fit yields the best result with an absorption corrected luminosity of (3.912-0.55) erg ! for the assumed distance of 97 Mpe., The combined PO+RS fit yields the best result with an absorption corrected luminosity of $\pm$ erg $^{-1}$ for the assumed distance of 97 Mpc.565matter fraction at z=0 and Hubble constant.,matter fraction at $z = 0$ and Hubble constant.566" The galaxy redshift distribution in the ith tomographic bin is an average of the underlyingn;(z) three-dimensional galaxy distribution over angles, and the mean surface density n; is the total number of galaxies per steradian in bin i."," The galaxy redshift distribution $n_i(z)$ in the $i$ th tomographic bin is an average of the underlying three-dimensional galaxy distribution over angles, and the mean surface density $\bar{n}_i$ is the total number of galaxies per steradian in bin $i$."567" For WL, we use 10 bins evenly spaced between z—0 to 3.5, and for BAO, we use 30 bins from z=0.15 to 3.5 with bin width proportional to 14-z."," For WL, we use 10 bins evenly spaced between $z = 0$ to 3.5, and for BAO, we use 30 bins from $z = 0.15$ to 3.5 with bin width proportional to $1+z$."568" The lo statistical error of the mean power spectrum in the multipole range (¢,@+A?) is given by (t) = (0), (0) = (0) +OR, where is the Kronecker osdelta function, fa is the sky coverage,ος gis=1, and «qs,=0.18+0.0422."," The $1\sigma$ statistical error of the mean power spectrum in the multipole range $(\ell,\ell+\Delta\ell)$ is given by ) = ), ) = ) + where $\delta_{ab}^{\rm K}$is the Kronecker delta function, $f_{\rm sky}$ is the sky coverage, ${\rm g_{\rm rms}}\equiv 1$, and $\gamma_{\rm569rms} = 0.18+0.042z$."570" shows several examples of the galaxy auto and cross power spectra, and shows shear auto and cross power spectra."," shows several examples of the galaxy auto and cross power spectra, and shows shear auto and cross power spectra."571" Cross power spectra between foreground galaxy bins and background shear bins, PE. are not shown, but they have similar characteristic shapes of those in and2."," Cross power spectra between foreground galaxy bins and background shear bins, $P_{ij}^{{\rm g}\gamma}$, are not shown, but they have similar characteristic shapes of those in and."572". The amplitude of the galaxy power spectra and the strength of the BAO feature decreases with increasing rms error, because a larger rms means more smoothing in the radial direction."," The amplitude of the galaxy power spectra and the strength of the BAO feature decreases with increasing rms error, because a larger rms means more smoothing in the radial direction."573" The amplitude of the shear power spectrum increases as redshift increases, because higher redshift galaxies are lensed by more intervening matter and, hence, have stronger shear signal fluctuations."," The amplitude of the shear power spectrum increases as redshift increases, because higher redshift galaxies are lensed by more intervening matter and, hence, have stronger shear signal fluctuations."574 The amplitude of the galaxy cross power spectrum between two bins is very sensitive to the separation between the two bins in true-redshift space., The amplitude of the galaxy cross power spectrum between two bins is very sensitive to the separation between the two bins in true-redshift space.575" Hence, galaxy cross powerspectra can be used to calibrate the error distribution."," Hence, galaxy cross powerspectra can be used to calibrate the error distribution."576 We use the Fisher information matrix (Tegmark1997) toestimate the errors of the parameters of interest., We use the Fisher information matrix \citep{tegmark97b} toestimate the errors of the parameters of interest.577" In summary, the Fisher matrix of the parameter set (q is given by ="," In summary, the Fisher matrix of the parameter set $\{q_\alpha\}$ is given by ="578original fit values to produce a new data set.,original fit values to produce a new data set.579 These data are fit using the same model with the resulting re-fitted parameters recorded., These data are fit using the same model with the resulting re-fitted parameters recorded.580 The process of randomizing residuals. applying them to the original fit and refitting ts then repeated a large number of times ( 100000).," The process of randomizing residuals, applying them to the original fit and refitting is then repeated a large number of times $\sim$ 000)."581 This technique is statistically rigorous and produces à more accurate result than a simple model-fit to the given data (Efron1979)... allowing each fit parameter to be characterised by a distribution.," This technique is statistically rigorous and produces a more accurate result than a simple model-fit to the given data \citep{Efron:1979ly}, allowing each fit parameter to be characterised by a distribution."582 The inset values in Figure 9. are the mean and standard deviation for these bootstrapped parameter distributions. which suggest that the pulse has an initial velocity of «446 km s! with a negative acceleration observed in both spacecraft (a=-260=130 m s7 in STEREO--A and a=-260+ 150m s7 in STEREO--B).," The inset values in Figure \ref{fig:kinematics} are the mean and standard deviation for these bootstrapped parameter distributions, which suggest that the pulse has an initial velocity of $\sim$ 446 km $^{-1}$ with a negative acceleration observed in both spacecraft $a\simeq -260\pm130$ m $^{-2}$ in -A and $a\simeq -260\pm150$ m $^{-2}$ in -B)."583 The errors associated with the acceleration terms are quite large (in each case the acceleration is negative within the |-sigma error range)., The errors associated with the acceleration terms are quite large (in each case the acceleration is negative within the 1-sigma error range).584 However. the kinematics are consistent between both spacecraft. suggesting that the deceleration is real.," However, the kinematics are consistent between both spacecraft, suggesting that the deceleration is real."585 The initial velocity values given here are similar to previous estimates obtained by Longetal.(2008) and Veronigetal.(2008) who both studied this event. although in both of those cases a three-point Lagrangian interpolation technique was used to determine the kinematics of the CBF.," The initial velocity values given here are similar to previous estimates obtained by \citet{Long:2008eu} and \citet{Veronig:2008ud} who both studied this event, although in both of those cases a three-point Lagrangian interpolation technique was used to determine the kinematics of the CBF."586 Although that technique retains all of the data points through the use of an interpolation method. it has been observed to introduce artificial trends through the skewing of the interpolated edge points.," Although that technique retains all of the data points through the use of an interpolation method, it has been observed to introduce artificial trends through the skewing of the interpolated edge points."587 As a result. it is possible to misinterpret the derived velocity and acceleration plots.," As a result, it is possible to misinterpret the derived velocity and acceleration plots."588 In contrast. the bootstrapping technique used here is designed to determine the best-fit of a model to a given small data set. producing accurate estimates of the model parameters and quantifiable associated errors.," In contrast, the bootstrapping technique used here is designed to determine the best-fit of a model to a given small data set, producing accurate estimates of the model parameters and quantifiable associated errors."589 The variation with distance of the full width at half maximum (FWHM: Ar=2 V2In2e) of the Gaussian fit applied to the PBD intensity profile was studied to identify any evidence of pulse broadening., The variation with distance of the full width at half maximum (FWHM; $\Delta r = 2\sqrt{2\mathrm{ln}2}\sigma$ ) of the Gaussian fit applied to the PBD intensity profile was studied to identify any evidence of pulse broadening.590 The top two panels in Figure 12 show this variation for STEREO--A (left) and STEREO--B (right). with measurements from the differing. passbands indicated by the different symbols.," The top two panels in Figure \ref{fig:broadening} show this variation for -A (left) and -B (right), with measurements from the differing passbands indicated by the different symbols."591 It i5 clear from these plots that an increase in pulse width with distance Is present for both passbands as observed by both spacecraft., It is clear from these plots that an increase in pulse width with distance is present for both passbands as observed by both spacecraft.592 This is indicative of pulse broadeningand shows that the CBF spreads out spatially as it propagates., This is indicative of pulse broadeningand shows that the CBF spreads out spatially as it propagates.593ancl Naito&Takahara(1994).,and \citet{naito}.594. Thev analyzed (he prospects for detection of emission from nearby SNR that should be produced by the decavs of mesons born in collisions of shock accelerated protons with the nuclei of interstellar gas., They analyzed the prospects for detection of super-TeV emission from nearby SNR that should be produced by the decays of mesons born in collisions of shock accelerated protons with the nuclei of interstellar gas.595 The expected fluxes were shown to be detectable by the imaging Ceerenkov telescopes., The expected fluxes were shown to be detectable by the imaging erenkov telescopes.596 Moreover. (he EGRET (Espositoetal.1996) detected a lower energv GeV)) emission coinciding with some ealactie SNRs.," Moreover, the EGRET \citep{esp} detected a lower energy GeV ) emission coinciding with some galactic SNRs."597 The spectra also seemed consistent with the DSA predictions., The spectra also seemed consistent with the DSA predictions.598 One may even argue that the low enerev EGRET data verified one of the most difficult elements of the entire acceleration mechanism. (he so called injection.," One may even argue that the low energy EGRET data verified one of the most difficult elements of the entire acceleration mechanism, the so called injection."599 In essence. this is a selection process (not completely understood) whereby a small number of thermal particles become subject to further acceleration (see Gieselerοἱal.2000:Zank2001. [or the latest development ol the injection theory and Malkov&Drury2001. [or à review) and may be then (treated by standard means of the DSA theory that was designed to describe particles with velocities much higher than (he shock velocity.," In essence, this is a selection process (not completely understood) whereby a small number of thermal particles become subject to further acceleration (see \citealt{gjk00,zank01} for the latest development of the injection theory and \citealt{mdru01} for a review) and may be then treated by standard means of the DSA theory that was designed to describe particles with velocities much higher than the shock velocity."600 Therefore. what seemed. left for the theory was to continue the EGRET spectrum (that sets (he normalization constant. or injection rate) wilh some standard DSA slope (nearly or somewhat steeper) and to predict the 5--rav flux in the TeV range where it could be detected by the Ceerenkov telescopes.," Therefore, what seemed left for the theory was to continue the EGRET spectrum (that sets the normalization constant, or injection rate) with some standard DSA slope (nearly or somewhat steeper) and to predict the -ray flux in the TeV range where it could be detected by the erenkov telescopes."601 Unfortunately. despite the physical robustness of the arguments given by Druryetal.(1994):Naito&Takahara (1994).. no statistically significant signal that could be attributed to anv of the EGRET sources was detected.," Unfortunately, despite the physical robustness of the arguments given by \cite{dav,naito}, no statistically significant signal that could be attributed to any of the EGRET sources was detected."602 The further complication is that some critical energv band between GeV anc TeV energies is currently uncovered by available instruments., The further complication is that some critical energy band between GeV and TeV energies is currently uncovered by available instruments.603 Therelore. based on these observational results it was suggested Buckleyοἱal. 1998)) Chat there is probably a spectral break or even cutoff somewhere within (his band.," Therefore, based on these observational results it was suggested \citealt{buck98}) ) that there is probably a spectral break or even cutoff somewhere within this band."604 However (he spectrum above GeV energies remains an enigma., However the spectrum above GeV energies remains an enigma.605 This will be resolved perhaps with the launch of the GLAST mission and when the new generation of Ceerenkov telescopes with lower energv thresholds begin to operate., This will be resolved perhaps with the launch of the GLAST mission and when the new generation of erenkov telescopes with lower energy thresholds begin to operate.606 However. the discovery ol the LOO TeV emission from SNIA1006 (Tanimori1993)... as well as some other remnants not seen bv the EGRET at lower energies (see. Aharonianetal.2001:AllenKirk&ον2001) for à complete discussion). although almost universally identified with electrons cdiffusively accelerated to similar energies. is widelv interpreted as a strong support ol the mechanism itself.," However, the discovery of the 100 TeV emission from SNR1006 \citep{tanim98}, as well as some other remnants not seen by the EGRET at lower energies (see, \citealt{ahar01,allen01,kirkd01} for a complete discussion), although almost universally identified with electrons diffusively accelerated to similar energies, is widely interpreted as a strong support of the mechanism itself."607 The above suggests. however. that in realitv it might be not as robust as is ils simplified test particle version with enhanced turbulence ancl particle scattering.," The above suggests, however, that in reality it might be not as robust as is its simplified test particle version with enhanced turbulence and particle scattering."608 In (his paper we attempt to understand what may happen (ο the spectrum provided that the acceleration is indeed [ast enough to access the TeV energies over the life time ol SNRs in question., In this paper we attempt to understand what may happen to the spectrum provided that the acceleration is indeed fast enough to access the TeV energies over the life time of SNRs in question.609 Our starting point is that the fast acceleration also means that. the pressure of accelerated particles becomes significant in an early stage of Supernova evolution so that the shock structure is highly nonlinear., Our starting point is that the fast acceleration also means that the pressure of accelerated particles becomes significant in an early stage of Supernova evolution so that the shock structure is highly nonlinear.610 At the first glance this should not slow down, At the first glance this should not slow down611for E/S0s divided by absolute magnitude.,for E/S0s divided by absolute magnitude.612" The negative correlation with density is seen over a very wide luminosity range, and strongest at —21>M,—23."," The negative correlation with density is seen over a very wide luminosity range, and strongest at $-21>M_r>-23$."613" To briefly summarize at this point, we find colour gradients in E/S0s to (i) be greatest at intermediate luminosities and (ii) at a fixed luminosity to be anticorrelated with velocity dispersion and mass density, and that (iii) colour gradients tend to be flatter in BCGs and to show much less dependence on galaxy properties."," To briefly summarize at this point, we find colour gradients in E/S0s to (i) be greatest at intermediate luminosities and (ii) at a fixed luminosity to be anticorrelated with velocity dispersion and mass density, and that (iii) colour gradients tend to be flatter in BCGs and to show much less dependence on galaxy properties."614 We now investigate the effects of age., We now investigate the effects of age.615" For most galaxies in the E/SO sample (64774/70378) and the M,«-—22.5 non-BCG sample (16337/18225), we have mean stellar age (‘age50’) estimated from the spectral index analysis of Gallazzi et al. ("," For most galaxies in the E/S0 sample (64774/70378) and the $M_r<-22.5$ non-BCG sample (16337/18225), we have mean stellar age (`age50') estimated from the spectral index analysis of Gallazzi et al. ("6162005) performed on the 3.7x10? galaxies of the DR4.,2005) performed on the $3.7\times 10^5$ galaxies of the DR4.617 This catalog includes only part of the max-BCG list., This catalog includes only part of the max-BCG list.618" By position-matching the two lists we find there are stellar age estimates for almost half (2274/4919) the max-BCG sample, sufficient to investigate age-dependent effects."," By position-matching the two lists we find there are stellar age estimates for almost half (2274/4919) the max-BCG sample, sufficient to investigate age-dependent effects."619" These stellar ages will generally be less than the lookback time to galaxy formation, and as they are luminosity-weighted, the age estimate of a galaxy may be greatly reduced if it has experienced recent episodes of star-formation."," These stellar ages will generally be less than the lookback time to galaxy formation, and as they are luminosity-weighted, the age estimate of a galaxy may be greatly reduced if it has experienced recent episodes of star-formation."620" The stellar mean age (and scatter) are 6.92 (1.74) Gyr for all E/SOs, 7.04 (1.69) Gyr for the non-BCGs, 7.54 (1.39) Gyr for the max-BCGs and 8.12 (1.33) Gyr for the C4 BCGs."," The stellar mean age (and scatter) are 6.92 (1.74) Gyr for all E/S0s, 7.04 (1.69) Gyr for the non-BCGs, 7.54 (1.39) Gyr for the max-BCGs and 8.12 (1.33) Gyr for the C4 BCGs."621 The greater mean age of the C4 BCGs is due to these being a lower redshift sample than the others., The greater mean age of the C4 BCGs is due to these being a lower redshift sample than the others.622 Figure 16a shows the age distributions for the 3 deep samples; the modal ages are similar but there are fewer younger («6 Gyr) and more older ages (9-10 Gyr) ages amongst the BCGs., Figure 16a shows the age distributions for the 3 deep samples; the modal ages are similar but there are fewer younger $<6$ Gyr) and more older ages (9–10 Gyr) ages amongst the BCGs.623" We add the mean stellar ages to the lookback times at the redshifts of observation, giving lookback times to the mean stellar formation time, and then convert these back to ‘formation’ redshifts (not the redshift of galaxy formation but the mean redshift of luminosity-weighted star formation, which will be lower) for the galaxies."," We add the mean stellar ages to the lookback times at the redshifts of observation, giving lookback times to the mean stellar formation time, and then convert these back to `formation' redshifts (not the redshift of galaxy formation but the mean redshift of luminosity-weighted star formation, which will be lower) for the galaxies."624" Figure 16b shows the distribution of the mean redshift of star formation; there is more of a difference here between the all-E/SO sample and the BCGs and other high luminosity galaxies, as the more luminous galaxies tend to be observed at higher redshifts as well as being older."," Figure 16b shows the distribution of the mean redshift of star formation; there is more of a difference here between the all-E/S0 sample and the BCGs and other high luminosity galaxies, as the more luminous galaxies tend to be observed at higher redshifts as well as being older."625" The mean star-formation redshift is 1.32 for all E/SO, 1.56 for the non-BCGs, 1.53 Gyr for the C4 BCGs and 1.82 for the max-BCGs; mean lookback times are 8.47, 9.17, 9.22 and 9.85 Gyr."," The mean star-formation redshift is 1.32 for all E/S0, 1.56 for the non-BCGs, 1.53 Gyr for the C4 BCGs and 1.82 for the max-BCGs; mean lookback times are 8.47, 9.17, 9.22 and 9.85 Gyr."626" Figure 17 shows mean age against M,; the BCGs are not only older than other E/S0s but show less of an increase in age with luminosity.", Figure 17 shows mean age against $M_r$; the BCGs are not only older than other E/S0s but show less of an increase in age with luminosity.627" We investigate whether the redder colours and weaker colour gradients of BCGs are related to their greater ages, and in addition, whether the trends in (non-BCG) colour gradients with galaxy properties are caused by age effects."," We investigate whether the redder colours and weaker colour gradients of BCGs are related to their greater ages, and in addition, whether the trends in (non-BCG) colour gradients with galaxy properties are caused by age effects."628" A higher luminosity or lower c relative to the mean (o|M,.) relation, found (in Paper I and here) to be associated with a stronger colour gradient, is also correlated with a younger stellar age (Forbes and Ponman 1999; Bernardi et al."," A higher luminosity or lower $\sigma$ relative to the mean $\langle\sigma|M_r\rangle$ relation, found (in Paper I and here) to be associated with a stronger colour gradient, is also correlated with a younger stellar age (Forbes and Ponman 1999; Bernardi et al."629 2005; Gallazzi et al., 2005; Gallazzi et al.630 2006)., 2006).631" On Figure 18, for non-BCG E/S0s the colour gradients are maximum at stellar ages 4—5 Gyr and decrease by almost a factor of two to the highest ages."," On Figure 18, for non-BCG E/S0s the colour gradients are maximum at stellar ages 4–5 Gyr and decrease by almost a factor of two to the highest ages."632" On Figure 19 it can be seen that this peak in gradient vs. age is most pronounced for galaxies of intermediate luminosities of —23, weaker at both —20>M,—21 and —23>M, —24, and apparently absent at M,>—20 and M,<—24, following a rather symmetric luminosity trend."," On Figure 19 it can be seen that this peak in gradient vs. age is most pronounced for galaxies of intermediate luminosities of $-21>M_r>-23$ , weaker at both $-20>M_r>-21$ and $-23>M_r>-24$ , and apparently absent at $M_r>-20$ and $M_r<-24$, following a rather symmetric luminosity trend."633" The BCG colour gradients show no age dependence (at >4 Gyr), and again the BCGs have lower colour gradients than the non-BCG E/S0 sample, at all but the highest ages where the two converge."," The BCG colour gradients show no age dependence (at $>4$ Gyr), and again the BCGs have lower colour gradients than the non-BCG E/S0 sample, at all but the highest ages where the two converge."634" A potential problem with this analysis is that the stellar ages are ‘central’ ages estimated from the spectra in fixed apertures, and if colour gradients were typically driven by age gradients this would bias the age-gradient relation."," A potential problem with this analysis is that the stellar ages are `central' ages estimated from the spectra in fixed apertures, and if colour gradients were typically driven by age gradients this would bias the age-gradient relation."635" We are not able to investigate age gradients with the single aperture of SDSS spectroscopy, but Tamura and Ohta (2004) and La Barbera and Carvalho (2009) find that age gradients in early-types are typically very small or consistent with zero"," We are not able to investigate age gradients with the single aperture of SDSS spectroscopy, but Tamura and Ohta (2004) and La Barbera and Carvalho (2009) find that age gradients in early-types are typically very small or consistent with zero"636(Seaton1959:Cen1992).,"\cite{sea59,cen92}."637 Phe average energy raciated per recombination. then. is 15.0eV.," The average energy radiated per recombination, then, is $15.0~\ev$."638 Under the on-the-spot approximation. all Lyman series photons degrade into n=2*1 transitions. where n is the energy. level of an excited hydrogen atom.," Under the on-the-spot approximation, all Lyman series photons degrade into $n=2\rightarrow1$ transitions, where $n$ is the energy level of an excited hydrogen atom."639 Since. 10.2eV. is released. per ΕΞ+ transition. about 2/3 of the recombinationenergy is released in Lye or two-photon emission. and 1/3 of the recombination energv (and about 1/6 of the total emitted energv) is emitted. in other recombination lines and [ree-bound continuous emission.," Since $10.2~\ev$ is released per $n=2\rightarrow1$ transition, about $2/3$ of the recombination energy is released in $\lya$ or two-photon emission, and $1/3$ of the recombination energy (and about $1/6$ of the total emitted energy) is emitted in other recombination lines and free-bound continuous emission."640 We ignore those processes here: see Schaerer (2002) for a more complete computation of nebular emission lines and free-bound spectra., We ignore those processes here; see Schaerer \shortcite{sch02} for a more complete computation of nebular emission lines and free-bound spectra.641 The relative importance of Lya enüssion. compared to two-photon emission is determined. by the effective recombination coellicients. to. the 2p anc 28 states. respectively. and the collisional excitation rate from the 2s state to the 2p state.," The relative importance of $\lya$ emission compared to two-photon emission is determined by the effective recombination coefficients to the $2p$ and $2s$ states, respectively, and the collisional excitation rate from the $2s$ state to the $2p$ state."642 At P=3103IX. 0.75 Lya photons are emitted for every rrecombination (Osterbrock1989:Storey&Hummer1995):: we ignore the contribution of two-photon emission (about 12 »er cent of the total emission) to the spectrum.," At $T=3\times10^4~\kelvin$, $0.75$ $\lya$ photons are emitted for every recombination \cite{ost89,sto95}; we ignore the contribution of two-photon emission (about 12 per cent of the total emission) to the spectrum."643 Since Lya photons resonantlv scatter in neutral ivdrogen. they will not travel far in the LGAL until their requencies are shifted away from the resonant frequency. Moye.," Since $\lya$ photons resonantly scatter in neutral hydrogen, they will not travel far in the IGM until their frequencies are shifted away from the resonant frequency, $\nu_\lya$ ."644 Photons initially scattered blueward. of the. line resonance will eventually cosmologically redshift back into he resonance., Photons initially scattered blueward of the line resonance will eventually cosmologically redshift back into the resonance.645 The result is that the IGM ultimately scatters all Lye photons to the red side of the line. resonance (broadening from motions inside a halo doesn't contribute to the line profile)., The result is that the IGM ultimately scatters all $\lya$ photons to the red side of the line resonance (broadening from motions inside a halo doesn't contribute to the line profile).646 For à homogeneous. expanding. IGM. the resulting scattered line profile. O07.z). was simulated by Loeb Itixbicki (1999).. and we fit their result with This profile results in a strong. asvometric Lya emission line near 12201225A with a scattering tail extending to long wavelengths.," For a homogeneous, expanding IGM, the resulting scattered line profile, $\phi(\nu,z)$, was simulated by Loeb Rybicki \shortcite{loe99}, and we fit their result with This profile results in a strong, asymmetric $\lya$ emission line near $1220-1225~\ang$ with a scattering tail extending to long wavelengths."647 In the rreeion. recombinationssometimes produce more than one photon capable of ionizing Ili.," In the region, recombinationssometimes produce more than one photon capable of ionizing ."648 Reeombinations directly to. the n2 state produce BBalmer continuum photons. which are capable of ionizing lli.," Recombinations directly to the $n=2$ state produce Balmer continuum photons, which are capable of ionizing ."649 In addition. two-photon decay from the 2s state produces 1.42 Li--ionizine photons per decay 1989).," In addition, two-photon decay from the $2s$ state produces 1.42 -ionizing photons per decay \cite{ost89}."650. Since Lye is also capable of ionizing hydrogen. the mean number of hydrogen ionizations per rrecombination is where ae(lleu.7) is the recombination cocllicient for recombinations clirectly to the n—2 state ancl ostPelle11.1) is the elfective total recombination cocllicient to the 2/ state (Storey&απο1995).," Since $\lya$ is also capable of ionizing hydrogen, the mean number of hydrogen ionizations per recombination is where $\alpha_2(\heii,T)$ is the recombination coefficient for recombinations directly to the $n=2$ state and $\alpha_{2l}^{\mathrm651eff}(\heii,T)$ is the effective total recombination coefficient to the $2l$ state \cite{sto95}."652. Free-[ree radiation accounts for about 1/4 of the cooling in the nebula., Free-free radiation accounts for about $1/4$ of the cooling in the nebula.653 This energy is radiated in à continuous spectrum. where jl is the specific emission. coctlicicnt. 7 is the frequeney of emitted radiation and nz is the number density of ions of net charge Z (Ferland1980). A Gaunt [actor of 1.3 has been assumed: this results in an error of less than S per cent over optical andUV frequencies 1961)..," This energy is radiated in a continuous spectrum, where $j_\nu^{\mathrm ff}$ is the specific emission coefficient, $\nu$ is the frequency of emitted radiation and $n_Z$ is the number density of ions of net charge $Z$ \cite{fer80} A Gaunt factor of 1.3 has been assumed; this results in an error of less than 8 per cent over optical andUV frequencies \cite{kar61}. ."654 Theluminosityper solar mass of sstars [rom [ree-f£ree radiation. /L/. is Llere aand label the emission coelficients ancl volumes computed. for the aand," Theluminosityper solar mass of stars from free-free radiation, $l_\nu^{\mathrm ff}$ , is Here and label the emission coefficients and volumes computed for the and"655While the processes which may lead to AIBIT seed formation have been modeled cosiologically (Wiseetal.2008:Regan&Tachuelt 2009).. if is not vet possible to do so in simulations involviug volumes larger than a few Mpc?.,"While the processes which may lead to MBH seed formation have been modeled cosmologically \citep{Wise08,Regan09}, it is not yet possible to do so in simulations involving volumes larger than a few $^3$."656 Previous large-scale sinmlations have incorporated seed black hole formation in a simplistic way. Sijackietal," Previous large-scale simulations have incorporated seed black hole formation in a simplistic way. \citet{Sijacki07},"657.(2007).. DiMatteoetal.(2008). and Booth&Schave(2000) all employ simular methods. which involves runing an on-the-fiv halo finder on the simulation as it evolves. aud. planting seed. MDIIs iu particular halos bx iud.," \citet{DiMatteo08}658 and \citet{Booth09} all employ similar methods, which involves running an on-the-fly halo finder on the simulation as it evolves, and planting seed MBHs in particular halos by hand."659 The halos are chosen based on a mass threshold (ecnerally Afiaiy~LOM ALA). aud seeds are. plated if there is not already a black hole preseut.," The halos are chosen based on a mass threshold (generally $M_{\rm halo} \sim 10^{10}$ ), and seeds are planted if there is not already a black hole present."660 There is 10 lnctallicity criterion. so seed MDIIS can form: at any redshift if a halo meets the eligibility criteria.," There is no metallicity criterion, so seed MBHs can form at any redshift if a halo meets the eligibility criteria."661 The seeds are placed at the ceuter of the halo aud ecucrally fixed here throughout the remainder of the simulation (Booth&Schave 2009)., The seeds are placed at the center of the halo and generally fixed there throughout the remainder of the simulation \citep{Booth09}.662. However. there is no plysical motivation for a halo uass threshold for MBIT placement of 1029AL... aud in act seeds Likely form iu halos of mich lower mass (sec 877)).," However, there is no physical motivation for a halo mass threshold for MBH placement of $10^{10}$, and in fact seeds likely form in halos of much lower mass (see \ref{sect:history}) )."663 This method also prevents more than one seccl roni formune per halo. which may be wmurealistic - halos nay experience the formation of nuultiple Population III stars if fragmentation occurs (Turketal.2009:Stacyetal.2010:Greif2011:Clark 2011).," This method also prevents more than one seed from forming per halo, which may be unrealistic - halos may experience the formation of multiple Population III stars if fragmentation occurs \citep{Turk09,Stacy10,Greif11,Clark11}."664. Fisine he black hole at the halo ceuter is similarly iuadvisable - dvuaiical effects such as galaxy mergers or eravitational recoil may cause a black hole to temporarily vacate the exact center of its galaxy. but such a circumstance ds prohibited in these models.," Fixing the black hole at the halo center is similarly inadvisable - dynamical effects such as galaxy mergers or gravitational recoil may cause a black hole to temporarily vacate the exact center of its galaxy, but such a circumstance is prohibited in these models."665 The formalisu of ποσα MDII formation we describe in this paper is a more sophisticated. ancl realistic model because it makes πο assunptious about the underlving halo properties., The formalism of seed MBH formation we describe in this paper is a more sophisticated and realistic model because it makes no assumptions about the underlying halo properties.666" Our method relies ouly onthe prospective MDBITIs local environment to determine when and where the MDBIT formes. including a requirement for gero-netallicity eas,"," Our method relies only on the prospective MBH's local environment to determine when and where the MBH forms, including a requirement for zero-metallicity gas."667 Our scenario is broadly cousisteut with either proposed MDBII formation mechanisin. and also allows MDIIsS to evolve dvuamically in a realistic way.," Our scenario is broadly consistent with either proposed MBH formation mechanism, and also allows MBHs to evolve dynamically in a realistic way."668 It is crucial to study low variations in halo lass affect the frequency of formation when examine simulated seed ALIBI populations. aud so we include iu our sample a variety of halo masses in order to preseut a more coherent picture of high-: seed formation aud evolution as a function of cosmic environment.," It is crucial to study how variations in halo mass affect the frequency of formation when examining simulated seed MBH populations, and so we include in our sample a variety of halo masses in order to present a more coherent picture of $z$ seed formation and evolution as a function of cosmic environment."669 With this technique. we investigate where aud when MDIT seeds form at high redshift. aud discuss the nuplicatious for ealaxies at 7=0.," With this technique, we investigate where and when MBH seeds form at high redshift, and discuss the implications for galaxies at $z = 0$."670 We emplov the. N-Body. Tree code Gasoline (Stadel2001:Wadsleyetal.2001). which las been shown to produce realistic galaxies in cosmological siuulatious (e.g.Governatootal.2010:Stinsonctetal. 2011).," We employ the $N$ -Body Tree code Gasoline \citep{Stadel01,Wadsley04}, which has been shown to produce realistic galaxies in cosmological simulations \citep[e.g.][]{Governato10,671Stinson10,Pontzen10,Oh10,Brooks11}."672.. Gasoline iucludes a plivsicallv. motivated prescription for star formation as well as a recipe for the formation aud evolution of AIBIT seeds (described in 22)., Gasoline includes a physically motivated prescription for star formation as well as a recipe for the formation and evolution of MBH seeds (described in \ref{sect:form}) ).673 We use a EKroupa TIF (svoupa2001) and a WMAP2 cosinology (Sperecletal.2007).., We use a Kroupa IMF \citep{Kroupa} and a WMAP3 cosmology \citep{WMAP3}.674 We use a supernova feedback model which incorporates the Sedov solution to the blastwave equations (seeStinsonetal. 2006). and set the blastwave enerev to Ee=10°! ores.," We use a supernova feedback model which incorporates the Sedov solution to the blastwave equations \citep[see][]{Stinson06}, and set the blastwave energy to $E_{SN} =67510^{51}$ ergs."676 We nauploiieut a uniforin ionizing radiation background with an onset at redslift ;=9 (IJaardt&Madau1996)., We implement a uniform ionizing radiation background with an onset at redshift $z = 9$ \citep{Haardt96}.677. Our sinulatious inchide cooling through metal lines as well as a imodol for turbulent metal diffusion (Shenetal.2010): however. we do not iucluding cooling via molecular hydrogen iu the simulations prescuted here.," Our simulations include cooling through metal lines as well as a model for turbulent metal diffusion \citep{Shen10}; however, we do not including cooling via molecular hydrogen in the simulations presented here."678 To probe a large parameter space. we enmiplov a suite of simulations with differeut halo masses to test seed MBO formation.," To probe a large parameter space, we employ a suite of simulations with different halo masses to test seed MBH formation."679 It is crucial to study how variations in halo mass affect the frequeney of formation when exanuning simulated seed MDII populations., It is crucial to study how variations in halo mass affect the frequency of formation when examining simulated seed MBH populations.680 We have chosen a set of three fiducial halos selected frou a uniform resolution. 50° Alpe volume aud resunulated at high resolution using the volume renormalization technique (Natz&White1993).," We have chosen a set of three fiducial halos selected from a uniform resolution, 50 Mpc volume and resimulated at high resolution using the volume renormalization technique \citep{Katz93}."681. We sample the region of interest at high resolution. aud then sample more coarsely as the distance from the chosen object increases.," We sample the region of interest at high resolution, and then sample more coarsely as the distance from the chosen object increases."682 This technique results i a large dvianic range. where we can capture the detailed plysics of galaxy. formation on sanall scales iu the region of interest. and also the larec-scale tidal torques from structures several Mpe away.," This technique results in a large dynamic range, where we can capture the detailed physics of galaxy formation on small scales in the region of interest, and also the large-scale tidal torques from structures several Mpc away."683 The three chosen halos correspond broadly to a disk ealaxy (1603). a disk galaxy (1258). aud aelliptical galaxy (11) at +=0.," The three chosen halos correspond broadly to a disk galaxy $h603$ ), a disk galaxy $h258$ ), and a galaxy $hz1$ ) at $z = 0$."684 We present the details of our simulations in Table 1.., We present the details of our simulations in Table \ref{table:sims}.685 Our sinulations are extremely high resolution. with eas particles having masses between 209«410!ALL. and a force resultion of 173 - 260 pc. depending ou the simulation.," Our simulations are extremely high resolution, with gas particles having masses between $2686- 9 \times 10^4$, and a force resultion of 173 - 260 pc, depending on the simulation."687 The mass ratio of gas and dark matter particles (column 1) is of order unitv. which reduces the effects due to two-body scattering and is critical for keeping MDIIs in the ceuters of their galaxies.," The mass ratio of gas and dark matter particles (column 4) is of order unity, which reduces the effects due to two-body scattering and is critical for keeping MBHs in the centers of their galaxies."688 Since we are interested in the epoch of seed MDII formation. we have run our simulations to 2=5 rather than 2=0. in order to maximize resolution while using only modest compuationual resources.," Since we are interested in the epoch of seed MBH formation, we have run our simulations to $z = 5$ rather than $z = 0$, in order to maximize resolution while using only modest compuational resources."689 We use the Amica Talo Finder (Calletal.2001:I&nolliuaun&I&nebe2009) to identify galaxy halos based ou the overdensity criterion for a flat universe (Cross1997).," We use the Amiga Halo Finder \citep{Gill04,Knollmann09} to identify galaxy halos based on the overdensity criterion for a flat universe \citep{Gross97}."690 For each siaulated reeion. we analyze MDBII populations for the primary halo as well as every satellite with at least 61 particles.," For each simulated region, we analyze MBH populations for the primary halo as well as every satellite with at least 64 particles."691 Iu the case of P603 there are 5370 total halos iu our analysis: for h258 there are 5170: for 2:1 there are 2160., In the case of $h603$ there are 5370 total halos in our analysis; for $h258$ there are 5170; for $hz1$ there are 2160.692 Since the actual physical process of MDITI seed formation is unresolvable i cosmological simulations. we have developed a model which is broadly cousistcut with both of the proposed seed formation scenarios.," Since the actual physical process of MBH seed formation is unresolvable in cosmological simulations, we have developed a model which is broadly consistent with both of the proposed seed formation scenarios."693 The cohunon thread between the proposed scenarios is au ability for eas to collapse into a large ceutral mass. which requires a zero or near-zero metallicity (thoueh see Mayeretal.(2010). for an alternate mechamisiur).," The common thread between the proposed scenarios is an ability for gas to collapse into a large central mass, which requires a zero or near-zero metallicity (though see \citet{Mayer10} for an alternate mechanism)."694 Because our star formation prescription is already based on snülar plivsics Ge. cold. dense. collapsing gas results in the formation of star particles). we use this prescription witli the additional criterion of zero ietallicitv to form seed AEDITs.," Because our star formation prescription is already based on similar physics (i.e. cold, dense, collapsing gas results in the formation of star particles), we use this prescription with the additional criterion of zero metallicity to form seed MBHs."695 Our current star formation recipe is described iu detail in Stinsonctal. (2006).. and we sunmuuize it here.," Our current star formation recipe is described in detail in \citet{Stinson06}, , and we summarize it here."696 For a star to form iu our simulations. ifs parent gas," For a star to form in our simulations, its parent gas"697distance estimate 2008).. although this value is subject to several uncertainties 2008b).,"distance estimate , although this value is subject to several uncertainties ."698". At the time of its ciscovery. wwas detected. oat 210. Καν luminosities of (17)10""(D/TAkpc)?ergs 1986)."," At the time of its discovery, was detected at 2–10 keV luminosities of $\sim(1-7)\times10^{36}~\mathrm{(D/7.4~kpc)^2}~\lum$ ."699. Llowever. during the observation of 1980. several vears prior to the ddetections. it clisplavecdl a 0.510. keV luminosity. of M55107(οτιkpeyeres+1999)..," However, during the observation of 1980, several years prior to the detections, it displayed a 0.5–10 keV luminosity of $\sim 5 \times 10^{33}~\mathrm{(D/7.4~kpc)^2}~\lum$."700 The source can therefore be classified as a transient X-ray. binary., The source can therefore be classified as a transient X-ray binary.701 Nevertheless. such svstenis typically exhibit accretion outbursts that last only weeks to months1997).. whereas wwas persistently detected at luminosities ol ~1078(D/TAkpe)?ergs bv. various satellites or over 24 wears.," Nevertheless, such systems typically exhibit accretion outbursts that last only weeks to months, whereas was persistently detected at luminosities of $\sim10^{36-37}~\mathrm{(D/7.4~kpc)^2}~\lum$ by various satellites for over 24 years."702 Similar prolonged accretion episodes continuing for vears to decades have been observed. for a ew other systems. which are termed quasicpersistent X-ray rinaries2004).," Similar prolonged accretion episodes continuing for years to decades have been observed for a few other systems, which are termed quasi-persistent X-ray binaries."703. In 2008 AugustSeptember. observations with the Proportional Counter Array (PCA) onboard the (RNTI)) and Suwiff'ss N-aray Telescope (NIE) indicated that the X-ray lux of wavas declining2008a.," In 2008 August–September, observations with the Proportional Counter Array (PCA) onboard the ) and s X-ray Telescope (XRT) indicated that the X-ray flux of was declining."704b).. Optical and near-LR observations of the optical counterpart. UY Vol. performed in 2008 October showed that the optical emission. had also faceck compared to the brighter X-ray state2008).," Optical and near-IR observations of the optical counterpart, UY Vol, performed in 2008 October showed that the optical emission had also faded compared to the brighter X-ray state."705. These events indicated that the accretion was ceasing and that the svstenir was transitioning [rom outburst to quiescence., These events indicated that the accretion was ceasing and that the system was transitioning from outburst to quiescence.706 This is also illustrated by Fie. 1..," This is also illustrated by Fig. \ref{fig:asm},"707 which displavs the X-ray lighteurve of aas observed with the Alb-Sky Monitor. (ASAI) onboard ssince 1996., which displays the X-ray lightcurve of as observed with the All-Sky Monitor (ASM) onboard since 1996.708 The decrease in source activity is clearly seen around ~4600 cays., The decrease in source activity is clearly seen around $\sim4600$ days.709 oobservations carried out in 2008 mid-October (1.o.. alter. the transition to quiescence started) revealed an X-ray spectrum composed of a soft. thermal component joined by a hard powerlawtail that dominates the spectrum above ~2B keV hisis f requentansccenfipenenunenvsdb-Not raybinariesingnicscence," observations carried out in 2008 mid-October (i.e., after the transition to quiescence started) revealed an X-ray spectrum composed of a soft, thermal component joined by a hard powerlawtail that dominates the spectrum above $\sim2-3$ keV. This is frequently seen for neutron star X-ray binaries in quiescence."710leg.. thermalcomponentisusiuallywell filledbyasimplepowerlawwilhinderttypicalBog. 1996)., The non-thermal component is usually well-fitted by a simple powerlaw with index 1--2 .711 Phe fractionalcontributiono fthchardpowariiie l10RkcV X ray flierwidelgearicsamongstsourcesandpossiblyalscuwad val, The fractional contribution of the hard powerlaw tail to the 0.5--10 keV X-ray flux widely varies amongst sources and possibly also with changing luminosity.712veibaacurd Lhephysicalpro , The physical process that is responsible for the powerlaw spectral component remains elusive.713Although the soft spectral component has been ascribed to low-level accretion.1995).. it js most often interpreted as thermal surface radiation. from. the cooling neutron star1998).," Although the soft spectral component has been ascribed to low-level accretion, it is most often interpreted as thermal surface radiation from the cooling neutron star."714. According to this model. the aceretion of matter compresses the neutron star crust. which induces a series of electron. captures. neutron emissions and. pyenonuclear fusion reactions2007).," According to this model, the accretion of matter compresses the neutron star crust, which induces a series of electron captures, neutron emissions and pycnonuclear fusion reactions."715. The heat energy released in these processes is spread over the neutron star via thermal conduction., The heat energy released in these processes is spread over the neutron star via thermal conduction.716 The neutron star cools primarily via neutrino emissions from the stellar core. as well as photon radiation from the surface.," The neutron star cools primarily via neutrino emissions from the stellar core, as well as photon radiation from the surface."717 The former depends on the equation of state of cold nuclear matter and the central clensity of the neutron star2006)., The former depends on the equation of state of cold nuclear matter and the central density of the neutron star.718. The neutron star core reaches a thermal steady. state in ~101 vears. vielding an incandescent emission [rom the neutron star surface set by the time-averaged accretion rate of the system. as well as the rate of neutrino emissions [roni the stellar core2001).," The neutron star core reaches a thermal steady state in $\sim10^{4}$ years, yielding an incandescent emission from the neutron star surface set by the time-averaged accretion rate of the system, as well as the rate of neutrino emissions from the stellar core."719. When combined with estimates of the outburst history. observations of quiescent neutron stars can constrain the rate of neutrino emissions. thereby. providing insight into the interior properties of the neutron star2009).," When combined with estimates of the outburst history, observations of quiescent neutron stars can constrain the rate of neutrino emissions, thereby providing insight into the interior properties of the neutron star."720. Once the steady state is reached. the neutron star change appreciably during a sjpiotbd ohoherath bua! 0 αμπο of the crust. can be dramatically altered.," Once the steady state is reached, the neutron star core temperature will not change appreciably during a single outburst, but the temperature of the crust can be dramatically altered."721 In. regular transients that have a outburst. duration of weeks to months. the crustal prbbobReO Aill only cause a slight increase in the avagé 20peoon," In regular transients that have a typical outburst duration of weeks to months, the crustal heating processes will only cause a slight increase in the crust temperature."722 hier quspersistent N-rav. binaries the prolonged accretion episodes can cause a significant temperature. gradient. between the neutron star crust ancl core.," However, in quasi-persistent X-ray binaries the prolonged accretion episodes can cause a significant temperature gradient between the neutron star crust and core."723 Once the accretion ceases. the crust is expected to thermally relax on a time scale of vears. until equilibrium with the core is re-established2002).," Once the accretion ceases, the crust is expected to thermally relax on a time scale of years, until equilibrium with the core is re-established."724. During the initial stages of the quiescent phase the thermal emission. will therefore. be dominated bv the cooling crust. whereas eventually a quiescent base level is reached that is set by the thermal state of the core 2002).," During the initial stages of the quiescent phase the thermal emission will therefore be dominated by the cooling crust, whereas eventually a quiescent base level is reached that is set by the thermal state of the core ."725. This provides the special opportunity to separately probe the properties of the neutron star crust 2009).., This provides the special opportunity to separately probe the properties of the neutron star crust .726(incoherent) signals produced. by dillerent. annuli.,(incoherent) signals produced by different annuli.727 As the οονομα emissivitv indices highlight dillering regions of he disc. the coherence between the corresponding energy vances can drop.," As the different emissivity indices highlight differing regions of the disc, the coherence between the corresponding energy bands can drop."728" This is shown in Fig. δι,"," This is shown in Fig. \ref{input_coh},"729 where we plot he coherence for light curves with broad. ((2= 0.5) and narrow (CQ)= 10) signal PSDs. using dotted and dashed lines respectively.," where we plot the coherence for light curves with broad $Q=0.5$ ) and narrow $Q=10$ ) signal PSDs, using dotted and dashed lines respectively."730 As expected. broader signal PSDs produce a stronger drop.," As expected, broader signal PSDs produce a stronger drop."731 Notice however. that in either case this loss of coherence is at most a few per cent.," Notice however, that in either case this loss of coherence is at most a few per cent."732 Phe emissivity indices have a much stronger elfect on the coherence., The emissivity indices have a much stronger effect on the coherence.733 For Fig., For Fig.734 δ we used son=23 and μα= +., \ref{input_coh} we used $\gamma_{\rm soft}=3$ and $\gamma_{\rm hard}=4$ .735" Phe same plot calculated [or soon=3 and sad=5 would show a drop of up to 7 per cent. and for 54,4=3 ancl 5p=o. of 40 per cent. for light curves with Q=0.5."," The same plot calculated for $\gamma_{\rm736soft}=3$ and $\gamma_{\rm hard}=5$ would show a drop of up to 7 per cent, and for $\gamma_{\rm soft}=3$ and $\gamma_{\rm hard}=\infty$, of 40 per cent, for light curves with $Q=0.5$."737 The effect of damping is to reduce the amplitude of the luctuations as they propagate inward so that the farther hey eo the smaller they get., The effect of damping is to reduce the amplitude of the fluctuations as they propagate inward so that the farther they go the smaller they get.738 An immediate consequence is he reduction of variability power at all frequencies as seen in the PSD., An immediate consequence is the reduction of variability power at all frequencies as seen in the PSD.739 As this suppression is frequency dependent. damping can change the shape of the PSD.," As this suppression is frequency dependent, damping can change the shape of the PSD."740 Fig., Fig.741 9. shows he PSD [or four hard lisht curves (5.= 4) with equal disc parameters and input signal variability zumplitucles. out. affected: by. dilferent. degrees. of damping.," \ref{damp_per} shows the PSD for four hard light curves $\gamma=4$ ) with equal disc parameters and input signal variability amplitudes, but affected by different degrees of damping."742 The. most noticeable change is the overall reduction. of variability »ower., The most noticeable change is the overall reduction of variability power.743 However. as the input rms is arbitrary. this effect does not help to quantify the amount of damping allecting real data.," However, as the input rms is arbitrary, this effect does not help to quantify the amount of damping affecting real data."744 A more important effect is the change in the low requency slope. as damping can produce values of ep«/— 1. as seen in e.g. citepMellardy MCCC.," A more important effect is the change in the low frequency slope, as damping can produce values of $\alpha_{\rm L}<-1$ , as seen in e.g. \\citep{McHardyMCG}."745.. Fable 20 shows the values of ap for damping coclicients 2=0.0.5.1 and 2. for cilferent emissivity indices.," Table \ref{lfslope} shows the values of $\alpha_{\rm L}$ for damping coefficients $D =0,0.5,1$ and 2, for different emissivity indices."746" In all cases. the stronger camping Iattens the low frequeney PSD. by similar values regardless of 7. e.g. [or D=2 compared to D=0. Χαν~0.17 for 3=0 and Aa,~0.1 For 3=1."," In all cases, the stronger damping flattens the low frequency PSD, by similar values regardless of $\gamma$, e.g. for $D=2$ compared to $D=0$, $\Delta \alpha_{\rm L} \sim 0.17 $ for $\beta=0$ and $\Delta \alpha_{\rm L} \sim 0.1$ for $\beta=1$."747 Damping can also reduce the coherence. especially in the case of broad input signal PSDs.," Damping can also reduce the coherence, especially in the case of broad input signal PSDs."748 The cot-clashecl line in Fig. S..," The dot-dashed line in Fig. \ref{input_coh},"749 shows the coherence for light curves of emissivity indices .5 and να=4+ Q=0.5 and damping coellicient D—1.," shows the coherence for light curves of emissivity indices $\gamma_{\rm750soft}=3$ and $\gamma_{\rm hard}=4$, $Q=0.5$ and damping coefficient $D=1$."751 These light curves only dilfer from those corresponding to the dashed line on the same plot by the elect. of damping. which increases the loss of coherence at the highest. frequeney bin from 2 to 4 per cent.," These light curves only differ from those corresponding to the dashed line on the same plot by the effect of damping, which increases the loss of coherence at the highest frequency bin from 2 to 4 per cent."752 Again. a larger dillerence in emissivitv indices makes this drop more pronounced. and damping elfects increase these losses accordingly: for σα=3 and shart=5 the coherence at the highest frequency bin drops by ~12 per cent and. for soot=Ὁ and tad=x by GO per cent.," Again, a larger difference in emissivity indices makes this drop more pronounced, and damping effects increase these losses accordingly: for $\gamma_{\rm soft}=3$ and $\gamma_{\rm hard}=5$ the coherence at the highest frequency bin drops by $\sim 12$ per cent and for $\gamma_{\rm soft}=3$ and $\gamma_{\rm hard}=\infty$, by 60 per cent."753 Finally. damping can produce slightly steeper lag spectra. reducing thelags at high frequencies.," Finally, damping can produce slightly steeper lag spectra, reducing thelags at high frequencies."754" This effect is small:damping coellicients of D=0 and 2=| produce lags dilfering bv. a factor of <1.5 at frequencies close to the measured. break-frequcney f, ancl less at lower frequencies.", This effect is small:damping coefficients of $D=0$ and $D=1$ produce lags differing by a factor of $< 1.5$ at frequencies close to the measured break-frequency $f_{b}$ and less at lower frequencies.755Observations during the last decade have led on one hand (o a remarkable breakthrough in (he identilication of the progenitor svstem. ancl on the other hand (o the realization (hat even the alterelow physics is much more complex than suspected.,"Observations during the last decade have led on one hand to a remarkable breakthrough in the identification of the progenitor system, and on the other hand to the realization that even the afterglow physics is much more complex than suspected."756 Theoretical modeling of the different physical processes improved. and became more detailed aud accurate. but also more complicated.," Theoretical modeling of the different physical processes improved and became more detailed and accurate, but also more complicated."757 Addiüonallv. new ingredients were invoked to explain various specilie observations of specific bursts. significantly increasing (he mocel's degrees of [reedom. sometimes to the point where the models do not provide useful predictions anymore.," Additionally, new ingredients were invoked to explain various specific observations of specific bursts, significantly increasing the model's degrees of freedom, sometimes to the point where the models do not provide useful predictions anymore."758 The purpose of the open discussion that I led in the conference was to take a step back and evaluate where do we stand in (he grand. picture., The purpose of the open discussion that I led in the conference was to take a step back and evaluate where do we stand in the grand picture.759 I presented to the audience three Unfortunately no detailed notes were taken during the discussion so here I will summarize onlv the introduction that I gave. which includes my short answers to these «questions.," I presented to the audience three Unfortunately no detailed notes were taken during the discussion so here I will summarize only the introduction that I gave, which includes my short answers to these questions."760 I divide mv discussion into live different stages of the explosion: progenitor. central engine. outflow properties. prompt emission and afterglow.," I divide my discussion into five different stages of the explosion: progenitor, central engine, outflow properties, prompt emission and afterglow."761 For each. I give (a very partial) answer to the three questions.," For each, I give (a very partial) answer to the three questions."762 I touch only the physics of the explosions themselves without. considering the study of Che large scale environment (e.g.. host galaxies).," I touch only the physics of the explosions themselves without considering the study of the large scale environment (e.g., host galaxies)."763" I discuss long and short GRBs separately only in the ""progenitor section."," I discuss long and short GRBs separately only in the “progenitor"" section."764 The reason is (hat as [ar as we know (he rest of the stages max be (ancl possibly are) rather similar., The reason is that as far as we know the rest of the stages may be (and possibly are) rather similar.765 Finally. one of the comments during the discussion (bv. Chris Frver) was that we are much better at ruling out models than al proposing viable ones.," Finally, one of the comments during the discussion (by Chris Fryer) was that we are much better at ruling out models than at proposing viable ones."766 Nevertheless. I discuss here only models which max describe the physies properly. without mentioning those that were ruled out.," Nevertheless, I discuss here only models which may describe the physics properly, without mentioning those that were ruled out."767 Given the very broad scope of the discussion and the limited space of this proceeding. ] veerettully cannot give the appropriate credit to the large number of observational anc," Given the very broad scope of the discussion and the limited space of this proceeding, I regretfully cannot give the appropriate credit to the large number of observational and"768 Given the very broad scope of the discussion and the limited space of this proceeding. ] veerettully cannot give the appropriate credit to the large number of observational ancl," Given the very broad scope of the discussion and the limited space of this proceeding, I regretfully cannot give the appropriate credit to the large number of observational and"769around 2-7 Gyr not significant.,around 2-7 Gyr not significant.770" When + increases to 2.0. the upper boundary of Mines and Av of the stars with age approximately 5 Gyr too.however.. the distributions of v,,,... and Av of the stars about 2-415 Gyr move down slightly. which leads to the peak locationΜΗ."," When $\eta$ increases to 2.0, the upper boundary of $\nu_{max}$ and $\Delta\nu$ of the stars with age approximately 5 Gyr too, the distributions of $\nu_{max}$ and $\Delta\nu$ of the stars about 2-4.5 Gyr move down slightly, which leads to the peak location."771 Fig., Fig.772 6 shows the histograms of ρα and Av of CHeB stars of simulated BSP with +) = 0.5 and the distributions of the Vine and Ag as a function of stellar age., \ref{fbsp} shows the histograms of $\nu_{max}$ and $\Delta\nu$ of CHeB stars of simulated BSP with $\eta$ = 0.5 and the distributions of the $\nu_{max}$ and $\Delta\nu$ as a function of stellar age.773 The distributions of Mines ind Av are very similar to those of the SSP with the same 5., The distributions of $\nu_{max}$ and $\Delta\nu$ are very similar to those of the SSP with the same $\eta$.774 This is because that most binary stars evolved into CHeB stage are wide binary stars., This is because that most binary stars evolved into CHeB stage are wide binary stars.775 Binary interactions such as mass transfer and mass accretion do not take effect in these wide binary stars., Binary interactions such as mass transfer and mass accretion do not take effect in these wide binary stars.776 However. the lower panels of Fig.," However, the lower panels of Fig."777 6. show that the value of Wines and Av of some stars with age = 2 Gyr is larger than 40 ΓΗΣ and 4 (/Hz. respectively.," \ref{fbsp} show that the value of $\nu_{max}$ and $\Delta\nu$ of some stars with age $>$ 2 Gyr is larger than 40 $\mu$ Hz and 4 $\mu$ Hz, respectively."778 This is due to the fact that binary interactions lead to an obvious increase or decrease in the mass of the stars when they evolve into the CHeB stage., This is due to the fact that binary interactions lead to an obvious increase or decrease in the mass of the stars when they evolve into the CHeB stage.779 Consequently. the Minos ad Av of these stars are larger than those of the wide binary stars.," Consequently, the $\nu_{max}$ and $\Delta\nu$ of these stars are larger than those of the wide binary stars."780 However. the fraction of these stars is very small.," However, the fraction of these stars is very small."781 Thus the effect of binary interactions on the distributions of ρω and Av is not significant., Thus the effect of binary interactions on the distributions of $\nu_{max}$ and $\Delta\nu$ is not significant.782 The star formation rate of the Galaxy is not a constant., The star formation rate of the Galaxy is not a constant.783 Rocha-Pintoal.(000b). gave that the SFR of the Galaxy at 2-5 Gyr ago is about 2 times larger than an average SFR over the past15 Gyr and the SFR of the Galaxy at 7-9 Gyr ago is around 1.5 times larger than the average SFR., \citet{roch00b} gave that the SFR of the Galaxy at 2-5 Gyr ago is about 2 times larger than an average SFR over the past15 Gyr and the SFR of the Galaxy at 7-9 Gyr ago is around 1.5 times larger than the average SFR.784 Fig., Fig.785 7. shows the distributions of ως and Av of CHeB stars of the simulated SSP with a constant SFR and with the Galaxy SFR given by Rocha-Pintoetal.(2000b)., \ref{fsfr} shows the distributions of $\nu_{max}$ and $\Delta\nu$ of CHeB stars of the simulated SSP with a constant SFR and with the Galaxy SFR given by \citet{roch00b}.786. Although the number of simulated CHeB stars is changed. the distributions are similar.," Although the number of simulated CHeB stars is changed, the distributions are similar."787 In fact. for jj 2 0.5. Fig.," In fact, for $\eta$ = 0.5, Fig."788 5 shows that the value of ο and Av of the CHeB stars with age 7 2 Gyr is mainly located in the range of 10-30 (Hz and 1-4 ΤΗΣ. and gathers around 25 (Hz and 3 j#Hz. respectively.," \ref{faged} shows that the value of $\nu_{max}$ and $\Delta\nu$ of the CHeB stars with age $>$ 2 Gyr is mainly located in the range of 10-30 $\mu$ Hz and 1-4 $\mu$ Hz, and gathers around 25 $\mu$ Hz and 3 $\mu$ Hz, respectively."789 Thus increasing or decreasing the SER of stars at a certain age > 2 Gyr can increase or decrease the number of CHeB stars but affect the peak αι and Av of CHeB stars., Thus increasing or decreasing the SFR of stars at a certain age $>$ 2 Gyr can increase or decrease the number of CHeB stars but affect the peak $\nu_{max}$ and $\Delta\nu$ of CHeB stars.790 increasing the SFR of the young stars 0-2 Gyr can increase the number of stars with relatively high ρω and Av. unless the SFR is enhanced many times.," increasing the SFR of the young stars 0-2 Gyr can increase the number of stars with relatively high $\nu_{max}$ and $\Delta\nu$ , unless the SFR is enhanced many times."791 The of simulated CHeB stars have a dominant peak. but the distribution of mass," The of simulated CHeB stars have a dominant peak, but the distribution of mass"792"determine a SER of 1021 AL. | from the reddeuine-corrected IL, huuinositv E (120270.08) το crest.",determine a SFR of $\pm$ 1 $M_\odot$ $^{-1}$ from the reddening-corrected $_\alpha$ luminosity of $\pm$ $\times$ $^{42}$ erg $^{-1}$.793 The method (again in Ixeunicut 1998) which instead. uses the extiuction-corrected [Ou] haninosity. (8.00.7) «1072 ere Ἐν yields a πιο larger SER values 1102830 AL. |l," The method (again in Kennicutt 1998) which instead uses the extinction-corrected [O ] luminosity, $\pm$ $\times$ $^{42}$ erg $^{-1}$, yields a much larger SFR value, $\pm$ 30 $M_\odot$ $^{-1}$."794" This Ligh value. so different from that obtained using the IL, cutission. may be produced by the abovementioned uucertaintv iu the flux calibration of the spectrum at its blue edge (see Sect."," This high value, so different from that obtained using the $_\alpha$ emission, may be produced by the abovementioned uncertainty in the flux calibration of the spectrum at its blue edge (see Sect."795 3). so we should treat this latter SER estimate verv cautiously.," 3), so we should treat this latter SFR estimate very cautiously."796" Moreover. as stressed hy Ieunieutt (1998). the SER determination from the [O 1] line flux suffers from larger uncertainties with respect to that obtaimed from the IL, eiission line."," Moreover, as stressed by Kennicutt (1998), the SFR determination from the [O ] line flux suffers from larger uncertainties with respect to that obtained from the $_\alpha$ emission line."797" The otal reddening estimate along the line of sight inferred roni the optical spectrum corresponds. using the Clupirica formula of Predehl Schinitt (1995). toa Αιzm 541073Di P2, whichB is. ~ 1 times+ less than that measured by Sazonoy et al. ("," The total reddening estimate along the line of sight inferred from the optical spectrum corresponds, using the empirical formula of Predehl Schmitt (1995), to a $N_{\rm H} \approx$ $\times$ $^{21}$ $^{-2}$, which is $\sim$ 4 times less than that measured by Sazonov et al. ("7982005) fromChandra N-vay data.,2005) from X-ray data.799 Next. asstmune for ICR 1610 the best-fit X- spectrin as in Sazonov et al. (," Next, assuming for IGR $-$ 1610 the best-fit X-ray spectrum as in Sazonov et al. ("8002005). we can determine the 210 keV flux of the source.,"2005), we can determine the 2–10 keV flux of the source."801 This results in 2.2.10 £2 ere ?os 1., This results in $\times$ $^{-12}$ erg $^{-2}$ $^{-1}$.802 The comparison between the reddening-corrected. [O Π) A5007 emission flux and the 210 keV N-ray flux estimated above implies an N-ray/[O Us007 ratio of —2.5. which indicates tha this source is in the Conmptou-thick regime (see Bassani ct al.," The comparison between the reddening-corrected [O ] $\lambda$ 5007 emission flux and the 2–10 keV X-ray flux estimated above implies an X-ray/[O $_{\rm 5007}$ ratio of $\sim$ 2.5, which indicates that this source is in the Compton-thick regime (see Bassani et al."803 1999)., 1999).804 Although. as we said above. the |OΠΠ emission line flix estimate is affected by large uncertainties. the detection of [O uj. [O ru]. and IL; also allows us to infer the gaseous oxveen abundance in this ealaxy.," Although, as we said above, the [O] emission line flux estimate is affected by large uncertainties, the detection of [O ], [O ], and $_\beta$ also allows us to infer the gaseous oxygen abundance in this galaxy."805 Following Robuluickv et al. (, Following Kobulnicky et al. (8061999). the Ro; parameter. defined as the ratio between |O uf | ο iu and IL; line fhixes. gives 19 | log (O/II) z 8.5.,"1999), the $R_{\rm 23}$ parameter, defined as the ratio between [O ] + [O ] and $_\beta$ line fluxes, gives 12 + log (O/H) $\approx$ 8.5."807 Considering the intrinsic luminosity of the source (it las rest-frame absolute B-baud magnitude Mp = 21.27 mae: Pruguiel 2005) and its [O 11]|/|N | ratio (~3). all of this information points to a basically solar oxveen abundance.," Considering the intrinsic luminosity of the source (it has rest-frame absolute $B$ -band magnitude $_B$ = $-$ 21.27 mag; Prugniel 2005) and its [O ]/[N ] ratio $\sim$ 3), all of this information points to a basically solar oxygen abundance."808" A similar result is obtained if we use the[N TL, flux ratio method Usewley Dopita 2002).", A similar result is obtained if we use the[N $_\alpha$ flux ratio method (Kewley Dopita 2002).809 As moeutioued in Sect., As mentioned in Sect.810 3. we also acquired a spectruu of 2NIASX. 161032. the edec-ou galaxy located l' cast of LEDA17019.," 3, we also acquired a spectrum of 2MASX $-$ 1610432, the edge-on galaxy located $'$ east of LEDA170194."811 The spectrum. albeit noisy (sce Fig.," The spectrum, albeit noisy (see Fig."812" 2. upper right). shows the presence of prominent and narrow IL, amd [N 1 AAG518.6583. cliission lines at redshift 2 = (.O7L40.001."," 2, upper right), shows the presence of prominent and narrow $_\alpha$ and [N ] $\lambda\lambda$ 6548,6583 emission lines at redshift $z$ = $\pm$ 0.001."813 In this case also. the NaD doublet in absorption is detected at lis same redshift.," In this case also, the NaD doublet in absorption is detected at this same redshift."814 This imuplies a huninosity distance dp = 315 Mpc for this ealaxy. thus twice as far from Earth as LEDA 17019," This implies a luminosity distance $d_L$ = 345 Mpc for this galaxy, thus twice as far from Earth as LEDA 170194."815" The detection of [N uj and IL, iu the spectiiu of 2ATASN 1670132 also allows us to inter the eascous oxveen abundance in this galaxy.", The detection of [N ] and $_\alpha$ in the spectrum of 2MASX $-$ 1610432 also allows us to infer the gaseous oxygen abundance in this galaxy.816" The use of the [N n]/IL, ratio method (among those isted in I&ewley Dopita 2002) for the determination of the metallicity of this galaxy is indicated because it isthe least scusitive to. aud therefore not substantially influenced by. the lack of our knowledge of the absorption intrinsic to this galaxy."," The use of the [N $_\alpha$ ratio method (among those listed in Kewley Dopita 2002) for the determination of the metallicity of this galaxy is indicated because it isthe least sensitive to, and therefore not substantially influenced by, the lack of our knowledge of the absorption intrinsic to this galaxy."817 Tudeed. the two emissiou lines are so close to cach other that the differeutial iutriusic reddening is not significant.," Indeed, the two emission lines are so close to each other that the differential intrinsic reddening is not significant."818" We thus fiud that the [N /TL, ratio observed here iniplies 12 | log (O/II) = 9.", We thus find that the [N $_\alpha$ ratio observed here implies 12 + log (O/H) $\approx$ 9.819 Therefore. in this case also we fud an oxveen abundance which is consistent with the solar value.," Therefore, in this case also we find an oxygen abundance which is consistent with the solar value."820" The iuteusitv of the IL, ciission line of 2\LASN 1610132. once corrected for the Galactic absorption. can also be used to estimate the SFR in this ealaxy."," The intensity of the $_\alpha$ emission line of 2MASX $-$ 1610432, once corrected for the Galactic absorption, can also be used to estimate the SFR in this galaxy."821 Using Eq. (, Using Eq. (8222) of IXcunicutt (1998) we determine a SFR of 1.l040.15 AZ. +.,2) of Kennicutt (1998) we determine a SFR of $\pm$ 0.15 $M_\odot$ $^{-1}$.823 This should conservatively be considered as a lower limit to the SFR because the effect ofabsorption intrinsic to 2ATASN 1610132 was not accounted for., This should conservatively be considered as a lower limit to the SFR because the effect of absorption intrinsic to 2MASX $-$ 1610432 was not accounted for.824" The poor S/N of the spectrum, of his source docs uot allow us to deduce mach more abou the nature of this narrow IL, enussiou-lne galaxy.", The poor S/N of the spectrum of this source does not allow us to deduce much more about the nature of this narrow $_\alpha$ emission-line galaxy.825 We of course can exclude that it isa Sevfert 1 type AGN due to he absence of broad chussion lines., We of course can exclude that it is a Seyfert 1 type AGN due to the absence of broad emission lines.826 But. eiven its optical (aud mavbhe radio) activity. we cannot exclude that this object is co-responsible. together with LEDA 170191. for the N-rax chussion detected by as ICR 1610.," But, given its optical (and maybe radio) activity, we cannot exclude that this object is co-responsible, together with LEDA 170194, for the X-ray emission detected by as IGR $-$ 1610."827 Dowever. a quick look at à 3.3 ks Chandra observation (Seq.," However, a quick look at a 3.3 ks Chandra observation (Seq."828 Num:, Num.:829 7011758. Obs.," 701178, Obs."830 ID: 6276. PE BR.À. Sunyaev) acquired on July 25. 2005. does not show detectable X-rav cluission either at the 2\TASN 1610132 position or within the NVSS 161011 radio error circle. whereas X-ravs are clearly detected fron LEDA 170191 (see also Halperu 2005 aud Sazonov et al.," ID: 6276, PI: R.A. Sunyaev) acquired on July 25, 2005, does not show detectable X-ray emission either at the 2MASX $-$ 1610432 position or within the NVSS $-$ 161041 radio error circle, whereas X-rays are clearly detected from LEDA 170194 (see also Halpern 2005 and Sazonov et al."831 2005)., 2005).832 This implies that either the soft (10 keV) X-ray enission. if any. frou the former sources is heavily absorbed or. more likely that they are not. N-rav cimittine and that LEDA 170191 is solely respousible for the hard N-ravs detected byI," This implies that either the soft $<$ 10 keV) X-ray emission, if any, from the former sources is heavily absorbed or, more likely, that they are not X-ray emitting and that LEDA 170194 is solely responsible for the hard X-rays detected by."833"NTEGRAL, Iu this latter case. the galaxy 2\TASN 1610132 can be identified as a ealaxy."," In this latter case, the galaxy 2MASX $-$ 1610432 can be identified as a galaxy."834 Regarding the association between this galaxy and the nearby NVSS radio source. we note that. given the relatively ow S/N ratio of the radio detection. the NVSS position aud the correspoucding uncertainties nav uot be very accurate: so. they Ισ actually be positiouallv cousistcut with each other.," Regarding the association between this galaxy and the nearby NVSS radio source, we note that, given the relatively low S/N ratio of the radio detection, the NVSS position and the corresponding uncertainties may not be very accurate; so, they might actually be positionally consistent with each other."835 Thus. oulv detailed radio observations can give an answer to this open issue.," Thus, only detailed radio observations can give an answer to this open issue."836 The optical spectrum of ο 106 is reported in the ceutra left panel of Fie., The optical spectrum of SS 406 is reported in the central left panel of Fig.837 2., 2.838 The absence of He lines poiuts to a D-star classification for this object., The absence of He lines points to a B-star classification for this object.839 Moreover. the shape of the Daliier absorption lines. along with the detection of fainter absorption features produced by Ie AA 1026.1171 and bv light metals (such as Si ALL28. C A1267 iux Met ALISI). poiuts to a 1uain-sequence. mid-tvpe D star (nost likely D5) ideutification.," Moreover, the shape of the Balmer absorption lines, along with the detection of fainter absorption features produced by He $\lambda\lambda$ 4026,4471 and by light metals (such as Si $\lambda$ 4128, C $\lambda$ 4267 and Mg $\lambda$ 4481), points to a main-sequence, mid-type B star (most likely B5) identification."840" Finally. the presence of a strong IL, line im ΟΙΙΟ (possibly showing a P-Cre"," Finally, the presence of a strong $_\alpha$ line in emission (possibly showing a P-Cyg"841Current Search [or Extraterrestrial Intelligence (SETI) programs concentrate on the search for radio or optical laser pulses emissions (Tarler2001).,Current Search for Extraterrestrial Intelligence (SETI) programs concentrate on the search for radio or optical laser pulses emissions \citep{tarter2001}.842. We propose here an alternative approach lor a new SETI: considering (hat artilicial planet-size bodies mav exist around other stars. and that such objects always. transit in front of their. parent star for a given remote observer. we may (hus have an opportunitv to detect aud even characterize them bv the transit method. assuming these transits are distinguishable from a simple planetary," We propose here an alternative approach for a new SETI: considering that artificial planet-size bodies may exist around other stars, and that such objects always transit in front of their parent star for a given remote observer, we may thus have an opportunity to detect and even characterize them by the transit method, assuming these transits are distinguishable from a simple planetary"843companiols in our fields is thus à maximum of =9x107 stars/aresec-4 (at confidence).,companions in our fields is thus a maximum of $\approx$ $\times$$10^{-5}$ $\rm arcsec^{2}$ (at confidence).844 Our total companion search area at <3’ “radius 1s 2007 aresee-: we thus expect 0.035 false associatior sin our dataset. or equivalently à probability of a single false positive association.," Our total companion search area at $<$ radius is 2007 $\rm arcsec^2$; we thus expect 0.035 false associations in our dataset, or equivalently a probability of a single false positive association."845 This limit is correct for companions with AM<4 (those detectable in all our observations)., This limit is correct for companions with $\Delta M$$\lsim$ 4 (those detectable in all our observations).846 Our faintest candidate. SLW1548+0443Ab. has a somewhat lower formal association probability because it would not have been detected in our Palomar imaging. and we leave its confirmation for CPM followup (Section 4.6.0)).," Our faintest candidate, SLW1548+0443Ab, has a somewhat lower formal association probability because it would not have been detected in our Palomar imaging, and we leave its confirmation for CPM followup (Section \ref{1548_comp}) )."847 Eleven of the 71 observed stars are actually close binaries. leading to à raw binary fraction of 2% (with a binomial uncertainty estimation).," Eleven of the 71 observed stars are actually close binaries, leading to a raw binary fraction of ${16}^{+5}_{-3}$ (with a binomial uncertainty estimation)."848 However. our 163common-proper-motion wide binary selection imposes several biases on the binarity fraction of our target sample: 1) the selection of photometrically-clean. non-extended SDSS objects removes some binaries from the sample: 2) binaries are brighter than single stars at the same colors in unresolved photometry. and so our magnitude-limited sample selects binaries from a larger space volume than the single stars (?):: 3) those binaries found at larger distances are also likely to have decreased proper motion. and so may be removed from our target sample (?):: 4) the widely-separated components of high-order multiple systems containing close pairs can appear to be at different photometric distances in unresolved photometry. removing them from the sample. and 5) the most distant (and therefore the physically widest) binaries will be de-selected by our minimum-separation cut.," However, our common-proper-motion wide binary selection imposes several biases on the binarity fraction of our target sample: 1) the selection of photometrically-clean, non-extended SDSS objects removes some binaries from the sample; 2) binaries are brighter than single stars at the same colors in unresolved photometry, and so our magnitude-limited sample selects binaries from a larger space volume than the single stars \citep{Burgasser2003}; 3) those binaries found at larger distances are also likely to have decreased proper motion, and so may be removed from our target sample \citep{Law2008}; 4) the widely-separated components of high-order multiple systems containing close pairs can appear to be at different photometric distances in unresolved photometry, removing them from the sample, and 5) the most distant (and therefore the physically widest) binaries will be de-selected by our minimum-separation cut."849lit can in turn be expressed as a lint ou the Iuninositv eaM. which is also shown in Figure 11.,"limit can in turn be expressed as a limit on the luminosity $\epsilon_0 \dot M_o c^2$, which is also shown in Figure 14."850" The luninosity upper liuüt for the infinite efficiency disk is much smaller than for the Novikov-Thorue disk since P/r peaks at rji. while for the Novikov-Thorne disk h/r peaks at larecr radius. where its maguitude is smaller ρω occurs at r=2lr, fora. =0 ud c=Tr, for e.=1]."," The luminosity upper limit for the infinite efficiency disk is much smaller than for the Novikov-Thorne disk since $h/r$ peaks at $r_{ms}$, while for the Novikov-Thorne disk $h/r$ peaks at larger radius, where its magnitude is smaller $(h/r)_{max}$ occurs at $r=24r_g$ for $a_*=0$ and $r=7r_g$ for $a_*=1$ ]."851" If either eyM,tD or ελt{οP exceeds their. respective. limits..+ then the thiu-diskc approximation breaks down."," If either $\epsilon_0 \dot M_o c^2$ or $\Delta \epsilon852\dot M_o c^2$ exceeds their respective limits, then the thin-disk approximation breaks down."853 In addition. if hir is sinall. then the approximation of a flat disk iu the coniputation of the returuing radiation will be appropriate.," In addition, if $h/r$ is small, then the approximation of a flat disk in the computation of the returning radiation will be appropriate."854 To treat the interesting cases where L—Lg will require a 2-D solution of the disk equations. which is bevonud the scope of this work.," To treat the interesting cases where $L \sim L_{Edd}$ will require a 2-D solution of the disk equations, which is beyond the scope of this work."855 The returning radiation will not affect the disk height since it diffuses through the disk on a thermal timescale. so there is no uct flux duc to returuiug radiation (unless the disk is wvarped).," The returning radiation will not affect the disk height since it diffuses through the disk on a thermal timescale, so there is no net flux due to returning radiation (unless the disk is warped)."856 21un 2nuu We have generalized. the equations for au azimuthally sviunietric. ecolmetrically thin. time-steady accretion disk around a black hole to include the effects of a torque operating at the iuuner boundary. taken to be at rj.," 2mm 2mm We have generalized the equations for an azimuthally symmetric, geometrically thin, time-steady accretion disk around a black hole to include the effects of a torque operating at the inner boundary, taken to be at $r_{ms}$."857 Coustant non-zero torque at ry). causes several plysical consequences that change the fundamental properties of the accretion flow: 1) The dus cau be expressed as a stun of the usual Novikov Thorne expression plus a part due to the torque which scales roughly as &67?, Constant non-zero torque at $r_{ms}$ causes several physical consequences that change the fundamental properties of the accretion flow: 1) The flux can be expressed as a sum of the usual Novikov Thorne expression plus a part due to the torque which scales roughly as $r^{-7/2}$.858 2) The acerction cficiency has a fundamental upper lait due to the second law of black hole dynamics for as<0.36., 2) The accretion efficiency has a fundamental upper limit due to the second law of black hole dynamics for $a_* < 0.36$.859 For larger a... infiudute efüciency is possible iu principle.," For larger $a_*$, infinite efficiency is possible in principle."860" 3) The black hole spin cau reach au equilibrimim for a,m0.998 since the augular momentum reaching the hole is znaller.", 3) The black hole spin can reach an equilibrium for $a_* < 0.998$ since the angular momentum reaching the hole is smaller.861 Raciation cau also exert a significant torque on the black hole. which changes the value of the equilibriumspin efficiency.," Radiation can also exert a significant torque on the black hole, which changes the value of the equilibrium-spin efficiency."862 Above au efficiency. of €=0.36. the black hole mast always be spun down.," Above an efficiency of $\epsilon = 0.36$, the black hole must always be spun down."863 1) Since the extra cuussivity is peaked at the ner οσο of the disk. if rj). is stuall then eravity causes a large fraction of the radiation (up to 58%)) to return to the accretion disk.," 4) Since the extra emissivity is peaked at the inner edge of the disk, if $r_{ms}$ is small then gravity causes a large fraction of the radiation (up to ) to return to the accretion disk."864 The fux of returning radiation scales as r3i)> at lareo radius., The flux of returning radiation scales as $r^{-3}$ at large radius.865 Up to of the radiation cau be captured by the black hole., Up to of the radiation can be captured by the black hole.866 5) The extra heating within the disk will iucrease the height of the disk if it is radiation pressure-supported., 5) The extra heating within the disk will increase the height of the disk if it is radiation pressure-supported.867 This limits severely the luminosity at which the thiu-cdisk approxiuation is appropriate., This limits severely the luminosity at which the thin-disk approximation is appropriate.868 6) Doppler beaming aud relativistic beudiug are stronecst in the inner parts of the accretion disk where the extra flux peaks. so that for large a. aud e. the disk will be lmb-briehtenued.," 6) Doppler beaming and relativistic bending are strongest in the inner parts of the accretion disk where the extra flux peaks, so that for large $a_*$ and $\epsilon$, the disk will be limb-brightened."869 These cach have multiple observable cousequences: 1) The extra surface brightuess changes the locally radiated spectrum, These each have multiple observable consequences: 1) The extra surface brightness changes the locally radiated spectrum.870 Though the local surface brightuess is usually not directly observable. it may be possible to lap it using several devices: eclipse mapping (Baptista ct al.," Though the local surface brightness is usually not directly observable, it may be possible to map it using several devices: eclipse mapping (Baptista et al."871 1998). although no eclipsing black hole N-xav. binaries have been discovered vet: quasar mieroleusine (Agol Isvolik 1999): or reverberation mapping (Collier et al.," 1998), although no eclipsing black hole X-ray binaries have been discovered yet; quasar microlensing (Agol Krolik 1999); or reverberation mapping (Collier et al."872 1999)., 1999).873 If the Fe Ίνα enuissivitv is proportional to the local dissipation. these effects can strenethen the red wing of the hue. particularly when the spin is simall.," If the Fe $\alpha$ emissivity is proportional to the local dissipation, these effects can strengthen the red wing of the line, particularly when the spin is small."874 This effect may uudercut the argument that lines with stroug red wines came from disks around black holes with higher spin (e.g. Dabrowski et al., This effect may undercut the argument that lines with strong red wings came from disks around black holes with higher spin (e.g. Dabrowski et al.875 1997)., 1997).876 Several authors have used the Novikov-Thorne model to fit the soft N-rav spectra of galactic black hole candidates., Several authors have used the Novikov-Thorne model to fit the soft X-ray spectra of galactic black hole candidates.877 Their procedure was to estimate the effective radiating area required to cut the observed dhwunuinositv at the observed effective temperature., Their procedure was to estimate the effective radiating area required to emit the observed luminosity at the observed effective temperature.878 Ou tlie basis of these fits they inferred that some black holes have rather high spius because the effective radiatiug area of a Novikov-Thorue disk decreases with iucreasing spin (Zhang. Cui. Chenu 1997).," On the basis of these fits they inferred that some black holes have rather high spins because the effective radiating area of a Novikov-Thorne disk decreases with increasing spin (Zhang, Cui, Chen 1997)."879 ILowever. for fixed spin aud ceutral mass. a disk with large Ae has a sinaller effective radiating area than a Novikov-Thorne disk. mimicking the effect of ereater spin.," However, for fixed spin and central mass, a disk with large $\Delta \epsilon$ has a smaller effective radiating area than a Novikov-Thorne disk, mimicking the effect of greater spin."880 2) The outer parts of accretion disks can be unstable to warping due to iradiation from the center (Pringle, 2) The outer parts of accretion disks can be unstable to warping due to irradiation from the center (Pringle881"2010) and the structure of the interstellar magnetic field, it is more likely that small angular scale anisotropies are generated by some local effect.","2010) and the structure of the interstellar magnetic field, it is more likely that small angular scale anisotropies are generated by some local effect."882 Figure 4 represents the possible structure of the heliotail which arises from the solar magnetic field cycles 1979)., Figure \ref{structure1} represents the possible structure of the heliotail which arises from the solar magnetic field cycles (Parker 1979).883 The magnetic fields of the opposite polarities (Parkeremerge as the result of 11 year solar dynamo cycle., The magnetic fields of the opposite polarities emerge as the result of 11 year solar dynamo cycle.884" As the magnetic field is carried away by solar wind, the reversed magnetic field regions get accumulated in the magnetotail region."," As the magnetic field is carried away by solar wind, the reversed magnetic field regions get accumulated in the magnetotail region."885 This is where reconnection is expected to occur., This is where reconnection is expected to occur.886" Naturally, the actual heliotail is going to be turbulent, which is not represented by the idealized drawing in Figure 4.."," Naturally, the actual heliotail is going to be turbulent, which is not represented by the idealized drawing in Figure \ref{structure1}."887" As the Alfven speed is smaller than the Solar wind speed, magnetic reconnection does not change the overall magnetic field structure."," As the Alfven speed is smaller than the Solar wind speed, magnetic reconnection does not change the overall magnetic field structure."888" Nevertheless, as we discuss in 85, the effects of turbulence are very important from the point of view of magnetic reconnection and the particle acceleration that it entails."," Nevertheless, as we discuss in 5, the effects of turbulence are very important from the point of view of magnetic reconnection and the particle acceleration that it entails."889 The simulations of the magnetotail are extremely challenging Pogorelov et al., The simulations of the magnetotail are extremely challenging (see Pogorelov et al.890 2009ab) and have not been done with(see the sufficient resolution and extent., 2009ab) and have not been done with the sufficient resolution and extent.891" While we believe that future research will provide details necessary for quantitative modeling, the schematic representation of the magnetotail structure depicted in Figure 4 is true in terms of major features."," While we believe that future research will provide details necessary for quantitative modeling, the schematic representation of the magnetotail structure depicted in Figure \ref{structure1} is true in terms of major features."892" In what follows, it will be used for describing the scenario for the origin of the cosmic ray excess that we advocate in this paper."," In what follows, it will be used for describing the scenario for the origin of the cosmic ray excess that we advocate in this paper."893 Astrophysical plasmas are often highly ionized and highly magnetized (Parker 1970)., Astrophysical plasmas are often highly ionized and highly magnetized (Parker 1970).894" The evolution of the magnetic field in a highly conducting fluid can be described by a simple version of the induction equation where B is the magnetic field, vis the velocity field, and η is the resistivity coefficient."," The evolution of the magnetic field in a highly conducting fluid can be described by a simple version of the induction equation where $\vec{B}$ is the magnetic field, $\vec{v}$is the velocity field, and $\eta$ is the resistivity coefficient."895 Under most circumstances this is adequate for discussing the evolution of magnetic field in an astrophysical plasma., Under most circumstances this is adequate for discussing the evolution of magnetic field in an astrophysical plasma.896" When the dissipative term on the right hand side is small, as is implied by simple dimensional estimates, the magnetic flux through any fluid element is constant in time and the field topology is an invariant of motion."," When the dissipative term on the right hand side is small, as is implied by simple dimensional estimates, the magnetic flux through any fluid element is constant in time and the field topology is an invariant of motion."897" On the other hand, reconnection is observed in the solar corona and chromosphere et al."," On the other hand, reconnection is observed in the solar corona and chromosphere (Innes et al."898" 1997, Yokoyama Shibata 1995, Masuda et al."," 1997, Yokoyama Shibata 1995, Masuda et al."899" (Innes1994, Ciaravella Raymond 2008), its presence is required to explain dynamo action in stars and galactic disks (Parker 1970, and the violent relaxation of magnetic fields following a 1993),change in topology is a promising process for the First order Fermi acceleration of high energy particles in the universe (de Gouveia Dal Pino Lazarian 2003, henceforth GL03, 2005, Lazarian 2005, Drake et al."," 1994, Ciaravella Raymond 2008), its presence is required to explain dynamo action in stars and galactic disks (Parker 1970, 1993), and the violent relaxation of magnetic fields following a change in topology is a promising process for the First order Fermi acceleration of high energy particles in the universe (de Gouveia Dal Pino Lazarian 2003, henceforth GL03, 2005, Lazarian 2005, Drake et al."900" 2006, Lazarian Opher 2009, Drake et al."," 2006, Lazarian Opher 2009, Drake et al."901 2010)., 2010).902" Quantitative general estimates for the speed of reconnection start with two adjacent volumes with different large scale magnetic field directions (Sweet 1958, Parker 1957)."," Quantitative general estimates for the speed of reconnection start with two adjacent volumes with different large scale magnetic field directions (Sweet 1958, Parker 1957)."903" The speed of reconnection, i.e. the speed at which inflowing magnetic field is annihilated by Ohmic dissipation, is roughly 7/A, where A is the width of the transition zone (see Figure 5))."," The speed of reconnection, i.e. the speed at which inflowing magnetic field is annihilated by Ohmic dissipation, is roughly $\eta/\Delta$, where $\Delta$ is the width of the transition zone (see Figure \ref{fig:recon1}) )."904" Since the entrained plasma follows the local field lines, and exits through the edges of the current sheet at roughly the Alfven speed, Va, the resulting reconnection speed is reduced compared"," Since the entrained plasma follows the local field lines, and exits through the edges of the current sheet at roughly the Alfven speed, $V_A$ , the resulting reconnection speed is reduced compared"905Observations by have shown that the gas does not in fact cool below 2kkeV (the so-called temperature Iloor) begging us to identify a source of heating which can olfset the radiative cooling.,Observations by have shown that the gas does not in fact cool below keV (the so-called temperature floor) begging us to identify a source of heating which can offset the radiative cooling.906 At the present time. the main observational cliagnostics for radio-galaxy/LOAL interactions are the spatial mapping in the N-rav. band. of temperature. pressure. cooling time and metal abundances of the LOCAL across the interaction region.," At the present time, the main observational diagnostics for radio-galaxy/ICM interactions are the spatial mapping in the X-ray band of temperature, pressure, cooling time and metal abundances of the ICM across the interaction region."907 However. with current data we lack any direct probe of the dvnamies of the interaction. that is the velocities of the Iuid. disturbances. shocks and turbulence.," However, with current data we lack any direct probe of the dynamics of the interaction, that is the velocities of the fluid disturbances, shocks and turbulence."908 X major rope for the observatory was that it would open a direct window on the dvnamics of these interactions through he velocity. profiles of the ICM emission lines (Brigeeen. lloeft Ruszkowski 2005): of course. the ervostat ziblure and subsequent loss of the high-resolution N-rav Spectrometer (XIUS) prevents this goal from being realized.," A major hope for the observatory was that it would open a direct window on the dynamics of these interactions through the velocity profiles of the ICM emission lines (Brügggen, Hoeft Ruszkowski 2005); of course, the cryostat failure and subsequent loss of the high-resolution X-ray Spectrometer (XRS) prevents this goal from being realized."909 Barring the emergence of a new mission possessing ahigh-resolution X-ray spectrometer. this important science goal must now await the launch ofConsfellalion-N.," Barring the emergence of a new mission possessing ahigh-resolution X-ray spectrometer, this important science goal must now await the launch of."910 The current baseline design. for has a requirement on the spatial resolution of aaresec (hall-power diameter)., The current baseline design for has a requirement on the spatial resolution of arcsec (half-power diameter).911 Coupled with its superior spectral resolution. this is sullicient to allow detailed. mapping of radio-galaxy driven LOCAL dynamics through emission line profiles in the nearest svstems such as theVargo andPerseus clusters.," Coupled with its superior spectral resolution, this is sufficient to allow detailed mapping of radio-galaxy driven ICM dynamics through emission line profiles in the nearest systems such as the and clusters."912 For more distant. systems. however. the interaction region is contained within a small number of spatial resolution elements. hampering the ability of the observation to constrain robust dynamical signatures.," For more distant systems, however, the interaction region is contained within a small number of spatial resolution elements, hampering the ability of the observation to constrain robust dynamical signatures."913 In this paper. we discuss the complementary technique of probing these racio-galaxy/ICM interactions through absorption line spectroscopy of the N-rav luminous core of the central racio galaxy. looking for jet-induced changes to the strength and line profile of the absorption lines present even in a static ICM.," In this paper, we discuss the complementary technique of probing these radio-galaxy/ICM interactions through absorption line spectroscopy of the X-ray luminous core of the central radio galaxy, looking for jet-induced changes to the strength and line profile of the absorption lines present even in a static ICM."914 While detection and characterization of these absorption line features is more observationally demanding. it has the major advantage of probing the kinematic state of the eas through a well-defined: “core sample” of the galaxy cluster. even if ones X-ray observatory has only limited spatial resolution.," While detection and characterization of these absorption line features is more observationally demanding, it has the major advantage of probing the kinematic state of the gas through a well-defined “core sample” of the galaxy cluster, even if ones X-ray observatory has only limited spatial resolution."915 We focus on predictions of resonance Ix-shell absorption features of Le and Lloe-like oxygen and iron arising in the ICM around an active as well as inactive racio ealaxy. examining the velocity structure and strength of the lines.," We focus on predictions of resonance K-shell absorption features of H- and He-like oxygen and iron arising in the ICM around an active as well as inactive radio galaxy, examining the velocity structure and strength of the lines."916 These radio-galaxv/1€CM. interactions are. particularly interesting and important since they appear to be the loca and accessible examples of the more general phenomena of “AGN feedback and potentially crucial for some aspects of cosmological structure formation (e.g.. sce Benson e al.," These radio-galaxy/ICM interactions are particularly interesting and important since they appear to be the local and accessible examples of the more general phenomena of “AGN feedback” and potentially crucial for some aspects of cosmological structure formation (e.g., see Benson et al."917 2003 and references therein)., 2003 and references therein).918 However. there is another iniportant reason to study this phenomenon: it is of direc relevance to our observational understanding of the Warn-Lot Intergalactic Medium (WIAD).," However, there is another important reason to study this phenomenon; it is of direct relevance to our observational understanding of the Warm-Hot Intergalactic Medium (WHIM)."919 Phe filamentary WIL is believed. to be the repository for half of the barvons in the local. Universe. anc so is of obvious interest to cosmologists., The filamentary WHIM is believed to be the repository for half of the baryons in the local Universe and so is of obvious interest to cosmologists.920 It is possible to detect the WIIEM. by. wide field imaging of dilfuse EUV and soft. X-ray emission (see for example Ixaastra et al., It is possible to detect the WHIM by wide field imaging of diffuse EUV and soft X-ray emission (see for example Kaastra et al.921 2003)., 2003).922 LEowever. the most important diagnostic of the WILAL ave UV. and X-ray absorption lines of oxygen seen in the spectra of bright ACGN.," However, the most important diagnostic of the WHIM are UV and X-ray absorption lines of oxygen seen in the spectra of bright AGN."923 Narrow oxvecn OVIOVILL absorption lines at. a. redshift’ intermediate between zero and that of the AGN are taken to be good candidates for the WIILA, Narrow oxygen OVI–OVIII absorption lines at a redshift intermediate between zero and that of the AGN are taken to be good candidates for the WHIM.924L However. all AGN towards which such lines have been claimed are radio-oud AGN (see Section 5).," However, all AGN towards which such lines have been claimed are radio-loud AGN (see Section 5)."925 The question arises whether these absorption ine features are. in fact. outLlows associated with the jettod AGN.," The question arises whether these absorption line features are, in fact, outflows associated with the jetted AGN."926 While the possibility remains that these features could x due to outllows associated. with the central engine itself. we demonstrate that it is not. possible to associate WILILM-ike features with racio-galaxy/I1CM interactions.," While the possibility remains that these features could be due to outflows associated with the central engine itself, we demonstrate that it is not possible to associate WHIM-like features with radio-galaxy/ICM interactions."927 We predict the absorption. line properties οἱ radio-galaxv/1€M interactions by using the NSTAR photoionization cock to post-process the hydrodynamies simulations of Itevnolds. Heinz Begelmanὃν (2002).," We predict the absorption line properties of radio-galaxy/ICM interactions by using the XSTAR photoionization code to post-process the hydrodynamics simulations of Reynolds, Heinz Begelman (2002)."928 Section 2 of this paper describes the underlying hyvdrodynamic simulations of the radio galaxy we use., Section 2 of this paper describes the underlying hydrodynamic simulations of the radio galaxy we use.929 Section 3 explains he simulation of our spectra using ASTAR., Section 3 explains the simulation of our spectra using XSTAR.930 In Section 4. we summarize the properties (strengths ancl velocity structure) of the absorption lines that we predict. as a unction of the inclination with which we view the racio galaxy. the age of the radio-galaxy. and the properties of he ICM into which the raclio-galaxy is expanding.," In Section 4, we summarize the properties (strengths and velocity structure) of the absorption lines that we predict as a function of the inclination with which we view the radio galaxy, the age of the radio-galaxy, and the properties of the ICM into which the radio-galaxy is expanding."931" Section 5 argues that these interactions cannot produce 7MLIIEM-like"" eatures. thereby. strengthening the WIIEM. interpretation of these intermediate. redshift) absorption lines."," Section 5 argues that these interactions cannot produce “WHIM-like” features, thereby strengthening the WHIM interpretation of these intermediate redshift absorption lines."932 Section 6 ocuses on the ability of current future X-ray observatories o study these absorption lines and hence prove these radio-galaxv/ICM interactions., Section 6 focuses on the ability of current future X-ray observatories to study these absorption lines and hence prove these radio-galaxy/ICM interactions.933 Finally. Section 7 presents our conclusions.," Finally, Section 7 presents our conclusions."934 The hvdrodynamic simulation underlying this work is the canonical axisvmnmietric simulation of Reynolds. Heinz Degelman (2002): for convenience. we briclly summarize the relevant details of the simulation here.," The hydrodynamic simulation underlying this work is the canonical axisymmetric simulation of Reynolds, Heinz Begelman (2002); for convenience, we briefly summarize the relevant details of the simulation here."935 This) 2-cd hydrodynamic simulation explores the evolution of an activejet in a radio galaxy as it interacts with a surrounding ICM. and the resulting structure after the jet is shut down.," This 2-d hydrodynamic simulation explores the evolution of an active jet in a radio galaxy as it interacts with a surrounding ICM, and the resulting structure after the jet is shut down."936 Initially. we start with an isothermal spherical gascous atmosphere with a (number) density. profile given by The (fixed) gravitational potential is defined such that this atmosphere is in hydrostatic equilibrium.," Initially, we start with an isothermal spherical gaseous atmosphere with a (number) density profile given by The (fixed) gravitational potential is defined such that this atmosphere is in hydrostatic equilibrium."937 Back-to-back supersonic jets are injected in initial pressure balance with a density of mi=ro100 and Mach number (with respect to the internal jet sound. speed) of 10., Back-to-back supersonic jets are injected in initial pressure balance with a density of $n_{\rm jet}=n_0/100$ and Mach number (with respect to the internal jet sound speed) of 10.938 The half-opening angle of the jet is 15. and the starting point for the jets is a radius of ry 720.," The half-opening angle of the jet is $15^{\circ}$, and the starting point for the jets is a radius of $r_0/20$ ."939 As discussed in the Appendix of Reynolds.," As discussed in the Appendix of Reynolds,"940ambiguity in defining the center could. actually be part of the puzzle material co-rotating with respect to the center as defined in one way may be counter-rotating with respect to the center defined in another.,ambiguity in defining the center could actually be part of the puzzle – material co-rotating with respect to the center as defined in one way may be counter-rotating with respect to the center defined in another.941 Large-scale νοντοςνnamic interactions seem to be important: many of the counter-rotating or misaligned. nuclei in the present sample form when galaxy pairs arrive at their second passage still dressed in extensive eas disks., Large-scale hydrodynamic interactions seem to be important; many of the counter-rotating or misaligned nuclei in the present sample form when galaxy pairs arrive at their second passage still dressed in extensive gas disks.942 A better understanding of the roles of eravitational and hyvdrodynamic torques in the formation of counter-rotating nuclei may require caleulations with higher spatial resolution than available at. present., A better understanding of the roles of gravitational and hydrodynamic torques in the formation of counter-rotating nuclei may require calculations with higher spatial resolution than available at present.943 The results presented here invite comparisons with observations of gas disks in merger remnants., The results presented here invite comparisons with observations of gas disks in merger remnants.944 As mentioned in the introduction. NGC 7252 is an obvious example: the central disk of ionized ancl molecular gas in this svsteni is about the same size as the disks. produced. in. these experiments. and the peculiar kinematics of this disk may be due to à strong warp of the kind seen in many of these simulations.," As mentioned in the introduction, NGC 7252 is an obvious example; the central disk of ionized and molecular gas in this system is about the same size as the disks produced in these experiments, and the peculiar kinematics of this disk may be due to a strong warp of the kind seen in many of these simulations."945 In addition. velocities confirm. that gas is Falling back into the remnant from the tidal tails (Libbzux et al.," In addition, velocities confirm that gas is falling back into the remnant from the tidal tails (Hibbard et al."946 1994)., 1994).947 NGC 7252 has been modeled as the resul of a close and fairly direct. encounter between two disk galaxies of comparable mass (Llibhard Mihos 1995)., NGC 7252 has been modeled as the result of a close and fairly direct encounter between two disk galaxies of comparable mass (Hibbard Mihos 1995).948 Phis nmoclel nicely. reproduced the morphology ancl kinematics of the tails. but the calculations did not include a dissipative component which could. form a central disk.," This model nicely reproduced the morphology and kinematics of the tails, but the calculations did not include a dissipative component which could form a central disk."949 Its worth repeating this calculation with a combined. N-body/SP code: a single model reproducing both the tails the central disk would be a fairly impressive accomplishment. and might help constrain the gas content of NGC 7252's POSCHILOLS.," It's worth repeating this calculation with a combined N-body/SPH code; a single model reproducing both the tails the central disk would be a fairly impressive accomplishment, and might help constrain the gas content of NGC 7252's progenitors."950 The peculiar elliptical galaxy NCC 3656 (Balcells 1997. Baleells οἱ al.," The peculiar elliptical galaxy NGC 3656 (Balcells 1997, Balcells et al."951 2001) may be a second example., 2001) may be a second example.952 Fhis galaxy contains shells ancl a pair of faint tails which suggest a merger of two disk galaxies., This galaxy contains shells and a pair of faint tails which suggest a merger of two disk galaxies.953 Ht also has a star-forming cust ane which corresponds to an extended: disk of1., It also has a star-forming dust lane which corresponds to an extended disk of.954 This disk is visibly warped. and the outer edge of the disk. is kinematically contiguous with gas at larger radii which may »e falling in from the tidal tails.," This disk is visibly warped, and the outer edge of the disk is kinematically contiguous with gas at larger radii which may be falling in from the tidal tails."955 The estimated gas accretion rate in NGC 3656 is an order of magnitude lower than in GC. 7252: this may imply that NGC 3656 is at a later stage in its evolution than NGC 7252., The estimated gas accretion rate in NGC 3656 is an order of magnitude lower than in NGC 7252; this may imply that NGC 3656 is at a later stage in its evolution than NGC 7252.956 NGC 5128. an elliptical galaxy with an active nucleus and an extended cisk of dust ancl gas. may represent an even later stage in the evolution of gas-vich merger remnants.," NGC 5128, an elliptical galaxy with an active nucleus and an extended disk of dust and gas, may represent an even later stage in the evolution of gas-rich merger remnants."957 Although the warped disk in this svstem. has been interpreted as the remains of an accreted. gas-rich satellite (e.g. Malin. Quinn. Graham. 1983). there are some grounds to suspect that this galaxy is the result. of a fairly major merger (e.g. Schweizer 1998).," Although the warped disk in this system has been interpreted as the remains of an accreted gas-rich satellite (e.g. Malin, Quinn, Graham 1983), there are some grounds to suspect that this galaxy is the result of a fairly major merger (e.g. Schweizer 1998)."958 These include the ripples visible in deep optical images (Malin et al., These include the ripples visible in deep optical images (Malin et al.959 1983). the fragments surrounding the galaxy (Schiminovich et al.," 1983), the fragments surrounding the galaxy (Schiminovich et al."960 1994). and the misaligned rotation revealed by kinematics of planetary. nebulae (ΕΠ et al.," 1994), and the misaligned rotation revealed by kinematics of planetary nebulae (Hui et al."961 1995)., 1995).962 eas disks ancl rings seem {ο be very common in merging galaxies selected by infrared Luminosity (e.g. Downes Solomon 1998)., gas disks and rings seem to be very common in merging galaxies selected by infrared luminosity (e.g. Downes Solomon 1998).963 In remnants with single nuclei the eas tends to rotate more rapidly than the stars. indicating some degree of kinematic decoupling (Genzel et al.," In remnants with single nuclei the gas tends to rotate more rapidly than the stars, indicating some degree of kinematic decoupling (Genzel et al."964 2001). though it’s not clear if the extreme decoupling in remnant POL 1:1 € has vet been seen in the observations.," 2001), though it's not clear if the extreme decoupling in remnant POL 1:1 C has yet been seen in the observations."965 Disks in svstems which appear to still have nuclei are not so casily explained., Disks in systems which appear to still have nuclei are not so easily explained.966 \lolecular-lince observations of Arp 220 and NGC 6240 have been interpreted in terms of a rotating disk located: the nuclei (Scoville. Yun. Bevant 1997: Tacconi et al.," Molecular-line observations of Arp 220 and NGC 6240 have been interpreted in terms of a rotating disk located the nuclei (Scoville, Yun, Bryant 1997; Tacconi et al."967 1999: οσα et al., 1999; Tecza et al.968 2000)., 2000).969 In the simulations. gas driven inward before the galaxies merge alwavs accumulates in disks or bars around the individual nuclei: moreover. an incipient disk between the nuclei would be torn apart by gravitational fields as the orbits of the nuclei decay.," In the simulations, gas driven inward before the galaxies merge always accumulates in disks or bars around the individual nuclei; moreover, an incipient disk between the nuclei would be torn apart by gravitational fields as the orbits of the nuclei decay."970 An inter-nuclear disk might be stable if it was more massive than either nucleus (Leeza et al., An inter-nuclear disk might be stable if it was more massive than either nucleus (Tecza et al.971 2000). but the amount of eas required seems quite extravagant. and earlier reports of a peak in the stellar velocity dispersion between he nuclei of NCC 6240 (Lester CGallnev 1994: Dovon et al.," 2000), but the amount of gas required seems quite extravagant, and earlier reports of a peak in the stellar velocity dispersion between the nuclei of NGC 6240 (Lester Gaffney 1994; Doyon et al."972 1994) are not supported by recent LIST observations (CGorssen et al., 1994) are not supported by recent HST observations (Gerssen et al.973 2001)., 2001).974" On the other iud. the interpretation of Arp 220's nuclei as bright spots in πα warped molecular gas disk"" (Lckart Downes 2001: Downes Solomon 1993) seems entirely consistent with the esent numerical results."," On the other hand, the interpretation of Arp 220's nuclei as bright spots in “a warped molecular gas disk” (Eckart Downes 2001; Downes Solomon 1998) seems entirely consistent with the present numerical results."975 Simulations indicate that extended: gas. disks can form in mergers of spiral galaxies., Simulations indicate that extended gas disks can form in mergers of spiral galaxies.976" These disks contain between 20 ancl GO percent of the total gas in the initial galaxies. aud may extend to several times the remnant half-lieht ασ,"," These disks contain between $20$ and $60$ percent of the total gas in the initial galaxies, and may extend to several times the remnant half-light radii."977 Its worth noting that the initial galaxy. mocels used in these experiments were not particularly eas-rich. and that the gas started with the same distribution as the disk stars.," It's worth noting that the initial galaxy models used in these experiments were not particularly gas-rich, and that the gas started with the same distribution as the disk stars."978 In real spiral galaxies the atomic gas is more extended. than the stars: mergers of such svstems should form remnants with even larger and more massive σας disks., In real spiral galaxies the atomic gas is more extended than the stars; mergers of such systems should form remnants with even larger and more massive gas disks.979 Moreover. the gas content of disk galaxies is generally expected to increase with redshift: remnants with disks containing 10 to 15 percent of their luminous mass would. very likely result by doubling the eas fraction in the initial galaxy. models.," Moreover, the gas content of disk galaxies is generally expected to increase with redshift; remnants with disks containing $\sim 10$ to $15$ percent of their luminous mass would very likely result by doubling the gas fraction in the initial galaxy models."980 These disks. if subsequently. converted to stars. would be fairly hard to detect photometrically (Rix White 1990) unless viewed from a favorable orientation (e.g. Scorza Jonder 1990).," These disks, if subsequently converted to stars, would be fairly hard to detect photometrically (Rix White 1990) unless viewed from a favorable orientation (e.g. Scorza Bender 1990)."981 On the other hand. disks could. significantly inlluence the observed of merger remnants.," On the other hand, disks could significantly influence the observed of merger remnants."982 Disks are kinematicallv cold. so stars formed in such disks would acd narrow features to the stellar absorption-line spectra of earlv-tvpe. galaxies.," Disks are kinematically cold, so stars formed in such disks would add narrow features to the stellar absorption-line spectra of early-type galaxies."983 Moreover. disks rotate faster than pressure-supported: components. so these narrow features will be systematically olfset. [rom the broader. profiles due to the rest of the galaxy.," Moreover, disks rotate faster than pressure-supported components, so these narrow features will be systematically offset from the broader profiles due to the rest of the galaxy."984 Phe combined velocity profile of a disk plus spheroid thus appears asymmetric. with a steep prograde wing and a shallow retrograde wing (e.g. Franx," The combined velocity profile of a disk plus spheroid thus appears asymmetric, with a steep prograde wing and a shallow retrograde wing (e.g. Franx"985their analvsis in the laboratory.,their analysis in the laboratory.986 A great stride was made in this direction with the return. in 2006. of samples from Comet SIP/Wild2 (Saudford et al. (2006))).," A great stride was made in this direction with the return, in 2006, of samples from Comet 81P/Wild2 (Sandford et al. \cite{san06}) )."987 Preliminary results were discussed in several coutributious to the IAU Sxauposiuu 251 in Wong Itong Uswok aud Saudtford. (2008)))., Preliminary results were discussed in several contributions to the IAU Symposium 251 in Hong Kong (Kwok and Sandford \cite{kwo08}) ).988 Work ou these samples requires considerable techuical means and sophisticated instruments: it is sill actively going on., Work on these samples requires considerable technical means and sophisticated instruments; it is still actively going on.989 While these and oher simular efforts are beiug pursued. it is of interest to approach the problem from other sides and with different means.," While these and other similar efforts are being pursued, it is of interest to approach the problem from other sides and with different means."990 Several attempts have bee1 and are cine. nade to this end.," Several attempts have been, and are being, made to this end."991 The most popular is based on the use of ineasurecd «x computed polvevelie aromatic livdrocarbous (PATIs) (see Dauscilielher et al. (2009).," The most popular is based on the use of measured or computed polycyclic aromatic hydrocarbons (PAHs) (see Bauschlicher et al. \cite{bau09},"992.. Mattioda et al. (2009)))., Mattioda et al. \cite{mat09}) ).993 Oue of the latest exaiuples of such a line of attack is given by Doersina et al. (2010)., One of the latest examples of such a line of attack is given by Boersma et al. \cite{boe10}.994. It consists in drawiug spectra roni a erowing data bank of PATIs of differcut shapes aud sizes., It consists in drawing spectra from a growing data bank of PAHs of different shapes and sizes.995 In general. tjose are colmpact or irregular. ueutrals. anions or cations. mostly of large size.," In general, these are compact or irregular, neutrals, anions or cations, mostly of large size."996 Another type of investigation. opposite in a wav. was iutiated by Papoular et al. (1989).," Another type of investigation, opposite in a way, was initiated by Papoular et al. \cite{pap89}."997. It is based on the analogy between some of the UID xvectra with sole of the spectra of “standard” coals (see Charcosse (1980))) axd kerogcus (see Duraud (1980 )). collected in the rich data bauss established over the vears in France aud the USA. by istitutions dedicatcc to coal aix oil 1imniug. like the Cvoupement Fraucais dEude des Caroes (GPEC) aud the Institut Francais du Petrole (FP).," It is based on the analogy between some of the UIB spectra with some of the spectra of ""standard"" coals (see Charcosset \cite{cha}) ) and kerogens (see Durand \cite {dur}) ), collected in the rich data banks established over the years in France and the USA, by institutions dedicated to coal and oil mining, like the Groupement Francais d'Etude des Carbones (GFEC) and the Institut Francais du Petrole (IFP)."998 This analogy was xeviouslv. noticed w Iserridec (1987)., This analogy was previously noticed by Kerridge \cite{ker87}.999. Analogy between kerogen and meteoriic organdes was also noticed by anetologists (see e.g. Khare ct al. (1990))), Analogy between kerogen and meteoritic organics was also noticed by planetologists (see e.g. Khare et al. \cite{kha}) )1000 and petrologists (see Cataldo et . C900 1)))., and petrologists (see Cataldo et al. \cite{cat04}) ).1001" An incresting feature of this method is that a tight parallel can. lu sole cases, be craw between series ο terrestrial and astronomical spectra (see Papoular (2001) for the CTD stretchiue spectral region}."," An interesting feature of this method is that a tight parallel can, in some cases, be drawn between series of terrestrial and astronomical spectra (see Papoular \cite{pap01} for the CH stretching spectral region)."1002 Since laboratory aud field studies alowed the changes in he terrestrial spectra to be ascribed to definite plivsica zd Chemical modifications. some of these conclusions can. by iuaogv. be extended to the IS baud carriers. thus drawing a Imk )otwoeen otherwise unconnected spectra aud simultaucouslv clarifvine their structure and composition.," Since laboratory and field studies allowed the changes in the terrestrial spectra to be ascribed to definite physical and chemical modifications, some of these conclusions can, by analogy, be extended to the IS band carriers, thus drawing a link between otherwise unconnected spectra and simultaneously clarifying their structure and composition."1003 This approach also lighliehted the key role of heteroatoms iu generatiug a nunber of features: oxvecu. uitrogen. for instance. which. together with hbydogen aud carbon. form the CIION family. celebrated carlicr on by plauctologists.," This approach also highlighted the key role of heteroatoms in generating a number of features: oxygen, nitrogen, for instance, which, together with hydogen and carbon, form the CHON family, celebrated earlier on by planetologists."1004 While lis approach helps in the &ugerpriuts spectral region. 1.0. below about LO ju. where bands are clearly distinct. aud assigned to definite siial eroups of a few atoms cach. it is less powerful bevoud that.," While this approach helps in the fingerprints spectral region, i.e. below about 10 $\mu$ m, where bands are clearly distinct, and assigned to definite small groups of a few atoms each, it is less powerful beyond that."1005 Besides. even in the," Besides, even in the"1006Preprint After four vears of waiting since the detection of X-Ray Flash (XRF) 060218/SN 2006aj. another pair of low-Iunüunositv (LL) CRB - supernova (SN) association. ARF 100316D/SN 2010bh at redshift +=0.059 etal. 2OLO).. was captured by POOL) on March 16. 2010 (Stiuunatikosctal.2010:etal. 2010).. with a detection rate i consistent with the population studies of these nearby LL-CRD events (Coward2005:Soderbergetal.2006:Liane 2007).," After four years of waiting since the detection of X-Ray Flash (XRF) 060218/SN 2006aj, another pair of low-luminosity (LL) GRB - supernova (SN) association, XRF 100316D/SN 2010bh at redshift $z=0.059$ \citep{Vergani10}, , was captured by \citep{Gehrels04} on March 16, 2010 \citep{Stamatikos10,Wiersema10,Chornock10,Rau10}, with a detection rate fully consistent with the population studies of these nearby LL-GRB events \citep{Coward05,Soderberg06,Liang07,Guetta07}."1007.— Before this event. four pairs of nearby (5< 0.2) secure GRBCARF)-SN associatious have been identified.," Before this event, four pairs of nearby $z<0.2$ ) secure GRB(XRF)-SN associations have been identified."1008" These are GRB 980125/SN. 1998lw at 2.=0.0085 (ee.CGabuuaot 1998). GRDB 030329/SN 2003dh at :=0.168 (c.g.Tjorthetal.2003).. CRB 031203/SN. 20031 at υπ2=0.105 2006(e.g...Malesanietal.2001)... aud ARF SN at i=(00331 (οι,Campanactal. 2006)."," These are GRB 980425/SN 1998bw at $z=0.0085$ \citep[e.g.,][]{Galama98}, GRB 030329/SN 2003dh at $z=0.168$ \citep[e.g.,][]{Hjorth03}, GRB 031203/SN 2003lw at $z=0.105$ \citep[e.g.,][]{Malesani04}, and XRF 060218/SN 2006aj at $z=0.0331$ \citep[e.g.,][]{Campana06}."1009. The nature of the GRD/SN connection aud the interplaybetween the GRD aud the SN compoucuts are still poorly understood., The nature of the GRB/SN connection and the interplay between the GRB and the SN components are still poorly understood.1010 Iu this paper. we analyze and interpret theSwift BAT aud XRT data of XRF 100316D. paving special attention to the similarities ane. differences between the ARF οὐ1οον 201011 and NRE 060218/SN 2006aj.," In this paper, we analyze and interpret the BAT and XRT data of XRF 100316D, paying special attention to the similarities and differences between the XRF 100316D/SN 2010bh and XRF 060218/SN 2006aj."1011 GRB 100316D Suift/BAT at 12:11:50. UT (Tie) and the ΗΝNRTbegan observing the field. 137.7 secouds after the BAT tiieeerao (Stamatilkos et al.," GRB 100316D triggered /BAT at 12:44:50 UT $T_{\rm trig}$ ), and the XRT began observing the field 137.7 seconds after the BAT trigger (Stamatikos et al."1012 2010)., 2010).1013 A aight. steady ui-catalogued X-ray source was doetected.," A bright, steady un-catalogued X-ray source was detected."1014 Dv aualvziue the BAT survey data before and after he trigecr time. it is found that this event possibly started at ~(Tias1500) seconds (Staunatikoseal. 2010).," By analyzing the BAT survey data before and after the trigger time, it is found that this event possibly started at $\sim (T_{\rm trig} - 1500)$ seconds \citep{Stamatikos10}."1015. Towever. the0 eaunuierav flix of the source sept almost coustant from Dini.—1500 to Tie9200. before Increasing siguificautly at Diy.—500.," However, the gamma-ray flux of the source kept almost constant from $T_{\rm trig}-1500$ to $T_{\rm1016trig}-500$, before increasing significantly at $T_{\rm1017trig}-500$."1018 We therefore ake the starting time of this event as Tiye500 seconds., We therefore take the starting time of this event as $T_{\rm trig}-500$ seconds.1019 The host galaxy redshift is +=0.059 (Vereanictal.20]0)., The host galaxy redshift is $z=0.059$ \citep{Vergani10}.1020". Au associated supernova. SN 20100, was detectec by Wiersema oet al. ("," An associated supernova, SN 2010bh, was detected by Wiersema et al. ("10212010) ancl spectroscopically coufirned by Chornock et al. (,2010) and spectroscopically confirmed by Chornock et al. (10222010).,2010).1023 We extract the observed liehteurves aud spectra of NRF 100316D from the BAT aud NRT event data., We extract the observed lightcurves and spectra of XRF 100316D from the BAT and XRT event data.1024 The deails of our data reduction are presented in Zhangctal.(2007) and Liangetal.(2007b)., The details of our data reduction are presented in \citet{ZLZ07} and \citet{Liang07a}.1025". The joint BATΝΑΤ spectrum from Ti;|138 s to Tie|736 8 is well fit witha Cutoff Power-law inodel. with a power-law index of 1:32keVc0.03 and a spectral peak energv L, of 19.633 (\?/dof= 1320/1057). ly."," The joint BAT/XRT spectrum from $T_{\rm trig}+138$ s to $T_{\rm trig}+736$ s is well fit with a Cutoff Power-law model, with a power-law index of $1.32\pm10260.03$ and a spectral peak energy $E_{\rm p}$ of $19.6^{+3.3}_{-2.8}1027~{\rm keV}$ $\chi^2/{\rm dof}=1320/1057$ ), \ref{fig:BAT-XRT}) ."1028 This sugeests that the N-rav aud the σαλάτα eniss19115 are from the same emission compoucut., This suggests that the X-ray and the gamma-ray emissions are from the same emission component.1029 A time-resolved spectral analysis during the period from Zi;|138 sto Tis|036 8 showsthat Ej clearly evolves with tine. .Pou E⋅D τον- (froin. Tias|↽∙∖138 8 to Tis|.21 Qs} o 18.3139 vdkeV. (from Triv| 2108Tij| 7318).," A time-resolved spectral analysis during the period from $T_{\rm trig}+138$ s to $T_{\rm trig}+736$ s showsthat $E_{\rm p}$ clearly evolves with time, from $32.6^{+ 14.2}_{-8.5}$ keV (from $T_{\rm trig}+138$ s to $T_{\rm trig}+240$ s) to $18.3^{+ 3.9}_{-3.2}$ keV (from $T_{\rm trig}+240$ s to $T_{\rm trig}+734$ s)."1030" The Tui-absori lichteurvesand the temporalto evolution of ave shown in Figure 2.. aloug withthe data of€‘RB DL,(60215."," The un-absorbed lightcurvesand the temporal evolution of $E_{\rm p}$ are shown in Figure \ref{fig:XRT}, , along withthe dataof GRB 060218."1031 The late X-ray Lehtcurves of the other uewhy, The late X-ray lightcurves of the other nearby1032are the eigenfuuction expausious (the prime meaus that zero moces are uot included) where 4A=1.2.3.L5.6 aud by uPD). wk) the non-zero eigentiuctious of fy ancl i are denoted respectively: £=1.2.3.1. IulGr)—wtjuP(n). 1—56. IulGg)=SGuUCg). ICur)—wih)uir).,"are the eigenfunction expansions (the prime means that zero modes are not included) where $A=1,2,3,4,5,6$ and by $u^{(I)}(k)$, $u(k)$ the non-zero eigenfunctions of $K$ and $K^G$ are denoted respectively: $I=1,2,3,4$, $Ku^{(I)}(\vec{x})=\omega(\vec{k})u^{(I)}(\vec{x})$, $I=5,6$, $Ku^{(I)}(\vec{x})=\gamma(\vec{k})u^{(I)}(\vec{x})$, $K^Gu(\vec{x})=\omega(\vec{k})u(\vec{x})$."1033 Canonical quantization leads to the quautuim free Hamiltouiar aud the grouud state energy (all the modes non-occupied) of the topological solitous reads: where the star meaus that zero eigenvalues are uot accounted Lor., Canonical quantization leads to the quantum free Hamiltonian and the ground state energy (all the modes non-occupied) of the topological solitons reads: where the star means that zero eigenvalues are not accounted for.1034 Note that the ghost fields are static in this combined Weyl-backeround gauge aud their vacuum energy is one-half with respect to the case., Note that the ghost fields are static in this combined Weyl-background gauge and their vacuum energy is one-half with respect to the time-dependent case.1035 Only the Goldstone fluctuations around the vortices must be subtracted., Only the Goldstone fluctuations around the vortices must be subtracted.1036 The zero-point vacuum euergy renormalization provides the CastinirH euergy forH self-dualH (5>=4 1) semi-localH topologicalH solitous: Iu (24-1)-dimeusioual model ouly graphs with oue or two exterual lines are divergent iu tlie vacuum Sector., The zero-point vacuum energy renormalization provides the Casimir energy for self-dual $\kappa^2=1$ ) semi-local topological solitons: In (2+1)-dimensional model only graphs with one or two external lines are divergent in the vacuum Sector.1037 We choose the following counter-terms to cancel these divergences: Therefore must be added to the bare action (1)) to tame the divergences arising in one-loop order.," We choose the following counter-terms to cancel these divergences: Therefore, must be added to the bare action \ref{eq:act}) ) to tame the divergences arising in one-loop order."1038 This specilic choice fixes fiuite reuorimalizations according to the following criteria:, This specific choice fixes finite renormalizations according to the following criteria:1039In this paper. for the [first time we numerically investigate elfect of a gravitational scattering encounter between stars on a cir¢cumbinary disc.,"In this paper, for the first time we numerically investigate effect of a gravitational scattering encounter between stars on a circumbinary disc."1040 The paper is structured as. follows., The paper is structured as follows.1041 Phe details. on simulations are described in §2.., The details on simulations are described in \ref{simdetails}.1042 Alain results from the numerical work are presented in δι. , Main results from the numerical work are presented in \ref{results}. .1043In S4.. we discuss these results in the context of existing observations of Orion DN/IXL. and in 85 we summarize our conclusions and look forward to future steps.," In \ref{discussion}, we discuss these results in the context of existing observations of Orion BN/KL, and in \ref{conclusions} we summarize our conclusions and look forward to future steps."1044 In this work we consider an isolated. system comprising of three stars. arranged as an initial binary and a single star. and a cireumbinary disc. which we model with low-mass particles.," In this work we consider an isolated system comprising of three stars, arranged as an initial binary and a single star, and a circumbinary disc, which we model with low-mass particles."1045 Lvelrodynamic forces. in he disc are neglected. an approximation which allows us ο complete many thousancl scattering experiments in a reasonable timelrame.," Hydrodynamic forces in the disc are neglected, an approximation which allows us to complete many thousand scattering experiments in a reasonable timeframe."1046 Large numbers of experiments are necessary. due to the chaos inherent to the three-bocky oblem.," Large numbers of experiments are necessary, due to the chaos inherent to the three-body problem."1047 It is impossible to definitively trace backwards rom the current svstem. observed. with uncertainty. to an earlier. configuration: instead statistical studies must be »erformed to determine the likelihood. of similar outcomes o what we observe.," It is impossible to definitively trace backwards from the current system, observed with uncertainty, to an earlier configuration; instead statistical studies must be performed to determine the likelihood of similar outcomes to what we observe."1048 The svstem is evolved. using Aarseth’s code2003).. which utilises a fourth-order Llermite integrator.," The system is evolved using Aarseth's code, which utilises a fourth-order Hermite integrator."1049 While the code includes à IxS regularization scheme for handling close encounters between bodies. because of the atypical setup of this problem (three dominant bodies and a dise of low-mass particles). we suppress this feature and the equations of motion are integrated in standard Cartesian space.," While the code includes a KS regularization scheme for handling close encounters between bodies, because of the atypical setup of this problem (three dominant bodies and a disc of low-mass particles), we suppress this feature and the equations of motion are integrated in standard Cartesian space."1050 Because does not smooth the force between particles. there is in principle the possibility Chat undesired relaxation effects can occur in the disc: however in practice. over the short timescale of these simulations this is not à concern.," Because does not smooth the force between particles, there is in principle the possibility that undesired relaxation effects can occur in the disc; however in practice, over the short timescale of these simulations this is not a concern."1051 The two main elements of this. problem. are. the interaction between the disc anc the stars. and the eravitational scattering encounter. between the stars.," The two main elements of this problem are the interaction between the disc and the stars, and the gravitational scattering encounter between the stars."1052" Separately, cach of these elements. have been extensively explored: the theory of binarv-single scattering is laid out in(1975).. and a long history of numerical work includes(1920)... (1974)..(1975)..(1979).. (1983).. and (1993)."," Separately, each of these elements have been extensively explored: the theory of binary-single scattering is laid out in, and a long history of numerical work includes, , and ."1053. Disruptive clise-star encounters have likewise been studied numerically in a variety of contexts2006)., Disruptive disc-star encounters have likewise been studied numerically in a variety of contexts.1054. The combination of the two processes. is unstudied to ourknowledgel.. and. interesting in general.," The combination of the two processes is unstudied to our, and interesting in general."1055 'araneter space is vast. and it is fortunate that DN/INL srovicles a physical reference from which to take initial steps.," Parameter space is vast, and it is fortunate that BN/KL provides a physical reference from which to take initial steps."1056 We ran two sets of simulations. which we label in terms of the masses of the initial stellar configuration in.," We ran two sets of simulations, which we label in terms of the masses of the initial stellar configuration in."1057AL... One set of 4500. simulations. (10.3)|10. has à. low-mass member in the binary. while the other set of 7000. runs. (10.10)|10. has all three stars with equal masses.," One set of 4500 simulations, (10,3)+10, has a low-mass member in the binary, while the other set of 7000 runs, (10,10)+10, has all three stars with equal masses."1058 Each experiment begins with a circular binary with semi-major axis e=10 au surrounded: by a cireumbinary. cise with otal mass 0.1:M., Each experiment begins with a circular binary with semi-major axis $a = 10$ au surrounded by a circumbinary disc with total mass 0.1.1059.. Phe results of this paper are. strictly speaking. only applicable to discs that are in this non-sel eravitating. dynamically unimportant regime.," The results of this paper are, strictly speaking, only applicable to discs that are in this non-self gravitating, dynamically unimportant regime."1060 While he dvnamies of the actual scattering interaction between he stars is probably insensitive to the disc mass. a more massive dise could potentially alter the trajectory of the stars in the immediate leadup to the stellar interactions.," While the dynamics of the actual scattering interaction between the stars is probably insensitive to the disc mass, a more massive disc could potentially alter the trajectory of the stars in the immediate leadup to the stellar interactions."1061 ote that the estimated remnant disc mass of less than L2 is comfortably unimportant to the dynamics of a massive binary over short timescales., Note that the estimated remnant disc mass of less than 0.2 is comfortably unimportant to the dynamics of a massive binary over short timescales.1062 The simulated cise is runcated at the inner edge at the 3:1 resonance with the παν which is approxiniately 2.08«=20.8 au. extends to the outer edge at 106=100 au. ancl follows an arbitrary but plausible surface. density mwolile X(r)xr ," The simulated disc is truncated at the inner edge at the 3:1 resonance with the binary which is approximately $2.08a=20.8$ au, extends to the outer edge at $10a = 100$ au, and follows an arbitrary but plausible surface density profile $\Sigma(r) \propto r^{-1}$."1063Jocause we ignore hydrodynamic forces here is no need to resolve the disc vertically. and the disc xwticle are initially coplanar with the binary.," Because we ignore hydrodynamic forces there is no need to resolve the disc vertically, and the disc particle are initially coplanar with the binary."1064 The disc is composed. of 2048. particles. ancl tests with 1024 and 4096 xwlicles show that the results are essentially independent of disc resolution.," The disc is composed of 2048 particles, and tests with 1024 and 4096 particles show that the results are essentially independent of disc resolution."1065 ln order to place the (initially unbound) third star. the παν is treated as a single body at its center of mass. and a Ixeplerian. orbit is determined. for some relative velocity at infinity e and desired. periastron kc.," In order to place the (initially unbound) third star, the binary is treated as a single body at its center of mass, and a Keplerian orbit is determined for some relative velocity at infinity $v_{\infty}$ and desired periastron $r_p$."1066 Εις periastron is he unperturbed separation that occurs when the binary is reated as its center of mass: the actual orbits of the three »odies diverges from this path asthey approach cach other, This periastron is the unperturbed separation that occurs when the binary is treated as its center of mass; the actual orbits of the three bodies diverges from this path asthey approach each other1067index. G eravitational constant. M. and Q the mass and the angular velocity of theNS. respectively.,"index, $G$ gravitational constant, $M$ and ${\bm \Omega}$ the mass and the angular velocity of the, respectively."1068 The third. fourth. and. fifth terms on the rhs of Eq. (," The third, fourth and fifth terms on the rhs of Eq. ("10691) represent the Coriolis force. the centrifugal force and the gravity. respectively.,"1) represent the Coriolis force, the centrifugal force and the gravity, respectively."1070 Observationally the accretion rates in LMXDs change on a timescale of ~10°—1019., Observationally the accretion rates in LMXBs change on a timescale of $\sim10^3-10^4 s$.1071 This is much more than (the relaxation (or dynamic) ünmescale of the plasma al the inner disc radius (~10?— 7s)., This is much more than the relaxation (or dynamic) timescale of the plasma at the inner disc radius $\sim10^{-2}-10^{-3} s$ ).1072 So we can approximate the plasma to be in an equilibrium state. which is subject to small perturbations. as discussed in Benz (2002).Dw use of the current expression. where je is is vacuum magnetic conductivity. Eqs. (," So we can approximate the plasma to be in an equilibrium state, which is subject to small perturbations, as discussed in Benz (2002).By use of the current expression, where $\mu$ is is vacuum magnetic conductivity, Eqs. ("10731)-(3) can be transformed to be where the subscript 0 denotes variables in the equilibrium state.,1)-(3) can be transformed to be where the subscript 0 denotes variables in the equilibrium state.1074 The initial relative velocily (v9) is equal to zero in (he corotating reference svstem and can be expressed. as Qxry in the inertial reference svstem.," The initial relative velocity $ {\bm {v_0} }$ ) is equal to zero in the corotating reference system and can be expressed as ${{ {\bm1075\Omega }}\times { {\bm r}}_0 }$ in the inertial reference system."1076" Now we consider the MUIID equations for the plasma subject to small perturbations. where v=V9+v,Vs. B=BotBs. T—v9 Τε. p=fut ps. p=4+DP, with the subscript s denoting the perturbed quantities (ο<|Qxrp). D;« Dg. re«rg. poKfy. D,« Py) ancl with the superscript the variables after (he disturbance."," Now we consider the MHD equations for the plasma subject to small perturbations, where $\hat{\bm v} = {{\bm v}}{ }_{ {\bm 0}} + { {\bm v}}_{ {\bm s}}1077= { {\bm v}}_{{\bm s}}$, $\hat{\bm B} = { {\bm B}}{ }_{ {\bm 0}} +1078{{\bm B}}{ }_{{\bm s}}$, $\hat{\bm r} = {{\bm r}}{ }_{{\bm 0}} +1079{{\bm r}}{ }_{{\bm s}}$ , $\hat{\rho} = \rho_0 + \rho_s$ , $\hat{P} =1080P_0 + P_s$ with the subscript s denoting the perturbed quantities $v_s \ll \left| {{ {\bm \Omega }}\times { {\bm r}}_0 }\right|$, $B_s\ll B{ }_0$ , $r_s \ll r{ }_0$, $\rho_s\ll \rho_0$, $P_s\ll P_0$ ) and with the superscript $\hat{}$ the variables after the disturbance."1081 Combining Eqs. , Combining Eqs. (6)-(12)1082we eget the equations about the perturbedquantities inthe first orderapproximation.," we get the equations about the perturbedquantities inthe first orderapproximation,"1083N-ray data.,X-ray data.1084" The ROSAT PSPC image (Moly. Mathieseu Evrard 1999). shows isophotes clongated iu the NE-SW direction on the few πάτος, kpe scale. while ou the ew Mpe scales the elongation shifts to the E-W direction."," The ROSAT PSPC image (Mohr, Mathiesen Evrard 1999), shows isophotes elongated in the NE-SW direction on the few hundred kpc scale, while on the few Mpc scales the elongation shifts to the E-W direction."1085 Iu the picture proposed by. QRW the two clusters started colliding about | Cyr ago. with the ceutral cores coming ogether in the last 1-2 Cir.," In the picture proposed by QRW the two clusters started colliding about 4 Gyr ago, with the central cores coming together in the last 1-2 Gyr."1086 Incidentally a fit with a -)- to the PSPC radial profile (see Mohr. Mathieseu Evrard 1999) vields a huge value for the core radius. r=0.5 AIpe. as inight expected in the case of recent nerecr in the core of A3266.," Incidentally a fit with a $\beta$ -model to the PSPC radial profile (see Mohr, Mathiesen Evrard 1999) yields a large value for the core radius, $r_c=0.5$ Mpc, as might expected in the case of recent merger in the core of A3266."1087 The radial tempcratiure eradient found by ASCA aud BeppoSAX lends further streneth to the mereing scenario proposed by QRW., The radial temperature gradient found by ASCA and BeppoSAX lends further strength to the merging scenario proposed by QRW.1088 The map we present in figure 3 shows that the radial telmperature eradient is present in all sectors., The map we present in figure 3 shows that the radial temperature gradient is present in all sectors.1089 We also find an indication of au azimuthal tempcrature eradicut occurring in the annulus withbounding radii ['-sN/ (0.35 Alpe - 0.7 Alpe): the data suggestsOO that the eastern side of the cluster may be somewhat cooler than the western side., We also find an indication of an azimuthal temperature gradient occurring in the annulus with bounding radii $^\prime$ $^\prime$ (0.35 Mpc - 0.7 Mpc); the data suggests that the eastern side of the cluster may be somewhat cooler than the western side.1090 Verv recentlv Heurksen. Dounelly Davis (1999). from) the analysis of an ASCA observation of A3266 find evidence of a temperature gradient in the SW NE direction indicative of au ongoing merger.," Very recently Henriksen, Donnelly Davis (1999), from the analysis of an ASCA observation of A3266, find evidence of a temperature gradient in the $-$ NE direction indicative of an ongoing merger."1091 The azimuthal telperature eradieut found in our data corroborates the ASCA result., The azimuthal temperature gradient found in our data corroborates the ASCA result.1092" The average metal abundance we find from the NECS data, (2140.03. solar units. i in agreement with the value 140.2 derived by Mushotzky. (1981). using HEAOL À2 crite, and with the average metallicity. 0.2140.05. derived bv Allen Fabian (1998) for a sample of nou-cooliug flow cluster."," The average metal abundance we find from the MECS data, $0.21\pm10930.03$, solar units, is in agreement with the value $\pm$ 0.2 derived by Mushotzky (1984), using HEAO1 A2 data, and with the average metallicity, $0.21\pm 0.05$, derived by Allen Fabian (1998) for a sample of non-cooling flow cluster."1094 The radial abuudance profile (see the bottom panel of fleure 2). does not show anv stroug evidence of a decrease in the abundance with increasing radius.," The radial abundance profile (see the bottom panel of figure 2), does not show any strong evidence of a decrease in the abundance with increasing radius."1095 Thus. A3266 would sccm to conform to the general rule that nou-cooling flow cluster do not present mictallicity eracicuts.," Thus, A3266 would seem to conform to the general rule that non-cooling flow cluster do not present metallicity gradients."1096 We acknowledge support from the BeppoSAX Scicuce Data Center., We acknowledge support from the BeppoSAX Science Data Center.1097"Low-auass stars du the pre-nuün sequence (PAIS) stage are classified into four types based ou their infrared. (IR) to sub-niu spectral energy distributions (SEDs): molecular cloud cores through protostars are represented by class 0 aud I SEDs, while classical aud weal-line T-Tiuul stars (CTTSs and WTTSs) exhibit class IL aud III SEDs (SlnAdams.&Lizano1987:AudréMoutiierle199 de","Low-mass stars in the pre-main sequence (PMS) stage are classified into four types based on their infrared (IR) to sub-mm spectral energy distributions (SEDs); molecular cloud cores through protostars are represented by class 0 and I SEDs, while classical and weak-line T-Tauri stars (CTTSs and WTTSs) exhibit class II and III SEDs \citep{Shu1987,Andre1994}."1098 Tιο satellite first revealed that T-Tauri stars (TTSs = CTTSs aud WTTSs) emit iuteise and fine variahle N-ravs with occasional rapk flares (Feigelson&DeCaipli1981:Monut-mereotal. 1983).," The satellite first revealed that T-Tauri stars (TTSs = CTTSs and WTTSs) emit intense and time variable X-rays with occasional rapid flares \citep{Feigelson1981,Montmerle1983}."1099". These properties are cousistcut with a scenario of euhlanced solu-tvpe activity. attitable to maenetic dynamo processes,"," These properties are consistent with a scenario of enhanced solar-type activity, attributable to magnetic dynamo processes."1100 Successive N-rav satellites πιch asROSAT.ASCA. andChandra lave cetected N-vay emission fron many PAIS stars. even fre1u the voungest class 0 stage (Tsuboietal.20n). and sueeestOO that XN-raw Cluission from. low-nass PAIS stars ds caused by naeuctic activiY.," Successive X-ray satellites such as, and have detected X-ray emission from many PMS stars, even from the youngest class 0 stage \citep{Tsuboi2001}, and suggest that X-ray emission from low-mass PMS stars is caused by magnetic activity."1101 These results. however. rely nainly on shor duration observations: neither ong-terni nor ¢letailed behavior of the X-rav spectruni aud tiwine has been well studied.," These results, however, rely mainly on short duration observations; neither long-term nor detailed behavior of the X-ray spectrum and timing has been well studied."1102 This paper hence addresses three topics on TTSs: (1) he duty ratio «Mf the N-ray flare. (2) the lo18Ouern plasuia sticture and variability. aud (3) he chemical compositions iu the quiescent aud flare phases. (," This paper hence addresses three topics on TTSs; (1) the duty ratio of the X-ray flare, (2) the long-term plasma structure and variability, and (3) the chemical compositions in the quiescent and flare phases. ("11031) Stelzer.Neuliuser.&Tamibarvan(2000) compiledROSAT PSPC observations of the Taurus-Auriga-Perseus region and found that the flare rate may depend on the stellar age. rotation,"1) \citet{Stelzer2000} compiled PSPC observations of the Taurus-Auriga-Perseus region and found that the flare rate may depend on the stellar age, rotation"1104RJM and ΜΗ are funded by PPARC PDRAs.,RJM and MJJ are funded by PPARC PDRAs.1105 This publication makes use of the material provided in the FIRST and SDSS DRI surveys., This publication makes use of the material provided in the FIRST and SDSS DR1 surveys.1106 FIRST is funded by the National Astronomy Observatory (NRAO) and is a research facility of the U.S. National Science foundation and uses the NRAO Very Large Array., FIRST is funded by the National Astronomy Observatory (NRAO) and is a research facility of the U.S. National Science foundation and uses the NRAO Very Large Array.1107 Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation. the Participating Institutions. the National Aeronauties and Space Administration. the National Science Foundation. the U.S. Department of Energy. the Japanese Monbukagakusho. and the Max Planck Society.," Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society."1108" Further details of the SDSS survey can be found on http://www.sdss.org/,", Further details of the SDSS survey can be found on http://www.sdss.org/.1109"llowever. a very laree ratio VH indicates that the [it under hypothesis. £/,, is. significantly. worse than achievable. with a four-parameter model. and that //,, has to be rejected.","However, a very large ratio $\frac{\chi^{2}(H_n)}{\chi^{2}(H_1)}$ indicates that the fit under hypothesis $H_n$ is significantly worse than achievable with a four-parameter model, and that $H_n$ has to be rejected."1110 The virtue of this approach is that it is insensitive to the assumed. uncertainties in the surface brightness measurement: the normalizations of the uncertainties cancel out as we form a ratio of X7 values., The virtue of this approach is that it is insensitive to the assumed uncertainties in the surface brightness measurement; the normalizations of the uncertainties cancel out as we form a ratio of $\chi^{2}$ values.1111" However. a complication lies in linding the critical threshold. value for the X7 ratio bevond which we can rule out that the improvement in the quality of the fit results simply from the ability to fit random Iluctuations better with more free parameters. and is indeed indicative ofa true preference of Lf, over IH."," However, a complication lies in finding the critical threshold value for the $\chi^{2}$ ratio beyond which we can rule out that the improvement in the quality of the fit results simply from the ability to fit random fluctuations better with more free parameters, and is indeed indicative ofa true preference of $H_1$ over $H_n$."1112" To calibrate r,2 as à statistical measure of significance. we use a Monte. Carlo simulation. that generates mock surface brightness proliles conforming to the. tested hypothesis //,."," To calibrate $r_{\chi^{2}}$ as a statistical measure of significance, we use a Monte Carlo simulation that generates mock surface brightness profiles conforming to the tested hypothesis $H_n$."1113 “Phe mock profile is created as à superposition of individual surface brightness luctuations., The mock profile is created as a superposition of individual surface brightness fluctuations.1114 The spatial shape of cach profile is described by a Gaussian function with a fixed width., The spatial shape of each profile is described by a Gaussian function with a fixed width.1115 The number of Ductuations that are added up in cach row of the simulated. profile are determined [rom an assumed average amplitude and Poisson statistics., The number of fluctuations that are added up in each row of the simulated profile are determined from an assumed average amplitude and Poisson statistics.1116 The actual flux. of cach Gaussian. [Iuctuation is drawn from a Caussian probability. distribution. over logarithmic Lux with a pre-defined width., The actual flux of each Gaussian fluctuation is drawn from a Gaussian probability distribution over logarithmic flux with a pre-defined width.1117 In addition. we add Uuctuations with a fixed. small amplitude to simulate Iluctuations in the background.," In addition, we add fluctuations with a fixed, small amplitude to simulate fluctuations in the background."1118 The average amplitudo. spatial width. width of the Lux probability distribution and background Iuctuation amplitude are the parameters of this model: in addition. we have two parameters controlling the size of the simulated: uncertainties.," The average amplitude, spatial width, width of the flux probability distribution and background fluctuation amplitude are the parameters of this model; in addition, we have two parameters controlling the size of the simulated uncertainties."1119 We have run a series of simulations to find the parameters that recreate profile shapes most similar in appearance to the observed. ones by fitting broken-exponential models to cach mock. profile and then using three statistics to compare them. to the observations: 1) the ratio of the integrated residuals between 0.2 and 1.5 Io; to the total integrated. profile in that range., We have run a series of simulations to find the parameters that recreate profile shapes most similar in appearance to the observed ones by fitting broken-exponential models to each mock profile and then using three statistics to compare them to the observations: 1) the ratio of the integrated residuals between 0.2 and 1.5 $_{25}$ to the total integrated profile in that range.1120" 2) the cumulative residual. between 0.21155; and an outer radius .r. normalized by the total integrated residual between 0.2 Re, and 1.51154."," 2) the cumulative residual between $_{25}$ and an outer radius $x$, normalized by the total integrated residual between 0.2 $_{25}$ and $_{25}$."1121 3) the cumulative us between 0.2155 and an outer radius or. normalized by the total integrated v7 between 0.2 Res and 1.51155.," 3) the cumulative $\chi^{2}$ between $_{25}$ and an outer radius $x$, normalized by the total integrated $\chi^{2}$ between 0.2 $_{25}$ and $_{25}$."1122 The first of these statistics characterizes the magnitude of the Uuctuations relative to the total lux., The first of these statistics characterizes the magnitude of the fluctuations relative to the total flux.1123 The second characterizes the radial behavior of the size of Dluctuations., The second characterizes the radial behavior of the size of fluctuations.1124 The third characterizes the relative goocdness-of-It as a function of radius., The third characterizes the relative goodness-of-fit as a function of radius.1125 Of particular iniportance is that. both in the mocks and in the real profiles. contributions to X72 come primarily from intermediate raclii around rz0.7.— Le. 7 Deds particularly+ sensitiveme to the regionH where we locate the break in the composite profile.," Of particular importance is that, both in the mocks and in the real profiles, contributions to $\chi^{2}$ come primarily from intermediate radii around $r\approx 0.7$, i.e., $\chi^{2}$ is particularly sensitive to the region where we locate the break in the composite profile."1126" We determine the X72 ratio. r2. for a given hypothesis 1. and compare it to the scatter of 7r,» in our mock profiles."," We determine the $\chi^{2}$ ratio, $r_{\chi^{2}}$, for a given hypothesis $H_n$, and compare it to the scatter of $r_{\chi^{2}}$ in our mock profiles."1127" 1E is outside the range reproduced\ by of the Monte Carlor,2 realizations. we reject the hypothesis with confidence."," If $r_{\chi^{2}}$ is outside the range reproduced by of the Monte Carlo realizations, we reject the hypothesis with confidence."1128 For our analysis of the IIa. surface brightness profiles in the individual galaxies. we estimate the racial range of Lo detections in each galaxy.," For our analysis of the $\alpha$ surface brightness profiles in the individual galaxies, we estimate the radial range of $\alpha$ detections in each galaxy."1129 Although our fitting procedure includes the entire racial range between 0.2<Aefes«1.5. the position of the outermost detection additionally provides a lower limit on the position of the truncation in I».," Although our fitting procedure includes the entire radial range between $0.2 < R/R_{25} < 1.5$, the position of the outermost detection additionally provides a lower limit on the position of the truncation in $_2$."1130 To determine the radius of the outermost significant llo detection. we integrate the total Πα lux between a eiven radius # and the Largest radius considered in our analysis. 1.54755.," To determine the radius of the outermost significant $\alpha$ detection, we integrate the total $\alpha$ flux between a given radius $R$ and the largest radius considered in our analysis, $1.5 R_{25}$."1131" To determine the limiting radius £2);,,,. we then use two alternative definitions."," To determine the limiting radius $R_{lim}$, we then use two alternative definitions."1132 The first. {ης is the radius of the outermost Ha detection with 73c significance between J? and. 1.54255.," The first, $R_{lim}^{a}$, is the radius of the outermost $\alpha$ detection with $> 3\sigma$ significance between $R$ and $1.5 R_{25}$."1133 The second. νι is the smallest racius A? for which the Ha Εικ is <2e significant.," The second, $R_{lim}^{b}$, is the smallest radius $R$ for which the $\alpha$ flux is $< 2\sigma$ significant."1134 Phe first definition may include isolated Ho. regions at large racial distance from the galaxy center. while the second definition measures only the extent of contiguous Lla emission aud does not include isolated sources in the outer disk.," The first definition may include isolated $\alpha$ regions at large radial distance from the galaxy center, while the second definition measures only the extent of contiguous $\alpha$ emission and does not include isolated sources in the outer disk."1135" Llowever. in practice. both definitions vield very similar results. except in the case of ESO 478- COLL. where an Ho detection of low significance pushes the outermost radius to fi,8 1.45. while contiguous Ho emission is only detected out to 1.34."," However, in practice, both definitions yield very similar results, except in the case of ESO 478- G011, where an $\alpha$ detection of low significance pushes the outermost radius to $R_{lim}\approx1.45$ , while contiguous $\alpha$ emission is only detected out to $R_{lim}\approx1.34$ ."1136" We list £25, For all galaxies in our sample in Table 1..", We list $R_{lim}$ for all galaxies in our sample in Table \ref{tab_outermost}. .1137 We have not [found obvious correlations between the extent of the Ha. emission and. other global galaxy properties. such as the total Ho emission. apart from a very weak correlation between {πρ and the integrated Ho lux beyond ry. which. indicates that only galaxies with large Hu may have significant Lux in this second. outer-clisk exponential Component.," We have not found obvious correlations between the extent of the $\alpha$ emission and other global galaxy properties, such as the total $\alpha$ emission, apart from a very weak correlation between $R_{lim}$ and the integrated $\alpha$ flux beyond $r_0$, which indicates that only galaxies with large $R_{lim}$ may have significant flux in this second, outer-disk exponential component."1138 Phe La surface brightness. profiles of the individual ealaxies are shown in Fig., The $\alpha$ surface brightness profiles of the individual galaxies are shown in Fig.1139 6 (which consists of two parts) for the morphologically regular sample. and in Fi," \ref{fig_profiles1} (which consists of two parts) for the morphologically regular sample, and in Fig."1140 .2 [or the warped sample., \ref{fig_profiles3} for the warped sample.1141 By visual inspection. we identify a break in several of these ealaxies. while others appear consistentwith single exponentials.," By visual inspection, we identify a break in several of these galaxies, while others appear consistentwith single exponentials."1142 Llowever.in general. there are large," However,in general, there are large"1143the point-source sensitivity for future instruments therefore potentially become dependent on the assumed: unresolved extragalactic component (due to quasars ancl star-forming ealaxies) ancl its (ux distribution. as well as the elective collecting area and photon extraction radius (or resolution).,"the point-source sensitivity for future instruments therefore potentially become dependent on the assumed unresolved extragalactic component (due to quasars and star-forming galaxies) and its flux distribution, as well as the effective collecting area and photon extraction radius (or resolution)."1144 Due to uncertainties in the decomposition. of the X-ray xickerouncd into contributions from the galaxy. star-forming galaxies and quasars. ib is somewhat uncertain what the confusion. limit of these telescopes will be.," Due to uncertainties in the decomposition of the X-ray background into contributions from the galaxy, star-forming galaxies and quasars, it is somewhat uncertain what the confusion limit of these telescopes will be."1145 The projected olnt-source sensitivities for the next. generation. telescope »ealk for soft X-ray energies., The projected point-source sensitivities for the next generation telescope peak for soft X-ray energies.1146 At these. energies the soft hermal galactic component is a significant portion of he background (see.e.g.2).. however the contribution of unresolved. point sources may. also play an important role if the ambitious goals for the resolution are not achieved.," At these energies the soft thermal galactic component is a significant portion of the background \citep[see, e.g.][]{parmar1999}, however the contribution of unresolved point sources may also play an important role if the ambitious goals for the resolution are not achieved."1147 Since our models predict the contribution of quasars to this background. the expected. sensitivity of an instrument. [ike NEUS to high redshift sources may depend on the modelling of faint sources.," Since our models predict the contribution of quasars to this background, the expected sensitivity of an instrument like XEUS to high redshift sources may depend on the modelling of faint sources."1148 7. discuss the anticipated point source sensitivity of Constellation-N and NEUS., \citet{hasinger2006} discuss the anticipated point source sensitivity of Constellation-X and XEUS.1149 Assuming that the observations are confusion limited when there are fewer than 40 “beams” per source. combined with estimates of the background due to unresolved. extragalactic sources. galactic emission. and cosmic-ravs. ? project that à 1 Ms observation with NEUS will vield a point source sensitivity o£3.LO “eres tem 2 approximately 200 times better than that of NMM-Newton.," Assuming that the observations are confusion limited when there are fewer than 40 “beams” per source, combined with estimates of the background due to unresolved extragalactic sources, galactic emission and cosmic-rays, \citet{hasinger2006} project that a 1 Ms observation with XEUS will yield a point source sensitivity of $3 \times115010^{-18}$ erg $^{-1}$ $^{-2}$, approximately $200$ times better than that of XMM-Newton."1151 This estimate presumes that Χοος goal. resolution of 2 arcseconds will be achieved., This estimate presumes that XEUS's goal resolution of 2 arcseconds will be achieved.1152 Achieving this resolution is expected: to be extremely challenging. and the required resolution of 5 arcseconds is perhaps more realistic (?)..," Achieving this resolution is expected to be extremely challenging, and the required resolution of 5 arcseconds is perhaps more realistic \citep{hasinger2006}. ."1153 For this resolution the sensitivity degrades to ~210to erg  7 for a 1 Ms integration.," For this resolution the sensitivity degrades to $\sim 2 \times115410^{-17}$ erg $^{-1}$ $^{-2}$ for a $1$ Ms integration."1155 Similarly. for Constellation-N. 7. project. a point source sensitivityof," Similarly, for Constellation-X \citet{hasinger2006} project a point source sensitivityof"1156have been emitted in the absence of the dust cover.,have been emitted in the absence of the dust cover.1157 This total hypothetical mass is a factor 2. 1 lieher than the total mass released within 33 davs after the outburst. ie. the burst dust mass.," This total hypothetical mass is a factor 3 – 4 higher than the total mass released within 33 days after the outburst, i.e. the burst dust mass."1158 Thus. we can safely asstme that the insolation provided enoush euergv to start sublimation witlin the nucleus.," Thus, we can safely assume that the insolation provided enough energy to start sublimation within the nucleus."1159 Within our observing interval. the ass in the halo is approximately proportional to the observed flux density iu the beam. thus to the observing accuracy.," Within our observing interval, the mass in the halo is approximately proportional to the observed flux density in the beam, thus to the observing accuracy."1160 Therefore the relative accuracy from dav to day and of the dust production rate is quite good., Therefore the relative accuracy from day to day and of the dust production rate is quite good.1161 The absolute accuracy depends on our kuowledee of the absorption cocfiicicut & of cometary dust. whose uncertaiutv was estimated by Altenhoff et al. (20003) ," The absolute accuracy depends on our knowledge of the absorption coefficient $\kappa$ of cometary dust, whose uncertainty was estimated by Altenhoff et al. \cite{alt2}) )"1162to be about a factor 2., to be about a factor 2.1163 The accuracies of the mass determinations of the suall-erained dust (Sekauina 2007)) and of water (Combi 20073) have unfortunately not been reported., The accuracies of the mass determinations of the small-grained dust (Sekanina \cite{sek}) ) and of water (Combi \cite{com}) ) have unfortunately not been reported.1164 The unclear structure of comets 9P/Tempcl 1 aud l17P/ILlohues before the outbursts are xobablv alike. a densely packed dust cover (1 1) )low a laver of pure water ice (10 1). below amorplous dirty H2O ice. whose upper part is possibly free of highlv volatile 1nolecules.," The nuclear structure of comets 9P/Tempel 1 and 17P/Holmes before the outbursts are probably alike, a densely packed dust cover $\approx$ 1 m) below a layer of pure water ice $\approx$ 10 m), below amorphous dirty $_2$ O ice, whose upper part is possibly free of highly volatile molecules."1165 At ΟΡΤωρα P the impactor acted as the exothermic enerev source to blow off the pancake-shaped dust cover. as the scheme of Sckauina (2007)) suggests. πι it a type (c) split nucleus.," At 9P/Tempel 1 the impactor acted as the exothermic energy source to blow off the pancake-shaped dust cover, as the scheme of Sekanina \cite{sek}) ) suggests, making it a type (c) split nucleus."1166 The development for 17P/Iohues is different., The development for 17P/Holmes is different.1167 When Πο sublimation startec inside the JOTOlls uucleus. water vapor spread all over the imcleus. initiating more sublimation: deeper inside. aud even other molecular ices with lower sublimation poiuts were heated. stored there at lower temperatures.," When $_2$ O sublimation started inside the porous nucleus, water vapor spread all over the nucleus, initiating more sublimation; deeper inside, and even other molecular ices with lower sublimation points were heated, stored there at lower temperatures."1168 The effective. sublimating surface inside the micleus. estimated as excess over the Cluission after he burst on dav 35 when it was near equilibrium with insolation. was more than l1 times the wiclear surface. corresponding roughly to the nuclear size of comet Uale-Bopp.," The effective sublimating surface inside the nucleus, estimated as excess over the emission after the burst on day 35 when it was near equilibrium with insolation, was more than 14 times the nuclear surface, corresponding roughly to the nuclear size of comet Hale-Bopp."1169 The stm of the saturated partial oressures of all ice species was obviously breaking up he air fight dust mantle. allowing the cometary wind to start through the dust mantle aud lifting dust particles. viece by piece. iuto the halo.," The sum of the saturated partial pressures of all ice species was obviously breaking up the air tight dust mantle, allowing the cometary wind to start through the dust mantle and lifting dust particles, piece by piece, into the halo."1170" The break up of he dust nautle is hardly spectacular. compared with the ""ull start of cometary wind."," The break up of the dust mantle is hardly spectacular, compared with the full start of cometary wind."1171 Different versious exist of the developmen of the outburst., Different versions exist of the development of the outburst.1172 Sekanina (2008.. 2007)) refers to an explosion and a single exothermic event. aud Diver et al. (2008) )," Sekanina \cite{sek3}, \cite{sek}) ) refers to an explosion and a single exothermic event, and Biver et al. \cite{biv2}) )"1173 and others repor a water production rate. which almost ends after 3 days.," and others report a water production rate, which almost ends after 3 days."1174 Iu Fie., In Fig.1175 1 the continua observations are plotted. showing that the outburstaclated. increased continui enmisson lasted for about 30 dave. as did the increased nuclear maguitude mo.," 1 the continuum observations are plotted, showing that the outburst-related increased continuum emission lasted for about 30 days, as did the increased nuclear magnitude $m_2$."1176" Additional proof are the uuuerous photographs taken within the first mouth of the outburst: see e.g. Sekauina (2008)). in which the comet appears as a filled Pleriou rather than a shell. miplviug that the dust injection iuto the coma continued. after the ""explosion"" for quite soe time."," Additional proof are the numerous photographs taken within the first month of the outburst; see e.g. Sekanina \cite{sek3}) ), in which the comet appears as a filled Plerion rather than a shell, implying that the dust injection into the coma continued after the “explosion” for quite some time."1177 Production of gas-phase molecules is responsible for all he cometary activity., Production of gas-phase molecules is responsible for all the cometary activity.1178 It is predominantly the cometary wind of the Πο ο molecules. which lifts the dust particles roni the uucleus. so a correlation between II;O aud dust xoduction is expected.," It is predominantly the cometary wind of the $_2$ O molecules, which lifts the dust particles from the nucleus, so a correlation between $_2$ O and dust production is expected."1179" Usually the production of ciffercut uolecules shows a fixed ratio. so that one can. ο,οι, predict he Ποῦ production rate from ICN observations."," Usually the production of different molecules shows a fixed ratio, so that one can, e.g., predict the $_2$ O production rate from HCN observations."1180 Not so or comet Wolmes!, Not so for comet Holmes!1181 Figre 1 shows that W2O production. observed with the SWAN satellite. asts at least for a nouth as does the enhanced nuu coutiunumun cussion. while c.g. the spectral lines of TICN. CO. NIT; (Dralius et al. 2007::," Figre 1 shows that $_2$ O production, observed with the SWAN satellite, lasts at least for a month as does the enhanced mm continuum emission, while e.g. the spectral lines of HCN, CO, $_3$ (Drahus et al. \cite{dra};"1182 Biver et al. 2008:, Biver et al. \cite{biv2};1183 Monteu. 2007)) had a ye signal at the start. which appareutlv petered out dramatically after 3 days. as shown in Fig.," Menten, \cite{men}) ) had a big signal at the start, which apparently petered out dramatically after 3 days, as shown in Fig."1184 1., 1.1185 The reason navy © the temperature/depth structure of the uucleus. vecamse the near surface ice unight be free of volatile uolecules.," The reason may be the temperature/depth structure of the nucleus, because the near surface ice might be free of volatile molecules."1186 All derived masses ave collected in Table 2. where the otal nüuass and the mass of the dust cover have becu calculated with the model values.," All derived masses are collected in Table 2, where the total mass and the mass of the dust cover have been calculated with the model values."1187 The mass of the 2 pan sized dust. ceterumued inunediately after the outburst by Sekamina (2008)). is surprisingly high. compared to the otal mass of the dust laver!," The mass of the 2 $\mu$ m sized dust, determined immediately after the outburst by Sekanina \cite{sek3}) ), is surprisingly high, compared to the total mass of the dust layer!"1188 Even the total particulate dust mass (erains of size 0.2 to 7 nuu) is sinaller., Even the total particulate dust mass (grains of size 0.2 to 7 mm) is smaller.1189 luteerated over the 33 davs of increased cometary activity. he dust mass produced by the outburst is z of the colct’s total mass.," Integrated over the 33 days of increased cometary activity, the dust mass produced by the outburst is $\approx$ of the comet's total mass."1190 Surprisingly low is also he Πο mass uicased iu the first part of the outburst. when we would lave expected a mass comparable to the particulate dust uas.," Surprisingly low is also the $_2$ O mass released in the first part of the outburst, when we would have expected a mass comparable to the particulate dust mass."1191 The hvpothetical dust mass. calculated youn the radio Πο curve backwards. is about a factor of 3 l1 higher than the total mass in he burst.," The hypothetical dust mass, calculated from the radio light curve backwards, is about a factor of 3 -- 4 higher than the total mass in the burst."1192 The eunergv released in the burst cau be provided by the insolation before the burst. even allowing euergv losses through radiation by the dust cover.," The energy released in the burst can be provided by the insolation before the burst, even allowing energy losses through re-radiation by the dust cover."1193 The total accounted mass loss during this apparition (nass of 2 fr size dust. burst dust mass. burst TO mass} is iu total < of the total nuclear mass.," The total accounted mass loss during this apparition (mass of $2~\mu$ m sized dust, burst dust mass, burst $_2$ O mass) is in total $\le$ of the total nuclear mass."1194 If a bulk density of p = 0.5 is assed for the outer nucleus. this loss corresponds to a laver of 20 1n thickness which is the same order of magnitude as found from observations of other coimets.," If a bulk density of $\rho$ = 0.5 is assumed for the outer nucleus, this loss corresponds to a layer of 20 m thickness which is the same order of magnitude as found from observations of other comets."1195 Sekauina (2007..1982)) explained the outburst of comet 17PΠοιος as a splitting nucleus. whereby the secondary uucleus is a fragment of a jettisoned iusulation mantle of debris.," Sekanina \cite{sek}, \cite{sek1}) ) explained the outburst of comet 17P/Holmes as a splitting nucleus, whereby the secondary nucleus is a fragment of a jettisoned insulation mantle of debris."1196 The splitting starts with an exothermic event., The splitting starts with an exothermic event.1197 Tis model considers neither particulate dust with particles, His model considers neither particulate dust with particles1198"source. the region with a count rate larecr than the noulinemritv level for a uniform ilhuuination case (0,15 13) is of the size ~pixel (~0.711« 0.722).","source, the region with a count rate larger than the nonlinearity level for a uniform illumination case (0.15 ) is of the size $\sim 10 \times119915$pixel $\sim 0.''14 \times 0.''22$ )."1200 To compare this size with the wine of the poiut spread function (PSF: see section ??7)). we fud that in the mages of the PSFs scaled to have peak counts of 1.1. the region with counts larger than 0.15 is of the size ~5ν6 pixel," To compare this size with the wing of the point spread function (PSF; see section \ref{sec-res}) ), we find that in the images of the PSFs scaled to have peak counts of 1.4, the region with counts larger than 0.15 is of the size $\sim 5 \times 6$ pixel."

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