ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 We developed two dilferent methods. for the numerical treatment of time dependent ionization in the calculations., We developed two different methods for the numerical treatment of time dependent ionization in the calculations.3 Both have in common the method of finding xuhs from the ionizing source to the particles. along which he optical depth for the Lvman continuum photons can x caleulatec.," Both have in common the method of finding paths from the ionizing source to the particles, along which the optical depth for the Lyman continuum photons can be calculated."4 They. dilfer in the wav the ionization rate is determined given the radiation field., They differ in the way the ionization rate is determined given the radiation field.5 Method A uses the ormalism to calculate the divergence of the radiation field in Eq. 3.., Method A uses the formalism to calculate the divergence of the radiation field in Eq. \ref{eq:ionrate}.6 In. method D. we adopt a dillerent approach also used in grid. methods. where we derive the ionization rate rom the difference in the numbers of photons entering and caving a particle.," In method B we adopt a different approach also used in grid methods, where we derive the ionization rate from the difference in the numbers of photons entering and leaving a particle."7 First. we specify the position. the rate of ionizing photons Stor and σ (from Iq. 2))," First, we specify the position, the rate of ionizing photons $S_{\rm8tot}$ and $\bar{\sigma}$ (from Eq. \ref{eq:crossect}) )"9 of the source., of the source.10 For cach particle i we now proceed in the following way (sce Fig. 1)):, For each particle i we now proceed in the following way (see Fig. \ref{fig:construct}) ):11 Given the list of nearest neighbours of particle i. which has to be determined anyway for the formalism. we look for the particle j in the list. closest to the line of sight defined by the smallest angle O between the line connecting he particles i and j and the line ofsight.," Given the list of nearest neighbours of particle i, which has to be determined anyway for the formalism, we look for the particle j in the list, closest to the line of sight defined by the smallest angle $\Theta$ between the line connecting the particles i and j and the line of sight."12 We choose the angle tween. not the distance from. the line ofsight. since we are interested in controlling the error in the direction. towards he source.," We choose the angle between, not the distance from, the line of sight, since we are interested in controlling the error in the direction towards the source."13 This is not garanteed by the latter criterion., This is not garanteed by the latter criterion.14 We store this particle in à list and determine the evaluation point 5; as the projected. particle position on the ine of sight., We store this particle in a list and determine the evaluation point $_{\rm j}$ as the projected particle position on the line of sight.15 To determine the next evaluation point Sy even closer to the source we now repeat this method. using the neighbor list of particle j and so forth until we reach the SOULCC., To determine the next evaluation point $_{\rm k}$ even closer to the source we now repeat this method using the neighbor list of particle j and so forth until we reach the source.16 Now the path from the source to particle tis known. and the integration of Eq. (1))," Now the path from the source to particle i is known, and the integration of Eq. \ref{eq:LOS}) )"17 can be diseretized by using the evaluation points 5; , can be discretized by using the evaluation points $S_{\rm i}$.18Phe value for my can be estimated by using the smoothing formalism: where the sum runs over the particle corresponding to the evaluation point and. its nearest neighbours., The value for $n_{\rm H}$ can be estimated by using the smoothing formalism: where the sum runs over the particle corresponding to the evaluation point and its nearest neighbours.19 Wo is the weight factor for cach neighboring particle provided by the smoothing kernel., $W$ is the weight factor for each neighboring particle provided by the smoothing kernel.20 We calculate the optical depth along the line of sighte by applyinge the Trapezian Formula. until we reach particle i: with 5j being the position of the evaluation point on the line of sight.," We calculate the optical depth along the line of sight by applying the Trapezian Formula, until we reach particle i: with $s_{\rm k}$ being the position of the evaluation point on the line of sight."21 Note that this treatment neglects the ellects of scattering of the ionizing photons by recombination or dust., Note that this treatment neglects the effects of scattering of the ionizing photons by recombination or dust.22 The clistance between two successive evaluation points is smaller or equal to the local smoothing length. which determines the largest distance of the particles included in," The distance between two successive evaluation points is smaller or equal to the local smoothing length, which determines the largest distance of the particles included in"23'The result is what would be expected intuitively.,The result is what would be expected intuitively.24" The infrared and optical flux, which is produced almost exclusively in the circumbinary disc, does not vary greatly from what is expected from a standard thin disc."," The infrared and optical flux, which is produced almost exclusively in the circumbinary disc, does not vary greatly from what is expected from a standard thin disc."25" However, the flux drops precipitously below wavelengths of AX;3000. (v>10!?Hz in the figure)."," However, the flux drops precipitously below wavelengths of $\lambda\ltsim 3000 \AA$ $\nu>10^{15}\Hz$ in the figure)."26" This is in stark contrast to most unobscured quasars thought to be powered by ~105Μο SMBHs, which have theirbrightest emission in the rest-frame near-UV near their Lyman-a line."," This is in stark contrast to most unobscured quasars thought to be powered by $\sim 10^{8-9}\Msol$ SMBHs, which have their emission in the rest-frame near-UV near their $\alpha$ line."27" The bolometric luminosity of the accreting PTA source is roughly O.03Lgaa (ie., L/Lgaa10 ?m) for P=1yr, and ~107?Lgaa for P=0.1yr."," The bolometric luminosity of the accreting PTA source is roughly $\sim 0.03L_{\rm Edd}$ (i.e., $L/L_{\rm Edd}\sim 10^{-2}\dot{m}$ ) for $P=1\yr$, and $\sim 10^{-3}L_{\rm Edd}$ for $P=0.1\yr$."28" The optical and infrared emission is dominated by the circumbinary disc, whereas the UV and X-rays are produced by circum-secondary accretion fueled by leakage of circumbinary gas into the cavity."," The optical and infrared emission is dominated by the circumbinary disc, whereas the UV and X-rays are produced by circum-secondary accretion fueled by leakage of circumbinary gas into the cavity."29" As the binary evolves toward shorter periods, the circum-secondary disc is depleted — the viscous time at the Hill radius is typically a few hundred years, shorter than the time to binary merger — and as a result, less gas is able to leak into the cavity, decreasing the UV and X-ray emission."," As the binary evolves toward shorter periods, the circum-secondary disc is depleted — the viscous time at the Hill radius is typically a few hundred years, shorter than the time to binary merger — and as a result, less gas is able to leak into the cavity, decreasing the UV and X-ray emission."30 The degree to which the UV and X-ray emission is suppressed depends on the model parameters (in particular ficax) and on the binary period., The degree to which the UV and X-ray emission is suppressed depends on the model parameters (in particular $f_{\rm leak}$ ) and on the binary period.31 Note that the system may still be luminous in hard X-rays due to inverse Compton scattering by a coronal electron plasma (Sesanaetal.2011)., Note that the system may still be luminous in hard X-rays due to inverse Compton scattering by a coronal electron plasma \citep[][]{Sesana+11}.32. We also note that the downturn in the near-UV flux at AS300nm could help distinguish PTA sources from single-SMBH AGN., We also note that the downturn in the near-UV flux at $\lambda \ltsim 300~{\rm nm}$ could help distinguish PTA sources from single-SMBH AGN.33" This feature will be observable in the optical if the source redshift is high; e.g., at z=1 it will be in theV band."," This feature will be observable in the optical if the source redshift is high; e.g., at $z=1$ it will be in the band."34" Hence, even in the optical, this source will have an unusual color: it will appear fainter in theU andB bands than a typical AGN."," Hence, even in the optical, this source will have an unusual color: it will appear fainter in the and bands than a typical AGN."35 The downturn could be distinguished from reddening due to dust obscuration through the deviation from the power-law spectral shape of dust reddening., The downturn could be distinguished from reddening due to dust obscuration through the deviation from the power-law spectral shape of dust reddening.36" We propose that once an individually resolved PTA source is detected and its error box determined, searching for AGN with weak UV emission lines (e.g., Ly a) and/or weak soft X-ray emission is a promising method to narrow the field of interlopers."," We propose that once an individually resolved PTA source is detected and its error box determined, searching for AGN with weak UV emission lines (e.g., Ly $\alpha$ ) and/or weak soft X-ray emission is a promising method to narrow the field of interlopers."37" AGN whose soft X-ray fluxes are weaker by more than a factor of 10 compared to the average have indeed been detected, and are estimated to constitute at most c1% of the general AGN population (e.g.,Brandt,Schneider2008;Wuetal.2011,andrefs. therein).."," AGN whose soft X-ray fluxes are weaker by more than a factor of 10 compared to the average have indeed been detected, and are estimated to constitute at most $\sim 1\%$ of the general AGN population \citep[e.g.,][and38refs. therein]{Brandt+00, Leighly+07a, Gibson+08, Wu+11}."39" There have also been observations of quasars with exceptionally weak lines (Diamond-Stanicetal.2009);; these objects have infrared and optical emission consistent with those of typical luminous AGN, and also tend to be X-ray weak (Shemmeretal. 2009)."," There have also been observations of quasars with exceptionally weak lines \citep{Diamond+09}; these objects have infrared and optical emission consistent with those of typical luminous AGN, and also tend to be X-ray weak \citep{Shemmer+09}."40". That X-ray weak AGN are so rare suggests that it will be possible to narrow the number of interlopers in a typical PTA error box by a factor of z100, i.e. either to a handful of objects, or yielding a unique EM counterpart candidate."," That X-ray weak AGN are so rare suggests that it will be possible to narrow the number of interlopers in a typical PTA error box by a factor of $\approx 100$, i.e. either to a handful of objects, or yielding a unique EM counterpart candidate."41" It is possible, furthermore, that some of these rare X-ray weak AGN are in fact the SMBH binaries that PTAs will be detecting."," It is possible, furthermore, that some of these rare X-ray weak AGN are in fact the SMBH binaries that PTAs will be detecting."42" Our results also strongly suggest that AGN counterparts to PTA sources should draw from optically selected surveys, as their nature makes them likely to be missed by X-ray searches (see, however, Sesanaetal.2011,, who investigate the possible X-ray searches of PTA source binaries that have not yet decoupled)."," Our results also strongly suggest that AGN counterparts to PTA sources should draw from optically selected surveys, as their nature makes them likely to be missed by X-ray searches (see, however, \citealt{Sesana+11}, who investigate the possible X-ray searches of PTA source binaries that have not yet decoupled)."43 The should be able to detect all of the optically luminous AGN in the PTA error box within z~1., The should be able to detect all of the optically luminous AGN in the PTA error box within $z\sim 1$.44" It may be possible to follow up candidates individually, but comparing the optical data with that of wide-field X-ray surveys such asMAXB,, or would greatly facilitate the multi-wavelength search for counterpart candidates inside the error box."," It may be possible to follow up candidates individually, but comparing the optical data with that of wide-field X-ray surveys such as, or would greatly facilitate the multi-wavelength search for counterpart candidates inside the error box."45 Additional follow-up studies of candidates may further corroborate the identification of a counterpart., Additional follow-up studies of candidates may further corroborate the identification of a counterpart.46" For example, the gas that leaks radially into the cavity can shock-heat the outer edge of the circum-secondary (or circum-primary) disc and produce hot spots."," For example, the gas that leaks radially into the cavity can shock-heat the outer edge of the circum-secondary (or circum-primary) disc and produce hot spots."47" The viscously dissipated luminosity of a circum-secondary disc is roughly Lais»SM»M3/ (1/2)GRisco,2, where M»XfiaM(RÀ) is the mass supply rate of the circum-secondary disc and Zisco, is the radius of innermost stable circular orbit around the secondary."," The viscously dissipated luminosity of a circum-secondary disc is roughly $L_{\rm disc2}\ltsim (1/2)GM_{2}\dot{M}_{2}/R_{\rm ISCO,2}$ , where $\dot{M}_{2}\le f_{\rm leak}\dot{M}(R_{\lambda})$ is the mass supply rate of the circum-secondary disc and $R_{\rm ISCO,2}$ is the radius of innermost stable circular orbit around the secondary."48 The time-averaged power per unit mass of the hot spots is limited by the amount of kinetic energy the flow can deposit at the outer edge of the circum-secondary, The time-averaged power per unit mass of the hot spots is limited by the amount of kinetic energy the flow can deposit at the outer edge of the circum-secondary49at different wavelengths dominated by +>2 sources will be a powerful tool for studving the evolution of the Iargc-scale structure of infrared galaxies.,at different wavelengths dominated by $z>2$ sources will be a powerful tool for studying the evolution of the large-scale structure of infrared galaxies.50 The effect of the Iecorrection ensures that each of these maps (at different wavelengths) are dominated by particular high-redshift ranges., The effect of the K-correction ensures that each of these maps (at different wavelengths) are dominated by particular high-redshift ranges.51 Alethods of indepeudent component separation based on the correlation matrix between these maps (e.g“Se7) should allow us to extract maps and power spectra for a nmnuuber of redshift ranges equal to the ummber of maps.," Methods of independent component separation based on the correlation matrix between these maps \citep [e.g.,][] {2003MNRAS.346.1089D} should allow us to extract maps and power spectra for a number of redshift ranges equal to the number of maps."52 This last step will fulfill the main objective of this work., This last step will fulfill the main objective of this work.53 It will allow the study of the evolution o| the IR ealaxy clistering at hieli redshifts by means of the IOWCYL προςπα analysis of CIB anisotropics., It will allow the study of the evolution of the IR galaxy clustering at high redshifts by means of the power spectrum analysis of CIB anisotropies.54 These maps may also be used to help us understand the contribution of hieh-z IR galaxies both to the CIB and the star-formation history., These maps may also be used to help us understand the contribution of high-z IR galaxies both to the CIB and the star-formation history.55We consider as acceptable templates all stellar populations for which the resulting VCDoe ds increased+ by less than Ay? = 9.21. correspoudiug to a coufideuce level of (~ 2.6 o level) for 2 parameters.,"We consider as acceptable templates all stellar populations for which the resulting $\chi^2$ is increased by less than $\Delta\chi^2$ = 9.21, corresponding to a confidence level of $\sim$ 2.6 $\sigma$ level) for 2 parameters."56 In the case of 3C 310. the error derived setting \2=Lis 9.5 1015 ere 7 st d.," In the case of 3C 310, the error derived setting $\chi^2_{\rm r}57= 1$ is 9.5 $10^{-18}$ erg $^{-2}$ $^{-1}$ $^{-1}$."58 This value is similar to he rus of he stellar subtracted spectrum. which is 8.2 10IS cre D27s + Al.," This value is similar to the rms of the stellar subtracted spectrum, which is 8.2 $10^{-18}$ erg $^{-2}$ $^{-1}$ $^{-1}$."59 AdoptingB the procedure outlined- above. he best fit is the only statistically acceptable model.," Adopting the procedure outlined above, the best fit is the only statistically acceptable model."60 Nonetheless. in Fig.," Nonetheless, in Fig."61 6 we show the two models with the ower values of 4c. differing by one step in moetalieitv or age.," \ref{3c310} we show the two models with the lower values of $\chi^2_{\rm r}$, differing by one step in metalicity or age."62 Iu both cases siguificant residuals already. οσο iu he contimmuu subtracted spectrum. coufinüug visually he poorer quality of he fit.," In both cases significant residuals already emerge in the continuum subtracted spectrum, confirming visually the poorer quality of the fit."63 We re-cstimate the ITJ fiux resulting frou these two templates., We re-estimate the $\beta$ flux resulting from these two templates.64 They differ bv ~ with respect to our original rieasureimmenut using the best fit ποσο]., They differ by $\sim$ with respect to our original measurement using the best fit model.65 This is lower than the statistical error associated Xspocfit to his measurement. which is 7.," This is lower than the statistical error associated by to this measurement, which is 7."66. We couclude hat. iu the case of 3€ 9180. a possible template nmis-natch does not have a significant effect on he IL) fiux ueasurement.," We conclude that, in the case of 3C 310, a possible template mis-match does not have a significant effect on the $\beta$ flux measurement."67 The infence of the stellar templates on the other Cluission lines (aud im particular on the Πα line that is © 3 brighter that I) while they have the same EW in absorption) is negligible because the stellar population las not significant absorption features at the me wavelengths., The influence of the stellar templates on the other emission lines (and in particular on the $\alpha$ line that is $\gtrsim$ 3 brighter that $\beta$ while they have the same EW in absorption) is negligible because the stellar population has not significant absorption features at the line wavelengths.68 For the majority of the sources. oulv models differing by at most one step in age or iu metalicity are cousistcut with the data. without a significant impact ou he ID> flux.," For the majority of the sources, only models differing by at most one step in age or in metalicity are consistent with the data, without a significant impact on the $\beta$ flux."69 Oulv for the lower quality spectra can this rauge be broader., Only for the lower quality spectra can this range be broader.70 For low SNR spectra. coupled with emission lines of low equivalent widths. the resulting measurement error is dominated by the choice of the stellar template.," For low SNR spectra, coupled with emission lines of low equivalent widths, the resulting measurement error is dominated by the choice of the stellar template."71 In these cases we adopted the following strategy: we subtracted the acceptable template with the lowest age (6. with the strongest IL) absorption) and considered he nuneasurenient as an upper lit., In these cases we adopted the following strategy: we subtracted the acceptable template with the lowest age (i.e. with the strongest $\beta$ absorption) and considered the measurement as an upper limit.72 One of such galaxy is 3C 173.1. already: shown in the bottom right panel of Fie. 9," One of such galaxy is 3C 173.1, already shown in the bottom right panel of Fig. \ref{ssp}."73 Finally we note that our approach o the starlight subtraction is limited to the choice of he best fitting single population model. similarly to the method adopted o» RKauffhiaunetal.(2003) for the SDSS sample of cluission liue galaxies aud by Tadhuuteretal.(1993) or the 2Jv sample of radio galaxies.," Finally we note that our approach to the starlight subtraction is limited to the choice of the best fitting single population model, similarly to the method adopted by \citet{kauffmann03} for the SDSS sample of emission line galaxies and by \citet{tadhunter93} for the 2Jy sample of radio galaxies."74 Nonetheless. this oxocedure docs not erasp the complexity of the emission orocesses dn radio galaxies.," Nonetheless, this procedure does not grasp the complexity of the emission processes in radio galaxies."75 As clearly shown bv several studies iu the literature the stellar content of radio galaxies is often composed by populatious of different ages. and here are also evideuce for internal absorption. particularly associated to the voung stars component (c.g. Wills 200[: Ranimauiaunoetal.2005: Tadluuteretal. 2005:: IToltetal. 2007)).," As clearly shown by several studies in the literature the stellar content of radio galaxies is often composed by populations of different ages, and there are also evidence for internal absorption, particularly associated to the young stars component (e.g. \citealt{wills04}; ; \citealt{raimann05}; \citealt{tadhunter05}; \citealt{holt07}) )."76 The preseuce of stars of various ages is iucdirectlv coufirmed by the fact that the best ft stellar population found iu the rreeious is on average ον one step iu age vouueer than for he ybaud., The presence of stars of various ages is indirectly confirmed by the fact that the best fit stellar population found in the regions is on average $\sim$ one step in age younger than for the band.77 This can be explained by the higher relative contribution of vounug stars at shorter wavelengths that dives the age of the SSP to lower values with respect to he redder part of of the spectimim., This can be explained by the higher relative contribution of young stars at shorter wavelengths that drives the age of the SSP to lower values with respect to the redder part of of the spectrum.78 Furthermore. other uechamisius contribute to the contiuuun enüsson. e.g. iebular coutinmiun aud Πο frou the accretion disk (direct or wvia scattering).," Furthermore, other mechanisms contribute to the continuum emission, e.g. nebular continuum and light from the accretion disk (direct or via scattering)."79 Towever. a full separation of these contributions requires a detailed analysis on a object-by-object basis iu order to properly measure the rather laree set of free paraincters describing the various coniponeuts.," However, a full separation of these contributions requires a detailed analysis on a object-by-object basis in order to properly measure the rather large set of free parameters describing the various components."80 Moreover this fitting requires spectra with a lüeh signal ο noise. aud this is not always achieved im our data.," Moreover this fitting requires spectra with a high signal to noise, and this is not always achieved in our data."81 We defer such a study to a forthcoming paper., We defer such a study to a forthcoming paper.82 There is however a 8general consideration that can be drawn already at this stage., There is however a general consideration that can be drawn already at this stage.83 Nuclear emission aud nebular continui are expected to be in ecueral bluer tha the starlight., Nuclear emission and nebular continuum are expected to be in general bluer that the starlight.84 This biases the derived age of the sinele stellar population that appears vounecr that it actually is., This biases the derived age of the single stellar population that appears younger that it actually is.85 It also dilutes the stellar absorption features. thus biasing the imetalicitv to lower values.," It also dilutes the stellar absorption features, thus biasing the metalicity to lower values."86 The main resulting effect is an overestimate of the correction of the Tine and. consequeuth. an over-rating of this line. particularly in the objects with brightest optical nuclei.," The main resulting effect is an overestimate of the correction of the line and, consequently, an over-rating of this line, particularly in the objects with brightest optical nuclei."87 We stunumarize here the data quality of our spectra from he point of view of the measurements of the emission liue intensities., We summarize here the data quality of our spectra from the point of view of the measurements of the emission line intensities.88 Onlv in three radio galaxies (namely 3€ 052. 3€ 129.1 and 3€ 130) did we fail to detect amy eiiission ine: dm all cases an accurate redshift ean be measured x fitting the stellar population models based on stellar absorption features.," Only in three radio galaxies (namely 3C 052, 3C 129.1 and 3C 130) did we fail to detect any emission line; in all cases an accurate redshift can be measured by fitting the stellar population models based on stellar absorption features."89 Leaving aside these objects. the line can be measured with a statistical accuracy of better han 10 with only three exceptions.," Leaving aside these objects, the line can be measured with a statistical accuracy of better than 10 with only three exceptions."90 In the red part of he spectrum. the [O TA6300 line is detected in 95 of he 3C galaxies.," In the red part of the spectrum, the [O $\lambda$ 6300 line is detected in 95 of the 3C galaxies."91 The completeness of the |S II] doublet neasurements is «iehtle lower (87 ) due to our choice o favor the inclusion of the |O I| in the high resolution spectra but. when covered by the data. the |S II] line is also measurable iu 95 of the objects.," The completeness of the [S II] doublet measurements is slightly lower (87 ) due to our choice to favor the inclusion of the [O I] in the high resolution spectra but, when covered by the data, the [S II] line is also measurable in 95 of the objects."92 In the blue part of the spectrmu there are ~ 15 of the sources in which aand/or |O IH] are not detected., In the blue part of the spectrum there are $\sim$ 15 of the sources in which and/or [O III] are not detected.93 All together. these measurements will enable us to ocate the vast majoritv of the sources in the diagnostic aues that compare pairs of emission line ratios and also o perforin a detailed analysis of the relationship betweeu he iuultivavelensth characteristics of the 3€ sources with heir cuussion line properties.," All together, these measurements will enable us to locate the vast majority of the sources in the diagnostic planes that compare pairs of emission line ratios and also to perform a detailed analysis of the relationship between the multiwavelength characteristics of the 3C sources with their emission line properties."94 Iu addition. a broad line is seenin 18 galaxies aud these are discussed m more detail in the Sect. 1.3.," In addition, a broad line is seenin 18 galaxies and these are discussed in more detail in the Sect. \ref{broad}. ."95extended to include significant numbers of compact. high surface-brightness galaxies that we think are background. in the stvle of the Drinkwater ct al. (,"extended to include significant numbers of compact, high surface-brightness galaxies that we think are background, in the style of the Drinkwater et al. ("962000a) Fornax survey.,2000a) Fornax survey.97 If the vast majority of such objects are indeed background objects. this would alleviate the concern. highlighted. in section 7 that the current sample is heavily incomplete due o us rejecting such galaxies.," If the vast majority of such objects are indeed background objects, this would alleviate the concern highlighted in Section 7 that the current sample is heavily incomplete due to us rejecting such galaxies."98 Such a project is now feasible with the advent of wide-field multi-objeet spectrographs on arge telescopes., Such a project is now feasible with the advent of wide-field multi-object spectrographs on large telescopes.99 More detailed: studies of the dwarf. galaxies found. in Virgo will also be of value in assessing the importance of he various physical processes at work during galaxy ancl cluster formations., More detailed studies of the dwarf galaxies found in Virgo will also be of value in assessing the importance of the various physical processes at work during galaxy and cluster formations.100 The following observations should be of xaurticular use: Ipt 1) measurement of colours ancl elemental abundances., The following observations should be of particular use: 1pt 1) measurement of colours and elemental abundances.101 These will constrain the star-formation histories of the Virgo dwarfs which in turn will allow a lower limit to be placed on the redshift at which gas was collected by small dark-matter halos., These will constrain the star-formation histories of the Virgo dwarfs which in turn will allow a lower limit to be placed on the redshift at which gas was collected by small dark-matter halos.102 This is of importance in the context of the squelching picture described above., This is of importance in the context of the squelching picture described above.103 This squelching picture also relies on the very existence of these dark-matter halos around the Virgo dwarks., This squelching picture also relies on the very existence of these dark-matter halos around the Virgo dwarfs.104 In the long term. studies of the kinematics of stars in the chwarfs will be required to verily this assertion: Ipt 2) LIE measurements of the cold gas content of the Virgo thwarts.," In the long term, studies of the kinematics of stars in the dwarfs will be required to verify this assertion; 1pt 2) HI measurements of the cold gas content of the Virgo dwarfs."105 Galaxies that have staved any appreciable time in the Vireo Cluster would lose their gas via ram-pressure stripping from the cluster A-ray halo., Galaxies that have stayed any appreciable time in the Virgo Cluster would lose their gas via ram-pressure stripping from the cluster X-ray halo.106 Therefore a large number of dwarfs in Virgo with LIE would suggest that many Virgo cbwarfs only recently entered the cluster., Therefore a large number of dwarfs in Virgo with HI would suggest that many Virgo dwarfs only recently entered the cluster.107 This would in turn argue against any physical process that requires Virgo cdwarfs to have formed in the cluster or at very early times in small groups which quickly. merged to form the eluster: Ipt 3) the location of dwarfs in the cluster., This would in turn argue against any physical process that requires Virgo dwarfs to have formed in the cluster or at very early times in small groups which quickly merged to form the cluster; 1pt 3) the location of dwarfs in the cluster.108 Do giants maintain their dwarf populations once they are in the Virgo Cluster. or do the dwarfs adopt orbits determined by the cluster potential?," Do giants maintain their dwarf populations once they are in the Virgo Cluster, or do the dwarfs adopt orbits determined by the cluster potential?"109 LW the former were true. chvarls would tend to cluster around giants.," If the former were true, dwarfs would tend to cluster around giants."110 I£ the latter were true. they would be smoothly clistributecl throughout the cluster. with a radial density profile similar to that of the giant galaxics.," If the latter were true, they would be smoothly distributed throughout the cluster, with a radial density profile similar to that of the giant galaxies."111 The answer to this question will provide constraints on the cark-matter structure of the Virgo Cluster., The answer to this question will provide constraints on the dark-matter structure of the Virgo Cluster.112 This work was based on observations made through the Isaac Newton Groups’ Wide Field. Camera Survey Programme with the Isaac Newton Telescope operated. on the. islam of La Palma by the Isaac. Newton Group in the Spanish Observatorio del oque de Ios Muchachos of the Instituto de Astrolisica de Canarias., This work was based on observations made through the Isaac Newton Groups' Wide Field Camera Survey Programme with the Isaac Newton Telescope operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias.113 Phe contributions of J. Davies anc t. AleAlahon in initiating the WES and the Virgo part of the survey are acknowledged in. particular., The contributions of J. Davies and R. McMahon in initiating the WFS and the Virgo part of the survey are acknowledged in particular.114 We are also eratefu o the referee. Dr. S. Phillipps. for detailed comments on the manuscript.," We are also grateful to the referee, Dr. S. Phillipps, for detailed comments on the manuscript."115 Fhis research has made use of the NASA/LPAC Extragalactic Database (NIZD) which is operated by the Je ropulsion Laboratory. Caltech. uncer agreement with the ational Aeronautics and Space Association.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, Caltech, under agreement with the National Aeronautics and Space Association."116aud a for different ratios of C'/E...,and $\alpha$ for different ratios of $_{\rm esc}$.117 A redshift dependent clumping factor such as that shown in Figure 2a and an escape fraction of 0.1 is shown as (he green lines., A redshift dependent clumping factor such as that shown in Figure 2a and an escape fraction of 0.1 is shown as the green lines.118 The red lines show the estimates lor a clumping factor of 30 ancl an escape fraction of 0.1., The red lines show the estimates for a clumping factor of 30 and an escape fraction of 0.1.119 The blue line show the estimates for a chunping factor of 10 and an escape fraction of 0.15 while the evan line is an extreme case with a clumping [actor of 10 and escape fraction of 0.5., The blue line show the estimates for a clumping factor of 10 and an escape fraction of 0.15 while the cyan line is an extreme case with a clumping factor of 10 and escape fraction of 0.5.120" A similar comparison between the redshift when reionization is complete ancl the IME slope is shown in Figure 6 for the ""low-V case."," A similar comparison between the redshift when reionization is complete and the IMF slope is shown in Figure 6 for the “low-V"" case."121" As can be seen for the ""high-V case. with a Salpeter IAIF (a=—2.3). unless the ratio of clumping factor to escape fraction is much smaller (han GO. the stus in z~6 galaxies cannot be responsible for reionization αἱ 2>7."," As can be seen for the “high-V"" case, with a Salpeter IMF $\alpha=-2.3$ ), unless the ratio of clumping factor to escape fraction is much smaller than 60, the stars in $z\sim6$ galaxies cannot be responsible for reionization at $z>7$."122" For our best estimate of the clumping factor and escape fraction. il 2,44,í,29./heslopea. =—1.65 while if οι. a=—1.5."," For our best estimate of the clumping factor and escape fraction, if $=9$ , the slope $\alpha=-1.65$ while if $=11$, $\alpha=-1.5$ ."123 In either scenario. only of all barvons at ze6 are processed through stars and are responsible for ionizing the IGM.," In either scenario, only of all baryons at $z\sim6$ are processed through stars and are responsible for ionizing the IGM."124 Our conclusions are therefore broadly consistent with the results of (2004) who argue against the need for a verv massive star population at high redshilt., Our conclusions are therefore broadly consistent with the results of \citet{Tum:04} who argue against the need for a very massive star population at high redshift.125 Their arguments are based on the metal abundance ratios in metal-poor Galactic halo stars., Their arguments are based on the metal abundance ratios in metal-poor Galactic halo stars.126 llowever. we disagree with Chem in the sense that a Salpeter IAIF even with a low mass οσο. is inadequate for producing sufficient ionizing photons to ionize the IGM while fitting the observed ultraviolet anc optical luminosity density al 2~6 (Figure 3).," However, we disagree with them in the sense that a Salpeter IMF even with a low mass cutoff, is inadequate for producing sufficient ionizing photons to ionize the IGM while fitting the observed ultraviolet and optical luminosity density at $z\sim6$ (Figure 3)."127 such a low mass turnover in the stellar IAIF maa indeed be present al 2< as has been argued by Dave(2008) ancl vanDokkum(2008).., Such a low mass turnover in the stellar IMF may indeed be present at $z<2$ as has been argued by \citet{Dave} and \citet{PvD}.128 The lormer attempted to reconcile specific star-formalion rate estimates between models and observations while the latter fit the color evolution of massive earlv-tvpe cluster galaxies al 2<I., The former attempted to reconcile specific star-formation rate estimates between models and observations while the latter fit the color evolution of massive early-type cluster galaxies at $z<1$.129 We attempted to fit the observed luminosity densities at z~6 with the IME proposed by Davé(2008).., We attempted to fit the observed luminosity densities at $z\sim6$ with the IMF proposed by \citet{Dave}.130 We assumed that the (1+2)? evolution in the mass turnover of the stellar IMFE. extends up to z~ 4. similar to the evolution proposed by vanDokkum(20," We assumed that the $(1+z)^2$ evolution in the mass turnover of the stellar IMF extends up to $z\sim4$ , similar to the evolution proposed by \citet{PvD}."131"05), Thus. the IME was assumed to have a slope of —1.3 at 0.1 «M12 MM. and —2.3 at 20022 M212 MM, ."," Thus, the IMF was assumed to have a slope of $-1.3$ at $0.1<$ $<12$ $_{\sun}$ and $-2.3$ at $>$ $>12$ $_{\sun}$ ."132" For the ""high-V case. I find that this IME produces 3.25 ionizing photons per baryon. an age of the stellar population of 90 Myr and a stellar mass of 1.2x LO"" nun."," For the “high-V"" case, I find that this IMF produces 3.25 ionizing photons per baryon, an age of the stellar population of 90 Myr and a stellar mass of $\times$ $^7$ ."133 Unlike the Salpeter IMEs with a sharp cutoff shown in Figure 3. such an IMF can account for the the “late” reionization historv of Trac&Cen(2007) and would maintain an ionized IGM at zX7.," Unlike the Salpeter IMFs with a sharp cutoff shown in Figure 3, such an IMF can account for the the “late"" reionization history of \citet{Trac:07}134 and would maintain an ionized IGM at $z\lesssim 7$."135" However. for the ""low-V case. this [MIF would produce only 1.0 ionizing photons per barvon. result in a stellar age of45Myr and a stellar mass of x 3.710°nun.. about a factor of 3lower thantheminimum required to maintain an ionized IGM at of."," However, for the “low-V"" case, this IMF would produce only 1.0 ionizing photons per baryon, result in a stellar age of45Myr and a stellar mass of $\times$ $^6$, about a factor of 3lower thantheminimum required to maintain an ionized IGM at $z<7$ ."136"all euergev is released from f=τι to f=rz;|τι, lu",all energy is released from $t = \tau_i$ to $t = \tau_i + \tau_b$.137" many problems. where 7,τι. this can be further siuplified to burniug occuring ouly at £—7;."," In many problems, where $\tau_b \ll \tau_i$, this can be further simplified to burning occurring only at $t = \tau_i$."138 Where such an approximation (often called ‘high asvinptotics) holds. it ereatly simplifies may problems of burning or ignition. reducing the region of burning in a flame to an infinitesimally thin ‘flamelet’ (7) surface. for iusance. or the structure of a detonation to a “square wave (7).," Where such an approximation (often called `high activation-energy asymptotics') holds, it greatly simplifies many problems of burning or ignition, reducing the region of burning in a flame to an infinitesimally thin `flamelet' \citep{mandm}139 surface, for instance, or the structure of a detonation to a `square wave' \citep{erpenbeck}."140" Where this approximation does not hold such as if slow decay processes are Προστit as bottlenecks for reactions to proceed(6.9... pep buruiie or the CNO evele) the simplification of burniie happeniug only over τεςf€7;|7, often remains uscef1."," Where this approximation does not hold – such as if slow $\beta$ -decay processes are important as bottlenecks for reactions to proceed, $p$ $p$ burning or the CNO cycle) the simplification of burning happening only over $\tau_i \le t \le \tau_i +141\tau_b$ often remains useful."142 Even for the simple case of one zone. ignition delay. times are rclevaut for iuvestigation of ignition iu SNIun because it sets a uininmuu time scale over which an initial local positive teurerature perturbation (hot spot) can successfully ignite and lauuch a combustion wave: other timescales. such as turbulent disruption of the hotspot. or diffsive timescales. must be larger than this for iguition to successfully occur.," Even for the simple case of one zone, ignition delay times are relevant for investigation of ignition in SNIa because it sets a minimum time scale over which an initial local positive temperature perturbation (hot spot) can successfully ignite and launch a combustion wave; other timescales, such as turbulent disruption of the hotspot, or diffusive timescales, must be larger than this for ignition to successfully occur."143 Tf the burning occurs in an ideal eas. or iu a 1iaterial with some other simple equation of state. it is fairly casy to write down approximate ignition times for various burning laws.," If the burning occurs in an ideal gas, or in a material with some other simple equation of state, it is fairly easy to write down approximate ignition times for various burning laws."144 In a white dwarf. however. where the material is partially degeuerate or relativistic and the equation of state is quite complicated (?2).. no such closed-form expression cau be written.," In a white dwarf, however, where the material is partially degenerate or relativistic and the equation of state is quite complicated \citep{eos}, no such closed-form expression can be written."145 In 822 we nuuerically follow the abundance aud thermodvuamiic evolution of a zone of white dwarf material in order to measure the iguition times as a function of the initial temperature. density. aud composition.," In 2 we numerically follow the abundance and thermodynamic evolution of a zone of white dwarf material in order to measure the ignition times as a function of the initial temperature, density, and composition."146 We follow both coustaut-deusitv and coustaut-pressure trajectories., We follow both constant-density and constant-pressure trajectories.147 The results are πανΊος by simple. moderately accurate. fitting fonnula.," The results are summarized by simple, moderately accurate, fitting formula."148 In 33 we cousider the ignition of a detonation through a localized οποιον release producing a Sedov blast wave. aud estimate the amount of energy that must be released for the detonation to successfully ignite.," In 3 we consider the ignition of a detonation through a localized energy release producing a Sedov blast wave, and estimate the amount of energy that must be released for the detonation to successfully ignite."149 In 511 we consider our results in liebt of likely temperature fluctuation svectra during the simuucring. couvective phase.," In 4 we consider our results in light of likely temperature fluctuation spectra during the simmering, convective phase."150 We performed a series of l-zone calculations for the purposes of measuiug ignition times iu mixtures., We performed a series of 1-zone calculations for the purposes of measuring ignition times in carbon-oxygen mixtures.151 For cach of these two burning conditions burning at constant νοιό aud coustaut pressure, For each of these two burning conditions – burning at constant volume and constant pressure152unification schemes since hotter dust is visible in BLIRGs due to their viewing angle.,unification schemes since hotter dust is visible in BLRGs due to their viewing angle.153 ISOCAAL was more than an order of magnitude more sensitive thanIRAS at MIB wavelengths., ISOCAM was more than an order of magnitude more sensitive than at MIR wavelengths.154 It was able to detect the hot dust component in 3C sources at much. higher redshifts and therefore expand comparison of tvpe 1 and 2 AGNs to a much wider range of Iuninosities., It was able to detect the hot dust component in 3C sources at much higher redshifts and therefore expand comparison of type 1 and 2 AGNs to a much wider range of luminosities.155 In this letter. we report first results [rom a program to investigate the hot dust component in 38C sources based on ISOCAM observations.," In this letter, we report first results from a program to investigate the hot dust component in 3C sources based on ISOCAM observations."156 A detailed discussion of our data reduction ancl modelling procecures as well as a full presentation of our data are delerred to a later paper (Siebenmoreenοἱal.2003a.hereafterSEIRIL., A detailed discussion of our data reduction and modelling procedures as well as a full presentation of our data are deferred to a later paper \citep[][hereafter SFKH]{mainpaper}.157. The goal of our program is (o investigate dust emission in AGNs. and in particular to search for differences related to the orientation of the central engine.," The goal of our program is to investigate dust emission in AGNs, and in particular to search for differences related to the orientation of the central engine."158 Constraining the SED enussion of both hot and cold dust requires infrared photometry covering the full wavelength range from jan to min wavelengths., Constraining the SED emission of both hot and cold dust requires infrared photometry covering the full wavelength range from $\mu$ m to mm wavelengths.159 We used the 3CR catalogue as our parent catalogue because its 178. MlIZ selection should be essentially unbiased in its inclusion of AGNs ad different. viewing angles., We used the 3CR catalogue as our parent catalogue because its 178 MHz selection should be essentially unbiased in its inclusion of AGNs at different viewing angles.160 In addition. photometric measurements at far infrared (FUR) and nm wavelengths are available for large subsamples of the catalogue (e.g.Meisenheimneretal.2001:Laaset 2003b).," In addition, photometric measurements at far infrared (FIR) and mm wavelengths are available for large subsamples of the catalogue \citep[e.g.][]{Meise,Haas}."161. By complementing these available data with our new MIR neasuremenis. we were able to construct SEDs [rom optical to mm wavelengths.," By complementing these available data with our new MIR measurements, we were able to construct SEDs from optical to mm wavelengths."162 We have compiled our sample by. eross-correlating (he position of 3€ sources with the ist of SOCAN observations available from the ISO post mission archive., We have compiled our sample by cross-correlating the position of 3C sources with the list of ISOCAM observations available from the ISO post mission archive.163 We then extracted all ISOCAM images which potentially include 3C sources., We then extracted all ISOCAM images which potentially include 3C sources.164 Each reduced image was visually inspected. and (he pointing of the spacecraft was verified using stellar ancl other sources contained in the images., Each reduced image was visually inspected and the pointing of the spacecraft was verified using stellar and other sources contained in the images.165 We report here on (hat subsample of our data For which the combined ISOCAM photometry. data from the literature and ISOPHOT photometry by (2003b) sullice to constrain the full infrared SED of our sources.," We report here on that subsample of our data for which the combined ISOCAM photometry, data from the literature and ISOPHOT photometry by \citet{Haas} suffice to constrain the full infrared SED of our sources."166 ISOCAM images in the ISO archive have been obtained [or a variety of programs and purposes and therefore vary in depth. resolution. used wavelength band and dither strategy.," ISOCAM images in the ISO archive have been obtained for a variety of programs and purposes and therefore vary in depth, resolution, used wavelength band and dither strategy."167 We re-processed (he raw data using a common set of procedures aimed al obtaining the, We re-processed the raw data using a common set of procedures aimed at obtaining the168Iu the past two decades ereat advances lave been iade in observational CORnology.,In the past two decades great advances have been made in observational cosmology.169 The most striking suele discovery Is the preseut acceleration of the universe expansion. now confrned by many iundepenudoenut experiments.," The most striking single discovery is the present acceleration of the universe expansion, now confirmed by many independent experiments."170 The inost peaverful probe of precision cosmology is the observatioIs of the cosmic microwave vackeround (?7)). which sceu to support the model of universe which at huge scales is flat. isotropic aud rolmogencous. as firmly precicted by inflation.," The most powerful probe of precision cosmology is the observations of the cosmic microwave background \cite{Bennett:2003bz,Hinshaw:2006ia}) ), which seem to support the model of universe which at large scales is flat, isotropic and homogeneous, as firmly predicted by inflation."171 However. at amore subtle level there seenis to be also hints of substantial auisotropy.," However, at more subtle level there seems to be also hints of substantial anisotropy."172 Such would mualy violation of 16 cosmoloeical principle. perlaps as strisine change of xuwadieni as the inroductiou of dark οwereyv.," Such would imply violation of the cosmological principle, perhaps as striking change of paradigm as the introduction of dark energy."173 Presently ιο evidence for anisotropy 1 Sebatable. but the bounds can definitely be expected to improve with the Plauck experiment.," Presently the evidence for anisotropy is debatable, but the bounds can definitely be expected to improve with the Planck experiment."174 Therefore it 1s extremely interesting to study theoretical links botween the acceleration and anisotropies. iun particular. the possibility to constrain them observationally ?," Therefore it is extremely interesting to study theoretical links between the acceleration and anisotropies, in particular, the possibility to constrain them observationally \cite{Copi:2010na}."175 Several distinct statically auisotropic features have been reported iu he cata analysis of he CMD. sky., Several distinct statically anisotropic features have been reported in the data analysis of the CMB sky.176 Amoue the most ctvous is the presence o hemispherical asvnuuetrv (7))., Among the most curious is the presence of hemispherical asymmetry \cite{Eriksen:2003db}) ).177 Recent mvoestigations exploiting the five-vear WALAP «ata lave ound that the evidence for this asviunetry Is Increasing and exends to much sinaller angular scales tha ]xeviouslv believed to TT).," Recent investigations exploiting the five-year WMAP data have found that the evidence for this asymmetry is increasing and extends to much smaller angular scales than previously believed to \cite{Hansen:2008ym,Hoftuft:2009rq}) )."178 Aligminent of the quadriupok| ad octupole. the so called Axis of Evi (?)) could also seeni an unuikcly result of statistically isotropic perulvations. even without taking into account that these multivoles happen also to be aligned to some extent with the dipole aud with the equinox.," Alignment of the quadrupole and octupole, the so called Axis of Evil \cite{Land:2005ad}) ) could also seem an unlikely result of statistically isotropic perturbations, even without taking into account that these multipoles happen also to be aligned to some extent with the dipole and with the equinox."179 In the CAD spectrtun. the angular correlation spectra seenis to be lacsing power at the largest scales.," In the CMB spectrum, the angular correlation spectrum seems to be lacking power at the largest scales."180 The aliguients seem to be statistically independent of the the lack of auelay power (?))., The alignments seem to be statistically independent of the the lack of angular power \cite{Rakic:2007ve}) ).181 For other studies. sec (Hg.," For other studies, see \cite{Prunet:2004zy,Gordon:2005ai}) )."182 It is natural to associate the appareut statistical anisotropy with dark energw. since the auonmalies occur at the largest scales. aud these cuter inside the horizon at the same epoch tvat the dark cuerey domuinauce beeius.," It is natural to associate the apparent statistical anisotropy with dark energy, since the anomalies occur at the largest scales, and these enter inside the horizon at the same epoch that the dark energy dominance begins."183 The paramoit characteristic of dark CLOITSM is its negativo pressire., The paramount characteristic of dark energy is its negative pressure.184 One may then coutcuplate whether this pressure nuelt vary with the direction., One may then contemplate whether this pressure might vary with the direction.185 Then also the universa acceleration becomes anisotropic. and one would iudeed see otherwise unexpecte effects.," Then also the universal acceleration becomes anisotropic, and one would indeed see otherwise unexpected effects."186 These would prestunably be strongest at the siuallest uultipoles of the CAD suce they describe tje large aneular scales which are most directly affected dirine the ate epochs of the universe., These would presumably be strongest at the smallest multipoles of the CMB since they describe the large angular scales which are most directly affected during the late epochs of the universe.187 Specifically. as the plotous ravel from the last scattering surface towards us. the heir temperature gets bπο and redshifted as μον fall iu aud climb out of the οravitational wells. respectively.," Specifically, as the photons travel from the last scattering surface towards us, the their temperature gets blue- and redshifted as they fall in and climb out of the gravitational wells, respectively."188 When the poteutials evolve. there ls a net effect in he temperature of the yhotous: this is the integrated Sachs-Wolte effect. CSW).," When the potentials evolve, there is a net effect in the temperature of the photons: this is the integrated Sachs-Wolfe effect (ISW)."189 Furthermore. if the average evolution of the poteutias was not the same in ciffereut directions of the sky. the effect. would )o anisotropic.," Furthermore, if the average evolution of the potentials was not the same in different directions of the sky, the effect would be anisotropic."190" ILowever. to explain the ack of larec-augle correlations. there should occur a cancellation with tιο Sachs-Wolte effect frou, the potentials last scatterine surface that typically coutribute to the| Jlayge aneles wih sinular order of magnitude as the ISW ?"," However, to explain the lack of large-angle correlations, there should occur a cancellation with the Sachs-Wolfe effect from the potentials last scattering surface that typically contribute to the large angles with similar order of magnitude as the ISW \cite{Afshordi:2008rd}."191 The potentials paraitering the portirbatious of the luetric. can be written iu the longitudiua gauge as ds? =," The potentials parameterising the perturbations of the metric, can be written in the longitudinal gauge as ds^2 ="192to a precision of 2.5 runs on average at the highest redshüfts of cach survey.,to a precision of 2.5 r.m.s on average at the highest redshifts of each survey.193 Iu 5 davs of calendar time. the total exposure times züunount to 56 and LO hours for the wide aud deep surveys respectively. which corresponds to SO of the 120 hours available.," In 5 days of calendar time, the total exposure times amount to 56 and 40 hours for the wide and deep surveys respectively, which corresponds to 80 of the 120 hours available."194 Both survevs are couducted in 7 bands. covering ο to TT (150<A1700 ni).," Both surveys are conducted in 7 bands, covering g to H $450<\lambda<1700\ \mathrm{nm}$ )."195 One mav arene that the bluest bands of the deep survey aud the reddest bauds of the wide survey are not strictly to neededmeasure supernovae distances. but since visible aud ΣΙ observations are assumed to happen in parallel. dropping blue visible bands for the deep survey. or red NIB. bauds for the wide survey docs uot save any observing time.," One may argue that the bluest bands of the deep survey and the reddest bands of the wide survey are not strictly needed to measure supernovae distances, but since visible and NIR observations are assumed to happen in parallel, dropping blue visible bands for the deep survey, or red NIR bands for the wide survey does not save any observing time."196 Note that we have not discussed the collection of a nearby sample (typically O(1000). events at 0.03<2« 0.1)., Note that we have not discussed the collection of a nearby sample (typically O(1000) events at $0.03<z<0.1$ ).197 As a baseline. and at variance with the SNAP project. we stick to a selfcoutained imaging survey. in order to realistically limit eross-calibration aud detection bias issues.," As a baseline, and at variance with the SNAP project, we stick to a self-contained imaging survey, in order to realistically limit cross-calibration and detection bias issues."198 SNe Ia exhibit reproducible rest fune colours (a(BV)zm(kl. see eg. Fig.," SNe Ia exhibit reproducible rest frame colours $\sigma(B-V) \simeq 0.1$ , see e.g. Fig."199 8 of ?7)) aud even more reproducible colour relations.," 8 of \citealt{Astier06}) ), and even more reproducible colour relations."200 For example. the rest frame U-band amplitude can be predicted. to better than ~ 0.01 from B- and V-bands (2)..," For example, the rest frame U-band amplitude can be predicted to better than $\sim$ 0.04 from B- and V-bands \citep{Astier06}."201 SNe Ia not oulv occupy narrow subspaces of imulticolour spaces. but also exhibit very reproducible elt curve shapes that permit to discriminate against imost of the core-collapse events (277)..," SNe Ia not only occupy narrow subspaces of multicolour spaces, but also exhibit very reproducible light curve shapes that permit to discriminate against most of the core-collapse events \citep{Poznanski02,JohnsonCrotts06,Rodney09}."202 Studies of the photometric selection of SNe Ia have heen conducted ou the SNLS data aud their preliminary conclusions are eucouragiue (22)...," Studies of the photometric selection of SNe Ia have been conducted on the SNLS data and their preliminary conclusions are encouraging \citep{RipochePHD,BazinPHD}."203 These studies rely on host galaxy photometric redshift aud inost of their identification-⋅⋅: failuresDa are duc to wrong assumed redshifts., These studies rely on host galaxy photometric redshift and most of their identification failures are due to wrong assumed redshifts.204""" The7 supernova survey we are considering. here would be in a more favourable situation than these studies: it measures To bands (whilst SNLS has at most d.) aud we assmuue that host galaxy spectroscopic redshifts will be available (whilst SNLS studies used host galaxy photometric redshifts)."," The supernova survey we are considering here would be in a more favourable situation than these studies: it measures 7 bands (whilst SNLS has at most 4), and we assume that host galaxy spectroscopic redshifts will be available (whilst SNLS studies used host galaxy photometric redshifts)."205 Iu order to estimate he contannuatio1 by core-collapse supernova ina sample selected im colour-colour subspaces. we would need a laree enough sample of niulti-baud measurements of such eveuts," In order to estimate the contamination by core-collapse supernova in a sample selected in colour-colour subspaces, we would need a large enough sample of multi-band measurements of such events."206 These should soon be available. thanks at least to the SDSS supornova survey. the Lick Observatory Supernova Search. iud tle Palomar Trausieut Factory. but they are not available vet.," These should soon be available, thanks at least to the SDSS supernova survey, the Lick Observatory Supernova Search, and the Palomar Transient Factory, but they are not available yet."207 Therefore. in order to bound the impact of core-collapse contanunation onthe SNe Ta distanee-redshitt relation. we resort to studying how clipping around the IIubble line rejects other supernuova types. following ?..," Therefore, in order to bound the impact of core-collapse contamination on the SNe Ia distance-redshift relation, we resort to studying how clipping around the Hubble line rejects other supernova types, following \cite{Conley10}."208 Core collapse supernovae are classified ia Type Ib aud To. and Type IL," Core collapse supernovae are classified in Type Ib and Ic, and Type II."209 Type Ib and Ic events are often merecd iuto a “The type (see e.g. 2?)).," Type Ib and Ic events are often merged into a “Ibc” type (see e.g. \citealt{Richardson02, Li10II}) )."210 Tvpo IL eveuts are about 3 times more frequent than Toc (? Fig., Type II events are about 3 times more frequent than Ibc \citealt{Li10II} Fig.211 9). but two thirds of those are Type I-plateau (IT-p) which are easily identified from their very flat light curves," 9), but two thirds of those are Type II-plateau (II-p) which are easily identified from their very flat light curves."212 Other Type EH events represent about the same rate as Toc. but their ποτ curves rise dn a few days. whereas SNe Ta rise iu more than 15 davs.," Other Type II events represent about the same rate as Ibc, but their light curves rise in a few days, whereas SNe Ia rise in more than 15 days."213 So. Type II eveuts might add a simall coutribution to Tbe interlopers. and we will now couceutrate on evaluating the impact of a The contaminationin the ITubble diagram.," So, Type II events might add a small contribution to Ibc interlopers, and we will now concentrate on evaluating the impact of a Ibc contaminationin the Hubble diagram."214 We model the Ibe population absolute magnitude distribution as a Caussian offset by A. from the Ia population with r.urs 25., We model the Ibc population absolute magnitude distribution as a Gaussian offset by $\Delta_{bc}$ from the Ia population with r.m.s $\sigma_{bc}$ .215 We expect the rate of Ibe to be proportional to the star formation rate. that we take from 2.. aud the amount of Inc eveuts follows from f5.(:=0). he ratio of the Thc rate to the Ia rate at :=0.," We expect the rate of Ibc to be proportional to the star formation rate, that we take from \cite{Hopkins06}, and the amount of Ibc events follows from $f_{bc}(z=0)$, the ratio of the Ibc rate to the Ia rate at $z=0$."216 The adopted Ia rate was preseuted in 82. 1.., The adopted Ia rate was presented in \ref{sec:cadence_and_coverage}. .217 We smiulate a iux of Ta aud Toe eveuts. fit the Wnbblediagram with a sixth degree polvuonial. aud iteratively," We simulate a mix of Ia and Ibc events, fit the Hubblediagram with a sixth degree polynomial, and iteratively"218that the value6=0.25 given by von Zeipel’s law is appropriate only as expected.,that the value $\beta=0.25$ given by von Zeipel's law is appropriate only as expected.219 In this figure we also plotted measurements of the gravity darkening exponent P derived by interferometry for several rapidly rotating stars. Altair (2).. a Cephei (?).. 8 Cassiopeiae. and a Leonis (?)..," In this figure we also plotted measurements of the gravity darkening exponent $\beta$ derived by interferometry for several rapidly rotating stars, Altair \citep{Monnier2007}, , $\alpha$ Cephei \citep{Zhao2009}, $\beta$ Cassiopeiae, and $\alpha$ Leonis \citep{Che2011}."220 We can see that there is good (Altair. aw Leo) or fair (a Cep) agreement between the observed values and the model.," We can see that there is good (Altair, $\alpha$ Leo) or fair $\alpha$ Cep) agreement between the observed values and the model."221 The discrepancy for 6 Cas may come from its small inclination angle (~20 deg). so we see the star near pole-on. which makes determining of its flattening more difficult and makes the results depend more on the model used for gravity darkening and limb darkening.," The discrepancy for $\beta$ Cas may come from its small inclination angle $\sim20$ deg), so we see the star near pole-on, which makes determining of its flattening more difficult and makes the results depend more on the model used for gravity darkening and limb darkening."222 Observing that the energy radiated by a star is produced almost entirely in the stellar core and that most stars may be considered close to a steady state. we have noted that the energy flux is essentially a divergence-free vector field.," Observing that the energy radiated by a star is produced almost entirely in the stellar core and that most stars may be considered close to a steady state, we have noted that the energy flux is essentially a divergence-free vector field."223 We also observed from full two-dimensional models of rotating stars that the direction of this vector is always very close to that of the effective gravity. so it only depends on the mass distribution.," We also observed from full two-dimensional models of rotating stars that the direction of this vector is always very close to that of the effective gravity, so it only depends on the mass distribution."224 Following this picture. the physical conditions 1n. the outer layers can only affect the flux radiated outside the star very slightly as long as a gray atmosphere can be assumed.," Following this picture, the physical conditions in the outer layers can only affect the flux radiated outside the star very slightly as long as a gray atmosphere can be assumed."225 Thus. we proposed that the energy flux in the envelope of a rotating star be approximated by F=-asFU.gar.," Thus, we proposed that the energy flux in the envelope of a rotating star be approximated by $\vec F = -\frac{L}{4\pi GM}F_\omega(r,\theta)\vec g_{\rm eff}$."226 We have shown how the non-dimensional function FG.0) can be evaluated by assuming that mass distribution is represented by the Roche model.," We have shown how the non-dimensional function $F_\omega(r,\theta)$ can be evaluated by assuming that mass distribution is represented by the Roche model."227 We then demonstrated that the latitudinal variation of the effective temperature only depends on a single parameter «o=OR;/GM.," We then demonstrated that the latitudinal variation of the effective temperature only depends on a single parameter $\omega=228\sqrt{\Omega^2R_e^3/GM}$."229 Such a model is very appropriate to interpreting the interferometric observations of rotating stars since. unlike von Zeipel's or Lucy’s laws. it is valid for high rotation rates (up to breakup) and depends only on a single parameter (the B-exponent is removed).," Such a model is very appropriate to interpreting the interferometric observations of rotating stars since, unlike von Zeipel's or Lucy's laws, it is valid for high rotation rates (up to breakup) and depends only on a single parameter (the $\beta$ -exponent is removed)."230 Adjustment of the observed surface flux would thus only require variations in w and i. the inclination of the rotation axis on the line of sight.," Adjustment of the observed surface flux would thus only require variations in $\omega$ and $i$, the inclination of the rotation axis on the line of sight."231 As previously mentioned. this model fits the fully 2D models well using a gray atmosphere with a rigidly rotating surface (but with interior differential rotation).," As previously mentioned, this model fits the fully 2D models well using a gray atmosphere with a rigidly rotating surface (but with interior differential rotation)."232 This dynamical feature of the models may not be very realistic. so we tested a surface rotation Q(4)=OI-0.1cos?4) inspired by observations (?)..," This dynamical feature of the models may not be very realistic, so we tested a surface rotation $\Omega(\theta)=\Omega_{\rm eq}(1-0.1\cos^2\theta)$ inspired by observations \citep{CC07}."233 The difference is hardly perceptible. therefore the proposed model of gravity darkening looksquite robust.," The difference is hardly perceptible, therefore the proposed model of gravity darkening looksquite robust."234 Future improvement of two-dimensional models will of course be used to confirm this robustness., Future improvement of two-dimensional models will of course be used to confirm this robustness.235Nova Centauri 1986 (later designated VWsd2 Cen) was discovered. on 22 November 1986 at V—5.6 and two cays later reached maximum at V=4.6 (AleNaught 1986).,Nova Centauri 1986 (later designated V842 Cen) was discovered on 22 November 1986 at $V = 5.6$ and two days later reached maximum at $V = 4.6$ (McNaught 1986).236 lt was a moderately fast nova. with decav time ο=48 d and developed an obscuring dust. shell starting 37 d after maximum and reaching greatest optical thickness 74 d after maximum. placing it in the group LL category of nova light. curves as defined by Duerbeck (1981). which is essentially of DQ Herculis tvpe.," It was a moderately fast nova, with decay time $t_3 = 48$ d and developed an obscuring dust shell starting 37 d after maximum and reaching greatest optical thickness 74 d after maximum, placing it in the group II category of nova light curves as defined by Duerbeck (1981), which is essentially of DQ Herculis type."237 The pre-eruption brightness was estimated by AleNaught (1986) to be in the range D-18.018.6., The pre-eruption brightness was estimated by McNaught (1986) to be in the range $B \sim 18.0 - 18.6$.238 Fifteen vears alter eruption it was at le15.5 (Downes Duerbeck 2000: Woudt Warner 2003). and our latest measurements give V.—16.3.," Fifteen years after eruption it was at $V \sim 15.8$ (Downes Duerbeck 2000; Woudt Warner 2003), and our latest measurements give $V \sim 16.3$."239 Therefore 22 vears after maximum it is still about two magnitudes above its pre-nova brightness. which could be a result. of irradiation enhanced mass transfer. caused. by a still very hot white dwarf primary. especially if it is relatively massive (see Warner (2002) for a discussion of anomalous post-nova luminosities arising from this elfect). though Ixato (2008) reports that there is nothing in the eruptive behaviour to suggest à mass much greater than 0.7 M...," Therefore 22 years after maximum it is still about two magnitudes above its pre-nova brightness, which could be a result of irradiation enhanced mass transfer caused by a still very hot white dwarf primary, especially if it is relatively massive (see Warner (2002) for a discussion of anomalous post-nova luminosities arising from this effect), though Kato (2008) reports that there is nothing in the eruptive behaviour to suggest a mass much greater than 0.7 $_{\odot}$."240 DO Ler. also à moderately fast nova (£a=94 d). but of low mass (0.60 AL.: Lorne. Welsh Wade 1993). reached its maximum optical thickness dust. obseurecl phase 101 d after maximum light: i was at mug~148 prior to its 1934 eruption (Robinson 1975) and vet had returned to that level less than 20 vears later (Walker 1956).," DQ Her, also a moderately fast nova $t_3 = 94$ d), but of low mass (0.60 $_{\odot}$: Horne, Welsh Wade 1993), reached its maximum optical thickness dust obscured phase 101 d after maximum light; it was at $m_{pg} \sim 14.8$ prior to its 1934 eruption (Robinson 1975) and yet had returned to that level less than 20 years later (Walker 1956)."241 Sekiguchi ct al. (, Sekiguchi et al. (2421989) give. distance estimates tha average 1.0 kpe from strengths of interstellar Na D lines and the 2200 feature. ancl a reddening of E(B-V) = 0.55.,"1989) give distance estimates that average 1.0 kpc from strengths of interstellar Na D lines and the 2200 feature, and a reddening of E(B-V) = 0.55."243 Gill O'Brien (1998) in 1995 detected an ejecta shell of diameter 1.5 arcsec in direct imaging., Gill O'Brien (1998) in 1995 detected an ejecta shell of diameter $\sim 1.5$ arcsec in direct imaging.244 There is nothing in the nova development of VS42 Cen that marks it as in anv way peculiar., There is nothing in the nova development of V842 Cen that marks it as in any way peculiar.245 Phe UV. and. sof X-lav turn-oll times are normal (CGonzallez-Riestra. Orio Gallagher 1998) and the abundances (including a high carbon content typical of dusty ejecta) are within the norma ranges CXndrea. Drechsel Starrlield 1994). except. tha Iben (1992) found that WS42 Cen was the only nova with a Lefll ratio falling between two groups having solar anc hall-solar values.," The UV and soft X-Ray turn-off times are normal (Gonzállez-Riestra, Orio Gallagher 1998) and the abundances (including a high carbon content typical of dusty ejecta) are within the normal ranges (Andrea, Drechsel Starrfield 1994), except that Iben (1992) found that V842 Cen was the only nova with a He/H ratio falling between two groups having solar and half-solar values."246 Recent spectra. however. show peculiarities that are relevant to the present state of the primary in. Wsd2 Cen.," Recent spectra, however, show peculiarities that are relevant to the present state of the primary in V842 Cen."247 An optical spectrum. obtained in 2008 (Sehmidtohreick et al., An optical spectrum obtained in 2003 (Schmidtobreick et al.248 2005) shows a strong blue continuum with weak high ionization lines. e.g. UV. probably coming from the nova," 2005) shows a strong blue continuum with weak high ionization lines, e.g. IV, probably coming from the nova"249several eround-hased near-intrarecl imagers thal have recently come online provide sullicient skv coverage and sensilivily (hat significant. portions of nearby clusters ean be mapped efficientlv.,Several ground-based near-infrared imagers that have recently come online provide sufficient sky coverage and sensitivity that significant portions of nearby clusters can be mapped efficiently.250 For example. Flamingos on Gemini South has a FOV = 2.6'x2.6' and reaches my=20.3 in 300 see (5o point source detection) (Elston1998).," For example, Flamingos on Gemini South has a FOV = $2.6'{\times}2.6'$ and reaches $m_J=20.3$ in 300 sec $\sigma$ point source detection) \citep{re}."251. WIRC-21Ix InfraRed Camera) on Palomar has a FOV = 9’x9' and reaches m.=19.9 in 300 see (5o point source detection) (Wilsonοἱal.2002)., WIRC-2K (Wide-field InfraRed Camera) on Palomar has a FOV = $9'{\times}9'$ and reaches $m_J = 19.9$ in 300 sec $\sigma$ point source detection) \citep{wils}.252. With (hese and other current ancl forthcoming instruments it will be possible to probe 3 to 5 magnitudes deeper (han past survevs. detect peak D. and better determine the behavior of the very. low-1ass/substellar mass function.," With these and other current and forthcoming instruments it will be possible to probe 3 to 5 magnitudes deeper than past surveys, detect peak D, and better determine the behavior of the very low-mass/substellar mass function."253 Although mid-infrared observations vield the best possibilities for detection of peak D and troughs B and E (Figure 10). ground-based. mid-inlrared observations are difIicult because atmospheric water bands limit sensitivity.," Although mid-infrared observations yield the best possibilities for detection of peak D and troughs B and E (Figure 10), ground-based mid-infrared observations are difficult because atmospheric water bands limit sensitivity."254 The forthcoming Space Infrared Telescope Facility (SIRTF) mission will probe to sullicient sensitivilies to detect these intermediate predicted cluster laminosity Function features., The forthcoming Space Infrared Telescope Facility (SIRTF) mission will probe to sufficient sensitivities to detect these intermediate predicted cluster luminosity function features.255 Figure 10 shows our M-band (4.8jm) models with the SIRTF Infralted Array Camera (RAC) Channel 2 (4.5jm) sensitivity limits 13) for a cluster at 200 pc., Figure 10 shows our $M$ -band $4.8~\mu$ m) models with the SIRTF InfraRed Array Camera (IRAC) Channel 2 $4.5~\mu$ m) sensitivity limits $M_M = 13$ ) for a cluster at 200 pc.256" The predicted sensitivity limit corresponds to a 5a detection of ama,=19 point source in 200 seconds (oraetal.2002).", The predicted sensitivity limit corresponds to a $\sigma$ detection of a $m_M=19$ point source in 200 seconds \citep{hora}.257. For this sensitivity ancl distance both troughs D and E will be easily detected by IRAC/SIRTF. as will be peaks D and C for clusters vounger than ~200 Mvr.," For this sensitivity and distance both troughs B and E will be easily detected by IRAC/SIRTF, as will be peaks D and C for clusters younger than $\sim$ 200 Myr."258 We have combined theoretical evolutionary. (racks for low-mass dwarls with empirical bolometric correction measurements to identily several characteristic features in the predicted hunimosity [unetions al £N AM-bands., We have combined theoretical evolutionary tracks for low-mass dwarfs with empirical bolometric correction measurements to identify several characteristic features in the predicted luminosity functions at $IJHKM$ -bands.259 The most significant. feature that we find is an intermediate peak in clusters (hat are tens to hundreds of millions of vears old (denoted as peak D) corresponding to high-mass brown dwarls (hat have ceased to burn deuterium., The most significant feature that we find is an intermediate peak in clusters that are tens to hundreds of millions of years old (denoted as peak D) corresponding to high-mass brown dwarfs that have ceased to burn deuterium.260 This feature appears to be present in observations of the voung cluster IC 2391. and the peaks location in absolute magnitude suggests an age of 35 Myr.," This feature appears to be present in observations of the young cluster IC 2391, and the peak's location in absolute magnitude suggests an age of 35 Myr."261 The location of this feature may prove a useful chronometer for voung clusters., The location of this feature may prove a useful chronometer for young clusters.262 Strong spectral features. such as the onset of methane absorption at the L/T transition. produce significant effects in the A-band luminosity function.," Strong spectral features, such as the onset of methane absorption at the L/T transition, produce significant effects in the $K$ -band luminosity function."263 Both these features ancl peak D should be detectable in nearby open clusters with the current generation of wide-field inlhared imagers and SIRTE. allowing comparison of the model predictions with clusters spanning a much wider range in age.," Both these features and peak D should be detectable in nearby open clusters with the current generation of wide-field infrared imagers and SIRTF, allowing comparison of the model predictions with clusters spanning a much wider range in age."264 PRA. acknowledges support by a grant made under the auspices of the NASA/NSF, P.R.A. acknowledges support by a grant made under the auspices of the NASA/NSF265The morphological types of the two nearby companions of the Cartwheel are cistinetly cüfferent.,The morphological types of the two nearby companions of the Cartwheel are distinctly different.266 The eastern colupanion (G2) is an early type SO galaxy., The eastern companion (G2) is an early type S0 galaxy.267 No is directly associated. wih this ealaxy (Ihedou 1996)) aud no star formation activity las been detectec (Amman ct al. 1998))., No is directly associated with this galaxy (Higdon \cite{jim_hi}) ) and no star formation activity has been detected (Amram et al. \cite{amram}) ).268 On the coutrary. the western companion (C1). a late type galaxy wihn somewhat imregular spiral aruis. has 2.7xl0? of and cmits strongly in Ho. which indicates the presence of a voung stellar population (see Fis.l of Amman et al. 1998)).," On the contrary, the western companion (G1), a late type galaxy with somewhat irregular spiral arms, has $^{9}$ $_{\odot}$ of and emits strongly in $\alpha$, which indicates the presence of a young stellar population (see Fig.1 of Amram et al. \cite{amram}) )."269 The IW2 fluxes of the two ealaxies. presented in Table 2.. are weak but very sluuilar.," The LW2 fluxes of the two galaxies, presented in Table \ref{phot}, are weak but very similar."270 This is not true thoughC» for the LAWS emission where oulv the western companion (CL). is detected.," This is not true though for the LW3 emission where only the western companion (G1), is detected."271 This coutrast on the LW? over LW2 ciission frou the two companion galaxies is in agreement with our current muaderstauding of the typical wid-IR signature iu ealaxies (Vieroux 19973)., This contrast on the LW2 over LW3 emission from the two companion galaxies is in agreement with our current understanding of the typical mid-IR signature in galaxies (Vigroux \cite{vigroux}) ).272 Tn carly type galaxies the iicl-IB enüssion is dominated by the old stellar populalon., In early type galaxies the mid-IR emission is dominated by the old stellar population.273 The LW2 fux is typically 6 times stronger than the LW2 enüssion which is consistent with temperaures of ~HOOO I& (Madden 1997))., The LW2 flux is typically 6 times stronger than the LW3 emission which is consistent with temperatures of $\sim$ 5000 K (Madden \cite{madden}) ).274 Since the LW? cussion o 2 is just ~ 1 ties strouger iui the rms noise of the ΤΑΝ2 enüssion any expected stellar contribution to the ΤΑV3 filter measurement would ο below our detection linüt., Since the LW2 emission of G2 is just $\sim$ 4 times stronger than the rms noise of the LW3 emission any expected stellar contribution to the LW3 filter measurement would be below our detection limit.275 The ratio of LW3 to LW? flux of Glas ~1.ἐν a value found in most normal late type galaxies.," The ratio of LW3 to LW2 flux of G1 is $\sim$ 1.4, a value found in most normal late type galaxies."276 Ow ISOCAN observations reveal that one-third. of the circunfereuce of Cartwheels outer ring is detected in the TW? baud (see Fig., Our ISOCAM observations reveal that one-third of the circumference of Cartwheel's outer ring is detected in the LW2 band (see Fig.277 laa)., \ref{lw23}a a).278 The detected segment of the rue corresponds to the brightest part of the optical rine as clefined by the Πα mise of Που (1995))., The detected segment of the ring corresponds to the brightest part of the optical ring as defined by the $\alpha$ image of Higdon \cite{jim_ha}) ).279 More than of the total Πα flux frou the Cartwheel outer ring originates from its south-castern scement., More than of the total $\alpha$ flux from the Cartwheel outer ring originates from its south-eastern segment.280 About of the LW2 flux from the outer rine comes from a hot-spot in the ring. unresolved to our observations. which corresponds to the position of two lage region couplexes (Tigdon 1995)).," About of the LW2 flux from the outer ring comes from a hot-spot in the ring, unresolved to our observations, which corresponds to the position of two large region complexes (Higdon \cite{jim_ha}) )."281 Tuterestinely. in the LW3 filter. ouly the hot-spot is detected in the outer ring (Fie.," Interestingly, in the LW3 filter, only the hot-spot is detected in the outer ring (Fig."282 1bb). aud the only other chussion from the Cartwheel is from the ceutral regions (see Table 2) and discussion below).," \ref{lw23}b b), and the only other emission from the Cartwheel is from the central regions (see Table \ref{phot} and discussion below)."283 The LAW3/TW2 flux ratio. often used as a diagnostic of the iueusitv of the radiation field. of the hot-spot is 5.2. a value which is among the highest detected in all interacting galaxies of the CAMACTIV sample.," The LW3/LW2 flux ratio, often used as a diagnostic of the intensity of the radiation field, of the hot-spot is 5.2, a value which is among the highest detected in all interacting galaxies of the CAMACTIV sample."284 For comparison. the highest TW3/TAWW2 flux ratio detected in the interacting ealaxy L1038/39 1s only 2.6 aud it is found in the region where the two disks overlap (Vieroux ct al. 1996:," For comparison, the highest LW3/LW2 flux ratio detected in the interacting galaxy 4038/39 is only 2.6 and it is found in the region where the two disks overlap (Vigroux et al. \cite{chef};"285 Mirabel et al. 1998))., Mirabel et al. \cite{felix}) ).286 The lack of detection in the LAV filter from regions of the ring other hau the hot-spot is iutziguing., The lack of detection in the LW3 filter from regions of the ring other than the hot-spot is intriguing.287 What is clear is that if the rest of the ring had the same ratio of TAW3/TAWW2 Gc. 5.2) as the hot-spot. then it would have been casily deteced in the LW3 band.," What is clear is that if the rest of the ring had the same ratio of LW3/LW2 (i.e. 5.2) as the hot-spot, then it would have been easily detected in the LW3 band."288 Therefore. using," Therefore, using"289to move to the region where there are no stable orbits around both stars.,to move to the region where there are no stable orbits around both stars.290 Hf a mass particle passes a certain point. then the only possibility for it is (o move (to one of the stars.," If a mass particle passes a certain point, then the only possibility for it is to move to one of the stars."291 In this wav. streams of material originating in the CD disk connect to the circumstellar disks (hese are disks inside the hole surrounding each star). acting as a feeding mechanism (CArtvmowiczGünther&Ixlev 2002).," In this way, streams of material originating in the CB disk connect to the circumstellar disks (these are disks inside the hole surrounding each star), acting as a feeding mechanism \citep{Artymowicz2,Gunther1}."292. Observations by Jensenetal.(2007) suggest the presence of streams from the CB disk towards the binary svstem UZ Tau E. Also. for the binary SC 34. Πατetal.(2005) argue that the presence of dust in the dynamically leared hole is consistent with having some accretion flow onto the central binary. (2004)..," Observations by \citet{Jensen} suggest the presence of streams from the CB disk towards the binary system UZ Tau E. Also, for the binary St 34, \citet{Hartmann} argue that the presence of dust in the dynamically cleared hole is consistent with having some accretion flow onto the central binary. \citet{Gunther2},"293 described the emission of DQ Tau and Alx Sco using a hydrodynamical simulation. where all the disks and the streams connecting them are taken into account.," described the emission of DQ Tau and AK Sco using a hydrodynamical simulation, where all the disks and the streams connecting them are taken into account."294 Also. binaries of very low mass are expected to be fed by a circumbinary disk (Robbertoοἱal.," Also, binaries of very low mass are expected to be fed by a circumbinary disk \citep{Robberto}."2952003).. From all these studies. we can argue that (he Gauncation radius £24 strongly depends on the viscosity in a non trivial wav.," From all these studies, we can argue that the truncation radius $R_{cb}$ strongly depends on the viscosity in a non trivial way."296 For an e>0 system. Artvmowiez&Lubow(1994). show clear differences between {το lor changes in the Revnolds number. which depends on the viscosity.," For an $e>0$ system, \citet{Artymowicz1} show clear differences between $R_{cb}$ for changes in the Reynolds number, which depends on the viscosity."297 llowever. in order to take a more realistic model. we decided to use the results of taken lrom Fieure 3. arguing that {ήν does not depend strongly on the stars mass ratio. and a Reynolds number of 10/.," However, in order to take a more realistic model, we decided to use the results of \citet{Artymowicz1} taken from Figure 3, arguing that $R_{cb}$ does not depend strongly on the stars mass ratio, and a Reynolds number of $10^{4}$."298 An important ingredient lor the calewlation of the radiation flix [rom the stars impinging on the circumbinary. wall is (he distance between each star and a specilic point on (he wall. liq» (see Figures 1606 and 17)).," An important ingredient for the calculation of the radiation flux from the stars impinging on the circumbinary wall is the distance between each star and a specific point on the wall, $R_{1,2}$ (see Figures \ref{fig-configuracion} and \ref{fig-configuracion-i}) )."299" In order to caleulate these distances. (he first step is lo give the locations of the stars in the configuration where one star is al the focus and (he other is on a point on ils corresponding, ellipse."," In order to calculate these distances, the first step is to give the locations of the stars in the configuration where one star is at the focus and the other is on a point on its corresponding ellipse."300 The second step is to place a circular CB disk with ils origin al the point corresponding to the eccentricity of the binary svstem 2008).. e.g. the center of mass of the binary svstem if e=0. with a radius given by Arivinowiez&Lubow(1994).," The second step is to place a circular CB disk with its origin at the point corresponding to the eccentricity of the binary system \citep[see][]{Pichardo1,Pichardo2}, e.g., the center of mass of the binary system if $e=0$, with a radius given by \citet{Artymowicz1}."301. Finally. the third step is to caleulate the distances between the coordinates of the stars and each surface element of the wall given in this reference svstem (see Figures 16 and 17)).," Finally, the third step is to calculate the distances between the coordinates of the stars and each surface element of the wall given in this reference system (see Figures \ref{fig-configuracion} and \ref{fig-configuracion-i}) )."302 The inner wall of the CD clisk is assumed to be vertical. ie.. (he vector normal to its surface is parallel to the disk micplane.," The inner wall of the CB disk is assumed to be vertical, i.e., the vector normal to its surface is parallel to the disk midplane."303 In order to relax the assumption of verticallitv one requires a wav to describe the shape of the wall. which is not possible to the best of our knowledge. because the observations and simulations do not have the necessary resolution.," In order to relax the assumption of verticallity one requires a way to describe the shape of the wall, which is not possible to the best of our knowledge, because the observations and simulations do not have the necessary resolution."304 Alore complicated shapes have been assumed for the dust to gas transition at (he dust destruction radius in the inner disk by (aking into account the dependence of the sublimation, More complicated shapes have been assumed for the dust to gas transition at the dust destruction radius in the inner disk by taking into account the dependence of the sublimation305enerev as the gas jumps above the photosphere would produce an effective local κιατω ol 2 to 3. compared with the true radiative ωμήτω of more like 15 at the same places. a reduction in local cooling time bv a [actor of between 5 and 8!,"energy as the gas jumps above the photosphere would produce an effective local $t_{\rm cool}/P_{\rm rot}$ of 2 to 3, compared with the true radiative $t_{\rm cool}/P_{\rm rot}$ of more like 15 at the same places, a reduction in local cooling time by a factor of between 5 and 8!"306" While /,,/7,4 is not vet small enough to induce fragmentation. the BOT simulations do evolve denser structure and stronger shocks than our simulation bevond this time. ancl it is easv to imagine that the Ganunie fragmentation criterion could then be satisfied."," While $t_{\rm cool}/P_{\rm rot}$ is not yet small enough to induce fragmentation, the B07 simulations do evolve denser structure and stronger shocks than our simulation beyond this time, and it is easy to imagine that the Gammie fragmentation criterion could then be satisfied."307 This estimate for the excess cooling introduced by the BOT temperature reset in the optically thin regions is highlv suggestive that [ast cooling in (he DOT scheme is artificial., This estimate for the excess cooling introduced by the B07 temperature reset in the optically thin regions is highly suggestive that fast cooling in the B07 scheme is artificial.308 Unfortunately. when we tried to implement the BOF BCs in our own code. we encountered numerical instability in the radiative heating and cooling terms near the photosphere. and we could not test this hypothesis anv further.," Unfortunately, when we tried to implement the B07 BCs in our own code, we encountered numerical instability in the radiative heating and cooling terms near the photosphere, and we could not test this hypothesis any further."309 50 what then is the status of disk instability as a gas giant. formation mechanism?, So what then is the status of disk instability as a gas giant formation mechanism?310 Although Maveretal.(2007) agree with Boss that disk instability may work under some conditions inside 40 AU. simulations by other groups (Stamatellos&Whitworth2008:Forganetal.2009) and analvtic treatments (Ralikoy2005.2007) support our arguments against Iragmentation in the inner disks of voung solar-tvpe stars.," Although \cite{mayer07} agree with Boss that disk instability may work under some conditions inside 40 AU, simulations by other groups \citep{stam08,forg09} and analytic treatments \citep{rafikov05,rafikov07} support our arguments against fragmentation in the inner disks of young solar-type stars."311 On (he other hand. recent. hycrodvnamics simulations (Stamatellos&Whitworth2009:DBolevx.Boleyetal.2009:Vorobvov&Basu2010:Havfield2010) and other arguments (Clarke2009:RafikovDocson-Robinsonetal.2009) now suggest that disk instability may very well operate in the outer regions of large (> LOO AU) massive disks during the early accretion phase to produce super-Jupilers. brown dwarls. and low-mass stars.," On the other hand, recent hydrodynamics simulations \citep{stam09,boley09,boleyetal09,vb10,hayfield10} and other arguments \citep{clarke09,rafikov09,dodson09} now suggest that disk instability may very well operate in the outer regions of large $>$ 100 AU) massive disks during the early accretion phase to produce super-Jupiters, brown dwarfs, and low-mass stars."312 We think these lragimentation results in outer massive disks are equite plausible and wortliv of further studs;, We think these fragmentation results in outer massive disks are quite plausible and worthy of further study.313 The disk instability mechanism (hat Boss has championed since 1997 may vel prove to be an important formation channel for eas giants. but in a different place than he originally proposed.," The disk instability mechanism that Boss has championed since 1997 may yet prove to be an important formation channel for gas giants, but in a different place than he originally proposed."314pressure. can be computed from the expression (see e.g. AlacLareu et al.,"pressure, can be computed from the expression (see e.g. MacLaren et al."315 1055) where Ae; is the observed line FWOAL, 1988) where $\Delta v_{1/2}$ is the observed line FWHM.316" In our case we took Ary from the CoII((65) lues. for the sake of cousistency with the derivation of O9, which las also been obtaiucd youn the CIT4CSII((65) emission."," In our case we took $\Delta v_{1/2}$ from the (6–5) lines, for the sake of consistency with the derivation of $\Theta_{\rm s}$ which has also been obtained from the (6–5) emission."317 The values of Aa are given in Table 9.., The values of $M_{\rm vir}$ are given in Table \ref{mvir-mcd}.318 Tn the same table we also eive the IT volume densities estimated. from Meg aud the chuup diameters. D.," In the same table we also give the $_2$ volume densities estimated from $M_{\rm CD}$ and the clump diameters, $D$."319 The uucertainties ou these values have been calculated by caus of the propagation of statistical errors. but they do not take iuto account distance uncertainties.," The uncertainties on these values have been calculated by means of the propagation of statistical errors, but they do not take into account distance uncertainties."320 All masses are equal to a few ~LOPAL... with the sole exception of 220126. for which themass is X107AL...," All masses are equal to a few $\sim 10^{3} M_{\odot}$, with the sole exception of 20126, for which themass is $\la10^2~M_\odot$."321 The IL volue deusities rauge from 107 to 5109 5m ereater than the πιααι critical deusitv of the transitions observed (see the discussion iu Sect. 3.13).," The $\rm H_{2}$ volume densities range from $10^{5}$ to $5~10^{6}$ $^{-3}$, greater than the maximum critical density of the transitions observed (see the discussion in Sect. \ref{sthtr}) )."322 Iu view of the Ποιος that will be illustrated in Sect. L..," In view of the findings that will be illustrated in Sect. \ref{sdisc},"323 before proceeding further it is important to diseuss the relevance of optical depth effects. temperature eradicuts. and density eracdicuts on the mass (aud hence deusitv) estimates obtained so far.," before proceeding further it is important to discuss the relevance of optical depth effects, temperature gradients, and density gradients on the mass (and hence density) estimates obtained so far."324 As explained in Sect. 3.2..," As explained in Sect. \ref{stempden},"325 we believe that the bulk of the observed. eenission is optically thin., we believe that the bulk of the observed emission is optically thin.326 However. one cannot exclude the existence of an optically thick region close to the ceutre of the clump. where the density is presumably higher thui average.," However, one cannot exclude the existence of an optically thick region close to the centre of the clump, where the density is presumably higher than average."327 Iudeed. deuse molecular cores with diameters of U.l pe are known to exist in most of our clumps.," Indeed, dense molecular cores with diameters of $\sim$ 0.1 pc are known to exist in most of our clumps."328 Hence the question is how auch of the clump mass is contained in the optically thick region., Hence the question is how much of the clump mass is contained in the optically thick region.329 Oue can calculate the size of the region over which 7:L for the lines., One can calculate the size of the region over which $\tau\ge1$ for the lines.330" For N44,=21075 7. T OK. Ae»=5 . aud a deusitv profile 7κAR7?D one Buds that the optica depth of e.g. the CIT4CSII((65) A —0 trausition along the line of sight passing through the ceutre of the chumpIs equal OTsOUR1). where Ry is the radius of the chuup aud δὲ the distance from the ceutre at which the optical depth 7 is achieved."," For $N_{\rm tot}=2~10^{15}$ $^{-2}$, $T=40$ K, $\Delta v_{1/2}=5$ , and a density profile $n\propto R^{-2}$ one finds that the optical depth of e.g. the (6–5) $K$ =0 transition along the line of sight passing through the centre of the clumpis equal to $\tau\simeq 0.1 (R_0/R-1)$, where $R_0$ is the radius of the clump and $R$ the distance from the centre at which the optical depth $\tau$ is achieved."331 By posing 7=1 one fiuds the radius of the optically thick core. R/Ryzz10.," By posing $\tau=1$ one finds the radius of the optically thick core, $R/R_0\simeq10$."332 Iu conclusion. the geas becomes optically thick at a distance from the centre equal to ~10% of the clump radius: this coutains 10 of the total mass.," In conclusion, the gas becomes optically thick at a distance from the centre equal to $\sim$ of the clump radius: this contains $\sim$ of the total mass."333 We conclude that optical depth effects are negligible for the estimate of Ac., We conclude that optical depth effects are negligible for the estimate of $M_{\rm CD}$.334 lucideutallv. it is worth notiug that in case of large optical depths the values of Neo aud ου are to be taken as lower limits.," Incidentally, it is worth noting that in case of large optical depths the values of $N_{\rm tot}$ and $M_{\rm CD}$ are to be taken as lower limits."335 Ou the coutrarv. Mg becomes au upper huit. because for rὃν1 line broadening occurs. which increases the value of ου used in Eq. (6)).," On the contrary, $M_{\rm vir}$ becomes an upper limit, because for $\tau\gg1$ line broadening occurs, which increases the value of $\Delta v_{1/2}$ used in Eq. \ref{emvir}) )."336 As a consequence. the ratio MopMy which one cau derive from Table 9 has to be regarded as a lower limit.," As a consequence, the ratio $M_{\rm CD}/M_{\rm vir}$ which one can derive from Table \ref{mvir-mcd} has to be regarded as a lower limit."337 This fact reinforces the conclusion attained later i Sect. L1.., This fact reinforces the conclusion attained later in Sect. \ref{sstab}.338 The hypotheses under which he method of the populatiou diagrams in Sect., The hypotheses under which the method of the population diagrams in Sect.339 3.2. is applied are that the gas is optically thin. in LTE. anc isothermal.," \ref{stempden} is applied are that the gas is optically thin, in LTE, and isothermal."340 The frst two assuniptious have already been discussed: as for the atter. Mauersberger et al. (," The first two assumptions have already been discussed; as for the latter, Mauersberger et al. ("3411988) have treated he case of a spherically saunietrie clo with temperature aud density laving power-law dependence on the distance roni the centre.,1988) have treated the case of a spherically symmetric cloud with temperature and density having power-law dependence on the distance from the centre.342 These authors demonstrate that in this case a linear correlation should hold between lnAN; aud uZZ; unlike the isothermal case where the correlation is expected with £;.," These authors demonstrate that in this case a linear correlation should hold between $\ln N_i$ and $\ln E_i$, unlike the isothermal case where the correlation is expected with $E_i$."343 We stress tha he differeut dependence on £; is causedondy by the temperature eradicut: in fact. or an isothermal cloud the usual lear relation between uN; and £; Is recovered. no lmatter what the density eradicut is.," We stress that the different dependence on $E_i$ is caused by the temperature gradient: in fact, for an isothermal cloud the usual linear relation between $\ln N_i$ and $E_i$ is recovered, no matter what the density gradient is."344 Following the example of Mauersberger ct al. (, Following the example of Mauersberger et al. (3451988). in Fie.,"1988), in Fig."346 5 we present “mocified” population diagranus for our sources. plotting InN; versus ln£;.," \ref{fmodbol} we present “modified” population diagrams for our sources, plotting $\ln N_i$ versus $\ln E_i$."347 Clearly. the linear fits to the points are mich less satisfactory than those in Fie. L..," Clearly, the linear fits to the points are much less satisfactory than those in Fig. \ref{grafico1}."348 We conclude that the eenission arises from a region m which the temperature profile 1uust be very shallow. hence making T=coust.," We conclude that the emission arises from a region in which the temperature profile must be very shallow, hence making $T$ =const."349 a suitable approximation., a suitable approximation.350" For an isothermal and optically thin cloud. the source averaged column deusitv cannot depend on the wav the eas is distributed inside the cloud: heuce. density eradieuts leave he estimates of N,,, aud Mog unaffected."," For an isothermal and optically thin cloud, the source averaged column density cannot depend on the way the gas is distributed inside the cloud: hence, density gradients leave the estimates of $N_{\rm tot}$ and $M_{\rm CD}$ unaffected."351 Ou the other haud. the virial mass depends ou the density profile. as demoustrated e.c. by MacLaren et al. (," On the other hand, the virial mass depends on the density profile, as demonstrated e.g. by MacLaren et al. ("3521988).,1988).353" For a power-Luv deusitv distribution of the type vy,x R. the virial mass obtained from Eq. (6))"," For a power-law density distribution of the type $n_{\rm H_2}\propto R^p$ , the virial mass obtained from Eq. \ref{emvir}) )"354" ust be ultiplied by the factor 2222 which is <1 for pox0,"," must be multiplied by the factor $\frac{3}{5}\frac{5+2p}{3+p}$, which is $\le$ 1 for $p\le0$."355 Therefore. the values of Aa of eapTable 9 derived: for. a homogeneous cloud are to be taken as upper lite.," Therefore, the values of $M_{\rm vir}$ of Table \ref{mvir-mcd} derived for a homogeneous cloud are to be taken as upper limits."356" lu couclusion. the ratio ALey/ALG, derived by us represents απ underestimate if the constant density asstuuption is released in the computation of AM."," In conclusion, the ratio $M_{\rm CD}/M_{\rm vir}$ derived by us represents an underestimate if the constant density assumption is released in the computation of $M_{\rm vir}$ ."357 This result is important iu view of the discussion in Sect. [.1.., This result is important in view of the discussion in Sect. \ref{sstab}. .358 Iu the previous section. we have demonstrated that the values of the plivsical paraiucters derived by us are robust.," In the previous section, we have demonstrated that the values of the physical parameters derived by us are robust."359 Iu particular. we have shown that: (1) optical deptheffects can be ueelected for the Ines: Gi) temperature eradieuts. if present. must be very," In particular, we have shown that: (i) optical deptheffects can be neglected for the lines; (ii) temperature gradients, if present, must be very"360and hence more of a problem for WD-MS systems.,and hence more of a problem for WD–MS systems.361 This ts because the majority of detectable systems (Fig. 1)), This is because the majority of detectable systems (Fig. \ref{fig:period}) )362 have short orbital periods. P©30d. and at these short periods. tidal effects due to the WD are significant (note that the code follows tidal effects in detail).," have short orbital periods, $P \la 30$ d, and at these short periods, tidal effects due to the WD are significant (note that the code follows tidal effects in detail)."363 Synchronization of the MS star's rotation with the orbital period 1s rapid if the MS star has a convective envelope (Mus= 1.6M..) and P=10d. This is the case for about half of the transits in principle detectable withKepler (see Fig. 2)).," Synchronization of the MS star's rotation with the orbital period is rapid if the MS star has a convective envelope $M_{\rm{MS}} \la 1.6364\rm{M}_\odot$ ) and $P \la 10$ d. This is the case for about half of the transits in principle detectable with (see Fig. \ref{fig:mass}) ),"365 and most forEddington., and most for.366 Rapidly rotating type stars display increased starspot activity (Messina.Rodond&Guinan 2001).. and hence greater photometric variability. due to these starspots rotating into and out of sight.," Rapidly rotating late-type stars display increased starspot activity \citep{mes01}, and hence greater photometric variability, due to these starspots rotating into and out of sight."367 Individual spots persist for only a few rotational cycles. meaning that over time the variability has a random nature.," Individual spots persist for only a few rotational cycles, meaning that over time the variability has a random nature."368 The amplitude of this variability scales with rotational period approximately as 0κ P1?.," The amplitude of this variability scales with rotational period approximately as $\sigma \propto369P_{\rm{rot}}^{-1.5}$ ."370 Examination of the power spectrum of solar irradiance variations (Frohlichetal.1997;Jenkins2002) shows that this starspot noise is present up to frequencies 25/Pro...," Examination of the power spectrum of solar irradiance variations \citep{fro97,jen02} shows that this starspot noise is present up to frequencies $\sim 25/P_{\rm{rot,}\odot}$."371 Of most importance in transit detection is the stellar noise on the timescale of a transit. since one can in principle filter out variability on other timescales (e.g. Jenkins2002... Aigrainetal. 2002)).," Of most importance in transit detection is the stellar noise on the timescale of a transit, since one can in principle filter out variability on other timescales (e.g. \citealt{jen02}, \citealt{aig02}) )."372" For WD-MS systems with P~1— 10d. we have Ty~ |-3h. If à WD transit has P4~Pap<257,4. then the fractional starspot noise will be z1077. and will drown out almost all WD transit signals. even if there are many transits during the survey lifetime."," For WD–MS systems with $P373\sim 1-10$ d, we have $T_{\rm{tr}} \sim 1-3$ h. If a WD transit has $P_{\rm{rot}} \sim P_{\rm{orb}} < 25 T_{\rm{tr}}$, then the fractional starspot noise will be $\gg 10^{-4}$, and will drown out almost all WD transit signals, even if there are many transits during the survey lifetime."374 Effectively this places a lower limit on the orbital period of detectable systems (of around P~ 2d for typical systems)., Effectively this places a lower limit on the orbital period of detectable systems (of around $P \sim 2$ d for typical systems).375 At frequencies >25/P.. the noise is governed by convective (super-)granulation. has an amplitude ~$«I0? in the Sun on the appropriate transit timescale. and is likely unaffected by rotation rate.," At frequencies $\ga 25/P_{\rm{rot}}$, the noise is governed by convective (super-)granulation, has an amplitude $\sim 5 \times 10^{-5}$ in the Sun on the appropriate transit timescale, and is likely unaffected by rotation rate."376 This noise will however reduce detectability of faint transits., This noise will however reduce detectability of faint transits.377 We approximate all convective stars as solar in these respects. and in Table | show the sizable effeets of adding this variability noise upon the detection rates forKepler. again requiring 8o detections.," We approximate all convective stars as solar in these respects, and in Table 1 show the sizable effects of adding this variability noise upon the detection rates for, again requiring $8 \sigma$ detections."378 The mission. which will observe in two colors. may be able to use the color signature of the stellar variability to enhance detection probabilities (Bordéetal. 2003).. in which case our predictions without stellar variability may be more appropriate.," The mission, which will observe in two colors, may be able to use the color signature of the stellar variability to enhance detection probabilities \citep{bor03}, in which case our predictions without stellar variability may be more appropriate."379 More massive MS stars have radiative. envelopes. and smaller |Af| for transits. since their radii are larger.," More massive MS stars have radiative envelopes, and smaller $\vert \Delta f\vert$ for transits, since their radii are larger."380 For these stars. there is less detailed literature available on stellar microvariability. so we do not attempt to calculate its expected impact on the WD-MS detection rate.," For these stars, there is less detailed literature available on stellar microvariability, so we do not attempt to calculate its expected impact on the WD–MS detection rate."381 We do however note that these stars can be intrinsically quite variable., We do however note that these stars can be intrinsically quite variable.382 Since the radiative tide is weaker. systems are most likely asynchronous. but Zaqarashvili.Javakhishvili&Belvedere(2002) suggest that in this case. tides can excite the fundamental mode of pulsation of the star. potentially leading to oscillations on roughly the timescale of a transit.," Since the radiative tide is weaker, systems are most likely asynchronous, but \citet{zaq02} suggest that in this case, tides can excite the fundamental mode of pulsation of the star, potentially leading to oscillations on roughly the timescale of a transit."383 We note also that some observations (e.g. Dempseyetal. 1993)) suggest that stars in close binaries display greater activity than single stars of the same rotation rate., We note also that some observations (e.g. \citealt{dem93}) ) suggest that stars in close binaries display greater activity than single stars of the same rotation rate.384 Thus. the numbers of detectable systems given in Table | for MS primaries with radiative envelopes are likely to be reduced by variability. but we have not attempted to quantify this reduction.," Thus, the numbers of detectable systems given in Table 1 for MS primaries with radiative envelopes are likely to be reduced by variability, but we have not attempted to quantify this reduction."385 Extensive data on these topics may only be acquired once space-based transit searches fly., Extensive data on these topics may only be acquired once space-based transit searches fly.386 An additional source of microvariability in the lghtcurve may be from flickering as the WD accretes at a low level from the MS star's wind., An additional source of microvariability in the lightcurve may be from flickering as the WD accretes at a low level from the MS star's wind.387" To have an accretion luminosity Lwp 107L.... the WD needs to accrete at a rate ~10ΜΙΑ, cf."," To have an accretion luminosity $L_{\rm{WD}} \sim 10^{-4} \rm{L}_\odot$ , the WD needs to accrete at a rate $\sim 10^{-13} \rm{M}_\odot388\rm{yr}^{-1}$, cf."389" the solar mass loss rate ~107M,yr7!."," the solar mass loss rate $\sim 10^{-14} \rm{M}_\odot390\rm{yr}^{-1}$."391 Only a fraction of the mass lost from the MS star will be accreted by the WD. though we note that stellar wind mass loss may be enhanced in close binaries.," Only a fraction of the mass lost from the MS star will be accreted by the WD, though we note that stellar wind mass loss may be enhanced in close binaries."392 As can be seen from Figure ].. it is unlikely that WDs will provide a significant source of spurious earth-like transits at P~ years. given the dearth of WD-MS systems and (hypothesized) large numbers of terrestrial planets at these periods.," As can be seen from Figure \ref{fig:period}, it is unlikely that WDs will provide a significant source of spurious earth-like transits at $P \sim$ years, given the dearth of WD–MS systems and (hypothesized) large numbers of terrestrial planets at these periods."393 In addition. primary transits will be microlensing events at these periods.," In addition, primary transits will be microlensing events at these periods."394 Although WD companions are easily distinguished using radial-velocity observations. it is useful to know that terrestrial planet signatures will not be swamped by those of WDs.," Although WD companions are easily distinguished using radial-velocity observations, it is useful to know that terrestrial planet signatures will not be swamped by those of WDs."395 A large sample of close WD-MS systems would enable useful tests of binary star evolution theories (such as common envelope evolution)., A large sample of close WD–MS systems would enable useful tests of binary star evolution theories (such as common envelope evolution).396 Also. if the mass and radius of the WD can be separately determined. then the WD mass-radius relation could be tested.," Also, if the mass and radius of the WD can be separately determined, then the WD mass-radius relation could be tested."397 The transiting WDs cannot. in general. be observed other than by their dimming effect upon the MS star. so that all properties must be inferred.," The transiting WDs cannot, in general, be observed other than by their dimming effect upon the MS star, so that all properties must be inferred."398 This approach is however independent of WD atmosphere modeling., This approach is however independent of WD atmosphere modeling.399 TheKepler and missions easily have the sensitivity to produce high-quality lighteurves of short-period WD-MS systems. since they are designed to search for longer-period terrestrial planets.," The and missions easily have the sensitivity to produce high-quality lightcurves of short-period WD–MS systems, since they are designed to search for longer-period terrestrial planets."400 The inclusion of stellar variability may affect this somewhat., The inclusion of stellar variability may affect this somewhat.401 However. there tis still a population of 250 WD-MS systems in principle detectable with each. given adequate signal processing power.," However, there is still a population of $\ga 50$ WD–MS systems in principle detectable with each, given adequate signal processing power."402 Many of these systems display both detectable primary and secondary transits. which can be distinguished using their transit profiles.," Many of these systems display both detectable primary and secondary transits, which can be distinguished using their transit profiles."403 WD transits are longer in duration than. and shaped differently from. grazing MS-MS transits of the same depth.," WD transits are longer in duration than, and shaped differently from, grazing MS–MS transits of the same depth."404 Radial velocity measurements should eliminate blending (dilution of a larger transit depth to the expected WD-MS level due to the presence of a brighter star within the same resolution element) às a source of confusion., Radial velocity measurements should eliminate blending (dilution of a larger transit depth to the expected WD–MS level due to the presence of a brighter star within the same resolution element) as a source of confusion.405 This i5 simpler than in the planetary case. since WDs induce larger radial velocity variations on the orbital timescale.," This is simpler than in the planetary case, since WDs induce larger radial velocity variations on the orbital timescale."406 If variability or ellipsoidal modulation of the MS star flux. radial velocity variations. or characteristic accretion lummosity from the WD. should draw attention to a system as a candidate close WD-MS system. then transit searches could also be targeted towards these sources. since thegeometric," If variability or ellipsoidal modulation of the MS star flux, radial velocity variations, or characteristic accretion luminosity from the WD, should draw attention to a system as a candidate close WD–MS system, then transit searches could also be targeted towards these sources, since thegeometric"407Llere. we present a justification of the method of analytical expansions near the folds used in Section 3.1.following (2003).,"Here, we present a justification of the method of analytical expansions near the folds used in Section \ref{ss2.1} following \cite{alzh_03}."408.. This problem is not trivial. because the Jacobian of the lens mapping on the Fold is equal to zero.," This problem is not trivial, because the Jacobian of the lens mapping on the fold is equal to zero."409" We sav that an analvtie function. f(/) has order & at /=O. f£,]7(0)20 for;0.1...1 and [""'(0)z0 1978)): then we write f(/)=Ot)."," We say that an analytic function $f(t)$ has order $k$ at $t=0$, if $f^{(i)}(0)=0$ for $i=0,1,...,k-1$, and $f^{(k)}(0)\neq0$ \citealt{Poston_78}) ); then we write $f(t)= O(t^k)$."410 Lt is well known that the mapping of two-dimensional manifolds (e.g.. mapping (3))) ina neighborhood of the fold (a=y 0) can be reduced by the coordinate transformations to the normalform (see. e.g.. Poston&StewartLOTS:Pettersctal. 2001)): We suppose the initial mapping CX1)) is analytic: then transformations CX2)) are also analytic.," It is well known that the mapping of two-dimensional manifolds (e.g., mapping \ref{eq3}) )) in a neighborhood of the fold $\bmath x=\bmath y=0$ ) can be reduced by the coordinate transformations to the normalform (see, e.g., \citealt411{Poston_78,Petters}) ): We suppose the initial mapping \ref{eq0a}) ) is analytic; then transformations \ref{eq2a}) ) are also analytic."412 Let the source move along a parameterized curve v(f) such that (4) es(/)20 for f= 0. ο(0)= 0: (ii) the functions," Let the source move along a parameterized curve ${\rm {\bmath v}}\left( t413\right)$ such that (i) $v_2 \left( t \right) > 0$ for $t \ne 0$ , ${\rm {\bmath v}}\left( 0 \right) = 0$ ; (ii) the functions"414"noise level on 7« pixel boxes as approximate PSF areas and devided bv the peak intensity,",noise level on $7 \times 7$ pixel boxes as approximate PSF areas and devided by the peak intensity.415 We compare the observed dynamic ranges with expected flux ratios for possible conipanious of different masses (calculated following Burrows ct al., We compare the observed dynamic ranges with expected flux ratios for possible companions of different masses (calculated following Burrows et al.416 1997) next to TD 358623 (Fig., 1997) next to HD 358623 (Fig.417 5)., 5).418 The AIPE speckle eiunera. SILARP-I clearly gives the best dynamic ranec., The MPE speckle camera SHARP-I clearly gives the best dynamic range.419" In the SITARDP. images. we should have detected all stellar companions above ~0.1 AL... outside of ~0.5"", "," In the SHARP images, we should have detected all stellar companions above $\sim 0.1$ $_{\odot}$ outside of $\sim 0.5^{\prime \prime}$."420"Browndwarf compamious with ~25 Mj, would have been detecable at ~3” separations. nore Massive oues at smaller separatious (between ~0.5 aud 3/f "," Brown dwarf companions with $\sim 25$ $_{\rm Jup}$ would have been detectable at $\sim 3^{\prime \prime}$ separations, more massive ones at smaller separations (between $\sim 0.5$ and $3^{\prime \prime}$ )."421Next. we can compute the luuinositics of the four objects. studied.," Next, we can compute the luminosities of the four objects studied."422 Wo assume the Uipparcos distance towards IID 19913 À (17.752.1 pe) for all four objects., We assume the Hipparcos distance towards HD 199143 A $47.7 \pm 2.4$ pc) for all four objects.423 From the known VII color indices and/or known spectral vpes. we can estimate the effective temperatures aud bolometric corrections D.C. (taken from Iyeuvou Tartumaun 1995).," From the known $V-H$ color indices and/or known spectral types, we can estimate the effective temperatures and bolometric corrections B.C. (taken from Kenyon Hartmann 1995)."424" Temperatures. D.C.. and buninosities are listed in Table {,"," Temperatures, B.C., and luminosities are listed in Table 4."425 We placed the stars into the ILR diagram and coupared their locations with theoretical tracks aud isochrones by Palla Stabler (1999) aud Daraffe et al. (, We placed the stars into the H-R diagram and compared their locations with theoretical tracks and isochrones by Palla Stahler (1999) and Baraffe et al. (4261998) to estimate masses and ages.,1998) to estimate masses and ages.427" Rough values are eiven in Table {,", Rough values are given in Table 4.428 All four stars appear to be co-eval with an age of ~20 Alvis., All four stars appear to be co-eval with an age of $\sim 20$ Myrs.429 Iu conclusion. we find that all our data are cousisteu with ND 199113 aud ΠΟ 358623 cach having an carly-AI type stellar companion.," In conclusion, we find that all our data are consistent with HD 199143 and HD 358623 each having an early-M type stellar companion."430 In the case of ΠΟ 358623 D. the spectral type is confirmed by JIS colors al a spectzunn and companiouship is also confined bx conunnon proper motion.," In the case of HD 358623 B, the spectral type is confirmed by JHK colors and a spectrum and companionship is also confirmed by common proper motion."431 For HD. 199113 D. spectru as well as proper motion aud. hence. companionship. stil have to be confine.," For HD 199143 B, spectrum as well as proper motion and, hence, companionship, still have to be confirmed."432were obtained by applving a two-dimensional Gaussian taper falling to at 1. MÀ in both the w and ivedirections of the visibility data.,were obtained by applying a two-dimensional Gaussian taper falling to at 1 $\lambda$ in both the - and -directions of the visibility data.433 The rms noises in these images are 28.3 and 26.8 μον beam! for the northern and southern components. respectively.," The rms noises in these images are 28.3 and $26.8~\mu$ Jy $^{-1}$ for the northern and southern components, respectively."434 Table 2 shows the results of fitting Gaussian models to the observed. source spatial profiles using the AIPS++ task -IMAGEFITTER., Table 2 shows the results of fitting Gaussian models to the observed source spatial profiles using the AIPS++ task “IMAGEFITTER”.435 The source names. redshifts. and positions are listed in columns 1. 2. 3. and 4. respectively.," The source names, redshifts, and positions are listed in columns 1, 2, 3, and 4, respectively."436 Columns 5 and 6 show the peak and total [ιν densities of the fitted. Gaussian functions., Columns 5 and 6 show the peak and total flux densities of the fitted Gaussian functions.437 Columns 7. 8. and 9 list the fitted hall-power ellipse axes and (he position augles.," Columns 7, 8, and 9 list the fitted half-power ellipse axes and the position angles."438" The intrinsic brightness temperatures (corresponding to a vest-lrame frequency of 8 GIIz) of (he radio continuum structures seen in the northern and southern components of BRI 12020725 are (2.240.3)x10! and (1.7220.3)x10! Is. respectively,"," The intrinsic brightness temperatures (corresponding to a rest-frame frequency of 8 GHz) of the radio continuum structures seen in the northern and southern components of BRI 1202–0725 are $(2.2 \pm 0.3) \times 10^4$ and $(1.7 \pm 0.3) \times 10^4$ K, respectively."439" The total radio [τις density of BRI 12020725 at 1.4 Gllz. i. ο. the sum of the [αν densities of both components. is 565—5n µ.]ν. as measured with our sensitive VLBI array,"," The total radio flux density of BRI 1202–0725 at 1.4 GHz, i. e., the sum of the flux densities of both components, is $565 \pm 55~\mu$ Jy, as measured with our sensitive VLBI array."440" No conünuum emission is detected at (he the full resolution of the array. which is 29xT mas (PA=—6""). indicating the absence of anv radio continuum emission with Πας densities of >4o~40 pJy |."," No continuum emission is detected at the the full resolution of the array, which is $29 \times 7$ mas $-6^{\circ}$ ), indicating the absence of any radio continuum emission with flux densities of $\geq 4\sigma \simeq 40~\mu$ Jy $^{-1}$."441 This implies an upper limit to the intrinsic brightness temperature of 6.7xLO? IX [or anv compact radio source in BRI 12020725., This implies an upper limit to the intrinsic brightness temperature of $6.7 \times 10^5$ K for any compact radio source in BRI 1202–0725.442 We emphasize again that our coherence tests during these observations using two VLBI calibrators show that the lack of a strong point source in BRI 12020725 at the full resolution of the array cannot be due to the phase referencing procedure., We emphasize again that our coherence tests during these observations using two VLBI calibrators show that the lack of a strong point source in BRI 1202–0725 at the full resolution of the array cannot be due to the phase referencing procedure.443 Figures 4 and 5 show (he two components in BRI 12020725 at the highest. angular resolutions for which there are at least 4o detections., Figures 4 and 5 show the two components in BRI 1202–0725 at the highest angular resolutions for which there are at least $\sigma$ detections.444 These detections correspond (to resolutions of81x57 mas in position angle —11 for the northern component (Figure 4:6=25.7 idv +). and 111x82 mas in posilion angle —37° for the southern component (Figure 5: Jv 1)|).," These detections correspond to resolutions of $81 \times 57$ mas in position angle $-11^{\circ}$ for the northern component (Figure 4; $\sigma=25.7~\mu$ Jy $^{-1}$ ), and $111 \times 83$ mas in position angle $-37^{\circ}$ for the southern component (Figure 5; $\sigma=25.5~\mu$ Jy $^{-1}$ )."445 These images were obtained by applying (vo-dimensional Gaussian (apers falling to al 2.6 and 1.8 AIA in both the w and idirections of the visibility data. for the northern and southern components. respectively.," These images were obtained by applying two-dimensional Gaussian tapers falling to at 2.6 and 1.8 $\lambda$ in both the - and -directions of the visibility data, for the northern and southern components, respectively."446 The intermediate angular resolution image of the northern component (Figure +) shows a single continuum structure al 4.30 level., The intermediate angular resolution image of the northern component (Figure 4) shows a single continuum structure at $4.3\sigma$ level.447 However. (he southern component (Figure 5) seenis to be resolved into (wo continuum structures separated by 320 mas.," However, the southern component (Figure 5) seems to be resolved into two continuum structures separated by 320 mas."448 The stronger continuum source in Figure 5 is detected at 4.80 level. and the weaker at 3.70.," The stronger continuum source in Figure 5 is detected at $4.8\sigma$ level, and the weaker at $3.7\sigma$."449 The derived intrinsic brightness temperature of the continuum structure in Figure 4 (northern component) is (0.52:1.9)x107. and of the structures in Figure 5 (southern component) are (4.541.0)x107 and (3.531.0)x10! K. The extent of these continuum sources range between 0.4 and 1 kpe on the plane of the sky.," The derived intrinsic brightness temperature of the continuum structure in Figure 4 (northern component) is $(7.5 \pm 1.9) \times 10^4$, and of the structures in Figure 5 (southern component) are $(4.5 \pm 1.0) \times 10^4$ and $(3.5 \pm 1.0) \times 10^4$ K. The extent of these continuum sources range between 0.4 and 1 kpc on the plane of the sky."450 The single. resolved structure of the northern component. ancl the double structure of the southern component seen in (hese VLBI images are similar to the," The single, resolved structure of the northern component, and the double structure of the southern component seen in these VLBI images are similar to the"451eeonmeltric (ransil probability (e.g... Sackett(1999))) and the low number of stars within this volume (between 300-400. mainly. M-cdiwarls).,"geometric transit probability (e.g., \citet{sackett99}) ) and the low number of stars within this volume (between 300-400, mainly M-dwarfs)."452 In order to monitor Chousands to hundreds of thousands of stars. à (vpical search volume in our Galaxy. extends [rom hundreds to a lew thousand parsecs.," In order to monitor thousands to hundreds of thousands of stars, a typical search volume in our Galaxy extends from hundreds to a few thousand parsecs."453 In such cases. Af would only be a few ~10.!—107? seconds.," In such cases, $\Delta t$ would only be a few $\sim 10^{-1} -10^{-2}$ seconds."454 The nominal (photon noise) l-o timing error for transits of a 9th magnitude star over (he mission lifetime of is approximately 1 seconds. for.Ivepler il is approximately 1 second.," The nominal (photon noise) $\sigma$ timing error for transits of a 9th magnitude star over the mission lifetime of is approximately 7 seconds, for it is approximately 1 second."455 For a 12th magnitude star. Corot's errors are in excess of a minute. whilefvepler allains approximately 6 seconds.," For a 12th magnitude star, 's errors are in excess of a minute, while attains approximately 6 seconds."456 A rough estimate of the timing precision that can be achieved for transit measurements wilh a larger instrument can be made assuming photon counting statistics., A rough estimate of the timing precision that can be achieved for transit measurements with a larger instrument can be made assuming photon counting statistics.457" (2005) point out that the uncertainty in thecentral transit tme (assuming the ingress and egeress are well sampled) o. for a transit of duration / is given bv: where 2, and 2, correspond to the planet and star radius respectively."," \citet{holman05} point out that the uncertainty in the transit time (assuming the ingress and egress are well sampled) $\sigma$, for a transit of duration $t$ is given by: where $R_p$ and $R_*$ correspond to the planet and star radius respectively."458 E 15 the actual photon count rate of the star in the given instrument., $\Gamma$ is the actual photon count rate of the star in the given instrument.459 While this assumes the simplest direct approach to estimating the central transit (nme il nonetheless provides an indication of the inherent noise level., While this assumes the simplest direct approach to estimating the central transit time it nonetheless provides an indication of the inherent noise level.460 For a (vpical 8m class telescope D. is estimated to be ft) , For a typical 8-m class telescope $\Gamma$ is estimated to be $\sim 5.5 \times 10^{10} \cdot 10^{-0.4(V-12)}$ $^{-1}$.461Thus. for transit durations of the order of a few hours and a. stellar magnitude V.= 12. (hen measurements of a Jupiter sized. planet (assuming a solar mass parent star) would vieldl σ~0.32 seconds and an Earth-sized planet &~10 seconds.," Thus, for transit durations of the order of a few hours and a stellar magnitude $V=12$ , then measurements of a Jupiter sized planet (assuming a solar mass parent star) would yield $\sigma \sim 0.3$ seconds and an Earth-sized planet $\sigma \sim 10$ seconds."462 Thus. while Al's of the level described are bevond the sensitivity of andNepler. for 8-1 class instruments ancl Jupiter sized planets around 12(h magnitude or brighter stars they begin to enter the realm of practical measurement.," Thus, while $\Delta t$ 's of the level described are beyond the sensitivity of and, for 8-m class instruments and Jupiter sized planets around 12th magnitude or brighter stars they begin to enter the realm of practical measurement."463 The above estimates are also based upon straightforwzuxd estimation of the transit curve imine., The above estimates are also based upon straightforward estimation of the transit curve timing.464 Observational strategies which exploit an iterative approach (ie. making use of (he prior knowledge of previous transil evenis) to determine the optimal start and end of a given exposure may improve on this somewhat., Observational strategies which exploit an iterative approach (i.e. making use of the prior knowledge of previous transit events) to determine the optimal start and end of a given exposure may improve on this somewhat.465 However. we also note that other svstematics. including orbital precession. stellar variation. and instrumental variations (e.g. thermal conditions) will need to be νουν well understood to reach the level of ~0.1 second timing precision.," However, we also note that other systematics, including orbital precession, stellar variation, and instrumental variations (e.g. thermal conditions) will need to be very well understood to reach the level of $\sim 0.1$ second timing precision."466 If the exoplanet-star plane is normal to thatof (he solar svstem (Figure Ib) (hen parallax, If the exoplanet-star plane is normal to thatof the solar system (Figure 1b) then parallax467by Galloetal(2003) had a correlation iudex of 0.70+0.20. whereas the addition of the A0620-00 point refined the index to (0.58+0.16 dex(Calloetal.2006).,"by \citet{Gal03} had a correlation index of $0.70\pm0.20$, whereas the addition of the A0620-00 point refined the index to $0.58\pm0.16$ \citep{Gal06}."468. Corbeletal.(2008). found an of (0.51+0.06 for Viol Cre taken although such a shallow correlationis ruled outby However.Bboth the A0620-00 detection aud our new data.," \citet{Cor08} found an index of $0.51\pm0.06$ for V404 Cyg taken alone, although such a shallow correlation is ruled out by both the A0620-00 detection and our new data."469 the lack of low-TIuiiuositv radio detections prevents the correlation iudex being better coustrained. with this uneccrtainty leading to aut variation in the predicted radiobrighitucsscs of signifiquiescent sources.," However, the lack of low-luminosity radio detections prevents the correlation index being better constrained, with this uncertainty leading to significant variation in the predicted radio brightnesses of quiescent sources."470" Furthermore, it is not clear whether the imeasured relationship between XN-aav and radio Iuinositv should indeed be a siugle. nbrosen aw down to the lowest quiesccut Iuuinosities. as powerhasbeen from the simnltancous radio and N-vay detection of interredANG20-00 (Galloetal.00601"," Furthermore, it is not clear whether the measured relationship between X-ray and radio luminosity should indeed be a single, unbroken power law down to the lowest quiescent luminosities, as has been inferred from the simultaneous radio and X-ray detection of A0620-00 \citep{Gal06}."471 Advectiou-dominated accretion flow models predict a gradual softening of ie power-law photou iudex in the X-ray baud as a source decays to quiesceuce (c.g.Esinetal.1997)., Advection-dominated accretion flow models predict a gradual softening of the power-law photon index in the X-ray band as a source decays to quiescence \citep[e.g.][]{Esi97}.472. That softening necessarily leads to chauges iu the fraction of 1ο accretion lunünositv arising ia even N-rav baud., That softening necessarily leads to changes in the fraction of the accretion luminosity arising in a given X-ray band.473 if» ∙ ∙. accretion∙ Evenimuünositv theaud jet underlyingcmission were relationshipindeed a betweensingle power-aw. these changes would cause a slight. deviation.+ fromBH law in any. enrpirical micasirenment 6 hethat indererelationship.," Even if the underlying relationship between accretion luminosity and jet emission were indeed a single power-law, these changes would cause a slight deviation from the underlying power law in any empirical measurement of that relationship."474 powerEqually. as discussed in Sectiou 1.. svuclirotrou-douinated jet models predict a steepeniug of the relation below a critical N-rav Duniuositv (Yuan&Cui 2005)... although such models do not secu to be iplicable to either A0620 5m he con00οι ofA ∙," Equally, as discussed in Section \ref{sec:intro}, synchrotron-dominated jet models predict a steepening of the relation below a critical X-ray luminosity \citep{Yua05}, although such models do not seem to be applicable to either A0620-00 or Sgr $^{\ast}$."475"∙∙n Tu STUN’, uation,source disance, Uucertaiutiesand «quiesceutcorrelation Acayindex. is Imuinosity.sufficient to explain the non-detections of all three of tle sources we observed with the EVLA."," In summary, the combination of uncertainties in quiescent X-ray luminosity, source distance, and correlation index is sufficient to explain the non-detections of all three of the sources we observed with the EVLA."476 As Fig., As Fig.477 shows. when considering these uucertaiuties. none of theH1. three sources observed with the EVLAi fall siguificautly below the radioκ⊸ταν correlation : by GalloAoetal⋅⊀E∙00). ∙ the A0620-00 fitteddetection to constrain the correlationwhich usedindex down to the lowest luminosities.," \ref{fig:lrlx} shows, when considering these uncertainties, none of the three sources observed with the EVLA fall significantly below the radio/X-ray correlation fitted by \citet{Gal06}, which used the A0620-00 detection to constrain the correlation index down to the lowest luminosities."478 Thus our upper limits cannot be used to robustly constrain models for the radio/X-ray coupling iu quiesceuce., Thus our upper limits cannot be used to robustly constrain models for the radio/X-ray coupling in quiescence.479" At the sensitivity levels achieved in our EVLA observations. we are reaching the regime where stellar radio cinission might be expected to contribute to wo Bnradio nabriehtuess akof paniwosonrees,"," At the sensitivity levels achieved in our EVLA observations, we are reaching the regime where stellar radio emission might be expected to contribute to the radio brightness of the sources."480 CHidel(2002 ο... evrosvuchnrotron cuussion from stellar radio flares cau reach LOS PII 3 for RS CVns EK Com aud Algol systems.," \citet{Gue02} showed that the peak luminosities of incoherent gyrosynchrotron emission from stellar radio flares can reach $10^{18}$ $^{-1}$ $^{-1}$ for RS CVn, FK Com and Algol systems."481 At the distance of the closest of our systems. NTE JLLLS|4NO (Table 23). this corresponds to a radio fux density of 3.," At the distance of the closest of our systems, XTE J1118+480 (Table \ref{tab:final}) ), this corresponds to a radio flux density of $^{-1}$."482 Hosever. our lack of detections suggests that any stellar eiuissiou is significautly faiuter than this.," However, our lack of detections suggests that any stellar emission is significantly fainter than this."483 The spectral types of the donor stars for our tarect sourees are listed in Table 2., The spectral types of the donor stars for our target sources are listed in Table \ref{tab:final}.484 The mass donors for the three sources observed with the EVLA are all ↕⋜↧↑↸∖≓↑⋅↖↴⋉∖⋯⋜↧↕∐↴∖↴↸∖≺∣⋯∖∐↸⊳↸∖∪↥⋅↴∖↴∏↴⋝∶, The mass donors for the three sources observed with the EVLA are all late-type main sequence or subgiant stars.485↴∙⊾↕⋜⋯↑↴∖↴↑⋜∐⋅↴∖↴∙⊺⋅↖↴≻↕↸⊳⋜↧↕ radio huniosities for carly/imid K stars are du the pueC» dos101 3&10!’ ssDIMI 32M with the higher cud of the range corresponding to the most rapid rotators.," Typical radio luminosities for early/mid K stars are in the range $1\times10^{14}$ $3\times10^{15}$ $^{-1}$ $^{-1}$, with the higher end of the range corresponding to the most rapid rotators."486 Although our N-ray binary targets will be tidally locked (with orbital periods iu the range 62.9hh). there is little correlation. between rotation rate and radio fux deusitv iu dwarf stars (Ciudel 2002).. so the high rotation speeds should have minimal effect.," Although our X-ray binary targets will be tidally locked (with orbital periods in the range h), there is little correlation between rotation rate and radio flux density in dwarf stars \citep{Gue02}, so the high rotation speeds should have minimal effect."487 We would therefore expect radio flux deusities of «2.5 i (d/lkpe) romthedonorstersM," We would therefore expect radio flux densities of $<2.5$ $\mu$ $(d/{\rm 1\,kpc})$ from the donor stars."488 hileforastromctriepurposc. f ragluiminositycorrelationshouldwepushdecperonthesequiescecitsgs," While for astrometric purposes the source of the radio emission is unimportant, stellar emission could begin to confuse the radio/X-ray luminosity correlation should we push deeper on these quiescent systems."489 We have inade deep radio observations of the fields anmroundiue Etwelve coufinued or candidate black hole binaries., We have made deep radio observations of the fields surrounding twelve confirmed or candidate black hole X-ray binaries.490 The ouly source detected was Swiftate J1753.5-Tpomη at a level 1.," The only source detected was Swift J1753.5-0127, at a level of $^{-1}$ ."491 The typical 30 upper limits on ofthe 0.1radio brielituesses of the remaiuiue, The typical $3\sigma$ upper limits on the radio brightnesses of the remaining492temperature equal to the effective radiation temperature Eq. (16),temperature equal to the effective radiation temperature Eq. )493)) to high until z200., to high until $z\sim 200$.494" In the thermodynamic equilibrium, entropy is conserved, and we can to calculate the energy losses due to the expansion of the Universe without referring toComptonization, since it is not really important which physical process is responsible for the equilibrium."," In the thermodynamic equilibrium, entropy is conserved, and we can to calculate the energy losses due to the expansion of the Universe without referring to, since it is not really important which physical process is responsible for the equilibrium."495" We start with the thermodynamic relation, where S is the entropy, U total thermal energy, P pressure, V volume, M chemical potential, and N the number of particles."," We start with the thermodynamic relation, where $S$ is the entropy, $U$ total thermal energy, $P$ pressure, $V$ volume, $M$ chemical potential, and $N$ the number of particles."496" T, is the common temperature of photons, ions, and electrons."," $T_{\gamma}$ is the common temperature of photons, ions, and electrons."497 We can ignore the last term for ions and electrons since their number is fixed after big bang nucleosynthesisend£]., We can ignore the last term for ions and electrons since their number is fixed after big bang nucleosynthesis.498" For photons as well we can ignore the last term if we assume that their chemical potential is 0, which is true in full thermodynamic equilibrium."," For photons as well we can ignore the last term if we assume that their chemical potential is $0$, which is true in full thermodynamic equilibrium."499" With this assumption we can write the equation for total entropy per baryon, σ=s/np, where s is the entropy density and ng is the baryon numberHere ag is the radiation constant and N the number of non-relativistic particles per baryon, so that Nng=περ."," With this assumption we can write the equation for total entropy per baryon, $\sigma=s/\nB$, where $s$ is the entropy density and $\nB$ is the baryon number $\aR$ is the radiation constant and $N$ the number of non-relativistic particles per baryon, so that $N\,\nB =\neb $."500 We now integrate Eq., We now integrate Eq.501" towhere C is an arbitrary constant of integration, which is not important for our calculation."," to $C$ is an arbitrary constant of integration, which is not important for our calculation."502" The first term above is the contribution from photons, and it is clear that in the absence of second term, T,ος(1+z)."," The first term above is the contribution from photons, and it is clear that in the absence of second term, $T_{\gamma}\propto (1+z)$."503 The second term is the contribution of non-relativistic particles and causes the temperature to drop slightly faster in order to conserve entropy per baryon., The second term is the contribution of non-relativistic particles and causes the temperature to drop slightly faster in order to conserve entropy per baryon.504" We can write T,=T(1*D, where T=Ti(1+z)/(1zi) is thebackground temperature proportional to 1+z, and {«1 is the fractional deviation from this law."," We can write $T_{\gamma}=T(1+t)$, where $T=\Ti (1+z)/(1+z_{i})$ is thebackground temperature proportional to $1+z$, and $t\ll 1$ is the fractional deviation from this law."505 We can take the initial deviation t(z;)=O at initial redshift z; and then calculate the subsequent energy losses to adiabatic cooling of baryons at later redshifts., We can take the initial deviation $t(z_i)=0$ at initial redshift $z_i$ and then calculate the subsequent energy losses to adiabatic cooling of baryons at later redshifts.506 Thus we can write as, Thus we can write as507 lu order to define a vaid DN sample for statistical investigation. one needs to separate a higher uuuber of nearby PNe with the same norphological characteristics.," In order to define a valid PN sample for statistical investigation, one needs to separate a higher number of nearby PNe with the same morphological characteristics."508 We define such a sample (USNO-PN) in the subsection below (Section 5.1)., We define such a sample (USNO-PN) in the subsection below (Section 5.1).509 Towever. the quality of radio-contimuni observations today simply is uot sensitive enough to provide a high nuniber of similar PNe distances.," However, the quality of radio-continuum observations today simply is not sensitive enough to provide a high number of similar PNe ."510"ereater than any other active pixel within a racius of 90, of the peak.",greater than any other active pixel within a radius of $2 {\theta_c}$ of the peak.511 We then eroup together all active poiuts in this area to create a unique cluster candidate., We then group together all active points in this area to create a unique cluster candidate.512 Ifa single active pixel is a peak by defaul ic. there are no other active pixels within 20. of it we disregard this peak aud do not include it in our final analysis., If a single active pixel is a peak by default – i.e. there are no other active pixels within $2 \theta_c$ of it – we disregard this peak and do not include it in our final analysis.513 Such isolated peaks are uulikely to be caused by real clusters., Such isolated peaks are unlikely to be caused by real clusters.514 Finally. we compute the weighted mean of all active points grouped. together as a single cluster to determine the most likely cluster candidate ceutroid.," Finally, we compute the weighted mean of all active points grouped together as a single cluster to determine the most likely cluster candidate centroid."515 Iu addition to deteriuniug our detection thresholds. our Monte Carlo παπαος were used to estimate our detection efficiency. and the effective area of our cluster search.," In addition to determining our detection thresholds, our Monte Carlo simulations were used to estimate our detection efficiency and the effective area of our cluster search."516 This is an iuportant aspect of any cosmological survey as it allows us to determine the volume sampled by the EDCCID and thus measure the space density of clusters from the catalogue., This is an important aspect of any cosmological survey as it allows us to determine the volume sampled by the EDCCII and thus measure the space density of clusters from the catalogue.517 Previous applications of automated optical clusterfinding algorithms have based their efficiency measurements on them ability to detect clusters within artificial galaxw data., Previous applications of automated optical cluster–finding algorithms have based their efficiency measurements on their ability to detect clusters within artificial galaxy data.518 For example. Postman et al. (," For example, Postman et al. ("5191999) and Kepner ot al. (,1999) and Kepner et al. (5201999) created artificial ealaxy catalogues with the same statistical properties as real galaxy catalogues they matched the surface deusity of galaxies and/or the lareescale clusteriug properties of the ealaxies.,1999) created artificial galaxy catalogues with the same statistical properties as real galaxy catalogues they matched the surface density of galaxies and/or the large–scale clustering properties of the galaxies.521 They then added artificial clusters to such simulated galaxy data., They then added artificial clusters to such simulated galaxy data.522 We however found this aethod imadequate since such simulated galaxy catalogues cannot fully reproduce he hierarchical structure of galaxies in the universe., We however found this method inadequate since such simulated galaxy catalogues cannot fully reproduce the hierarchical structure of galaxies in the universe.523" Specifically, the effects of superpositions of structures. variations in the field couuts. or largescale structures are ward to include in these simulations."," Specifically, the effects of superpositions of structures, variations in the field counts, or large–scale structures are hard to include in these simulations."524" Therefore. for cach combination of RT aud LT (see Tables 1 and 2)). we added a total 6000 artificial clusters (20 at a time) at random iu he EDSCC and computed our success rate in detecting hese artificial clusters above our thresholds (an artificial cluster was considered detected if a candidate cluster was found bv our algorithm within 20, of the original coordinate of the artificial cluster: 0. was evaluated at redshift of the artificial cluster)."," Therefore, for each combination of RT and LT (see Tables \ref{richthreshtable} and \ref{likethreshtable}) ), we added a total 6000 artificial clusters (20 at a time) at random in the EDSGC and computed our success rate in detecting these artificial clusters above our thresholds (an artificial cluster was considered detected if a candidate cluster was found by our algorithm within $2 \theta_c$ of the original coordinate of the artificial cluster; $\theta_c$ was evaluated at the redshift of the artificial cluster)."525" Iu Figure 2.. we present our average detection efficienci as a function of input redshift aud richness (using 10""«105test area discussed in Section 2. aud in Figu 1)."," In Figure \ref{efficiencies}, we present our average detection efficiencies as a function of input redshift and richness (using the $10^{\circ}\times10^{\circ}$test area discussed in Section \ref{EDSGC} and in Figure \ref{edsgc10x10}) )."526 The error bars are the standard deviation observed between the different trials of 20 clusters added to EDSGCC data at any one time., The error bars are the standard deviation observed between the different trials of 20 clusters added to the EDSGC data at any one time.527 Iu FigureOo ?3.. we show an example of the aneOo dependence of our detection efficiency.," In Figure \ref{LSSinterfere}, we show an example of the angular dependence of our detection efficiency."528 It is iuteresti o note that our efficicney is stronely correlated with areescale structure (LSS) in the Universe., It is interesting to note that our efficiency is strongly correlated with the large–scale structure (LSS) in the Universe.529" This effec uost prominent for lower richness clusters. while for vicher systems (A,2 200). it is insignificant (except he highest redshifts probed by our simulations where effect resulted in a loss of z15% of artificial clusters)."," This effect is most prominent for lower richness clusters, while for the richer systems $R_m\ge 200$ ), it is insignificant (except at the highest redshifts probed by our simulations where the effect resulted in a loss of $\simeq15\%$ of artificial clusters)."530 Iu addition to this interference. our efficiency in detectiis ow richness. hieh redshift clusters was hiudered bv 1ο uaenitude linüt of the EDSQGC€ data since the nuuber of potentially visible galaxies in these artificial clusters decreases below the noise in the backerounud.," In addition to this interference, our efficiency in detecting low richness, high redshift clusters was hindered by the magnitude limit of the EDSGC data since the number of potentially visible galaxies in these artificial clusters decreases below the noise in the background."531 This problem of LSS iutertercuce appears to citect both over aud underdense regious of the EDSCC data (νο compared to the mean surface density of galaxies)., This problem of LSS interference appears to effect both over– and under–dense regions of the EDSGC data (when compared to the mean surface density of galaxies).532 For example. the most striking example of this LSS interference is seen at coordinates 2.0. 7.5Γ in Figure 1.. where a eroup of clusters Abell 2730. :2]. 2719. 2755. and 128 ias reduce our detection cficiency to alinost zero (seo Figure 3. )," For example, the most striking example of this LSS interference is seen at coordinates 2.0, 7.5 in Figure \ref{edsgc10x10}, where a group of clusters – Abell 2730, 2721, 2749, 2755, and 12S – has reduce our detection efficiency to almost zero (see Figure \ref{LSSinterfere} )."533 Iu contrast. we also observe iu Figure 3. a low detection efficiency. near coordinates 8.0. 6.0 which coincides with a underdense region in Fietye L..," In contrast, we also observe in Figure \ref{LSSinterfere} a low detection efficiency near coordinates 8.0, 6.0 which coincides with a under–dense region in Figure \ref{edsgc10x10}."534 We believe this effect is caused by our assuptiou o a flat. unifoiiu Oogalaxy backeroundC» as used in the matched filter algorithui (see Section 23.1)).," We believe this effect is caused by our assumption of a flat, uniform galaxy background as used in the matched filter algorithm (see Section \ref{matched}) )."535 In both over axd nuderdense regions. our assuniptiou of a flat backegrouxd with the mean surface deusity of the EDSQGC is poOr and therefore. we suppress the overall likelihood of t10 cluster detection the model of the backeround arouxd," In both over– and under–dense regions, our assumption of a flat background with the mean surface density of the EDSGC is poor and therefore, we suppress the overall likelihood of the cluster detection the model of the background around"536The SDSS photometric data indicate that the SLACS enses are only slightly. more likely to be associated with companion galaxies than comparable lenses selected. from he SDSS. though the uncertainty of the photometric identilications makes the dillerence negligible.,"The SDSS photometric data indicate that the SLACS lenses are only slightly more likely to be associated with companion galaxies than comparable lenses selected from the SDSS, though the uncertainty of the photometric identifications makes the difference negligible."537 We therefore conclude that SLACS lenses do lie in tvpical environments roth globally anc locally., We therefore conclude that SLACS lenses do lie in typical environments both globally and locally.538 However. we also find that lens systems with steeper than isothermal density slopes are orclerentially associated with companion galaxies compared o lenses with shallower density slopes.," However, we also find that lens systems with steeper than isothermal density slopes are preferentially associated with companion galaxies compared to lenses with shallower density slopes."539 N-body simulations suggest that interactions with neighbouring galaxies can induce a steepening in the density. slope (Dobkeοἱal.2007) and there are other lens systems. with companion ealaxics that are found to be best modelled. with steeper [jan isothermal profiles (e.g...Rusinetal.2002Augerct 2007b).," N-body simulations suggest that interactions with neighbouring galaxies can induce a steepening in the density slope \citep{dobke} and there are other lens systems with companion galaxies that are found to be best modelled with steeper than isothermal profiles \citep[e.g.,][]{rusin,augerb}."540. The interaction-induced. steepening is a transient ellect and the density. profile of the galaxy. will return to isothermal approximately 0.5-2 ον after the encounter with the neighbour (Dobkeetal.200, The interaction-induced steepening is a transient effect and the density profile of the galaxy will return to isothermal approximately 0.5-2 Gyr after the encounter with the neighbour \citep{dobke}.541" This may account. for the large range of IN, for lenses with nearly isothermal profiles: isothermal lenses with a companion may be in the relaxed state before or after an encounter with the neighbour galaxy.", This may account for the large range of $N_w$ for lenses with nearly isothermal profiles; isothermal lenses with a companion may be in the relaxed state before or after an encounter with the neighbour galaxy.542 ‘his stripping mechanism may also account for local observations of dark matter deficient. galaxies (e.g...lto-manowskyetal.2t)03:Proctorct 2005).," This stripping mechanism may also account for local observations of dark matter deficient galaxies \citep[e.g.,][]{romanowsky,proctor}."543. We note that one stecper-prolile system. SDSSJ1250|0523. is not photometrically associated with any neighbouring galaxies: this perhaps illustrates the limits of using photometry to find. perturbing companions or demonstrates that other factors also inlluence the slope of the density. profile.," We note that one steeper-profile system, SDSSJ1250+0523, is not photometrically associated with any neighbouring galaxies; this perhaps illustrates the limits of using photometry to find perturbing companions or demonstrates that other factors also influence the slope of the density profile."544 Additionally. we have not found. an environmental bias to account. for. the shallower lenses.," Additionally, we have not found an environmental bias to account for the shallower lenses."545 Llowever. if a lens galaxy is embedded. in a cluster. the joint. profile of the cluster ancl galaxy would tend to be moclelled with a shallower than isothermal profile if a single-component power law is used (ignoring interactions between he two halos).," However, if a lens galaxy is embedded in a cluster, the joint profile of the cluster and galaxy would tend to be modelled with a shallower than isothermal profile if a single-component power law is used (ignoring interactions between the two halos)."546 This ellect is dependent on the location of he lens with respect to the centre of the cluster. which our photometric analysis is unable to address.," This effect is dependent on the location of the lens with respect to the centre of the cluster, which our photometric analysis is unable to address."547 More complete 101 spectroscopy would. better characterize the globa environments of these lens systems anc allow correlations »tween the mass slopes and. cluster centre ollsets t0. be investigated., More complete field spectroscopy would better characterize the global environments of these lens systems and allow correlations between the mass slopes and cluster centre offsets to be investigated.548 Furthermore. a spectroscopic investigation of he local environments of the complete sample of SLACS earlv-tvpe lenses would confirm the correlation. indicate » our photometric analvsis ancl provide strong evidence or truncation caused by ealaxy interactions.," Furthermore, a spectroscopic investigation of the local environments of the complete sample of SLACS early-type lenses would confirm the correlation indicated by our photometric analysis and provide strong evidence for truncation caused by galaxy interactions."549 The author would like to thank Chris Fassnacht and Ami Choi. as well as the anonymous referee. for providing useful comments on the manuscript.," The author would like to thank Chris Fassnacht and Ami Choi, as well as the anonymous referee, for providing useful comments on the manuscript."550 This work has made extensive use of the SDSS database., This work has made extensive use of the SDSS database.551 Funding for the SDSS and SDSS-LL has been provided by the Alfrecl P. Sloan Foundation. the Participating Institutions. the National Science Foundation. theU.S. Department of Energy. the National Aeronautics and Space Administration. theJapanese Monbukagakusho. the Max Planck Society. and the Higher Education Funding Council for England.," Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England."552 Part. of this work was supported by the European Communitys Sixth Framework. Marie Curie Research Training Network Programme. Contract No.," Part of this work was supported by the European Community's Sixth Framework Marie Curie Research Training Network Programme, Contract No."553" AIREN-C""I-2004-505183. “ANGLES”"," MRTN-CT-2004-505183 “ANGLES""."554x* for Ep(Il) = 600 IK to 47 x ! for Ep(Il) = 350 K. A similar increase is seen in (he atomic deuterium abundance.,$\times$$^{-7}$ for $E_D$ (H) = 600 K to 4.7 $\times$ $^{-4}$ for $E_D$ (H) = 350 K. A similar increase is seen in the atomic deuterium abundance.555 The column density of H doesnt change but (hat of D increases. from. 6.5. x 15 »? to 3.5. x LO! »7., The column density of H doesn't change but that of D increases from 6.5 $\times$ $^{15}$ $^{-2}$ to 3.5 $\times$ $^{16}$ $^{-2}$.556 The changes in Il and D abundance lead to changes in the abundances and column densities of other gaseous molecules. (, The changes in H and D abundance lead to changes in the abundances and column densities of other gaseous molecules. (557The column densiües of molecules in the ice mantles do not show much variation with the change in £p(I1)).,The column densities of molecules in the ice mantles do not show much variation with the change in $E_D$ (H)).558 In the gas phase. with Ly) U1) 2350 IX NCOIISQ) increases by ~ a factor of 10. as does NOHISCO). NOIDCO) and NCNIIS).," In the gas phase, with $E_D$ (H) =350 K $_2$ O) increases by $\sim$ a factor of 10, as does $_2$ CO), N(HDCO) and $_2$ )."559" The largest increase is for. CLL,OLL whose column density increases from 6.5 x 10 em7 to 2.2 x 10* 7.", The largest increase is for $_3$ OH whose column density increases from 6.5 $\times$ $^5$ $^{-2}$ to 2.2 $\times$ $^7$ $^{-2}$.560 The reduction in binding enerev also changes some cleuteration levels. most noliceably — increases to 3.4 (up from 0.9 in Model D).," The reduction in binding energy also changes some deuteration levels, most noticeably $^+$ $^+$ increases to 3.4 (up from 0.9 in Model B)."561 Using the lower value of Ep(II) in Model D changes the eoliunn densities even more., Using the lower value of $E_D$ (H) in Model D changes the column densities even more.562 There are large reductions in column densities of [actors between 5 and 10 for many of the deuterated molecules. including DeQ. all of the deuterated isotopomers of CII; and NIL. DCN. and CoD. For example. N(DsQ) is reduced from 7 x 107 ? to 87 x LOM and N(DCN) from 1.2 x LO’ ? to L2 x 10! ?.," There are large reductions in column densities of factors between 5 and 10 for many of the deuterated molecules, including $_2$ O, all of the deuterated isotopomers of $_4$ and $_3$, DCN, and $_2$ D. For example, $_2$ O) is reduced from 7 $\times$ $^{12}$ $^{-2}$ to 8.7 $\times$ $^{11}$ and N(DCN) from 1.2 $\times$ $^{12}$ $^{-2}$ to 1.2 $\times$ $^{11}$ $^{-2}$."563 1 also a factor of 4 higher than inModel D. The choice of H-atom binding energy can therefore have quite an effect on the calculated abundances., $^+$ $^+$ is also a factor of 4 higher than inModel D. The choice of H-atom binding energy can therefore have quite an effect on the calculated abundances.564 We have chosen (he larger value since (he results lor a few species are in better agreement with (he observations: in particular using ZZp(II) = 350 Ix produces ratios which are much higher than observed., We have chosen the larger value since the results for a few species are in better agreement with the observations: in particular using $E_D$ (H) = 350 K produces $^+$ $^+$ ratios which are much higher than observed.565 The radial distribution of molecular column densities. as calculated in Model D. are shown in Figure 14..," The radial distribution of molecular column densities, as calculated in Model B, are shown in Figure \ref{fig:rad_cd_a}."566 N(CO) is fairly flat. with radius for /2 > LOO AU. but shows a sharp rise al 22 = 50 AU. where thermal desorption of this molecule is efficient throughout the verlical extent of the disk.," N(CO) is fairly flat with radius for $R$ $>$ 100 AU, but shows a sharp rise at $R$ = 50 AU, where thermal desorption of this molecule is efficient throughout the vertical extent of the disk."567 Ne has a more gradual increase because of its slightly. lower binding enerev which means that it can desorb at larger radii., $_2$ has a more gradual increase because of its slightly lower binding energy which means that it can desorb at larger radii.568 As cliseussecl in the previous section. the presence of these volatile species in (he eas can lead to the production of other molecules. which are not themselves thermally desorbed.," As discussed in the previous section, the presence of these volatile species in the gas can lead to the production of other molecules, which are not themselves thermally desorbed."569 For example. photocdissociation of No and COcan release the elements required (o form molecules such as HCN and and the column densities of these species increase towards (he central star.," For example, photodissociation of $_2$ and COcan release the elements required to form molecules such as HCN and $^+$ , and the column densities of these species increase towards the central star."570" N(D, ) Falls off with decreasing radius. because of the increase in temperature which inhibits its formation from IL,."," $_3^+$ ) falls off with decreasing radius, because of the increase in temperature which inhibits its formation from $_3^+$."571 The relative levels of deuteration in some other molecules such as . H4O. ICN and NID; also decrease with decreasing 2.," The relative levels of deuteration in some other molecules such as $^+$ , $_2$ O, HCN and $_3$ also decrease with decreasing $R$ ."572 CII; and H3CO clo, $_4$ and $_2$ CO do573"High-energy y--ray emission is expected from starburst galaxies, as well as normal galaxies, through cosmic ray interactions with the interstellar medium.","High-energy -ray emission is expected from starburst galaxies, as well as normal galaxies, through cosmic ray interactions with the interstellar medium."574" Cosmic rays, which are mainly made up of charged, light hadrons and leptons that are accelerated in these galaxies, interact with the ambient interstellar medium, creating pions which will in turn decay into y--rays for the hadrons, and bremsstrahlung and inverse Compton radiations for the leptons (seee.g.?).."," Cosmic rays, which are mainly made up of charged, light hadrons and leptons that are accelerated in these galaxies, interact with the ambient interstellar medium, creating pions which will in turn decay into -rays for the hadrons, and bremsstrahlung and inverse Compton radiations for the leptons \citep[see e.g.][]{2010ApJ...722L..58S}."575" Based on a simple model, ? suggest there is a relationship between the expected y--ray luminosity, the supernova rateRsn,, and the total gas mass ffor a given galaxy: where is the supernova rate in the MilkyWay and d the distance of the considered galaxy."," Based on a simple model, \citet{2001ApJ...558...63P} suggest there is a relationship between the expected -ray luminosity, the supernova rate, and the total gas mass for a given galaxy: where is the supernova rate in the MilkyWay and $d$ the distance of the considered galaxy."576" Figure D] shows the y--ray luminosity measured with ffor the different known high-energy emitting starburst and normal galaxies, against RsNXMeas, and the upper limits at 2σ confidence level on the fluxes derived from the sources in our sample."," Figure \ref{fig-relation} shows the -ray luminosity measured with for the different known high-energy emitting starburst and normal galaxies, against $\rsn \times \mgas$, and the upper limits at $\sigma$ confidence level on the fluxes derived from the sources in our sample."577" For comparison, we added the 22 galaxy 11068, also detected at high energies (?).."," For comparison, we added the 2 galaxy 1068, also detected at high energies \citep{2010A+A...524A..72L}."578" This object encompasses a starburst region in its centre as well, but we argued in ? that the associated high-energy y--ray emission is more likely due to the central AGN activity than from the starburst activity only, owing to its relatively high y--ray Iuminosity."," This object encompasses a starburst region in its centre as well, but we argued in \citet{2010A+A...524A..72L} that the associated high-energy -ray emission is more likely due to the central AGN activity than from the starburst activity only, owing to its relatively high -ray luminosity."579" As can be seen in Fig. D],"," As can be seen in Fig. \ref{fig-relation},"580" the upper limit reported by the ccollaboration on 333 (?) is very stringent compared to expectation, if the high-energy flux is indeed directly connected to RsX Μαις, that is, if the high energy emission is solely due"," the upper limit reported by the collaboration on 33 \citep{2010A+A...523L...2A} is very stringent compared to expectation, if the high-energy flux is indeed directly connected to $\rsn \times \mgas$ , that is, if the high energy emission is solely due"581mules away from the former two stations.,miles away from the former two stations.582 WINX signal iun Pune (8723 UN . 7357 19E) did not respond to this CRB.," VTX signal in Pune $18^\circ$ 34'N , $73^\circ$ 49'E) did not respond to this GRB."583 This station is located im the west of Iolkata by 1000 iles., This station is located in the west of Kolkata by $1000$ miles.584 The vertical line in Fig., The vertical line in Fig.585 Lb is the satellite detection time of the CRB (in UT ou 21th. April. 2009).," 4 is the satellite detection time of the GRB (in UT on 24th April, 2009)."586 The upper panel is the RIIESSI satellite observation (background subtracted)., The upper panel is the RHESSI satellite observation (background subtracted).587 The VLF response occurred after a few secouds., The VLF response occurred after a few seconds.588 The signals in all the three places were found to be correlated uuediatelv after the eveut for tens of seconds., The signals in all the three places were found to be correlated immediately after the event for tens of seconds.589 The Earth's ionosphere is a large. open and noisy detector. but it is free of production and naintenance costs.," The Earth's ionosphere is a large, open and noisy detector, but it is free of production and maintenance costs."590 We reported here two observations where we showed how the responded to one Soft Cama Bay and oue Ganuna Ray Burst., We reported here two observations where we showed how the ionosphere responded to one Soft Gamma Ray repeater and one Gamma Ray Burst.591 Civeu that oulyionosphere a handful of such detections are present m therepeater literature. our observation certainly is a significant addition.," Given that only a handful of such detections are present in the literature, our observation certainly is a significant addition."592 From a careful analysis. significant kuowledge about the interaction of Ligh euergy radiation with 1e Earth's ionosphere can be obtained.," From a careful analysis, significant knowledge about the interaction of high energy radiation with the Earth's ionosphere can be obtained."593 The analysis is in progress aud will be reported elsewhere., The analysis is in progress and will be reported elsewhere.594 Frou the astrobiological point of view. the burning of the iouosphiere by the coustant bombardinenuts of the extra-terrestrial energetic events is of great concern.," From the astrobiological point of view, the burning of the ionosphere by the constant bombardments of the extra-terrestrial energetic events is of great concern."595 It has been reported by Than et al. [, It has been reported by Inan et al. [59622] iat the otherwise neutral atmosphere at a height of 30 km was ionized by the SCR 1806-20.,22] that the otherwise neutral atmosphere at a height of $30$ km was ionized by the SGR 1806-20.597 If iv such event happened in our own galaxy. the effect could have been aud the whole would have been converted. to for a short-wliledevastating at least. exposing us to rostile mosphereradiations from space.," If any such event happened in our own galaxy, the effect could have been devastating and the whole atmosphere would have been converted to ionosphere for a short-while at least, exposing us to hostile radiations from space."598 In this seuse. ionosphereour study is linked to the exobiologv. aud more such studies are encouraged in this direction.," In this sense, our study is linked to the exobiology and more such studies are encouraged in this direction."599 SISAL acknowledges the support froma CSIR/JRE Fellowship., SKM acknowledges the support from a CSIR/JRF Fellowship.600 SS acknowledge the support frou ISRORESPOND project., SS acknowledge the support from ISRORESPOND project.601gyrokinetic theory. that they are passive to the Alfvénnic fluctuations. but have no parallel cascade along the exact magnetic field lines (Schekochihinetal.2009).,"gyrokinetic theory, that they are passive to the Alfvénnic fluctuations, but have no parallel cascade along the exact magnetic field lines \citep{schekochihin09}."602. This may explain why there is a compressive cascade in the solar wind: the compressive fluctuations are expected to be damped at a rate proportional to their parallel wavenumber >~& (Barnes1966;Schekochihinetal.2009;Klein2012). but if is very small then they are not heavily damped and ean ΚΕcascade nonlinearly.," This may explain why there is a compressive cascade in the solar wind: the compressive fluctuations are expected to be damped at a rate proportional to their parallel wavenumber $\gamma\sim k_\parallel$ \citep{barnes66,schekochihin09,klein12} but if $k_\parallel$ is very small then they are not heavily damped and can cascade nonlinearly."603 An alternative explanation is that the less anisotropic compressive fluctuations are generated but are quickly damped. leaving the highly elongated structures to be observed.," An alternative explanation is that the less anisotropic compressive fluctuations are generated but are quickly damped, leaving the highly elongated structures to be observed."604 This work was supported by NASA contract NNNOOAADIC. NASA grant NNXOO9AEAIG and the Leverhulme Trust Network for Magnetized Plasma Turbulence.," This work was supported by NASA contract NNN06AA01C, NASA grant NNX09AE41G and the Leverhulme Trust Network for Magnetized Plasma Turbulence."605 Ulysses data was obtained from CDAWeb nasa.gov)., Ulysses data was obtained from CDAWeb ).606 , 607In the deeply embedded phase of low-mass star formation. it is often only possible to trace the dynamics of gas in a young stellar object (YSO) by analysing resolved emission-line profiles.,"In the deeply embedded phase of low-mass star formation, it is often only possible to trace the dynamics of gas in a young stellar object (YSO) by analysing resolved emission-line profiles."608 The various dynamical processes include infall from the surrounding envelope towards the central protostar. molecular outflows caused by jets ejected from the central object. and strong turbulence induced within the inner parts of the envelope by small-scale shocks(??).," The various dynamical processes include infall from the surrounding envelope towards the central protostar, molecular outflows caused by jets ejected from the central object, and strong turbulence induced within the inner parts of the envelope by small-scale shocks."609. One of the goals of the Water In Star-forming regions with. (WISH) key programme Is to use water as a probe of these processes and determine its abundance in the various components as a function of evolution (van Dishoeck et al., One of the goals of the Water In Star-forming regions with (WISH) key programme is to use water as a probe of these processes and determine its abundance in the various components as a function of evolution (van Dishoeck et al.610 in prep.)., in prep.).611 Spectrally resolved observations of the H»O Γιο line at 557 GHz with ODIN and SWAS towards low-mass star-forming regions have revealed it to be broad. «20 ss'. indicative of an origi in shocks?).," Spectrally resolved observations of the $_2$ O $_{10}$ $_{01}$ line at 557 GHz with ODIN and SWAS towards low-mass star-forming regions have revealed it to be broad, $\sim$ 20 $^{-1}$, indicative of an origin in shocks."612. Within the large beams aand 4)). where both the envelope and the entire outflow are present. outflow emission most likely dominates.," Within the large beams and ), where both the envelope and the entire outflow are present, outflow emission most likely dominates."613 Observations and subsequent modelling of the more highly excited HO lines with ISO-LWS were unable to distinguish between an origin in shocks or an infalling envelope222)., Observations and subsequent modelling of the more highly excited $_2$ O lines with ISO-LWS were unable to distinguish between an origin in shocks or an infalling envelope.614. -HIFI has a much higher sensitivity. higher spectral resolution. and smaller beam than previous space-based missions. thus is perfectly suited. to addressing this question.," -HIFI has a much higher sensitivity, higher spectral resolution, and smaller beam than previous space-based missions, thus is perfectly suited to addressing this question."615 Complementary CO data presented by üre used to constrain the role of the envelope and determine outflow temperatures and densities., Complementary CO data presented by are used to constrain the role of the envelope and determine outflow temperatures and densities.616 NGC1333 is a well-studied region of clustered. low-mass star formation at a distance of 235 pe(?).," NGC1333 is a well-studied region of clustered, low-mass star formation at a distance of 235 pc."617. In particular. the three deeply embedded. low-mass class 0 objects IRAS2A. IRASAA. and IRASA4B have been observed extensively with ground-based submillimetre telescopes and interferometers 22).," In particular, the three deeply embedded, low-mass class 0 objects IRAS2A, IRAS4A, and IRAS4B have been observed extensively with ground-based submillimetre telescopes and interferometers ."618. All sources have strong outflows extending over aremin scales (>15 0000 AU)., All sources have strong outflows extending over arcmin scales $>$ 000 AU).619 Both IRAS4A and 4B consist of multiple protostars?)., Both IRAS4A and 4B consist of multiple protostars.620. Because of the similarities between the three sources in terms of luminosity (20. 5.5. and 3.8 L.). envelope mass and presumably also age. they provide ideal grounds for comparing YSOs in the same region.," Because of the similarities between the three sources in terms of luminosity (20, 5.8, and 3.8 $L_\odot$ ), envelope mass and presumably also age, they provide ideal grounds for comparing YSOs in the same region."621 Three sources in NGCI333. IRAS2A. ΠΑΡΑ. and IRASAB. were observed with HIFI on on March 15. 2010 in dual beam switch mode in bands |. 3. 4. and 5 with a nod of3'.," Three sources in NGC1333, IRAS2A, IRAS4A, and IRAS4B, were observed with HIFI on on March 3--15, 2010 in dual beam switch mode in bands 1, 3, 4, and 5 with a nod of."622". Observations detected several transitions of H:O and HO in the range E,,/Kg ~50-250 K (Table 2. in the online appendix).", Observations detected several transitions of $_2$ O and $_2^{18}$ O in the range $E_{\rm u}/k_{\rm B}$$\approx$ 50–250 K (Table \ref{tab:h2o_line} in the online appendix).623 Diffraction-linited beam sizes were in the range ((4500-9500 AU)., Diffraction-limited beam sizes were in the range (4500–9500 AU).624" In general. the calibration is expected to be accurate to ~20% and the pointing to -2"".."," In general, the calibration is expected to be accurate to $\sim$ and the pointing to $\sim$."625 Data were reduced with HIPE 3.0., Data were reduced with HIPE 3.0.626 À main-beam efficiency of 0.74 was used throughout., A main-beam efficiency of 0.74 was used throughout.627 Subsequent analystSs was performed in CLASS., Subsequent analysis was performed in CLASS.628 The rms was in the range 3-150 nK in 0.5 ss! bins., The rms was in the range 3–150 mK in 0.5 $^{-1}$ bins.629" Linear baselines were subtracted from all spectra. except around 750 GHz (corresponding to the H»O 2,,-2os transition) where higher-order polynomials are required."," Linear baselines were subtracted from all spectra, except around 750 GHz (corresponding to the $_2$ O $_{11}$ $_{02}$ transition) where higher-order polynomials are required."630 A difference 1 ms was always seen between the H- and V-polarizations. with the rms in the H-polarization being lower.," A difference in rms was always seen between the H- and V-polarizations, with the rms in the H-polarization being lower."631 In cases where the difference exceededand qualitative differences appear 1 the line profile. the V-polarization was discarded. otherwise the spectra were averaged.," In cases where the difference exceededand qualitative differences appear in the line profile, the V-polarization was discarded, otherwise the spectra were averaged."632is then considered TPAGB enriched.,is then considered TPAGB enriched.633 It should be noted that regardless of the orbital period. TPAGB enrichment during the second mass transfer episode is negligible in all systems.," It should be noted that regardless of the orbital period, TPAGB enrichment during the second mass transfer episode is negligible in all systems."634 The second part of Table 1 revisits the case A/Br of systems below 75 days. but now under the assumption that stable mass transfer occurs semi-conservatively (6= 0.5).," The second part of Table 1 revisits the case A/Br of systems below 75 days, but now under the assumption that stable mass transfer occurs semi-conservatively $\beta=0.5$ )."635 In that case. due to the angular momentum loss. systems below 3.8 days will merge.," In that case, due to the angular momentum loss, systems below 3.8 days will merge."636 Obviously. the non-conservatism also means that matter is ejected beforehand. 0.3 to 2.0 M. of which 0.1 to 0.5 M. helium.," Obviously, the non-conservatism also means that matter is ejected beforehand, 0.3 to 2.0 $_{\odot}$ of which 0.1 to 0.5 $_{\odot}$ helium."637 Above 3.8 days. the system survives and this amount increases to 2.3 M; of which 0.7 M.helium. plus 0.5 to 2.7 M. of which 0.2 to 0.8 Μι helium during the mass transfer episode of the secondary.," Above 3.8 days, the system survives and this amount increases to 2.3 $_{\odot}$ of which 0.7 $_{\odot}$helium, plus 0.5 to 2.7 $_{\odot}$ of which 0.2 to 0.8 $_{\odot}$ helium during the mass transfer episode of the secondary."638 This episode results in a merger and happens between 92 and 102 Myr. depending on whether accretion ϱuring the first mass transfer occurs through direct impact or an ο”nsecretion disk (the latter causing accretion induced full mixing.σι rejuvinating the secondary and thus delaying its evolution).," This episode results in a merger and happens between 92 and 102 Myr, depending on whether accretion during the first mass transfer occurs through direct impact or an accretion disk (the latter causing accretion induced full mixing, rejuvinating the secondary and thus delaying its evolution)."639 The last two parts of Table | again focus on the period range where mass transfer is stable (with 8=1 in part 3 and p=0.5 in part 4). but now for a companion with an initial mass of 3.6 M...," The last two parts of Table 1 again focus on the period range where mass transfer is stable (with $\beta=1$ in part 3 and $\beta=0.5$ in part 4), but now for a companion with an initial mass of 3.6 $_{\odot}$."640 In the conservative case. 1o matter Is lost during the first mass transfer phase. but depending on the initial orbital period. the second mass transfer phase is either a case B merger. a survivable case B. or a survivable case C. Here. more matter Is lost in the case B than in the case C as the core has grown larger in the latter case. and there is no possibility for case BB towards a WD to compensate for this.," In the conservative case, no matter is lost during the first mass transfer phase, but depending on the initial orbital period, the second mass transfer phase is either a case B merger, a survivable case B, or a survivable case C. Here, more matter is lost in the case B than in the case C as the core has grown larger in the latter case, and there is no possibility for case BB towards a WD to compensate for this."641 In the semi-conservative case. systems below 11 days will merge during the first mass transfer phase. while those above will survive and cause a second (merging) case A or B mass transfer phase.," In the semi-conservative case, systems below 11 days will merge during the first mass transfer phase, while those above will survive and cause a second (merging) case A or B mass transfer phase."642 Table 2 summarizes our population synthesis simulations for a population consisting of intermediate mass close binaries with primary mass Mj. between 3 M and 10 M. and initial chemical composotion (X... Y. Z) = (0.26. 0.24. 0.0001).," Table 2 summarizes our population synthesis simulations for a population consisting of intermediate mass close binaries with primary mass $M_1$ between 3 $M_{\odot}$ and 10 $M_{\odot}$ and initial chemical composotion (X, Y, Z) = (0.26, 0.24, 0.0001)."643 The evolution of every binary in the population code is followed as illustrated in Table I., The evolution of every binary in the population code is followed as illustrated in Table 1.644 Our simulations lead to the following conclusions., Our simulations lead to the following conclusions.645 Figure 2 shows the ejecta of a population of intermediate mass close binaries as a function of time since starburst., Figure 2 shows the ejecta of a population of intermediate mass close binaries as a function of time since starburst.646 It includes the total ejected mass. the ejected mass in helium (both primordial and newly synthesized) and the ejected mass of TPAGB-enriched matter.," It includes the total ejected mass, the ejected mass in helium (both primordial and newly synthesized) and the ejected mass of TPAGB-enriched matter."647 It has been calculated using the models with a flat g-distribution. B= 1. at=| and 0.1. y=2.92.3 (meaning mass loss through £2).," It has been calculated using the models with a flat $q$ -distribution, $\beta = 1$ , $\alpha\lambda = 1$ and 0.1, $\eta = 2.3$ (meaning mass loss through $L_2$ )."648 Note that with this, Note that with this649that the PLE model describes the data adequately.,that the PLE model describes the data adequately.650 This rejection of the PLIZ model is not found when testing case A. a vellection of the fact that it is the high luminosity. quasars which are responsible for the effect.," This rejection of the PLE model is not found when testing case A, a reflection of the fact that it is the high luminosity quasars which are responsible for the effect."651 This is not surprising: the PLE model was developed to fit the fainter AAT data. which we continue to use in this analysis. plus the brighter data of Schmidt Green (1983). which we have previously argued. is significantly. incomplete (CioldschimicdtaL. 1992) and which we have replaced by the Edinburgh quasar survey.," This is not surprising: the PLE model was developed to fit the fainter AAT data, which we continue to use in this analysis, plus the brighter data of Schmidt Green (1983), which we have previously argued is significantly incomplete (Goldschmidt, 1992) and which we have replaced by the Edinburgh quasar survey."652 We should therefore expect to see the most significant dillerences between the model and the IExlinburgh data., We should therefore expect to see the most significant differences between the model and the Edinburgh data.653 We can extend our analysis to test whether the data can be fitted byany evolving power-law model in which the power-law index remains constant., We can extend our analysis to test whether the data can be fitted by evolving power-law model in which the power-law index remains constant.654 Both the PLE mocel. at magnitudes more luminous than the BSP break. and the Llawkins Vérron model are examples of this class of model.," Both the PLE model, at magnitudes more luminous than the BSP break, and the Hawkins Vérron model are examples of this class of model."655 In this section we fit a single power-law model to the data in cach redshift slice. but only at. luminositics higher than the BSP break luminosity.," In this section we fit a single power-law model to the data in each redshift slice, but only at luminosities higher than the BSP break luminosity."656 We then test the null hypothesis that these bright-end power-law indices in each redshift slice have the same value., We then test the null hypothesis that these bright-end power-law indices in each redshift slice have the same value.657 The best-lit values for the indices are calculated: using niaximum likelihood assuming a single power-law fit to the data more luminous than the BSP break in each recdshift slice. where ἆ is the estimate of the index of the power law.," The best-fit values for the indices are calculated using maximum likelihood assuming a single power-law fit to the data more luminous than the BSP break in each redshift slice, where $\hat{\alpha}$ is the estimate of the index of the power law."658" ‘The faintest absolute magnitude in each redshift slice used to fit the model to the data is calculated by taking a comoving space density p=10""!Mpe (qo=0.5) and [inding the corresponding absolute magnitude in the DSP model at different redshifts.", The faintest absolute magnitude in each redshift slice used to fit the model to the data is calculated by taking a comoving space density $\rho=10^{-6.4} {\rm Mpc}^{-3}$ $q_0 = 0.5$ ) and finding the corresponding absolute magnitude in the BSP model at different redshifts.659 Under the null hypothesis this should give a constant index for all redshifts., Under the null hypothesis this should give a constant index for all redshifts.660 Phe answer should not be overly dependent on the value of p chosen. although too low a value will result in too little relevant data being used to fit he mocel. thereby. reducing the statistical significance.," The answer should not be overly dependent on the value of $\rho$ chosen, although too low a value will result in too little relevant data being used to fit the model, thereby reducing the statistical significance."661 Too ugh a value will result in some of the data from the flat part of the luminosity function being used. again underestimating he true significance.," Too high a value will result in some of the data from the flat part of the luminosity function being used, again underestimating the true significance."662 Errors on 4 are caleulated by assuming a X7 distribution or SSao=2og(LíLuaus) where L is the likelihood unction., Errors on $\hat{\alpha}$ are calculated by assuming a $\chi^{2}$ distribution for $S-S_{max}= -2 log(L/L_{max})$ where $L$ is the likelihood function.663 For quy=0.5 the best-fit power-law index increases rom à=2.7 in the lowest redshift slice to ἆ=41 in he highest slice. 33)., For $q_{0}=0.5$ the best-fit power-law index increases from $\hat{\alpha} = 2.7$ in the lowest redshift slice to $\hat{\alpha} = 4.1$ in the highest slice 3).664 Asingle. value for à is ruled out at a significance level of 0.14., Asingle value for $\hat{\alpha}$ is ruled out at a significance level of $0.1 \%$.665 For qu=0.1 the index increases [rom à=2.6 to ἆ=3.6 and a single value for à is unacceptable at a significance level of 2%., For $q_{0}=0.1$ the index increases from $\hat{\alpha}=2.6$ to $\hat{\alpha}=3.6$ and a single value for $\hat{\alpha}$ is unacceptable at a significance level of $2\%$.666 The evidence presented in this section. shows clearly that the high-Iuminosity part of the luminosity function does evolve according to the expectations of pure luminosity evolution., The evidence presented in this section shows clearly that the high-luminosity part of the luminosity function does evolve according to the expectations of pure luminosity evolution.667 The slope of the luminosity function at 26 displays significant steepening with redshift., The slope of the luminosity function at $M_B \la -26$ displays significant steepening with redshift.668 The analysis in the previous section showed that the power-law index changed with redshift for quasars more Iuminous than the break., The analysis in the previous section showed that the power-law index changed with redshift for quasars more luminous than the break.669 We have also previously remarked that there appears to be no evidence for a break in the luminosity function at low redshift., We have also previously remarked that there appears to be no evidence for a break in the luminosity function at low redshift.670 In this section we fit a single power- model to απο quasars brighter than Mg=23 in the lowest redshift bin with 0.3<2.< 0.7:, In this section we fit a single power-law model to the quasars brighter than $M_{B}=-23$ in the lowest redshift bin with $0.3 \le z \le 0.7$ ;671In the original self-similar CDAF. the viscous heating is balanced by the convective enerev (ransporl.,"In the original self-similar CDAF, the viscous heating is balanced by the convective energy transport."672 However. radiative heating can be dominant in some region of (he flow.," However, radiative heating can be dominant in some region of the flow."673 The importance of each heating aud cooling process is measured by the corresponding timescale., The importance of each heating and cooling process is measured by the corresponding timescale.674 The radiative heating timescale is given by the ratio between the internal οποιον of the electron.earepsilon... and the Compton heating rate. The radiative cooling timescale is similarly given bv The timescale for convective enerev (ransport al a certain. two-dimensional position is nol trivial to evaluate because (he two-dimensional velocity profile including the convective motion is needed.," The radiative heating timescale is given by the ratio between the internal energy of the electron, and the Compton heating rate, The radiative cooling timescale is similarly given by The timescale for convective energy transport at a certain two-dimensional position is not trivial to evaluate because the two-dimensional velocity profile including the convective motion is needed."675 Although two-dimensional numerical simulations (Stone. Pringle. Begelman 1999: lovwmenshchey. Abramowicz. Narayan 2000: leumenshehey Abramowicz 2000) suggest verv limited flow motion near the pole as in self-similar two-dimensional ADAF (CNaravan Yi 1995). the (vo-dimensional flow motion has vet to be expressed in simple analvtie form.," Although two-dimensional numerical simulations (Stone, Pringle, Begelman 1999; Igumenshchev, Abramowicz, Narayan 2000; Igumenshchev Abramowicz 2000) suggest very limited flow motion near the pole as in self-similar two-dimensional ADAF (Narayan Yi 1995), the two-dimensional flow motion has yet to be expressed in simple analytic form."676" IHere. we adopt the height-averaged convective energy flux given by NIA. where is the convection coefficient analogous to the usual Shakura Sunyaev)..cc, the isothermal sound speed. the Ixeplerian angular velocity. and the entropy of the flow."," Here, we adopt the height-averaged convective energy flux given by NIA, where is the convection coefficient analogous to the usual Shakura Sunyaev, the isothermal sound speed, the Keplerian angular velocity, and the entropy of the flow."677" In a sell-sinmilar solution that is marginally stable to convection. a, 244 is expected when 0.05 (NIA)."," In a self-similar solution that is marginally stable to convection, _c = 3 is expected when 0.05 (NIA)."678 The cooling rate per volume cue to this convective enerev flux is then Corresponding timescale is now, The cooling rate per volume due to this convective energy flux is then Corresponding timescale is now679burst 1. 3. and d: and four detectors during burst 2).,"burst 1, 3, and 4; and four detectors during burst 2)."680 We calculated. one Fourier power spectrum of the whole burst for each burst. anc power spectra of 1. 2. aud Ls intervals covering the whole duration of the bursts.," We calculated one Fourier power spectrum of the whole burst for each burst, and power spectra of 1, 2, and 4-s intervals covering the whole duration of the bursts."681 In the case of the 1. 2. and ds power spectra. to mucrease the sensitivity to trausieut oscillations lasting less than the time leneth of the power spectra. we set the beeimuime of the interval for which we calculated the power spectrum 1.125. 0.250. aud. 0.5 s respectively. after the beeiuniug of the previous interval (the power spectra were therefore rot iudependoenut).," In the case of the 1, 2, and 4-s power spectra, to increase the sensitivity to transient oscillations lasting less than the time length of the power spectra, we set the beginning of the interval for which we calculated the power spectrum 0.125, 0.250, and 0.5 s, respectively, after the beginning of the previous interval (the power spectra were therefore not independent)."682 For the brightest burst (no., For the brightest burst (no.683 3). we also calculated power spectra of 1) 0.25 aud 0.5 s intervals. shifting bv 1/32 aud 1/16 s at a time: aud 2) for >6 keV photons oulv.," 3), we also calculated power spectra of 1) 0.25 and 0.5 s intervals, shifting by 1/32 and 1/16 s at a time; and 2) for $>6$ keV photons only."684 We did not find auv significant oscillation iu the 1.0251021 Uz frequeney range in anv of the bursts., We did not find any significant oscillation in the $0.025 - 1024$ Hz frequency range in any of the bursts.685distribution within four analogous cylindrical volumes drawn from a variety of radii are presented.,distribution within four analogous cylindrical volumes drawn from a variety of radii are presented.686" At larger radii, the steepness of the profiles are quite different; the LSB simulation has a constant slope while the MW simulation shows a transition from a steep to a shallow density distribution."," At larger radii, the steepness of the profiles are quite different; the LSB simulation has a constant slope while the MW simulation shows a transition from a steep to a shallow density distribution."687" Thus the MW simulation cannot be characterized by a single exponential or sech? component in the vertical direction, as we show explicity in the following Section and Figure 5.."," Thus the MW simulation cannot be characterized by a single exponential or $^2$ component in the vertical direction, as we show explicity in the following Section and Figure \ref{f:density_profiles}."688 It is this double-component nature which first led to the identification of the thick disk (Gilmore&Reid1983);; we have thus shown that this feature need represent nothing more than internal evolution of the Milky Way., It is this double-component nature which first led to the identification of the thick disk \citep{Gilmore1983}; we have thus shown that this feature need represent nothing more than internal evolution of the Milky Way.689" In the following section, we compare SDSS observations with the MW simulation to demonstrate its usefulness as a model for understanding the Milky Way thick disk."," In the following section, we compare SDSS observations with the MW simulation to demonstrate its usefulness as a model for understanding the Milky Way thick disk."690" Here we study the stellar mass distribution, rotational velocity and metallicity as functions of distance from the Galactic plane, and galactocentric cylindrical radius, R."," Here we study the stellar mass distribution, rotational velocity and metallicity as functions of distance from the Galactic plane, $|z|$, and galactocentric cylindrical radius, $R$."691" We draw qualitative|z|, comparisons between the datasets by examining their mass weighted metallicity and kinematic distributions in this R-|z| space.", We draw qualitative comparisons between the datasets by examining their mass weighted metallicity and kinematic distributions in this $R$ $|z|$ space.692" The observed Milky Way disk is best fit by a 2-component model that is exponential both in the R and z directions Table 10, bias-corrected results, "," The observed Milky Way disk is best fit by a 2-component model that is exponential both in the $R$ and $z$ directions (see Table 10, bias-corrected results, J08)."693The top panel of (seeFigure 6 shows the mass weighted J08).density distribution of the entire MW simulation at its final timestep., The top panel of Figure \ref{f:color_maps} shows the mass weighted density distribution of the entire MW simulation at its final timestep.694 This distribution is in qualitative agreement in both the R and z directions with J08 for up to ~2.5 kpc above the disk's plane and ~15 kpc from the galactic center., This distribution is in qualitative agreement in both the $R$ and $z$ directions with J08 for up to $\sim$$2.5$ $\kpc$ above the disk's plane and $\sim$$15$ $\kpc$ from the galactic center.695for incorporation into phototonization codes such as Cloudy. which can be used to robustly determine correction factors for the abundances of unobserved tons. enabling much more accurate abundance determinations of n--capture elements in ionized nebulae than previously possible.,"for incorporation into photoionization codes such as Cloudy, which can be used to robustly determine correction factors for the abundances of unobserved ions, enabling much more accurate abundance determinations of -capture elements in ionized nebulae than previously possible."696 Until our study. the photoionization and recombination properties of Se tons had received very little attention. aside from a handful of PI studies of neutral Se(???).. which Is a trace species in ionized nebulae such as PNe.," Until our study, the photoionization and recombination properties of Se ions had received very little attention, aside from a handful of PI studies of neutral Se\citep{manson79, gibson86, chen94}, which is a trace species in ionized nebulae such as PNe."697 We therefore utilize recently conducted experimental PI. cross-section measurements of Se tons (222?) às a comparison and benchmark to our calculations.," We therefore utilize recently conducted experimental PI cross-section measurements of Se ions \citep{esteves09, sterling11, esteves10, esteves11b} as a comparison and benchmark to our calculations."698 We estimate the uncertainties in our computed data by utilizing three different. configuration-interaction (CI) expansions for each ion. and test the sensitivity of our results to various internal parameters of the AUTOSTRUCTURE code.," We estimate the uncertainties in our computed data by utilizing three different configuration-interaction (CI) expansions for each ion, and test the sensitivity of our results to various internal parameters of the AUTOSTRUCTURE code."699 Upon completion of our theoretical study of Se. Kr. and Xe ions we will use Monte Carlo simulations with photoionization codes to reveal the effects of these atomic data uncertainties on nebular abundance determinations. illustrating which systems and atomic processes require further theoretical and/or experimental analysis.," Upon completion of our theoretical study of Se, Kr, and Xe ions we will use Monte Carlo simulations with photoionization codes to reveal the effects of these atomic data uncertainties on nebular abundance determinations, illustrating which systems and atomic processes require further theoretical and/or experimental analysis."700 This paper is organized as follows: in Sect. ??..," This paper is organized as follows: in Sect. \ref{calcs},"701 we provide details of our calculations. and in Sect.," we provide details of our calculations, and in Sect."702 2? we present our computed PI cross sections and RR and DR rate coefficients along with their associated uncertainties., \ref{results} we present our computed PI cross sections and RR and DR rate coefficients along with their associated uncertainties.703 Finally. in Sect.," Finally, in Sect."704 ?? we summarize our results and provide concluding remarks., \ref{summ} we summarize our results and provide concluding remarks.705" We have computed the electronic structure and MCBP distorted-wave PI cross sections for Se""-Se with the AUTOSTRUCTURE code (??)."," We have computed the electronic structure and MCBP distorted-wave PI cross sections for $^0$ $^{5+}$ with the AUTOSTRUCTURE code \citep{badnell86, badnell97}."706 Higher charge states were not considered since they are negligibly populated in. PNe. whose central stars are not sufficiently hot (generally<2x10°K:??). to significantly ionize species with ionizatior thresholds greater than 100 eV. RR and DR rate coetlficients were determined from the direct and resonant portions of the PI cross sections. respectively. using detailed balance.," Higher charge states were not considered since they are negligibly populated in PNe, whose central stars are not sufficiently hot \citep[generally $<2\times10^5$~K;][]{napiwot99, stang02} to significantly ionize species with ionization thresholds greater than 100 eV. RR and DR rate coefficients were determined from the direct and resonant portions of the PI cross sections, respectively, using detailed balance."707 We usec the independent processes approximation (?). to treat RR anc DR separately., We used the independent processes approximation \citep{pindzola92} to treat RR and DR separately.708 The theoretical background of such calculations have beer provided elsewhere (??).. and we refer the reader to those references for a full discussion.," The theoretical background of such calculations have been provided elsewhere \citep{badnell03, badnell06b}, and we refer the reader to those references for a full discussion."709 In each of the following subsections. we discuss details of calculations for each process considered.," In each of the following subsections, we discuss details of calculations for each process considered."710" To test the sensitivity of our results to the adopted CI expansions. we constructed three different CI expansions for each Se ton (designated ""small “medium.” and large"")."," To test the sensitivity of our results to the adopted CI expansions, we constructed three different CI expansions for each Se ion (designated “small,” “medium,” and “large”)."711 The medium configuration sets provide the best compromise between accuracy and computational expense. and are adopted as the basis for our PI.RR. and DR results.," The medium configuration sets provide the best compromise between accuracy and computational expense, and are adopted as the basis for our PI,RR, and DR results."712 The CI expansions for each ion are listed in Table 4., The CI expansions for each ion are listed in Table \ref{ciexp}.713 In all cases. the electronic. structure was computed with Thomas-Fermi-Dirac-Amaldi model potentials. adopting intermediate-coupling and &-averaged relativistic wavefunctions (?)..," In all cases, the electronic structure was computed with Thomas-Fermi-Dirac-Amaldi model potentials, adopting intermediate-coupling and $\kappa$ -averaged relativistic wavefunctions \citep{cowan76}."714 We Schmidt orthogonalized the «-averaged orbitals 1n. our calculations. but compared our PI and RR results with those obtained without foreing the radial orbitals to be orthogonal.," We Schmidt orthogonalized the $\kappa$ -averaged orbitals in our calculations, but compared our PI and RR results with those obtained without forcing the radial orbitals to be orthogonal."715 The average of the LS term energies were optimizec by varying the orbital radial scaling parameters so as to best reproduce experimental energy levels and ionizatior potentials. from. NIST (?).., The average of the LS term energies were optimized by varying the orbital radial scaling parameters so as to best reproduce experimental energy levels and ionization potentials from NIST \citep{NIST}.716 The radial sealing parameters adopted for each CI expansion ofeach ton are given m Table 5.., The radial scaling parameters adopted for each CI expansion ofeach ion are given in Table \ref{lambdas}.717 A comparison of selected calculated and experimental energies is shown in Table 6.., A comparison of selected calculated and experimental energies is shown in Table \ref{ecomp}.718 We also compare our Einstein A-coefficients with those available in the literature in Table 8.., We also compare our Einstein A-coefficients with those available in the literature in Table \ref{acomp}.719 The structure computed with the medium configuration sets generally reproduce experimental energies and ionization potentials to a good degree of accuracy. lending credence to our use of orthogonalized radial orbitals.," The structure computed with the medium configuration sets generally reproduce experimental energies and ionization potentials to a good degree of accuracy, lending credence to our use of orthogonalized radial orbitals."720" The ionization potentials are within of the NIST values for all ions with the exception of Se"". which is a very challenging system to model but fortunately is negligibly populated in ionized nebulae such as PNe."," The ionization potentials are within of the NIST values for all ions with the exception of $^0$, which is a very challenging system to model but fortunately is negligibly populated in ionized nebulae such as PNe."721 The energies of most levels tabulated by NIST are reproduced to within for Se—Se. and to within for the neutral case.," The energies of most levels tabulated by NIST are reproduced to within for $^+$ $^{6+}$, and to within for the neutral case."722 À few exceptions occur. particularly for levels within the ground configuration. where the discrepancies with experiment are as large as for Se 4s 44p? Di;;.," A few exceptions occur, particularly for levels within the ground configuration, where the discrepancies with experiment are as large as for $^+$ $^2$ $^3$ $^2$ $_{3/2}$."723 The accuracies of the level energiesaffects those of the Einstein A-coefficients Àj; we computed., The accuracies of the level energiesaffects those of the Einstein A-coefficients $A_{ij}$ we computed.724 For forbidden transitions within the ground configurations. the main comparisons are the theoretical studies of Biémmont Hansen in the 1980s (see Table 8 for references). who utilized the Hartree-Fock with relativistic corrections (HER) and Harteeplus-statistical-exchange (HXR) methods.," For forbidden transitions within the ground configurations, the main comparisons are the theoretical studies of Biémmont Hansen in the 1980s (see Table \ref{acomp} for references), who utilized the Hartree-Fock with relativistic corrections (HFR) and Hartee-plus-statistical-exchange (HXR) methods."725 In many cases. our computed Aj; agree with those previous studies to within50%.. but exceptions occur. particularly for the lowest-charge states.," In many cases, our computed $A_{ij}$ agree with those previous studies to within, but exceptions occur, particularly for the lowest-charge states."726 However. it is worth noting that our CI expansions generally are larger than those used by Biémmont Hansen. and henee include configuration mixing not incorporated in their calculations.," However, it is worth noting that our CI expansions generally are larger than those used by Biémmont Hansen, and hence include configuration mixing not incorporated in their calculations."727 MCBP distorted-wave PI cross sections. were computed in intermediate coupling for neutral Se and each of the first five ions., MCBP distorted-wave PI cross sections were computed in intermediate coupling for neutral Se and each of the first five ions.728 Cross sections were calculated for each level in the ground configurations. as thevast majority of ions in photoionized plasmas reside in the ground configuration.," Cross sections were calculated for each level in the ground configurations, as thevast majority of ions in photoionized plasmas reside in the ground configuration."729 The Cl expansions listed in Table + were used in these calculations. with radial scaling parameters from Table 5..," The CI expansions listed in Table \ref{ciexp} were used in these calculations, with radial scaling parameters from Table \ref{lambdas}. ."730 In the cases of Se-. Se. Se. and Se. experimental PI cross sections measured at the Advanced Light Source (ALS) synchrotron," In the cases of $^+$ $^{2+}$ , $^{3+}$ , and $^{5+}$ , experimental PI cross sections measured at the Advanced Light Source (ALS) synchrotron"731LIX Peg is a chwarl nova with an orbital period of 4.52 h and frequent outbursts (recurrence time ~ 30 d) taking it from a minimum V — 16.5 to maximum at V — 12.9.,HX Peg is a dwarf nova with an orbital period of 4.82 h and frequent outbursts (recurrence time $\sim$ 30 d) taking it from a minimum V $\sim$ 16.5 to maximum at V $\sim$ 12.9.732 We have observed. parts of three outbursts Clable 1)., We have observed parts of three outbursts (Table 1).733 Late in the July 2002 outburst there is à clear signal (run $6475) with a period of 112.3 (+ 1.1) s and mean amplitude 2.6 mmag. illustrated in Fig. 21..," Late in the July 2002 outburst there is a clear signal (run S6475) with a period of 112.3 $\pm$ 1.1) s and mean amplitude 2.6 mmag, illustrated in Fig. \ref{lc6475}."734 On the late rise of the October 2002. outburst (run S6584) a similar signal is seen in the first half of the run: L148 (4 0.7) s with an amplitude of 3.1 mmag., On the late rise of the October 2002 outburst (run S6584) a similar signal is seen in the first half of the run: 114.8 $\pm$ 0.7) s with an amplitude of 3.1 mmag.735 Just after maximum of the November 2002 outburst (run 86656) there is an obvious moclulation in the light curve (Fig. 22)), Just after maximum of the November 2002 outburst (run S6656) there is an obvious modulation in the light curve (Fig. \ref{lc6656}) )736 which the ET shows is the result of two oscillations. with periods 83.4 (— 0.5) s and 111.1 (= 1.0) s with amplitudes of 4.8 and 4.2 mmag. respectively.," which the FT shows is the result of two oscillations, with periods 83.4 $\pm$ 0.5) s and 111.1 $\pm$ 1.0) s with amplitudes of 4.8 and 4.2 mmag, respectively."7371 The beat 1period. between these oscillations is 335 s and there is a peak in the EE near this period., The beat period between these oscillations is 335 s and there is a peak in the FT near this period.738 The above properties. especially the similarity of the periods near 112 s despite a brightness range of ~ 2.7 mag (Table 1) suggests an IpDNO. in which the true period is ~s3s. andthe ~ 112 s is a reprocessed signal.," The above properties, especially the similarity of the periods near 112 s despite a brightness range of $\sim$ 2.7 mag (Table 1) suggests an lpDNO, in which the true period is $\sim$ 83 s, and the $\sim$ 112 s is a reprocessed signal."739 In none of the runs ciscussed so far is there any sien of a normal DNO., In none of the runs discussed so far is there any sign of a normal DNO.740 Llowever. runs 56646 and 86650. made at the maximum of the November 2002 outburst. have distinct signals at. 16.22 s ancl 16.39 s. with mean amplitudes of 1.7 mmag ancl 1.6 numag. respectively — seen in the PVs in Fig. 23..," However, runs S6646 and S6650, made at the maximum of the November 2002 outburst, have distinct signals at 16.22 s and 16.39 s, with mean amplitudes of 1.7 mmag and 1.6 mmag, respectively – seen in the FTs in Fig. \ref{fthxpeg}."741 Combine with the QPO observed a few davs later in the same outburs (sce above) we have 2 = 20.6., Combined with the QPO observed a few days later in the same outburst (see above) we have $R$ = 20.6.742 In addition to the suite of periodicities described above. LIX Peg frequently shows very large QPOs with periods over the range of 1400 1900 s. Lt was our early observations of these. combined with the ~ 112 s oscillations that we now recognise to be IpDNOs. that led us to combine these into the # value ~ 16 announced earlier (Warner Woucdt 2003).," In addition to the suite of periodicities described above, HX Peg frequently shows very large QPOs with periods over the range of 1400 – 1900 s. It was our early observations of these, combined with the $\sim$ 112 s oscillations that we now recognise to be lpDNOs, that led us to combine these into the $R$ value $\sim$ 16 announced earlier (Warner Woudt 2003)."743 We now consider that the  1800 s QPOs arise from a dilferent. physical mechanism. perhaps connected. to the rotation period of the outer edge of the accretion disc (see Section S.T of Warner 1995a)," We now consider that the $\sim$ 1800 s QPOs arise from a different physical mechanism, perhaps connected to the rotation period of the outer edge of the accretion disc (see Section 8.7 of Warner 1995a)."744 AAL CVn stars are very short orbital period. helium-transferring double degenerate CVs.," AM CVn stars are very short orbital period, helium-transferring double degenerate CVs."745 DNOs at 19 22 s ancl 26.2 s respectively have been. observed. in the high state of CR Boo (Patterson. private communication). ancl in AM CVn (Patterson et al.," DNOs at 19 – 22 s and 26.2 s respectively have been observed in the high state of CR Boo (Patterson, private communication) and in AM CVn (Patterson et al."746 1979. Patterson et al.," 1979, Patterson et al."747 1992)., 1992).748 In most of the AM CVn stars the dominant modulation is a non-sinusoidal superhump that has strong harmonics to 5th order or higher: this makes the detection of QPOs very cüllicult., In most of the AM CVn stars the dominant modulation is a non-sinusoidal superhump that has strong harmonics to 5th order or higher; this makes the detection of QPOs very difficult.749 Nevertheless. Patterson ct al. (," Nevertheless, Patterson et al. ("7502002) have observed,2002) have observed751effect of saturation.,effect of saturation.752" The fraction of IRDCs lying above this limit is4%,, very similar to the value estimated in PF09."," The fraction of IRDCs lying above this limit is, very similar to the value estimated in PF09."753" However, even for these saturated clouds only a small fraction of their area is above the saturation limit, only marginally affecting the averaged IRDC column density (and therefore any estimate of the cloud mass)."," However, even for these saturated clouds only a small fraction of their area is above the saturation limit, only marginally affecting the averaged IRDC column density (and therefore any estimate of the cloud mass)."754" But, the saturation has a much stronger effect on some fragments."," But, the saturation has a much stronger effect on some fragments."755" For this reason, in the analysis presented here IRDCs containing saturated pixels are considered, but fragments with saturated pixels are excluded."," For this reason, in the analysis presented here IRDCs containing saturated pixels are considered, but fragments with saturated pixels are excluded."756 Our estimated saturation limit is roughly twice as large as the one found by Vasyunina et al. (, Our estimated saturation limit is roughly twice as large as the one found by Vasyunina et al. (7572009) from millimetre emission in their study of particularly high column density IRDCs.,2009) from millimetre emission in their study of particularly high column density IRDCs.758 The discrepancy between the low absorption column densities Vasyunina et al., The discrepancy between the low absorption column densities Vasyunina et al.759" determined by assuming the minimum possible foreground emission, that due to the zodical light, and the high values they determined from millimeter dust continuum led Vasyunina et al."," determined by assuming the minimum possible foreground emission, that due to the zodical light, and the high values they determined from millimeter dust continuum led Vasyunina et al."760 to derive a relatively low saturation limit., to derive a relatively low saturation limit.761" However, the majority of their clouds do not in fact appear saturated as considerable substructure can be seen in the 8um extinction maps."," However, the majority of their clouds do not in fact appear saturated as considerable substructure can be seen in the $\mu$ m extinction maps."762 This study aims to statistically analyze the density and mass distributions of IRDCs and their fragments., This study aims to statistically analyze the density and mass distributions of IRDCs and their fragments.763 To derive such quantities we first need to know the angular size and column density distributions as measured on the column density maps constructed by PF09., To derive such quantities we first need to know the angular size and column density distributions as measured on the column density maps constructed by PF09.764" Figure 3 shows the distribution of the angular size and column density for the ~11,000 IRDCs and the ~50,000 fragments identified within them."," Figure \ref{NH2_distrib} shows the distribution of the angular size and column density for the $\sim11,000$ IRDCs and the $\sim 50,000$ fragments identified within them."765 Fragments with saturated pixels have been excluded (see Sec., Fragments with saturated pixels have been excluded (see Sec.766 2.1)., 2.1).767 In addition IRDCs which are not fragmented (~40% of the IRDC sample) have also been removed to maintain a clear definition of a fragment as a substructure within a cloud., In addition IRDCs which are not fragmented $\sim40\%$ of the IRDC sample) have also been removed to maintain a clear definition of a fragment as a substructure within a cloud.768" However, in practice keeping these single peak clouds has little effect on the results."," However, in practice keeping these single peak clouds has little effect on the results."769" It is important to note that the column densities we plot here are thesubstracted column densities, equivalent to the one obtained in theclipping option of the dendogram analysis of Rosolowsky et al. ("," It is important to note that the column densities we plot here are the column densities, equivalent to the one obtained in the option of the dendogram analysis of Rosolowsky et al. ("7702008).,2008).771" In the context of centrally concentrated structures, these column densities are the relevant ones when interested in the physical properties of the gas enclosed in a given radius."," In the context of centrally concentrated structures, these column densities are the relevant ones when interested in the physical properties of the gas enclosed in a given radius."772 Figure 3 clearly shows that the distributions are dominated by small structures of low column density, Figure \ref{NH2_distrib} clearly shows that the distributions are dominated by small structures of low column density.773" We can also clearly see the effect of incompleteness on the distributions with the decrease in the number of sources at low radius/column density, responsible for the formation of artificial peaks."," We can also clearly see the effect of incompleteness on the distributions with the decrease in the number of sources at low radius/column density, responsible for the formation of artificial peaks."774 The incompleteness in the sample and these distributions are discussed in Section 4.., The incompleteness in the sample and these distributions are discussed in Section \ref{sec:complete}.775" To calculate the density and mass of the clouds the distance of each IRDC is required, however this is not yet known for most of the 11,000 IRDCs."," To calculate the density and mass of the clouds the distance of each IRDC is required, however this is not yet known for most of the 11,000 IRDCs."776 For this analysis we have therefore adopted a statistical approach based on previous measurements of the distances to samples of IRDCs., For this analysis we have therefore adopted a statistical approach based on previous measurements of the distances to samples of IRDCs.777 Several studies have measured the distance, Several studies have measured the distance778the effects of a lew low-probabilitw detections of SDOs that are unstable. can explain the difference in the rate of objects leaving the scattered disk curing the first gigavear and the remainder of the simulation.,"the effects of a few low-probability detections of SDOs that are unstable, can explain the difference in the rate of objects leaving the scattered disk during the first gigayear and the remainder of the simulation."779 With (he above considerations. we conservatively report (he escape rate of SDOs as (1—2)x10 Ivy| where the former is the average over the last 1 Gyr and the latter is the average over the first | Gvr of the simulation.," With the above considerations, we conservatively report the escape rate of SDOs as $(1-2)\times10^{-10}$ $^{-1}$ where the former is the average over the last 1 Gyr and the latter is the average over the first 1 Gyr of the simulation."780 A small source of error in this estimate is introduced by our assumption (hat all the particles (hat had a close encounter wilh a planet will leave the scattered disk after their first such encounter., A small source of error in this estimate is introduced by our assumption that all the particles that had a close encounter with a planet will leave the scattered disk after their first such encounter.781 To estimate the magnitude of (his error. we note that ? found that only ~57€ of trans-Neptunian objects will persist on stable trans-Neptunian orbits alter their first Neptune encounter: therefore. (he error in our estimate due to this assumption is e5%. well within the range reported here.," To estimate the magnitude of this error, we note that \citet{duncan97} found that only $\sim5\%$ of trans-Neptunian objects will persist on stable trans-Neptunian orbits after their first Neptune encounter; therefore, the error in our estimate due to this assumption is $\sim5\%$, well within the range reported here."782 For comparison. our estimate of the SDO outfIux rate is nearly an order of magnitude larger than the 4xLOvy ! found by ο for the Classical Kuiper belt. confirming that the scattered. disk is a less stable population. than the Classical belt.," For comparison, our estimate of the SDO outflux rate is nearly an order of magnitude larger than the $4\times10^{-11}$ $^{-1}$ found by \citet{duncan95} for the Classical Kuiper belt, confirming that the scattered disk is a less stable population than the Classical belt."783 The rate for the scattered disk that we find here is an order of magnitude smaller (han the rate found bv ον a likely explanation for this difference is that those authors considered only initial orbits with perihelia interior to 37 AU. a sample that is overall less stable than one that includes the full range of perihelia representative of the scattered disk.," The rate for the scattered disk that we find here is an order of magnitude smaller than the rate found by \citet{emel04}; a likely explanation for this difference is that those authors considered only initial orbits with perihelia interior to 37 AU, a sample that is overall less stable than one that includes the full range of perihelia representative of the scattered disk."784 When only objects with initial q«37 AU in our simulation are considered. the fractional escape rate is (1.6—3.5)x10. Myrft which is a factor of ~3 smaller (han (he ? estimate.," When only objects with initial $q<37$ AU in our simulation are considered, the fractional escape rate is $(1.6-3.5)\times10^{-10}$ $^{-1}$, which is a factor of $\sim\!3$ smaller than the \citet{emel04} estimate."785 The rest of the discrepancy night be explained bv the difference in the number of observed SDOs contributing to each sample: based their population on 7 real SDOs. while our sample of real SDOs with q<37 AU is 43.," The rest of the discrepancy might be explained by the difference in the number of observed SDOs contributing to each sample; \citet{emel04} based their population on 7 real SDOs, while our sample of real SDOs with $q<37$ AU is 43."786 Our estimate is also smaller than that of 2.. who report that SDOs enter the Centaur population at a fractional rate of 5x10.10! |: the discrepancy here is likely due to the difference between their debiased semi-major axis distribution. which is sharply peaked αἱ smaller semi-major axes. and ours. which is more flat (see Section ??)).," Our estimate is also smaller than that of \citet{disisto07}, who report that SDOs enter the Centaur population at a fractional rate of $5\times10^{-10}$ $^{-1}$; the discrepancy here is likely due to the difference between their debiased semi-major axis distribution, which is sharply peaked at smaller semi-major axes, and ours, which is more flat (see Section \ref{ss:tp}) )."787 Our estimate is closer in agreement to the rate 2.7x10.I! vr.! obtained [rom the simulations of ?. (as reported by 7)): the latter was based on only 33 scattered disk.(vpe particles. compared with our sample ol 80.," Our estimate is closer in agreement to the rate $2.7\times10^{-10}$ $^{-1}$ obtained from the simulations of \citet{duncan97} (as reported by \citet{levison06}) ); the latter was based on only 33 scattered disk–type particles, compared with our sample of 80."788 Our estimate also agrees with the estimate of 2: however their work was based on 76 observed objects. only 49 of which meet the criteria for 5DOs as given in (liis paper.," Our estimate also agrees with the estimate of \citet{fernandez04}; however their work was based on 76 observed objects, only 49 of which meet the criteria for SDOs as given in this paper."789 The rate at which objects leave the scattered. disk can be used to estimate (he total number of SDOs by making (he assumption that the scattered. disk is in steady state wilh the JFC population., The rate at which objects leave the scattered disk can be used to estimate the total number of SDOs by making the assumption that the scattered disk is in steady state with the JFC population.790 Previous studies (??) have consistently shown," Previous studies \citep{levison97,fernandez04} have consistently shown"791position angle were first teutativelv proposed as possible large scale lobes (Martíetal. 20003).,position angle were first tentatively proposed as possible large scale lobes \cite{m2000}) ).792 Towever. no further evidence for a physical counection couk be found bevoud the mere ecometric aligunient.," However, no further evidence for a physical connection could be found beyond the mere geometric alignment."793" A closer search revealed later the existence of two possible ho PArot candidates (USC's) associated with Cyeuus. δν, tms sugeesting i analogy with Fanarofi-Riley type II (FR ID) radio HHealaxies (Martietal. 2005))."," A closer search revealed later the existence of two possible hot spot candidates (HSCs) associated with Cygnus X-3, thus suggesting an analogy with Fanaroff-Riley type II (FR II) radio galaxies \cite{m2005}) )."794" The apurvut hot spots were two faint radio sources wihn noitjiermal spectra aueular distances of n alc 1136 trou, Cyveuus. N-3.", The apparent hot spots were two faint radio sources with non-thermal spectra at angular distances of 07 and 36 from Cygnus X-3.795 The liue jomüug them was also within one deere O© the almost North-South YONIlon ange of the iuner arc-second radio jets (Martíetal. 20013)., The line joining them was also within one degre of the almost North-South position angle of the inner arc-second radio jets \cite{m2001}) ).796 Unfortunately. Allow up radio aud near ifrared. observations of both the IISC« and the Cyveuus N-3 nearby cuviromment did uot confinu the proposed ho spot nature and indicated that they were most Likely backernud or foreground objects (Alaitretal. 2006)).," Unfortunately, follow up radio and near infrared observations of both the HSCs and the Cygnus X-3 nearby environment did not confirm the proposed hot spot nature and indicated that they were most likely background or foreground objects \cite{m2006}) )."797 This fac left open again for Crenus X-3 the issue of seareune for sienatures of energv deposition from its relativistic jets iuo the ISM., This fact left open again for Cygnus X-3 the issue of searching for signatures of energy deposition from its relativistic jets into the ISM.798 Iu this context. the nudi purpose of this paper is to present new very deep radio nuages of Crenus 3 ancl its environnent obtained afer combining niulti-epoch archive iik survey data. together with our own observations.," In this context, the main purpose of this paper is to present new very deep radio images of Cygnus X-3 and its environment obtained after combining multi-epoch archive and survey data, together with our own observations."799 The resullug maps eable us to put very strong limits for aIv ONended or compact radio features that could be asrclater| to Cyeuus XN on a scale of a few pc., The resulting maps enable us to put very strong limits for any extended or compact radio features that could be associated to Cygnus X-3 on a scale of a few pc.800 We also reort on several exended radio features in the field with aypareit non-thermial spectra that were previously unknown., We also report on several extended radio features in the field with apparent non-thermal spectra that were previously unknown.801 TI1011 possible connection with the niücroquasar Is discussec from a skeptical point of view., Their possible connection with the microquasar is discussed from a skeptical point of view.802 Mos observations used im this paper come from archive data obtained with the Very Large Array (VLA) operated by the National Radio Astroioniv. Observatory (NRAO) in the USA., Most observations used in this paper come from archive data obtained with the Very Large Array (VLA) operated by the National Radio Astronomy Observatory (NRAO) in the USA.803 Dy similarity with the Cyvguus N-1 or Circiuus X-1 cases. the possible large scales radio features around Crems N-3 (if αν) are likev to be a few pc extended aud. therefore. with few arc-nuuute augular sizes.," By similarity with the Cygnus X-1 or Circinus X-1 cases, the possible large scales radio features around Cygnus X-3 (if any) are likely to be a few pc extended and, therefore, with few arc-minute angular sizes."804 The more compact D coufiguration of the VLA appears thus as fre best choice for our «ctection purposes at these angellay scales., The more compact D configuration of the VLA appears thus as the best choice for our detection purposes at these angular scales.805 Concerning waveleneth. we decided to use 6 cm data which provides a Full-Width Πα Ααπλάι (FWIIND primary beam of about 97.," Concerning wavelength, we decided to use 6 cm data which provides a Full-Width Half Maximum (FWHM) primary beam of about $9^{\prime}$."806" Thou. a 1 to 10 pe feature would be well covered wihi-—i the more seusitive, Inner part of he primary beam."," Then, a 1 to 10 pc feature would be well covered within the more sensitive, inner part of the primary beam."807" Surprisingly. the NRAO cata arcllive contains no VLA observations for Cveuus X-3? int 1C »ue D configuration of the array aid with more than a half hour of integration tie,"," Surprisingly, the NRAO data archive contains no VLA observations for Cygnus X-3 in the pure D configuration of the array and with more than a half hour of integration time."808 The archive data Lore closely imnatchiug «nur requirclcuts comes from the livDar DuC configuration., The archive data more closely matching our requirements comes from the hybrid DnC configuration.809 We also inchded one D configuration run for eulauced angular resoltion in addition to seusitivitv on arcauimute scales., We also included one B configuration run for enhanced angular resolution in addition to sensitivity on arc-minute scales.810 The xs of the observiig PULLS that we lewe colmbined iuto a single radio image is given in Table 1.., The log of the observing runs that we have combined into a single radio image is given in Table \ref{vlaobs}.811 The data were srocessedL using the ΑΠΣ software xickage. of NRAO following the standard procedures for coutiuuuni calibratio1 of interferometers., The data were processed using the AIPS software package of NRAO following the standard procedures for continuum calibration of interferometers.812 Moreover. theυ data were selt-calibrated in phase.," Moreover, the data were self-calibrated in phase."813 Due to the siguificaut Hiring variability of he Cveuus X-3? core during the observations the standard CLEAÀNiug methods could not © applied. directly., Due to the significant flaring variability of the Cygnus X-3 core during the observations the standard CLEANing methods could not be applied directly.814 Tustead. we proceeded following a nethod simular to tli outlined in et al. (," Instead, we proceeded following a method similar to that outlined in et al. ("8152000).,2000).816 Therefore. we had to πιbtract a time-variable point source roni he visibilitics iu thewer plane using differcut AIPS asks.," Therefore, we had to subtract a time-variable point source from the visibilities in the plane using different AIPS tasks."817 First of all. we split he data im blocks of less than uplitude variation.," First of all, we split the data in blocks of less than amplitude variation."818 Ex each block. we derived the flux density of the COipoient to be subtracted at the wicroquasar core position witli ask UVFIT.," For each block, we derived the flux density of the component to be subtracted at the microquasar core position with task UVFIT."819" Afterwawcls, we reuoved it using UVSUD aud all alinost variability free docks were finally reconibiued 1sing DBCON."," Afterwards, we removed it using UVSUB and all almost variability free blocks were finally recombined using DBCON."820 T1ο next step was to ¢λα]ne the data of the three xojects we are haudine., The next step was to combine the data of the three projects we are handling.821 At this point. the UVFIN task of AIPS had to be used to set a comnon phase ceuter to alow he appropriate couination oftje visibilities.," At this point, the UVFIX task of AIPS had to be used to set a common phase center to allow the appropriate combination of the visibilities."822 Finally. he DBCON task was 1sed again to meree all data scts iu Table 1 into a single «c file.," Finally, the DBCON task was used again to merge all data sets in Table \ref{obslog} into a single $uv$ file."823 This was performed with au appropriate woeleiting to cuhance the short baseli108 of DuC|D configuration that eive scusitivity to extened cuuission., This was performed with an appropriate weighting to enhance the short baselines of DnC+D configuration that give sensitivity to extended emission.824 A constai point source. with the average fAx density of Cyeuus X-3.H was finally added at the position of the subtracted core.," A constant point source, with the average flux density of Cygnus X-3, was finally added at the position of the subtracted core."825 We first mapped our nulti-coufiguration data without using the long baseline visibilites provided by the D configuration of the array., We first mapped our multi-configuration data without using the long baseline visibilites provided by the B configuration of the array.826 Fig., Fig.827 1 shows a conteUY plot of the field obtained using the IAIAGR task of AIPS with the ROBUST parameter set to five and, \ref{cd+d} shows a contour plot of the field obtained using the IMAGR task of AIPS with the ROBUST parameter set to five and828for ice.,for ice.829 Fig., Fig.830 1 shows their model and Equation (1)). 2010).," \ref{fig:me} shows their model and Equation \ref{eq:fragmass}) ).,"831 most of the laboratory experiucuts aud the lydrodvuamic nuuerical simulations of collisional disruption showed i. not to have a discontinuity at o=1 (Tousenctal.1991:Takagietal1981:Beng&Asphaneg 1999).," most of the laboratory experiments and the hydrodynamic numerical simulations of collisional disruption showed $m_{\rm e}$ not to have a discontinuity at $\phi = 1$ \citep{housen,takagi,benz99}."832. Therefore. Equation 619) includes. erosive colliious (a<1) more accurately.," Therefore, Equation \ref{eq:fragmass}) ) includes erosive collisions $\phi < 1$ ) more accurately."833 ↽∕∏∐∖↸⊳↥⋅↕↑↕↸⊳⋜↧↕↸∖∐↸∖↥⋅∶↴⋁∙↖↽∣⊋↻↕↴∖↴∶↴↜⊾↕↖↽↸∖∐↴⋝∙↖↽ where s and s are the radius and mass of a body. py is its density. aud G is the gravitatioual constant.," The critical energy $Q_{\rm D}^*$ is given by where $r$ and $m$ are the radius and mass of a body, $\rho_{\rm p}$ is its density, and $G$ is the gravitational constant."834" The first term on the right-hand side of Equation (5)) is dominant for rXtot 10 com. the secoud term describes Qi of &X coun, and the third term coutrols (Qj for the luger bodies."," The first term on the right-hand side of Equation \ref{eq:qd}) ) is dominant for $r \la83510^4$ $10^5$ cm, the second term describes $Q_{\rm D}^*$ of $r \la83610^7$ cm, and the third term controls $Q_{\rm D}^*$ for the larger bodies."837 Beng&Asphane(1999) performed the lyvdrocdvuamical sinuulations of collisional dispersion for 7=1 10% cm aud provided the values of Qui...Que. aud ἐν.," \citet{benz99} performed the hydrodynamical simulations of collisional dispersion for $r = 1$ $10^7\,$ cm and provided the values of $Q_{\rm 0s}, {\beta_{\rm s}}, Q_{\rm 0g}$, and ${\beta_{\rm g}}$."838 For rom lU'ecna Qj is purely determined by the eravitational binding enerex. Όσιο independent of material properties.," For $r \ga 10^7$ cm, $Q_{\rm D}^*$ is purely determined by the gravitational binding energy, being independent of material properties."839 The collisional simulation for eravitational agereeatesOUO vields Cu~10 (Stewart&Leinhardt2009).," The collisional simulation for gravitational aggregates yields $C_{\rm840gg} \sim 10$ \citep{stewart09}."841. Ouce a plauctary oenbrvo has erown larecr than the Moon. it acquires an atmosphere.," Once a planetary embryo has grown larger than the Moon, it acquires an atmosphere."842 It helps the accretion of plauetesimals or fragimoeuts outo an enmbrvo: small bodies are captured by the atinosplhiere of the embryo., It helps the accretion of planetesimals or fragments onto an embryo; small bodies are captured by the atmosphere of the embryo.843 Iuaba&Dsoma(2003) provided an analytical uodel for a density profile of the atinosphiere., \citet{inaba_ikoma03} provided an analytical model for a density profile of the atmosphere.844 We consider the atinosphere at a distance A. from an ubrvo center., We consider the atmosphere at a distance $R_{\rm e}$ from an embryo center.845 We asstune that 2. is nach smaller han that at the outer boundary of the atinosphiere id that its temperature is much higher than hat at the boundary., We assume that $R_{\rm e}$ is much smaller than that at the outer boundary of the atmosphere and that its temperature is much higher than that at the boundary.846 The atmospheric deusitv pa is then proportional to Ro? (Mizuno.1980:Stevenson 1982).," The atmospheric density $\rho_{\rm a}$ is then proportional to $R_{\rm e}^{-3}$ \citep{mizuno80,stevenson82}."847". Applving the temperature Zi. oxessure Py. and density pa, of the nebula in the disk: unidplane as those at the outer boundary of atmosphere. the density profile of the atmosphere around an embryo with mass A is given by where # is the opacity of the atmosphere aud asp ds the Stephanu-Boltzimanu constant."," Applying the temperature $T_{\rm neb}$, pressure $P_{\rm neb}$, and density $\rho_{\rm neb}$ of the nebula in the disk midplane as those at the outer boundary of atmosphere, the density profile of the atmosphere around an embryo with mass $M$ is given by where $\kappa$ is the opacity of the atmosphere and $\sigma_{\rm SB}$ is the Stephan-Boltzmann constant."848" The planetary luninosity L,. mainly comes from the accretion of bodies.", The planetary luminosity $L_{\rm e}$ mainly comes from the accretion of bodies.849 We approximate where Rois the enibrvo radius., We approximate where $R$ is the embryo radius.850 To validate the assumption of pax7. we will apply the complete model by Tnaba&Tsoma(2003). to our statistic sinulation and compare our analytical solutions with the statistical simulations iu Section l..," To validate the assumption of $\rho_{\rm a} \propto R_{\rm e}^{-3}$, we will apply the complete model by \citet{inaba_ikoma03} to our statistic simulation and compare our analytical solutions with the statistical simulations in Section \ref{sc:simulation}."851 When a body passes by a planetary enibryo with an atmosphere. the embryo can accrete the body without direct collision due to the atinosplere.," When a body passes by a planetary embryo with an atmosphere, the embryo can accrete the body without direct collision due to the atmosphere."852" The relative velocity between the body aud the embryo at infinity is determined by the eccentricity e of the small body: it is eiven by cee, with the Nheplerian velocity κ=VGALa AZ. beiug the mass of a central star."," The relative velocity between the body and the embryo at infinity is determined by the eccentricity $e$ of the small body; it is given by $e v_{\rm k}$ with the Keplerian velocity $v_{\rm k} =853\sqrt{GM_*/a}$ $M_*$ being the mass of a central star."854 The, The855 , 856We characterize the changes in the longitudinal photospheric magnetic fiekl during 38 and 39 M-class flares within 65° of disk-center using I-minute GONG magnetograms.,We characterize the changes in the longitudinal photospheric magnetic field during 38 X-class and 39 M-class flares within $65^{\circ}$ of disk-center using 1-minute GONG magnetograms.857 In all T7 cases we identilv al least one site in (he flaring active region where clear. permanent. stepwise field changes occurred.," In all 77 cases we identify at least one site in the flaring active region where clear, permanent, stepwise field changes occurred."858 The median duration of the field changes was about 15 minutes and was approximately equal for N-class and for M-class flares., The median duration of the field changes was about 15 minutes and was approximately equal for X-class and for M-class flares.859 The absolute values ol the field changes ranged from the detection limit of —10 C to as high as ~450 G in two exceptional cases., The absolute values of the field changes ranged from the detection limit of $\sim\!\!10$ G to as high as $\sim\!\!450$ G in two exceptional cases.860 The median value was 69 G. Field changes were significantly stronger for X-elass than lor M-class {flares ancl for limb flares than for clisk-center flares., The median value was 69 G. Field changes were significantly stronger for X-class than for M-class flares and for limb flares than for disk-center flares.861 Longitudinal field changes less than LOO G tended to decrease longitudinal field strengths. both close to disk-center and close to the limb. while field changes greater than 100 G showed no such pattern.," Longitudinal field changes less than 100 G tended to decrease longitudinal field strengths, both close to disk-center and close to the limb, while field changes greater than 100 G showed no such pattern."862 Likewise. longitudinal flux strengths tended to decrease during Hares.," Likewise, longitudinal flux strengths tended to decrease during flares."863 Flux changes. particularly net [lux changes near disk-center. correlated better (han local field changes with GOES peak X-ray Πας.," Flux changes, particularly net flux changes near disk-center, correlated better than local field changes with GOES peak X-ray flux."864 The strongest longitudinal field and fIux changes occurred in flares observed close to the limb., The strongest longitudinal field and flux changes occurred in flares observed close to the limb.865 We estimate the change of Lorentz force associated. with each flare and find that this is large enough in some cases {ο power seisnic waves., We estimate the change of Lorentz force associated with each flare and find that this is large enough in some cases to power seismic waves.866 We find (hat longitudinal field decreases would likely outinmber increases at all parts of the solar disk within 65° of disk-center. as in our observations. if photospheric fiekd (ilis increase curing flares as predicted by IIudson et al.," We find that longitudinal field decreases would likely outnumber increases at all parts of the solar disk within $65^{\circ}$ of disk-center, as in our observations, if photospheric field tilts increase during flares as predicted by Hudson et al."867 Solar flares are generally believed to be caused by strong. stressed. topologically complicated magnetic fields.," Solar flares are generally believed to be caused by strong, stressed, topologically complicated magnetic fields."868 The energy estimated (to power a solar [are can only come Irom the magnetic field and this field must be sulliciently stressed. to contain enough [ree energy {ο power the flare., The energy estimated to power a solar flare can only come from the magnetic field and this field must be sufficiently stressed to contain enough free energy to power the flare.869 The topology must be complicated enough to contain a magnetic null point for abrupt energy release to be possible (e.g. Priest Forbes 2000. Aschwanden 2004).," The topology must be complicated enough to contain a magnetic null point for abrupt energy release to be possible (e.g. Priest Forbes 2000, Aschwanden 2004)."870 Magnetic, Magnetic8711n this work. we focus on how well CO emission can trace the intrinsic CO column density.Neo... the ecolumn densityWH... as well as the total column density of hydrogen nuclei.γαι.,"In this work, we focus on how well CO emission can trace the intrinsic CO column density, the column density, as well as the total column density of hydrogen nuclei,."872 The intrinsic column densities can be easily calculated. cirecthy from the AID simulation by integrating the CO. πο hydrogen nuclei volume densities along a given axis.," The intrinsic column densities can be easily calculated directly from the MHD simulation by integrating the CO, and hydrogen nuclei volume densities along a given axis."873" Accordingly. for the radiative transfer caleulation we orient the simulation cube such that the €CO line is ""observed? along the same axis for which the column censitics are computed (c.g. along the z-axis)."," Accordingly, for the radiative transfer calculation we orient the simulation cube such that the CO line is “observed” along the same axis for which the column densities are computed (e.g. along the $\hat{z}$ -axis)."874 We have verified that our results are not sensitive to the choice of orientation., We have verified that our results are not sensitive to the choice of orientation.875 Jesdes the viewing geometry. the only other user defined parameter. required for the radiation transfer calculations is the microturbulent velocity erin see Eqns. 6--1]].," Besides the viewing geometry, the only other user defined parameter required for the radiation transfer calculations is the microturbulent velocity $v_{mtrb}$ [see Eqns. \ref{lineprof}- \ref{broaden}] ]."876 Extrapolating the observed linewiclth-size relationship (ος.7) down to the resolution of the ALD simulations  Ll pe. appropriate microturbulent velocities are in the range ~ 0.2 - 0.7 |.," Extrapolating the observed linewidth-size relationship \citep[e.g.][]{Larson81} down to the resolution of the MHD simulations $\sim$ 0.1 pc, appropriate microturbulent velocities are in the range $\sim$ 0.2 - 0.7 ."877 We have explored. this range in ο and have Found that the results are insensitive to the particular choice of mre.," We have explored this range in $v_{mtrb}$, and have found that the results are insensitive to the particular choice of $v_{mtrb}$."878 In the analysis presented here. Crip issct to 0.5," In the analysis presented here, $v_{mtrb}$ is set to 0.5."879" In our discussion. we will sometimes express aas an extinction«Εν, in order to allow for direct comparison with observational analyses."," In our discussion, we will sometimes express as an extinction, in order to allow for direct comparison with observational analyses."880 Observers are often. required to employ indirect. measures ofNia. sinceI»... à major constituent of iin MCS.. is) cliflieult το observe. directly.," Observers are often required to employ indirect measures of, since, a major constituent of in MCs, is difficult to observe directly."881 Extinction measurements provide estimates of the total amount of dust along the line of sight., Extinction measurements provide estimates of the total amount of dust along the line of sight.882 Using a 7reddening law.” (and an assumption for the dust-to-gas ratio). the total gaseous column follows cürectlv from the amount of extinction citepBohlinctalTs..," Using a “reddening law,” (and an assumption for the dust-to-gas ratio), the total gaseous column follows directly from the amount of extinction \\citep{Bohlinetal78}."883 In many Galactic molecular clouds. nearly all the hydrogen is molecular. so iis directly. proportional tozl.," In many Galactic molecular clouds, nearly all the hydrogen is molecular, so is directly proportional to."884.. In. other environments. however. such as those with lower metallicity or higher UV fields (e.g. Models. n300-Z01. and n300-UV2). there may be significant amounts of atomic hydrogen. soNy... and hence the extinction. is dependent on both the molecular and atomic column densities.," In other environments, however, such as those with lower metallicity or higher UV fields (e.g. Models n300-Z01 and n300-UV2), there may be significant amounts of atomic hydrogen, so, and hence the extinction, is dependent on both the molecular and atomic column densities."885 To allow for straightforward comparison between models and. observations. we use a simple conversion between aanddc where Z is the metallicity of the gas.," To allow for straightforward comparison between models and observations, we use a simple conversion between and: where ${\rm Z}$ is the metallicity of the gas."886 Phus. in. the comparison of CO emission with intrinsic cloud. properties. any discussion involving ccan be directly translated into total column clensity.," Thus, in the comparison of CO emission with intrinsic cloud properties, any discussion involving can be directly translated into total column density."887 The radiative transfer calculations. produce. spectral (position-position-velocity. or PPV) cubes of the CO 0) line.," The radiative transfer calculations produce spectral (position-position-velocity, or PPV) cubes of the CO (J=1-0) line."888" Phe 3D cube indicates the intensity. Z, in a given frequeney or velocity channel of width de at cach 2D position.", The 3D cube indicates the intensity $I_\nu$ in a given frequency or velocity channel of width $dv$ at each 2D position.889" We choose a sulliciently large range in frequencies such that all (linc-of-sight) velocities in the simulation are detected. so that all emission from the model MC is ""observed."""," We choose a sufficiently large range in frequencies such that all (line-of-sight) velocities in the simulation are detected, so that all emission from the model MC is “observed.”"890 In the comparison of CO) intensities. with intrinsic column densities. as well as the analysis of the [factor. the quantity of interest is the velocity integrated intensity. which is simply the PPV cube integrated over its velocity axis.," In the comparison of CO intensities with intrinsic column densities, as well as the analysis of the factor, the quantity of interest is the velocity integrated intensity, which is simply the PPV cube integrated over its velocity axis."891" The intensity £,. which has units of erg s7 > t sterto ean be expressed as the Planck function evaluated at a “brightness temperature” Zg. (1)."," The intensity $I_\nu$, which has units of erg $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$, can be expressed as the Planck function evaluated at a “brightness temperature” $T_B$, $B_\nu(T_B)$."892" Since the CO (J=1-0) line is located in the Ravleigh-Jeans part of the spectrum. Z,xYe."," Since the CO (J=1-0) line is located in the Rayleigh-Jeans part of the spectrum, $I_\nu \propto893T_B$."894 The intensity of CO line emission is thus often conveyed in Zp units., The intensity of CO line emission is thus often conveyed in $T_B$ units.895 We follow this convention and express the intensity as a velocity integrated brightness Lemperatiure: ‘This integrated intensity is thus a measure of total CO emission along the line of sight., We follow this convention and express the intensity as a velocity integrated brightness temperature: This integrated intensity is thus a measure of total CO emission along the line of sight.896 iis computed at all positions vielding a 2D map. which can be used with the 2D map of tto obtain the [factor through Equation 1..," is computed at all positions yielding a 2D map, which can be used with the 2D map of to obtain the factor through Equation \ref{Xfac}."897 To begin our investigation of CO emission. we assess the correlation of velocity integrated intensities wwith the total column density ool the model MICs.," To begin our investigation of CO emission, we assess the correlation of velocity integrated intensities with the total column density of the model MCs."898 Figure 1. shows aas a [function of the Που four of the simulations listed in Table 1.., Figure \ref{wvsav} shows as a function of the for four of the simulations listed in Table \ref{exptab}.899 The bottom abscissa showsNici. while the top abscissa shows the corresponding ((see Eqn. 11)).," The bottom abscissa shows, while the top abscissa shows the corresponding (see Eqn. \ref{extinction}) )."900 A number of features are readily apparent in Figure , A number of features are readily apparent in Figure \ref{wvsav}.901First. there is à general trend. of increasing intensity with increasing column density. though the slopes ciller between the various simulations.," First, there is a general trend of increasing intensity with increasing column density, though the slopes differ between the various simulations."902 For the high-density run (n1000) shown in Figure lee. the intensities do not) demonstrate any clear. trend with increasingely.," For the high-density run (n1000) shown in Figure \ref{wvsav}c c, the intensities do not demonstrate any clear trend with increasing."903 Rather. the vast majority of the intensities reach a threshold value of 7265 +2.," Rather, the vast majority of the intensities reach a threshold value of $\sim$ 65 ."904 his saturation of intensities is expected to occur at high densities since the [line becomes optically thick., This saturation of intensities is expected to occur at high densities since the line becomes optically thick.905 Saturation is also found. at the highest extinctions in the Alilky Way simulation (1300. Fie laa).," Saturation is also found at the highest extinctions in the Milky Way simulation (n300, Fig \ref{wvsav}a a)."906 Indeed. saturation at high. CO intensities has been observed in ALC's in the solar neighborhood (c.g.??)..," Indeed, saturation at high CO intensities has been observed in MCs in the solar neighborhood \citep[e.g.][]{Lombardietal06,Pinedaetal08}."907 Though model n300 qualitatively. reproduces the observed trends of increasing €CO intensity with increasing density. up to a threshold value. a detailed quantitive comparison is not appropriate here: the precise slope. scatter. minimum. and threshold intensity are dependent on additional physics not included in our models. such as additional heating clue to stars. outllows. and supernovac.," Though model n300 qualitatively reproduces the observed trends of increasing CO intensity with increasing density, up to a threshold value, a detailed quantitive comparison is not appropriate here; the precise slope, scatter, minimum, and threshold intensity are dependent on additional physics not included in our models, such as additional heating due to stars, outflows, and supernovae."908 Figure 1bb and d show that for clouds with lower metallicities (Modeln300-Z03) or densities (Model n100). there is a wider cistribution of intensities at [ow extinctions.," Figure \ref{wvsav}b b and d show that for clouds with lower metallicities (Modeln300-Z03) or densities (Model n100), there is a wider distribution of intensities at low extinctions."909 The saturation of the CO line is not easilyevident. compared to Model n1000., The saturation of the CO line is not easilyevident compared to Model n1000.910 Since CO, Since CO911We adopted an assumption that only very voung stellar populations formed Iron mixed eas of Ηλος and ISM in the LMC have unusually low [N/II] in the LAC: the accretion evenis of IIVCs onto the LAIC need to happen only recently.,We adopted an assumption that only very young stellar populations formed from mixed gas of HVCs and ISM in the LMC have unusually low [N/H] in the LMC: the accretion events of HVCs onto the LMC need to happen only recently.912" We have shown that the large number (>2500) of massive HIVCs (~10*ML, ) are required to exist within the LMC's orbital radius wilh respect to the Galactic center: the required total mass of IWCs (Αἱμνος) in the Galactic halo is about ~2.5xI0!7ML,. [or a reasonable set of model parameters."," We have shown that the large number $> 2500$ ) of massive HVCs $\sim 10^7 {\rm M}_{\odot}$ ) are required to exist within the LMC's orbital radius with respect to the Galactic center: the required total mass of HVCs $M_{\rm HVC, G}$ ) in the Galactic halo is about $\sim 2.5 \times 10^{10} {\rm M}_{\odot}$ for a reasonable set of model parameters."913 Although (he previous numerical simulations tried to predict the total mass of the Galactic HINC's. the predicted mass ranges widely from ~107M. (Peek οἱ al.," Although the previous numerical simulations tried to predict the total mass of the Galactic HVCs, the predicted mass ranges widely from $\sim 10^8 {\rm M}_{\odot}$ (Peek et al."914 2008) to ~2xLOMAL. (Maller Bullock 2004)., 2008) to $\sim 2 \times 10^{10} {\rm M}_{\odot}$ (Maller Bullock 2004).915 The required Mic to explain the observed low [N/IH] in the present scenario appears lo exceed (he predicted. Mie.," The required $M_{\rm HVC, G}$ to explain the observed low [N/H] in the present scenario appears to exceed the predicted $M_{\rm HVC, G}$."916 Given that the typical mass and 3D distribuGon of the Galactic ILVCs remains observationally unclear (e.g. Wakker 2004). the above inconsistency between the required total mass of IIVC's and the theoretically. predicted one does not rule out the present scenario.," Given that the typical mass and 3D distribution of the Galactic HVCs remains observationally unclear (e.g. Wakker 2004), the above inconsistency between the required total mass of HVCs and the theoretically predicted one does not rule out the present scenario."917 It is. however. reasonable (hat gas from other sources (e.g.. galaxies in the Local Group) can also play a role in diluting the ISM of the LAIC.," It is, however, reasonable that gas from other sources (e.g., galaxies in the Local Group) can also play a role in diluting the ISM of the LMC."918 Recently. Dekki Chiba (2007) have shown that the ISM stripped from the SAIC durug the LAIC-SAIC-Galaxy interaction for the past 2 Gyr can collide with the LAIC’s disk around 0.2 Gyr ago.," Recently, Bekki Chiba (2007) have shown that the ISM stripped from the SMC during the LMC-SMC-Galaxy interaction for the past 2 Gyr can collide with the LMC's disk around 0.2 Gyr ago."919" We thus suggest the following “SAICtransfer” scenario (or ""Magellanic squall”: Bekki Chiba 2007).", We thus suggest the following “SMC-transfer” scenario (or “Magellanic squall”; Bekki Chiba 2007).920 During the last 0.2 Gyr. the LAIC and the SAIC have interacted each other like a binary through their strong gravitational fields.," During the last 0.2 Gyr, the LMC and the SMC have interacted each other like a binary through their strong gravitational fields."921 As a result of this tidal interaction. gas with low [N/IH] in the SMIC can be transferred. efficiently to the LAIC sporadically. which induces star formation ancl thus creation of HII regions with low [N/II] in the LAIC.," As a result of this tidal interaction, gas with low [N/H] in the SMC can be transferred efficiently to the LMC sporadically, which induces star formation and thus creation of HII regions with low [N/H] in the LMC."922" Thus the observed low /Η of voung populations in the LAIC is a result of a close tidal interaction between. AICs in (he last 0.2 GAP Or κο,", Thus the observed low [N/H] of young populations in the LMC is a result of a close tidal interaction between MCs in the last 0.2 Gyr or so.923 We here suggest that this SAICtransler scenario has the following three advantages in explaining the observed low [N/IH]., We here suggest that this SMC-transfer scenario has the following three advantages in explaining the observed low [N/H].924 Firstly. the relative velocities between the infalling gas rom the SMC and the LMC's gas disk can be as small as ~60 kins !. because the relative velocity between the LMC and the SMC is ~ 60 kins |! for the last 200 Mrs (e.g. orbital nodels of the LAIC and the SAIC shown in Bekki Chiba 2005).," Firstly, the relative velocities between the infalling gas from the SMC and the LMC's gas disk can be as small as $\sim 60$ km $^{-1}$, because the relative velocity between the LMC and the SMC is $\sim$ 60 km $^{-1}$ for the last 200 Myrs (e.g., orbital models of the LMC and the SMC shown in Bekki Chiba 2005)."925 The relative velocity is Significantly smaller than the circular. velocities of the LMC (~80—120 km ! for a reasonable mass model of the LAIC: Bekki Chiba 2005) so that the infalling gas is highly ikelv to be (rapped by the gravitational potential of the LMC., The relative velocity is significantly smaller than the circular velocities of the LMC $\sim 80-120$ km $^{-1}$ for a reasonable mass model of the LMC; Bekki Chiba 2005) so that the infalling gas is highly likely to be trapped by the gravitational potential of the LMC.926 On the other hand. the relative velocities of the ΗΝCs and the LMC can be as large as (he velocity dispersion of the Galactic halo (~160 kins |) so that the infalling HVCs are less likely to be trapped by the LMC in comparison with the infalling SMC gas.," On the other hand, the relative velocities of the HVCs and the LMC can be as large as the velocity dispersion of the Galactic halo $\sim 160$ km $^{-1}$ ) so that the infalling HVCs are less likely to be trapped by the LMC in comparison with the infalling SMC gas."927must be more we do no know at this moment.,must be more we do not know at this moment.928 We think we should first express our thanks to the referee and editors of APJ and the authors of these papers for their informing us of these works. we then also would like to point out that. almost all these works contain the problem we pointed out in this paper. 1.6.. when writing down the basic equations to describe the evolution of the over-deuse regions. the assuuption that dark energv moves svuchronouslv with ordinary matters is made. but iu the basic equations when the dark eunergvs density is involved. the results from another different assumption is used.," We think we should first express our thanks to the referee and editors of APJ and the authors of these papers for their informing us of these works, we then also would like to point out that, almost all these works contain the problem we pointed out in this paper, i.e., when writing down the basic equations to describe the evolution of the over-dense regions, the assumption that dark energy moves synchronously with ordinary matters is made, but in the basic equations when the dark energy's density is involved, the results from another different assumption is used."929"Below, we determined the signal-to-noise ratios for an optimal orientation of detector and source.","Below, we determined the signal-to-noise ratios for an optimal orientation of detector and source."930" As detector sensitivity we used the present performance of one single LIGO instrument and the improved, possible future performance of an Advanced LIGO (AdvLIGO) detector (Shoemaker 2007)."," As detector sensitivity we used the present performance of one single LIGO instrument and the improved, possible future performance of an Advanced LIGO (AdvLIGO) detector \citep{Shoemaker:2007}."931. Note that Advanced LIGO possesses several adjustable frequency responses., Note that Advanced LIGO possesses several adjustable frequency responses.932" While the 'burst' selection provides a broad range of nearly maximum sensitivity (optimal for model s15g), a ""nsns'-tuned instrument is likely to be used for sources that radiate at lower frequencies (model s15h)."," While the 'burst' selection provides a broad range of nearly maximum sensitivity (optimal for model s15g), a 'nsns'-tuned instrument is likely to be used for sources that radiate at lower frequencies (model s15h)."933" The computed quadrupole wave amplitudes of models s15g and s15h are displayed in Figs. 3,,"," The computed quadrupole wave amplitudes of models s15g and s15h are displayed in Figs. \ref{fig3.eps},"934 4 and 6.., \ref{fig4.eps} and \ref{fig6.eps}.935 A quantitative overview of key properties andresults is given in Table 1 and Fig. 5.., A quantitative overview of key properties andresults is given in Table \ref{table:1} and Fig. \ref{fig5.eps}.936 Detector dependent quantities are summarised in Table 2 and Fig. 8.., Detector dependent quantities are summarised in Table \ref{table:2} and Fig. \ref{fig8.eps}.937" Model s15g undergoes a rotational core collapse as we assume an initial central angular velocity of Q=27 rad/s. During the stage of contraction, the core spins up massively while becoming oblate."," Model s15g undergoes a rotational core collapse as we assume an initial central angular velocity of $\Omega=2\pi$ rad/s. During the stage of contraction, the core spins up massively while becoming oblate."938 The collapse gets abruptly halted due to the stiffening of the EoS above nuclear densities., The collapse gets abruptly halted due to the stiffening of the EoS above nuclear densities.939 Afterwards the core rebounds and drives a hydrodynamical shock wave outwards., Afterwards the core rebounds and drives a hydrodynamical shock wave outwards.940 These conditions give rise to strong time-dependent variations in the physical quantities that affect the quadrupole tensor (see equ.7))., These conditions give rise to strong time-dependent variations in the physical quantities that affect the quadrupole tensor (see \ref{equ:2}) ).941 Its behaviour is then directly reflected by the model's gravitational wave signature., Its behaviour is then directly reflected by the model's gravitational wave signature.942" Since the core collapse proceeds nearly axisymmetrically, the only GW amplitude considerably driven by the rotationally induced large-scale asymmetries is À,j;."," Since the core collapse proceeds nearly axisymmetrically, the only GW amplitude considerably driven by the rotationally induced large-scale asymmetries is $_{+II}$."943" It exceeds the other s15g-wave trains A,j,, Ax; and Ay;; by 1-2 orders of magnitude, as one can see in Fig."," It exceeds the other s15g-wave trains $_{+I}$ $_{\times I}$ and $_{\times II}$ by 1-2 orders of magnitude, as one can see in Fig."944 3 for times around bounce., \ref{fig3.eps} for times around bounce.945" The initial GW signal is emitted just before core bounce, when the rapid infall of matter and the spin-up of the core is dominant."," The initial GW signal is emitted just before core bounce, when the rapid infall of matter and the spin-up of the core is dominant."946 It shows a prebounce rise., It shows a prebounce rise.947" Then, the instantaneous slowdown of matter at core bounce leads to a prominent negative peak, which is followed by a ring-down behaviour that lasts for the first few ms postbounce."," Then, the instantaneous slowdown of matter at core bounce leads to a prominent negative peak, which is followed by a ring-down behaviour that lasts for the first few ms postbounce."948 This generic type I wave characteristics is displayed in the lower left panel of Fig. 3.., This generic type I wave characteristics is displayed in the lower left panel of Fig. \ref{fig3.eps}.949" Performing a Fourier transform of the GW signal around bounce (—5 ms «f5 ms), we find a spectrum with a very narrow bandwidth peaking around 893 Hz."," Performing a Fourier transform of the GW signal around bounce $-5$ ms $<t<5$ ms), we find a spectrum with a very narrow bandwidth peaking around $893$ Hz."950" This is in good agreement with the recent findings from Mülleretal. (2004), but more than 150 Hz higher than in (Dimmelmeieretal.2007;Ott 2007).."," This is in good agreement with the recent findings from \citet{M2004}, , but more than $150$ Hz higher than in \citep{Dimmelmeier07PhysRev,2007CQGra..24..139O}. ."951So far our definition of au upper limit assmues that there are no wnkuown miüsauce parameters. and im particular that Ap is known.,"So far our definition of an upper limit assumes that there are no unknown nuisance parameters, and in particular that $\lamB$ is known."952 Unfortunately. the probabilities in EquationsLo and 6 cannot be computed if Ap is uuknown.," Unfortunately, the probabilities in Equations\ref{eq:alpha}953 and \ref{eq:power} cannot be computed if $\lamB$ is unknown."954 In this section we describe several strategies that cau be used in the more realistic situation when Ap is not known precisely., In this section we describe several strategies that can be used in the more realistic situation when $\lamB$ is not known precisely.955 The most conservative procedure cusures that the detection probability of the wpper limit is ercater than iin for anv possible value of Ap., The most conservative procedure ensures that the detection probability of the upper limit is greater than $\beta_{\rm min}$ for any possible value of $\lamB$.956 Conucrally speaking. the larger Ap is. the larger As ust be in order to be detected with a given probability. aud thus the lareer the upper Init.," Generally speaking, the larger $\lamB$ is, the larger $\lamS$ must be in order to be detected with a given probability, and thus the larger the upper limit."957 Thus. a useful upper lit requires a finite ranec. Xp. to be specified for Ap.," Thus, a useful upper limit requires a finite range, $\rangeLamB$, to be specified for $\lamB$."958 Given this range. aiit eau be defined as the sinallest Ag that satisfies (We use the toxinΠα (Gut) rather than to allow for the case when the mini may be ou the boundary of. but outside. the rauge of interest.," Given this range, a can be defined as the smallest $\lamS$ that satisfies (We use the term $\inf$ ) rather than to allow for the case when the minimum may be on the boundary of, but outside, the range of interest."959 It is the largest umber that is suialler than all the Παπους in the rage., It is the largest number that is smaller than all the numbers in the range.960 For instance. the minim of the range [Le2O} is undefined. but the στ is 0.)," For instance, the minimum of the range $\{x>0\}$ is undefined, but the infimum is $0$ .)"961 Unufortunatelv. uuless the range of values Ap is relatively precise. this upper limit will often be too large to be useful.," Unfortunately, unless the range of values $\rangeLamB$ is relatively precise, this upper limit will often be too large to be useful."962 Iu practice. there is better solution.," In practice, there is better solution."963 The backeround count provides information on the likely values of Ap that should be used when computing the upper luit., The background count provides information on the likely values of $\lamB$ that should be used when computing the upper limit.964 Iu particular the distribution of Ap2 05eiven ip cau be computed using standard Bayesian and used to evaluate the probability. where S*(a) is the smallest value such that The upper lint is then computed as the sinallest As that satisfies J(As)c ig ," In particular the distribution of $\lamB$ given $\nB$ can be computed using standard Bayesian and used to evaluate the , where $\thresh(\alpha)$ is the smallest value such that The upper limit is then computed as the smallest $\lamS$ that satisfies $\beta(\lamS) \geq \beta_{\rm min}$ ."965Unlike the upper Huit described in rets:def.. these calculations require data. in particular sp.," Unlike the upper limit described in \\ref{s:def}, these calculations require data, in particular $\nB$ ."966 For this reason. we call the smallest As that satisfies Equation & the düuit oyHiit.," For this reason, we call the smallest $\lamS$ that satisfies Equation \ref{eq:ULpost} the or."967 Au iuteriucdiate approach that is more practical than using Equation 7 but more conservative than using Equation & is to simply compute a high percentile of pCAg]np). perhaps its 95th percentile.," An intermediate approach that is more practical than using Equation \ref{eq:ULinf}968 but more conservative than using Equation \ref{eq:ULpost} is to simply compute a high percentile of $p(\lamB | \nB)$, perhaps its 95th percentile."969 The procedure for kuown Ap can then be used with this percentile treated as the known value of Ap., The procedure for known $\lamB$ can then be used with this percentile treated as the known value of $\lamB$.970 This is a couscrvative strategv in that it asstunes a nearly worst case scenario for the level of background contamination., This is a conservative strategy in that it assumes a nearly worst case scenario for the level of background contamination.971 Asan illustration. suppose the uncertainty in Ap eiven the observed backerouud counts eau be sunnuimidzed in the posterior distribution plottediu theleftpanel of Figure 7..," As an illustration, suppose the uncertainty in $\lamB$ given the observed background counts can be summarized in the posterior distribution plottedin theleftpanel of Figure \ref{fig:unknown}. ."972 This isa ooOgana posterior distribution, This isa gamma posterior distribution973trrue nature of these sources.,rue nature of these sources.974"mav be written as where yp=AL, /Al,. Ad, is the planetary mass. M; is the stellar mass. ry is the location of the planet from the host star. Ph is the disc scale height. and the normalization constant po is chosen so that the density at 2=Q0 corresponds to the unperturbed density (Following ?)).","may be written as where $\mu=M_p/M_*$ , $M_p$ is the planetary mass, $M_*$ is the stellar mass, $r_p$ is the location of the planet from the host star, $h$ is the disc scale height, and the normalization constant $\rho_0$ is chosen so that the density at $z=0$ corresponds to the unperturbed density (following \citet{js03}) )."975 Thus. the dust and gas distributions are perturbed by the presence of a planet. ancl consequently. the thermal and. density structures are alfected by the planet.," Thus, the dust and gas distributions are perturbed by the presence of a planet, and consequently the thermal and density structures are affected by the planet."976" We interpolate the density ancl temperature of. disces in the region. within the Hill radius ryzrj,M) following ? because the hydrostatic assumption breaks down within the region.", We interpolate the density and temperature of discs in the region within the Hill radius $r_H\approx r_p(M_p/M_*)^{1/3}$ following \citet{js05} because the hydrostatic assumption breaks down within the region.977 Lt is well known that the horseshoe structure and/or circumplanctary disces are. located: well inside the Lill radius (c.g.2). so that we consider this treatment to be very conservative since use of the Lll radius overestimates the effect.," It is well known that the horseshoe structure and/or circumplanetary discs are located well inside the Hill radius \citep[e.g.][]{cbkm09}, so that we consider this treatment to be very conservative since use of the Hill radius overestimates the effect."978 This interpolation is unlikely to be harmful for caleulating the Lindblad torque because the resonant positions are pushed away from their distance from. the planets 25/3zrg., This interpolation is unlikely to be harmful for calculating the Lindblad torque because the resonant positions are pushed away from their distance from the planets $2h/3\approx r_H$.979 The Lindblad torque. takes its maximum value at. the distance 25/3 from a planet due to the gas pressure ellect., The Lindblad torque takes its maximum value at the distance $2h/3$ from a planet due to the gas pressure effect.980 Η then decreases with increasing distance from the planet (?).., It then decreases with increasing distance from the planet \citep{ward97}.981 Thus. it is important to highly resolve the thermal structure of the region in the vicinity of planets in order to calculate the Lindblad: torque accurately.," Thus, it is important to highly resolve the thermal structure of the region in the vicinity of planets in order to calculate the Lindblad torque accurately."982 The thermal structure of dises. however. is determined globally since clises are heated by their central stars.," The thermal structure of discs, however, is determined globally since discs are heated by their central stars."983 Furthermore. the higher the resolution. the longer is the computational time.," Furthermore, the higher the resolution, the longer is the computational time."984 This is because a larger number of photons is required in finer grid systems to avoid random noise produced by the Monte Carlo methods., This is because a larger number of photons is required in finer grid systems to avoid random noise produced by the Monte Carlo methods.985 Vherefore. carefully constructed. grid svstenis. are required to optimize the computational time.," Therefore, carefully constructed grid systems are required to optimize the computational time."986 We adopt a logarithmic grid svstem in vertical direction., We adopt a logarithmic grid system in vertical direction.987 This ids sullicient because the inner region in which the torques provide the largest. contribution is eLlicienthy resolved. by finer grids., This is sufficient because the inner region in which the torques provide the largest contribution is efficiently resolved by finer grids.988 In the radial direction. however. more complicated. grid svstems are required.," In the radial direction, however, more complicated grid systems are required."989 We combine two types of erid systems., We combine two types of grid systems.990 Around the planet. we use an equally spaced. fine mesh system. while for the rest of the disc. we use a logarithmically spaced. coarser mesh.," Around the planet, we use an equally spaced, fine mesh system while for the rest of the disc, we use a logarithmically spaced, coarser mesh."991 The fine mesh is inserted into the vicinity of the planets with its size AR centered at the location of the planet. while the larger remaining region is resolved by the coarser mesh.," The fine mesh is inserted into the vicinity of the planets with its size $\bigtriangleup R$ centered at the location of the planet, while the larger remaining region is resolved by the coarser mesh."992 We fix the disc size as 50 au and the total number of grids as 180., We fix the disc size as 50 au and the total number of grids as 180.993 Also. the region AR is resolved by 65 grids for any case.," Also, the region $\bigtriangleup R$ is resolved by 65 grids for any case."994 Furthermore. we recalculate the density structure alter temperature calculations are completed by using à times finer svstems mentioned above.," Furthermore, we recalculate the density structure after temperature calculations are completed by using $n$ times finer systems mentioned above."995 This treatment enables one to reduce the computational time by a factor of 10 or so., This treatment enables one to reduce the computational time by a factor of 10 or so.996 We performed a convergence study for the Lindblad. torque in our eid system by changing AR and n. and found AR=6h and pὁ are sullicient (7)..," We performed a convergence study for the Lindblad torque in our grid system by changing $\bigtriangleup R$ and $n$, and found $\bigtriangleup R=6h$ and $n=6$ are sufficient \citep{h08}."997 Dust is important for heating the disc., Dust is important for heating the disc.998 We first summarise the results of our simulations of the thermal and. density structures of disces with a planet for the case of well mixed dust., We first summarise the results of our simulations of the thermal and density structures of discs with a planet for the case of well mixed dust.999 Fig., Fig.1000 1. shows the dust. gas and total (gas. | cust) density distributions of the MMSN disc models respectively on the top to bottom panels.," \ref{fig1} shows the dust, gas and total (gas + dust) density distributions of the MMSN disc models respectively on the top to bottom panels."1001 A LO AZ) planet is placed. without loss of generalitv. at 6 au in Fig. l..," A 10 $M_{\oplus }$ planet is placed, without loss of generality, at 6 au in Fig. \ref{fig1}."1002 Also. the zoomed-in versions of disc structures around the planet are shown in the right column.," Also, the zoomed-in versions of disc structures around the planet are shown in the right column."1003 Note that the disc temperatures shown on every panel are the same because they are defined as the mass-averaged temperature of dust with various grain sizes (Paper I)., Note that the disc temperatures shown on every panel are the same because they are defined as the mass-averaged temperature of dust with various grain sizes (Paper I).1004 Our numerical simulations confirm that the mid-plane region has much lower temperatures than the surface [aver for both disc models (Fig. 1))., Our numerical simulations confirm that the mid-plane region has much lower temperatures than the surface layer for both disc models (Fig. \ref{fig1}) ).1005 We found. for both disc models. that the temperature of the surface is =SO Ix at lau while that of the mic-plane is =25 Ix there.," We found, for both disc models, that the temperature of the surface is $\simeq 80$ K at 1 au while that of the mid-plane is $\simeq 25$ K there."1006 This arises because the surface laver is directly. heated by the central star while the mid-plane region is heated only by the thermal emission of dust (2).., This arises because the surface layer is directly heated by the central star while the mid-plane region is heated only by the thermal emission of dust \citep{cg97}.1007 Also. the direct heating of discs results in ecometrically [ared. structures. which are seen in both he dust and σας density distributions.," Also, the direct heating of discs results in geometrically flared structures, which are seen in both the dust and gas density distributions."1008 Furthermore. the oesence of the planet produces a low density region above it due to the compression by the gravitational force of the planet.," Furthermore, the presence of the planet produces a low density region above it due to the compression by the gravitational force of the planet."1009 Consequently. the region above a planet has higher emperatures while the mid-plane region. in turn. has lower emperatures.," Consequently, the region above a planet has higher temperatures while the mid-plane region, in turn, has lower temperatures."1010 The compression results in a small localized »ak in a temperature profile in the mid-plane (Fig. 2))., The compression results in a small localized peak in a temperature profile in the mid-plane (Fig. \ref{fig2}) ).1011 The dillerence between the MMSN and S07 disc models is that the temperature in the mid-plane region becomes ower in the former model. especially at the region around he planet and larger disc radii.," The difference between the MMSN and S07 disc models is that the temperature in the mid-plane region becomes lower in the former model, especially at the region around the planet and larger disc radii."1012 We found that the emperature dilference between two cise models is about 15 »er cent in the mid-plane region at. 10 au., We found that the temperature difference between two disc models is about 15 per cent in the mid-plane region at 10 au.1013 This is explained w the two combined. effects arising from a steeper slope ποιο., This is explained by the two combined effects arising from a steeper slope profile.1014 “Phe MMSN cise model has a denser inner region. and is therefore more optically thick.," The MMSN disc model has a denser inner region, and is therefore more optically thick."1015 La our disc setup. he MMSN model is about 4 times denser than the S07 model in the mid-plane region at 1 au.," In our disc setup, the MMSN model is about 4 times denser than the S07 model in the mid-plane region at 1 au."1016 As a result. the inner region in the MMSN disc. model prevents photons rom readily. penetrating the outer region. resulting in lower emperatures there.," As a result, the inner region in the MMSN disc model prevents photons from readily penetrating the outer region, resulting in lower temperatures there."1017 Ehe higher density in the inner region. in turn. results in a less dense. outer region even in the müd-plane region. since the total disc mass is fixed.," The higher density in the inner region, in turn, results in a less dense, outer region even in the mid-plane region, since the total disc mass is fixed."1018 The outer region consequently becomes optically thin even to the thermal emission of dust. resulting in lower temperatures.," The outer region consequently becomes optically thin even to the thermal emission of dust, resulting in lower temperatures."1019 Since the difference of the optical cepth is accumulated with increasing disc radius. the temperature difference far from the star becomes larger.," Since the difference of the optical depth is accumulated with increasing disc radius, the temperature difference far from the star becomes larger."1020 Lt is interesting that both disc models result in the almost identical temperature profiles at the mid-plane. especially at smaller cise radii. and are well represented by rBe (Fig. 2)).," It is interesting that both disc models result in the almost identical temperature profiles at the mid-plane, especially at smaller disc radii, and are well represented by $r^{-3/5}$ (Fig. \ref{fig2}) )."1021 We showed that cust settling. drives the dise structure from Uared to fatter shapes in Paper L Fig., We showed that dust settling drives the disc structure from flared to flatter shapes in Paper I. Fig.1022 3.shows that the same behaviour pertains to the MMSN model as well., \ref{fig3} shows that the same behaviour pertains to the MMSN model as well.1023 Dust settling arises from the size distribution of dust erains., Dust settling arises from the size distribution of dust grains.1024 In general. larger grains cannot be kept aloft even," In general, larger grains cannot be kept aloft even"1025"We have presented first results from the analysis of a long observation of IGR J16318—4848, the most extreme of the strongly absorbed --sources"".","We have presented first results from the analysis of a long observation of IGR $-$ 4848, the most extreme of the strongly absorbed -sources”."1026" As found in previous studies, the average ""INTEGRALspectrum of the source is consistent with a strongly absorbed exponentially cutoff power-law and strong flourescent line emission."," As found in previous studies, the average spectrum of the source is consistent with a strongly absorbed exponentially cutoff power-law and strong flourescent line emission."1027" In contrast to earlier studies, the power-law photon index was found to be considerably harder than before (AT from 0.67 up to 1.93)."," In contrast to earlier studies, the power-law photon index was found to be considerably harder than before $\Delta\Gamma$ from 0.67 up to 1.93)."1028" This result can be due to the significantly better signal to noise ratio in the energy band above kkeV compared to the earlier studies, which allows for a better determination of the high energy cutoff, the continuum parameters, and Ny than the earlier soft X-ray measurements, although an instrinsic change in the source is not ruled out."," This result can be due to the significantly better signal to noise ratio in the energy band above keV compared to the earlier studies, which allows for a better determination of the high energy cutoff, the continuum parameters, and $N_\mathrm{H}$ than the earlier soft X-ray measurements, although an instrinsic change in the source is not ruled out."1029 The soft excess below 2kkeV is probably due to a serendipitous source near IGR J16318—48482007)., The soft excess below keV is probably due to a serendipitous source near IGR $-$ 4848.1030 The considerable variability of the source can be explained as being due to variations in Ny., The considerable variability of the source can be explained as being due to variations in $N_\mathrm{H}$ .1031" As pointed out by(2004),, the general spectral characteristics derived from the fit are typical for accreting neutron stars2008)."," As pointed out by, the general spectral characteristics derived from the fit are typical for accreting neutron stars."1032". Note that this result does not mean that the neutron star nature of the compact object in IGR J16318—4848 is confirmed, which would require e.g. the detection of pulsations."," Note that this result does not mean that the neutron star nature of the compact object in IGR $-$ 4848 is confirmed, which would require e.g. the detection of pulsations."1033" A search for pulsations in the range between ss and kksec was negative, while shorter period pulsations are probably not detectable due to the smearing of pulsations by Compton scattering2005)."," A search for pulsations in the range between s and ksec was negative, while shorter period pulsations are probably not detectable due to the smearing of pulsations by Compton scattering."1034". Turning to the emission lines, we note that our fit requires a slight overabundance of iron with respect to the ISM values of(2000),, as one would expect for an evolved star."," Turning to the emission lines, we note that our fit requires a slight overabundance of iron with respect to the ISM values of, as one would expect for an evolved star."1035" Furthermore, the flux ratio of Fe and Ni also points towards a Ni overabundance by a factor of ~2.5 with respect to Fe."," Furthermore, the flux ratio of Fe and Ni also points towards a Ni overabundance by a factor of $\sim$ 2.5 with respect to Fe."1036" The ratio of the Fe Ka and Fe Kj line fluxes is given by n=(F(FeKBi)+F(FeK65)/(""(FeΚαι)Κα») 0.008."," The ratio of the Fe $\alpha$ and Fe $\beta$ line fluxes is given by $\eta=(F(\mathrm{Fe\ K}\beta_1)+F(\mathrm{Fe\1037 K}\beta_3))/(F(\mathrm{Fe\ K}\alpha_1)+F(\mathrm{Fe\ K}\alpha_2))=10380.086\pm 0.008$ ."1039" This flux ratio is formally slightly smaller than that found in theoretical calculations for neutral gas phase Fe atoms of0.121),,0.125), or 0.132(2)),, and it is also smaller than the value of 7 found in experimental measurements performed in solid Fe (e.g., 7=0.1307(7) found by and 2002))."," This flux ratio is formally slightly smaller than that found in theoretical calculations for neutral gas phase Fe atoms of, or , and it is also smaller than the value of $\eta$ found in experimental measurements performed in solid Fe (e.g., $\eta=0.1307(7)$ found by and )."1040" The difference between the different theoretical calculations is due to certain approximations made in solving the structure of the excited Fe ion after the K-shell photoabsorption, while for the latter measurements η is affected by internal absorption in the Fe crystal used to make the measurements as well as by the dependence of the emission probability of the photoelectron on orientation."," The difference between the different theoretical calculations is due to certain approximations made in solving the structure of the excited Fe ion after the K-shell photoabsorption, while for the latter measurements $\eta$ is affected by internal absorption in the Fe crystal used to make the measurements as well as by the dependence of the emission probability of the photoelectron on orientation."1041" The systematic uncertainty of 7 in theory and measurements is therefore probably as large as 0.02, which would make our measurement consistent with neutral Fe."," The systematic uncertainty of $\eta$ in theory and measurements is therefore probably as large as 0.02, which would make our measurement consistent with neutral Fe."1042 We note that our value for 7 is significantly smaller than the 7=0.207953 found in the EPIC-pn analysis of2003)mainBodyCitationEnd1859]Matt:03a.," We note that our value for $\eta$ is significantly smaller than the $\eta1043=0.20^{+0.02}_{-0.03}$ found in the EPIC-pn analysis of."1044 These authors speculated that this higher 7 could be due to the absorbing wind being moderately ionized., These authors speculated that this higher $\eta$ could be due to the absorbing wind being moderately ionized.1045" Given that the line ratio (and also the line energy) found in the higher resolution data are consistent with neutral Fe, we might be seeing a change in the ionization structure of the wind between the and the observations."," Given that the line ratio (and also the line energy) found in the higher resolution data are consistent with neutral Fe, we might be seeing a change in the ionization structure of the wind between the and the observations."1046" Alternatively, the larger value for 7 may be due to systematic effects in the analysis: WithSuzaku,, the Fe K line and the Fe K edge are easier to separate and the spectral continuum is better constrained in the present analysis than withXMM-Newton,, since spectral information is available above kkeV. Finally, despite the large column of the source, no significant evidence for the presence of a Compton shoulder is found in the spectrum, which is consistent with previous results."," Alternatively, the larger value for $\eta$ may be due to systematic effects in the analysis: With, the Fe $\beta$ line and the Fe K edge are easier to separate and the spectral continuum is better constrained in the present analysis than with, since spectral information is available above keV. Finally, despite the large column of the source, no significant evidence for the presence of a Compton shoulder is found in the spectrum, which is consistent with previous results."1047" This result is in contrast to the expectation for absorption in an homogeneous medium: As shown by(2002),, with this assumption the equivalent width of the Fe Ka line at the Ny of IGR J16318—484 should be much less than that observed here, and a strong Compton shoulder should be present, in line e.g. with the Compton shoulder observed by in301-2."," This result is in contrast to the expectation for absorption in an homogeneous medium: As shown by, with this assumption the equivalent width of the Fe $\alpha$ line at the $N_\mathrm{H}$ of IGR $-$ 484 should be much less than that observed here, and a strong Compton shoulder should be present, in line e.g. with the Compton shoulder observed by in."1048". As pointed out by e.g. and(2007),, the non-existence of the Compton shoulder could be due to a strongly inhomogeneous absorbing medium."," As pointed out by e.g. and, the non-existence of the Compton shoulder could be due to a strongly inhomogeneous absorbing medium."1049" Since the strength of the shoulder is strongly dependent on the assumed accretion geometry, further work using self-consistent modeling of the absorption, fluorescent line formation and Compton shoulder formation is required."," Since the strength of the shoulder is strongly dependent on the assumed accretion geometry, further work using self-consistent modeling of the absorption, fluorescent line formation and Compton shoulder formation is required."1050" We will present such self-consistent analyses, as well as a more detailed study of the variability of the source, in a future publication."," We will present such self-consistent analyses, as well as a more detailed study of the variability of the source, in a future publication."1051eccenlricily of (he Jovian planet. ancl (he eccentricity damping rate (through A) of (he rocky planet.,"eccentricity of the Jovian planet, and the eccentricity damping rate (through $K$ ) of the rocky planet."1052 Since (hese svstems are highly. chaotic. a laree ensemble of numerical experiments must be performed [ον each point in parameter space (~1000 independent realizations).," Since these systems are highly chaotic, a large ensemble of numerical experiments must be performed for each point in parameter space $\sim$ 1000 independent realizations)."1053 The main result from these simulations is (he fraction of the trials that end with the two planets colliding., The main result from these simulations is the fraction of the trials that end with the two planets colliding.1054 For a given migration rate. collisions represent the most common outcome provided that eccentricity damping is not (oo effective.," For a given migration rate, collisions represent the most common outcome provided that eccentricity damping is not too effective."1055 These results are depicted in Figures 1 and 2. which show the [fraction of collisions plotted. versus the parameter A. that sets the strength. of eccentricity damping for the rocky planet (equation [5]]).," These results are depicted in Figures \ref{fig:efraction} and \ref{fig:mfraction}, which show the fraction of collisions plotted versus the parameter $K$ that sets the strength of eccentricity damping for the rocky planet (equation \ref{edamp}] ])."1056 Figure 1. shows collision [fractions for four choices of starting eccentricity for the Lot Jupiter. from €=0 ilo e= 0.3.," Figure \ref{fig:efraction} shows collision fractions for four choices of starting eccentricity for the Hot Jupiter, from $e=0$ to $e=0.3$ ."1057" Figure 2. shows collision fractions for fixed starting eccentricity ο=0.2 and three choices for the Hot Jupiter mass. Mp/AM,= 0.5. 1. and 2."," Figure \ref{fig:mfraction} shows collision fractions for fixed starting eccentricity $e=0.2$ and three choices for the Hot Jupiter mass, $M_P/M_J=$ 0.5, 1, and 2."1058 In both Figures. each point shown corresponds to the fractions caleulated from NV1000 independent realizations of the starting conditions.," In both Figures, each point shown corresponds to the fractions calculated from $N\sim{1000}$ independent realizations of the starting conditions."1059 The error bars (1/ VN) provide a crude measure of the uncertainties., The error bars $\sim1/\sqrt{N}$ ) provide a crude measure of the uncertainties.1060 The results displaved in Figures 1 and 2 show a robust (rend: For sulliciently weak eccenlricily damping. A«Ac810. most simulations end with collisions between the planets.," The results displayed in Figures \ref{fig:efraction} and \ref{fig:mfraction} show a robust trend: For sufficiently weak eccentricity damping, $K<K_C\approx{10}$, most simulations end with collisions between the planets."1061 For stronger eccentricity damping. A>Ac. (he collision fraction becomes neelieible and nearly all of the svstenis survive (keeping both planets) over the entire range of integration limes.," For stronger eccentricity damping, $K>K_C$, the collision fraction becomes negligible and nearly all of the systems survive (keeping both planets) over the entire range of integration times."1062 Further. the critical level of eccentricity damping (Ac) depends on the starting eccenlricily and mass of the Hot Jupiter.," Further, the critical level of eccentricity damping $K_C$ ) depends on the starting eccentricity and mass of the Hot Jupiter."1063" Larger eccentricities (for the Jovian orbit) allow collisions to occur in the face of greater eccentricity damping. following a trend of the approximate form log,Ac2(32-106)/4 (from fitting)."," Larger eccentricities (for the Jovian orbit) allow collisions to occur in the face of greater eccentricity damping, following a trend of the approximate form $\log_{10}{K_C}\approx(3+10e)/4$ (from fitting)."1064 ILowever. larger eccentricities combined with smaller A values vield lower collision rates.," However, larger eccentricities combined with smaller $K$ values yield lower collision rates."1065 In this regime. collision events are replaced (primarily) bv accretion events (onto the star).," In this regime, collision events are replaced (primarily) by accretion events (onto the star)."1066 The lareer eccentricity of the Jovian planet provides (he rocky planet with greater opportunity (o pass by and enter the gravitational realm of the star., The larger eccentricity of the Jovian planet provides the rocky planet with greater opportunity to pass by and enter the gravitational realm of the star.1067 Similarly. larger masses for the Jovian planet allow collisions to occur lor larger values of the eccentricity damping parameter.," Similarly, larger masses for the Jovian planet allow collisions to occur for larger values of the eccentricity damping parameter."1068 In addition. larger masses combined with smaller ἐν values lead to lower collision rates.," In addition, larger masses combined with smaller $K$ values lead to lower collision rates."1069 Ii this case. the collision events are (again) replaced wilh accretion events.," In this case, the collision events are (again) replaced with accretion events."1070 The larger mass of the Jovian planet can scatter the rocky planet before impact. and the scattering alters the orbit of the rocky planet enough to send it into the star (or. more rarely. eject the planet).," The larger mass of the Jovian planet can scatter the rocky planet before impact, and the scattering alters the orbit of the rocky planet enough to send it into the star (or, more rarely, eject the planet)."1071 These results were obtained for a single migration rate: for faster (slower) migration. the outer planet is less (more) likely to lock into mean motion resonance and is more (less) likely to collide with the Jovian planet (INetehum οἱal.," These results were obtained for a single migration rate; for faster (slower) migration, the outer planet is less (more) likely to lock into mean motion resonance and is more (less) likely to collide with the Jovian planet (Ketchum etal."1072 2011)., 2011).1073 We have performed additional simulations with faster migration (not shown) to confirm these trends., We have performed additional simulations with faster migration (not shown) to confirm these trends.1074For further analysis we have prepared a system of relatively siualler programs for processing the Cherenkov LehtoO CORSISA output.,For further analysis we have prepared a system of relatively smaller programs for processing the Cherenkov light CORSIKA output.1075 These programsC» rescale the photon wavelengthC» range.Oo apply t," These programs rescale the photon wavelength range, apply the losses of photons in the atmosphere (which are wavelength, altitude and pathlength dependent; see for more details."1076he ↕∪↴∖↴↴∖↴↸∖↴∖↴∪↕⋟↻∐∪↑∪∐↴∖↴↕∐↑∐↸∖⋜↧↑↕⊔∪↴∖↴↻∐↸∖↥⋅↸∖≺↖↖⇁↕∐↸⊳∐⋜⋯∖↖↖⇁⋜↧↖↽↸∖↕↸∖∐∶↴∙⊾↑∐∙⋜↧↕↑↕↑∏≼∐∖⋜⋯≼⊔≻⋜↧↑↕∐, These programs calculate the detector mirror reflection probability according to (for the THEMISTOCLE mirrors).1077↸∖∐∶↴⋁↑∐≼∐∖↻↸∖∐≼∐∖∐↑∶↴∖↴↸∖↸∖ ⊺↕∐∖↻↥⋅∪∶↴⋁⋜⊔⊔↴∖↴∏↴∖↴↸∖∐↕≧↻↻↕↘⊽↻↥⋅∪↸⊳↸∖≺⊔∐⋅↸∖↴∖↴∪↕⋟≼⊲⊏↕⊰⋀∖⊽∎↴∖↴∣∣∪↙∶∣≖⋅∣∣⊲≀⇂∩⋜⋯≼⊔∐⋅∪∶↴⋁↥⋅⋜↧⋯↕⋟⋎∖↖⊽∏∶≩∫↖↖⇁⋜↧↴∖↴∏↴∖↴↸∖≺⊔∪ make most of prescuted eraphics., The photomultiplier quantum efficiency for bialkali cathode was taken from as for phototube PHILIPS The progams use HBOOK procedures of CERN's and program PAW was used to make most of presented graphics.1078 To demonstrate the role of 1unuous iu Cherenkov light detection as result of MlonteCarlo simulations we present an example of one shower observed by oue detector iu one place., To demonstrate the role of muons in Cherenkov light detection as result of Monte–Carlo simulations we present an example of one shower observed by one detector in one place.1079 EASs almost never look ‘the sanie because of fluctuations in their development and stochastic nature of most of physical processes involved., EASs almost never look 'the same' because of fluctuations in their development and stochastic nature of most of physical processes involved.1080 Average. mean nor ‘most probable values do not describe well the situation we are going to preseut.," Average, mean nor most probable' values do not describe well the situation we are going to present."1081 Many features are detector dependent aud full detailed analysis should be performed for specific experimental Iu our example we present result of MonteCarlo simulation ofEAS trigecred by vertical CR proton with cnerey 10 TeV. The detector is at the 1650 ii as...," Many features are detector dependent and full detailed analysis should be performed for specific experimental In our example we present result of Monte–Carlo simulation of EAS triggered by vertical CR proton with energy 10 TeV. The detector is at the 1650 m a.s.l.,"1082 altitude of Themus site (Freuch Pyrenees) or Baksan Neutrino Observatory (Russia). where the Narpet detector is now equipped with Chercukoy device.," altitude of Themis site (French Pyrenees) or Baksan Neutrino Observatory (Russia), where the Karpet detector is now equipped with Cherenkov device."1083 Our simulated detector is located in (αρ) position equal to (50 1. 0 11).," Our simulated detector is located in $x_{\qqq{d}}$ $y_{\qqq{d}}$ ) position equal to (–50 m, 0 m)."1084 We will consider two detector areas: 211 x 2r aud the circle of the area of 0.L1 1? both ceutred at Grp., We will consider two detector areas: 3 m x 3 m and the circle of the area of 0.44 $^{2}$ both centred at $x_{\qqq{d}}$ $y_{\qqq{d}}$ ).1085" OLL in? is the effective area of cach of the TUENUSTOCLE experiment detectorsTHEMISTOCLE. and heliostat mirrors in the CELESTE experiment would have effective area of & 31 ni? (of 5 Lin? for each mirror) The muon of euergv. 88.52 GeV has been fouud at position (53.07 1i. 3.59 1). with direction cosines e }.COS(O (cos(oyf)) equal to €.0.00196.0.00088) at AT, = 0.3 Iu the Figure 2aa we present the angular distribution of Chercukov photons which fall on the area Birx3iurceutredat (50 12. 0 11) and are produced iu the EAS described above."," 0.44 $^{2}$ is the effective area of each of the THEMISTOCLE experiment detectors, and heliostat mirrors in the CELESTE experiment would have effective area of $\approx$ 31 $^{2}$ (of 54 $^{2}$ for each mirror) The muon of energy 88.52 GeV has been found at position (–53.07 m, –3.59 m), with direction cosines $\alpha_{\qqq{x}}$ $\alpha_{\qqq{y}}$ )) equal to (–0.00196,0.00088) at $\Delta t_{\mu}$ = 0.3 In the Figure \ref{fig:all_ph}a a we present the angular distribution of Cherenkov photons which fall on the area 3 m x 3 m centred at (–50 m, 0 m) and are produced in the EAS described above."1086" This ia 2dimensional distribution for cos(a,.) aud costa) at the axes aud grey scale corresponding to the decimal logaritlin of the wmmber of photous iu the pixel (0.001 1uirad x 0.001 mrad).", This is a 2–dimensional distribution for $\alpha_{x}$ ) and $\alpha_{y}$ ) at the axes and grey scale corresponding to the decimal logarithm of the number of photons in the pixel (0.001 mrad x 0.001 mrad).1087 White colour indicates no photons iu the pixel., White colour indicates no photons in the pixel.1088 We have $y.=2/2.ay aud (5=7/2ay., We have $\theta_{\qqq{x}} = \pi /2 - \alpha_{\qqq{x}}$ and $\theta_{\qqq{y}} = \pi /2 - \alpha_{\qqq{y}}$.1089" For small zenith angle 0. costei) aud οσα) approximately correspoud to 0, aud Gy iu radiaus. the axes limits ave fron, 50 nirad to |50 nirad in x and v The result shown in the Figure 2aa corresponds to the theoretical registration of all photons produced in EAS by the imaging Cherenkov device at 501 away from the EAS core."," For small zenith angle $\theta$, $\alpha_{\qqq{x}}$ ) and $\alpha_{\qqq{y}}$ ) approximately correspond to $\theta_{\qqq{x}}$ and $\theta_{\qqq{y}}$ in radians, the axes limits are from –50 mrad to +50 mrad in x and y The result shown in the Figure \ref{fig:all_ph}a a corresponds to the theoretical registration of all photons produced in EAS by the imaging Cherenkov device at 50 m away from the EAS core."1090 The main pattern has an assvinetry which corresponds to the fact that most of photons are comine from the direction of the EAS axis (negative cos(os) values)., The main pattern has an assymetry which corresponds to the fact that most of photons are coming from the direction of the EAS axis (negative $\alpha_{\qqq{x}}$ ) values).1091 The perpendicular spread of photon directions (in this case along cosa) axis) is quite large. aud this fact is often used to discriminate between gamma aud ladron primary Around the point (0.015.0.015) there is an are of Cherenkov pliotous produced by the i1nnon.," The perpendicular spread of photon directions (in this case along $\alpha_{\qqq{y}}$ ) axis) is quite large, and this fact is often used to discriminate between gamma and hadron primary Around the point (0.015,0.015) there is an arc of Cherenkov photons produced by the muon."1092 The aneular radius of the are is about 20 iirad which indicates the relativistic particle. the ceutre of arc points to the direction of the mon mentioned above Iu the Figure 2bb we preseut the distribution of Clherenkov photons iu 2dimensional space of arrival time vs. height of photon cussion (these are the same photous as in the Fieure 2aa).," The angular radius of the arc is about 20 mrad which indicates the relativistic particle, the centre of arc points to the direction of the muon mentioned above In the Figure \ref{fig:all_ph}b b we present the distribution of Cherenkov photons in 2–dimensional space of arrival time vs. height of photon emission (these are the same photons as in the Figure \ref{fig:all_ph}a a)."1093 Most of plotous are produced 600 1i above the detector level or higher and arrive with Af— 2 us., Most of photons are produced 600 m above the detector level or higher and arrive with $\Delta t~>$ 2 ns.1094 The small spot at place Af — 0.5 1 ns aud production height 120 200 1m correspouds to the Cherenkov lelt enitted by the cucrectic umon., The small spot at place $\Delta t$ = 0.5 – 1 ns and production height 120 –200 m corresponds to the Cherenkov light emitted by the energetic muon.1095 The isolated spot with Af=6 7 us and production height 120 2001 corresponds to another. not very fast Iu the FieureOo 2cc we show the arrival time distribution (projection of the FigureOo 2bb on the time axis) The first peak in the histogram represents Chercukov plotous enutted by the energetic," The isolated spot with $\Delta t$ = 6 – 7 ns and production height 120 – 200 m corresponds to another, not very fast In the Figure \ref{fig:all_ph}c c we show the arrival time distribution (projection of the Figure \ref{fig:all_ph}b b on the time axis) The first peak in the histogram represents Cherenkov photons emitted by the energetic"1096"Civeu that Wü,~20dans| corresponds to Tuy.~10! K. the gas distribution would form an extended pressure-supported euvelope.","Given that $V_{\rm vir} \sim 20 \kms$ corresponds to $T_{\rm1097 vir} \sim 10^4$ K, the gas distribution would form an extended pressure-supported envelope."1098Depeudiug ou halo mass. this envelope would be barely confined by the halo potential. or even be completely evaporated from the halo (?)..,"Depending on halo mass, this envelope would be barely confined by the halo potential, or even be completely evaporated from the halo \citep{Barkana_Loeb99}."1099 Regardless. the gas component would not rena confined into the central disk. except perhaps near the very center where sclfshiclkling might be sufficient to preserve a tiny fraction of cold dense gas (?)..," Regardless, the gas component would not remain confined into the central disk, except perhaps near the very center where self-shielding might be sufficient to preserve a tiny fraction of cold dense gas \citep{Susa_Umemura04}."1100 Therefore. either the mergers would occur essenutiallv in a clissipationless regunue. as we have assmmed here. or they would contain a teuuous hot atmosphere which could be shock heated further as a result of the mereer (e.8..2) and become unbound thereafter in the absence of efficient cooling.," Therefore, either the mergers would occur essentially in a dissipationless regime, as we have assumed here, or they would contain a tenuous hot atmosphere which could be shock heated further as a result of the merger \citep[e.g.,][]{Kazantzidis_etal05} and become unbound thereafter in the absence of efficient cooling."1101 In either case. uo gas inflows and no appreciable triggered star formation would occur. contrary to what happens in mergers of massive galaxies (ee.7T). Ta sunuuuuw.," In either case, no gas inflows and no appreciable triggered star formation would occur, contrary to what happens in mergers of massive galaxies \citep[e.g,.][]{Barnes_Hernquist96}."1102 the effects of eas and star formation on the structure of the merecr renimnants are expected to be weak.," In summary, the effects of gas and star formation on the structure of the merger remnants are expected to be weak."1103 Therefore. although they will have to be coufirmed with lvdrodvuamical saulatious. the qualitative arguneuts outlined above suggest that our collisionless experiments have likely captured the esseuce of the morphological evolution iu mergers of diskv dawart ealanics.," Therefore, although they will have to be confirmed with hydrodynamical simulations, the qualitative arguments outlined above suggest that our collisionless experiments have likely captured the essence of the morphological evolution in mergers of disky dwarf galaxies."1104 The authors would lise to thank the referee. Matteo Alouelli. for constructive conuneuts on the mamuscript and the CLUES members S. CGottlóbbor. Y. Hoffiuan. and C. Yepes for providing the LG simulation.," The authors would like to thank the referee, Matteo Monelli, for constructive comments on the manuscript and the CLUES members S. Gottlöbber, Y. Hoffman, and G. Yepes for providing the LG simulation."1105 We also acknowledge stimulating discussions with Jürg Diecniuid. Alan MeConuachie. Chris Orban. and David Weinberg.," We also acknowledge stimulating discussions with Jürrg Diemand, Alan McConnachie, Chris Orban, and David Weinberg."1106 SAX. is supported bv the Center for Cosmology and Astro-Particle Plysics at The Ohio State University., S.K. is supported by the Center for Cosmology and Astro-Particle Physics at The Ohio State University.1107 This research was partially supported by the Polish National Scicuce Centre under eraut N N203 580910., This research was partially supported by the Polish National Science Centre under grant N N203 580940.1108" AJ. ds supported bv the Ministerio de Ciencia e lunovacion (MICTNN) in Spain through the Ramou v Cajal program. and further acknowledges support from erauts AYA 2009-13875-C03-02. AYA2009-12792-C'O03-03. aud CAM, $2009/ESP-1196."," A.K. is supported by the Ministerio de Ciencia e Innovacion (MICINN) in Spain through the Ramon y Cajal program and further acknowledges support from grants AYA 2009-13875-C03-02, AYA2009-12792-C03-03, and CAM S2009/ESP-1496."1109 This research was also. supported by the Ohio. Supercomputer Center (http://vww.osc.edu)., This research was also supported by the Ohio Supercomputer Center (http://www.osc.edu).1110"isochrones for metallicity Z=0.02, which are those used by R02 to obtain their mass values.","isochrones for metallicity $Z=0.02$, which are those used by R02 to obtain their mass values."1111" The physical parameters of the PMS member samples in and mentioned above, as well as the MS members obtained from the same photometric studies cited (Sharma et al."," The physical parameters of the PMS member samples in and mentioned above, as well as the MS members obtained from the same photometric studies cited (Sharma et al."1112" 2007; Walker 1956; R02), will be used as examples in the rest of the paper."," 2007; Walker 1956; R02), will be used as examples in the rest of the paper."1113 We now turn to a discussion of the results obtained for the physical parameters of the assigned member samples., We now turn to a discussion of the results obtained for the physical parameters of the assigned member samples.1114 The general results for the clusters are listed in Table 2.., The general results for the clusters are listed in Table \ref{t3}.1115 All the results in this table for which the use of models is involved are obtained with S00 models., All the results in this table for which the use of models is involved are obtained with S00 models.1116" As explained before (see DAYOT, and Sect."," As explained before (see DAY07, and Sect."1117" 2 above), the membership assignment to a particular star can happen with respect to several isochrones."," 2 above), the membership assignment to a particular star can happen with respect to several isochrones."1118" This provides several age and mass values for the particular candidate, and the averages and rms deviations of all these values are taken as the mass and age values, with uncertainties, of the particular PMS candidate member."," This provides several age and mass values for the particular candidate, and the averages and rms deviations of all these values are taken as the mass and age values, with uncertainties, of the particular PMS candidate member."1119 How do these values compare to previous results?., How do these values compare to previous results?.1120" In this section we refer to previous results on PMS masses, and discuss the differences between the values obtained using various evolutionary models."," In this section we refer to previous results on PMS masses, and discuss the differences between the values obtained using various evolutionary models."1121" To illustrate the results of our method to determine membership and the physical parameters of members, we discuss in this section the results on the masses of PMS members in2264.."," To illustrate the results of our method to determine membership and the physical parameters of members, we discuss in this section the results on the masses of PMS members in."1122" For this cluster the comparison to isochrones can be performed on the basis of a previously established member population by independent means, and it therefore offers a good possibility to test the reliability of our procedure."," For this cluster the comparison to isochrones can be performed on the basis of a previously established member population by independent means, and it therefore offers a good possibility to test the reliability of our procedure."1123" We extract two sets of mass values from each of the published works (R02, F06, D07)."," We extract two sets of mass values from each of the published works (R02, F06, D07)."1124" First, we have the three sets of published values, all of them obtained with S00 models, which we call P-values in the following."," First, we have the three sets of published values, all of them obtained with S00 models, which we call P-values in the following."1125 These three sets of P-values differ from each other., These three sets of P-values differ from each other.1126" Values from F06 and D07 do not differ systematically, although the dispersion of the difference is high (Mpo;—Mros=0.005+ 0.26)."," Values from F06 and D07 do not differ systematically, although the dispersion of the difference is high $M_\mathrm{D07}-M_\mathrm{F06}=0.005\pm0.26$ )."1127" The R02 values are systematically lower than both the F06 and D07 values (Mno2—Mpos=—0.14+0.26, Mao»—Mpo7 0.28)."," The R02 values are systematically lower than both the F06 and D07 values $M_\mathrm{R02}-M_\mathrm{F06}=-0.14\pm0.26$, $M_\mathrm{R02}-M_\mathrm{D07}=-0.22\pm0.28$ )."1128" A second set of mass values is obtained by shifting the transformed model isochrones in the CM V,(B—V) diagram, according to the published values of distance and absorption for every individual star, and reading its mass from the isochrone at nearest distance."," A second set of mass values is obtained by shifting the transformed model isochrones in the CM $V,(B-V)$ diagram, according to the published values of distance and absorption for every individual star, and reading its mass from the isochrone at nearest distance."1129 The mass values from this second procedure are named T-values., The mass values from this second procedure are named T-values.1130" P- and T-values should in principle be similar to each other, since they are obtained from the same observed colours, and are compared to the same isochrone models."," P- and T-values should in principle be similar to each other, since they are obtained from the same observed colours, and are compared to the same isochrone models."1131" However, they differ significantly."," However, they differ significantly."1132 The results attained by HO04 on star masses from different evolutionary models show that the masses of PMS stars predicted by all models are systematically lower than dynamical masses., The results attained by H04 on star masses from different evolutionary models show that the masses of PMS stars predicted by all models are systematically lower than dynamical masses.1133 In Fig., In Fig.1134 3 we reproduce the comparative plots produced by H04 in their Figs., \ref{MassPC} we reproduce the comparative plots produced by H04 in their Figs.1135 4 and 5., 4 and 5.1136" Here we plot the difference of T- minus P- values, versus P- values."," Here we plot the difference of T- minus P- values, versus P- values."1137 The plot shows that T-values are systematically higher than the P-values., The plot shows that T-values are systematically higher than the P-values.1138" On the other hand, the results by H04 show that the values from the 4-models (associated to our P-values here) are smaller than ""correct"" dynamical masses."," On the other hand, the results by H04 show that the values from the +models (associated to our P-values here) are smaller than ""correct"" dynamical masses."1139 The differences shown in our plots and in those by H04 are furthermore of the same order., The differences shown in our plots and in those by H04 are furthermore of the same order.1140" This means that the T-procedure produces higher mass values than the published P- values, and therefore mass results in better agreement with the dynamical values."," This means that the T-procedure produces higher mass values than the published P- values, and therefore mass results in better agreement with the dynamical values."1141the radiation pressure or hot eas pressure likely plaved this role iu 30 Doradus. decreasing the available mass to lake new stars and slowi19 star formation iu the region.,"the radiation pressure or hot gas pressure likely played this role in 30 Doradus, decreasing the available mass to make new stars and slowing star formation in the region."1142 Iu this paper. we have τιilized iiultiwavceleusth (radio. infrared. optical/UV. zx AX-oav) dnmaging to assess the role of several stellar fecback iiechanisiis im the eiaut ΠΠ region 30 Doradus in the LMC.," In this paper, we have utilized multi-wavelength (radio, infrared, optical/UV, and X-ray) imaging to assess the role of several stellar feedback mechanisms in the giant HII region 30 Doradus in the LMC."1143 In particular. we ave nieasured observatiouallv the pressures associated with possible sources of enerev and momentum to diive the dynamics of the region: the direct radiation rom stars. the dust-processed infrared radiation field. he warm ionized eas from massive stars. and the vot eas shock-heated Dy stellar winds aud superuovae.," In particular, we have measured observationally the pressures associated with possible sources of energy and momentum to drive the dynamics of the region: the direct radiation from stars, the dust-processed infrared radiation field, the warm ionized gas from massive stars, and the hot gas shock-heated by stellar winds and supernovae."1144" We have exploited the high-resolution. images of 30 Doradus to map these pressure components in LLL square regions, with dimensions of «/35.."," We have exploited the high-resolution images of 30 Doradus to map these pressure components in 441 square regions, with dimensions of $\times$."1145 We have ound that the direct radiation pressure from stars dominates at distances less than 75 pe from the central star cluster. whereas the wari ionized gas pressure dominates at lurger radi.," We have found that the direct radiation pressure from stars dominates at distances less than 75 pc from the central star cluster, whereas the warm ionized gas pressure dominates at larger radii."1146 By contrast. the ho eas pressure and the dust-processed radiation pressure do uot contribute significantly. indicating these compoucuts are not dvnuauicallv important.," By contrast, the hot gas pressure and the dust-processed radiation pressure do not contribute significantly, indicating these components are not dynamically important."1147 However. we cannot rule out that the hot eas pressure douunated at carly times and has become weaker with the WIT region expansion.," However, we cannot rule out that the hot gas pressure dominated at early times and has become weaker with the HII region expansion."1148 We have discussed two implications of our results: the partial confinement of the hot eas aud the dynamical role of radiation pressure in 30) Doradus., We have discussed two implications of our results: the partial confinement of the hot gas and the dynamical role of radiation pressure in 30 Doradus.1149 First. the weakness of the ταν gas pressure relative to the direct radiation pressure sugeests the hot eas is only partially confined and is leaking out of the pores iu the III shell.," First, the weakness of the X-ray gas pressure relative to the direct radiation pressure suggests the hot gas is only partially confined and is leaking out of the pores in the HII shell."1150 Secondly. the significant radiation pressure near the star cluster indicates that radiation pressure may have driven the expansion of the WIT shell at carly times.," Secondly, the significant radiation pressure near the star cluster indicates that radiation pressure may have driven the expansion of the HII shell at early times."1151 This result suggests observationally that radiation pressure nav be dynamically iu])ortant mb massive star clusters. reinforcing that radiation pressure is a viable mechanism to remove eas frou ΠΠ regions aud to regulate star formation.," This result suggests observationally that radiation pressure may be dynamically important in massive star clusters, reinforcing that radiation pressure is a viable mechanism to remove gas from HII regions and to regulate star formation."1152 Tudeed. if NGC 2070 was more massive. the radiation pressure coul even expel eas at hieh enough velocities to launch a ealactic wind (?)..," Indeed, if NGC 2070 was more massive, the radiation pressure could even expel gas at high enough velocities to launch a galactic wind \citep{wind}."1153 The work presented 100 qoa first step to MCASTIEC observationa∙dv the relaIve role of stellar feedback miechauisiis n stanford1 regions., The work presented here is a first step to measure observationally the relative role of stellar feedback mechanisms in star-forming regions.1154" Although we have applied our techuiques to One soluce, 30 Doradus. onr methods to extract dynamical information frou iuulti-wavelcngth nuages can be applied to other sources as well."," Although we have applied our techniques to one source, 30 Doradus, our methods to extract dynamical information from multi-wavelength images can be applied to other sources as well."1155" we plan fo perfor. these analyses ou allConsequently, the ΠΤΙ regious iu the LMC with available data to develop a road observational of these stellar feedback mechanisms aud understandingtheir role iu regulatiug star formation."," Consequently, we plan to perform these analyses on all the HII regions in the LMC with available data to develop a broad observational understanding of these stellar feedback mechanisms and their role in regulating star formation."1156 We would like to thank Bruce Draine. Naucder Ticlens. and Fernando Selman for helpful discussions.," We would like to thank Bruce Draine, Xander Tielens, and Fernando Selman for helpful discussions."1157 Also. we would like to thauk Jasmina Lazenucic-Calloway for eenerouslv providing the 3.5-cm ATCA image of 30 Doradus.," Also, we would like to thank Jasmina Lazendic-Galloway for generously providing the 3.5-cm ATCA image of 30 Doradus."1158 This work is supported by an AAUW American Dissertation Fellowship (LAL) aud by the National Science Foundation through grant NSE-ASTO0055836., This work is supported by an AAUW American Dissertation Fellowship (LAL) and by the National Science Foundation through grant NSF-AST0955836.1159 MRI acknowledges support from: an Alfred. P. Sloan Fellowship: NASA through ATFP erant NNNOQAB3LC: NASA as part of the Spitzer Theoretical. Research Program. through a coutract issued bv the JPL: the National Science Foundation through eraut AST-," MRK acknowledges support from: an Alfred P. Sloan Fellowship; NASA through ATFP grant NNX09AK31G; NASA as part of the Spitzer Theoretical Research Program, through a contract issued by the JPL; the National Science Foundation through grant AST-0807739."1160"SN remnant before the ejection to the ISM (N07), the grain-grain collision rate is enhanced in higher ng environments.","SN remnant before the ejection to the ISM (N07), the grain–grain collision rate is enhanced in higher $\nH$ environments."1161 The increase of grains with a<0.1um could efficiently affect the UV and optical extinction curves., The increase of grains with $a\la 0.1~\micron$ could efficiently affect the UV and optical extinction curves.1162 This point is quantitatively addressed in Section 3.2.., This point is quantitatively addressed in Section \ref{subsec:theor_extinc}.1163" Large grains with a>0.1um are marginally affected by shattering; namely, shattering of a small fraction of large grains can produce a large number of small grains."," Large grains with $a>0.1~\micron$ are marginally affected by shattering; namely, shattering of a small fraction of large grains can produce a large number of small grains."1164" In the case of HYOO, on the other hand, grains with a>0.1um are more shattered becauseabundant small grains in the MRN (Mathis,Rumpl,&Nordsieck1977) grain size distribution, which they assumed as the initial condition, enhance the grain-grain collision rate."," In the case of HY09, on the other hand, grains with $a>0.1~\micron$ are more shattered becauseabundant small grains in the MRN \citep{mathis77} grain size distribution, which they assumed as the initial condition, enhance the grain–grain collision rate."1165 The extinction curve of grains ejected from SNe II tends to be flat because small grains are efficiently destroyed in SNRs without escaping into the ISM (Hirashitaetal. 2008).., The extinction curve of grains ejected from SNe II tends to be flat because small grains are efficiently destroyed in SNRs without escaping into the ISM \citep{hirashita08}. .1166 Here we investigate, Here we investigate1167drives these shock oscillations?,drives these shock oscillations?1168 Surprisingly. the answer to this basic question remains somewhat uncertain.," Surprisingly, the answer to this basic question remains somewhat uncertain."1169 Indeed. two mechanisms have been proposed the advective-acoustic evcle (2?) ancl a purely acoustic mechanism (?) ancl distinguishing them remains a controversial issue.," Indeed, two mechanisms have been proposed – the advective-acoustic cycle \citep{foglizzo07} and a purely acoustic mechanism \citep{blondin06} – and distinguishing them remains a controversial issue."1170 The aim of this article is to clarify his question., The aim of this article is to clarify this question.1171 In Section 2.. we review previous works by explaining he two instability mechanisms and the arguments ooposed. to distinguish. them.," In Section \ref{sec:review}, we review previous works by explaining the two instability mechanisms and the arguments proposed to distinguish them."1172 On the basis of these works. the acdvective-acoustic evele is ecnerally favoured out. the purely. acoustic mechanism. cannot. be ruled out for realistic flow parameters.," On the basis of these works, the advective-acoustic cycle is generally favoured but the purely acoustic mechanism cannot be ruled out for realistic flow parameters."1173 In Section 3.. we oxopose a new method to extract a radial propagation ime from the frequency of the unstable modes.," In Section \ref{sec:radial_time}, we propose a new method to extract a radial propagation time from the frequency of the unstable modes."1174 For his purpose we compare the [frequencies of several iwnmonics that ciller only by their racial structure., For this purpose we compare the frequencies of several harmonics that differ only by their radial structure.1175 As he radial advective-acoustic and purely acoustic times are significantly. clilferent. this allows us to clistinguish »etween the two eveles and discards the purely acoustic mechanism.," As the radial advective-acoustic and purely acoustic times are significantly different, this allows us to distinguish between the two cycles and discards the purely acoustic mechanism."1176 1n Section +.. we describe a method. to compute purely acoustic modes. which are shown to oe stable.," In Section \ref{sec:acoustic_modes}, we describe a method to compute purely acoustic modes, which are shown to be stable."1177 In. Section r5.. we give a detailed: analysis of he frequency. spectra of advective-acoustic and. purely acoustic eveles. which confirms the validity of the method used in Section 3..," In Section \ref{sec:details}, we give a detailed analysis of the frequency spectra of advective-acoustic and purely acoustic cycles, which confirms the validity of the method used in Section \ref{sec:radial_time}."1178 Finally in Section 6.. we summarise our work and conclude that the instability mechanism behind SASL is the advective-acoustic evele.," Finally in Section \ref{sec:conclusion}, we summarise our work and conclude that the instability mechanism behind SASI is the advective-acoustic cycle."1179 SASL has been observed to grow in many cdillerent models of stellar collapse. which have very diverse degrees of complexity.," 	 SASI has been observed to grow in many different models of stellar collapse, which have very diverse degrees of complexity."1180 “Vhese include state of the art realistic numerical simulations (?7?).. as well as models where successive simplifications have been mace: approximate neutrino transport (7).. neglect of the stellar structure (???).. simple equation of state with the effect of neutrinos parameterised. by cooling/heating functions (227727277). adiabatic approximation (?).. ancl finally the most simplified of these models is the planar toy model (??)..," These include state of the art realistic numerical simulations \citep{marek09,burrows06}, as well as models where successive simplifications have been made: approximate neutrino transport \citep{scheck08}, neglect of the stellar structure \citep{ohnishi06,yamasaki07,iwakami08}, simple equation of state with the effect of neutrinos parameterised by cooling/heating functions \citep{blondin03,blondin06,foglizzo07,yamasaki08,fernandez09a,fernandez09b}, adiabatic approximation \citep{blondin07a}, and finally the most simplified of all these models is the planar toy model \citep{foglizzo09,sato09}."1181 Very often. allthe simpler the model. the deeper the ohvsical understanding.," Very often, the simpler the model, the deeper the physical understanding."1182 For example. models neglecting he heating have clearly shown that SASL is. distinct rom neutrino-driven convection (2)... a fact missed. by most realistic models.," For example, models neglecting the heating have clearly shown that SASI is distinct from neutrino-driven convection \citep{blondin03}, a fact missed by most realistic models."1183 Phe neglect of the stellar structure allows to set up an initial steady state. which can be studied with a perturbative analvsis (?22)..," The neglect of the stellar structure allows to set up an initial steady state, which can be studied with a perturbative analysis \citep{foglizzo07,yamasaki07}."1184 Finally. the simplest of these models. the planar tov mocdoel. allows for an analvtical treatment where the instability mechanism can be clearly demonstratect.," Finally, the simplest of these models, the planar toy model, allows for an analytical treatment where the instability mechanism can be clearly demonstrated."1185 Lt is widely accepted that the same instability is at work in all these models., It is widely accepted that the same instability is at work in all these models.1186 Therefore. although we base our investigation on the model of ?.. we expect our conclusions to hold more generally.," Therefore, although we base our investigation on the model of \citet{blondin06}, we expect our conclusions to hold more generally."1187 We choose this model because it is simple enough to perform a perturbative analvsis with clean boundary. conditions., We choose this model because it is simple enough to perform a perturbative analysis with clean boundary conditions.1188 This idealised flow was first described by ? and later studied by ? to describe SASL., This idealised flow was first described by \citet{houck92} and later studied by \citet{blondin06} to describe SASI.1189 It has then become a classical moclel of SASL and has been studied. also by 2?77..," It has then become a classical model of SASI and has been studied also by \citet{foglizzo07,fernandez09a,fernandez10,guilet10b}."1190 In this flow. he Uuiel is modelled as a non selt-gravitating perfect eas with an aciabatic index of 4/3.," In this flow, the fluid is modelled as a non self-gravitating perfect gas with an adiabatic index of 4/3."1191 The clleets of neutrinos are simply modelled by a cooling function. while neutrino jeating is neglected.," The effects of neutrinos are simply modelled by a cooling function, while neutrino heating is neglected."1192 The proto-neutron star is treated as a hard surface on which the cooling fluid. settles down., The proto-neutron star is treated as a hard surface on which the cooling fluid settles down.1193 The shock is stationary at a radius which can »e freely chosen by adjusting the normalisation of the cooling function., The shock is stationary at a radius which can be freely chosen by adjusting the normalisation of the cooling function.1194 This model and its near analysis will » extensively. used. in the remainder of the paper., This model and its linear analysis will be extensively used in the remainder of the paper.1195 We refer the reader to 7. for a description of the linear cigenmodes caleulation. and to Appendix A for a slightly different (but equivalent) formulation of the equations governing it. which has been used in this paper.," We refer the reader to \cite{foglizzo07} for a description of the linear eigenmodes calculation, and to Appendix \ref{sec:equations} for a slightly different (but equivalent) formulation of the equations governing it, which has been used in this paper."1196 The two mechanisms proposed to explain the growth of SASL are schematically illustrated by Figure 1.., 	 	 The two mechanisms proposed to explain the growth of SASI are schematically illustrated by Figure \ref{fig:lineaire:mecanismes}.1197 Phe advective-acoustic evele (left. panel) is a evele between two wavesc an advected wave composed of entropy ane vorticity and an acoustic wave propagating outward (2???7).," The advective-acoustic cycle (left panel) is a cycle between two waves: an advected wave composed of entropy and vorticity and an acoustic wave propagating outward \citep{foglizzo00,foglizzo02,blondin03,foglizzo07}."1198. Note that the acoustic wave also propagates in the transverse direction., Note that the acoustic wave also propagates in the transverse direction.1199 Phese waves are coupled via two coupling processes taking place at the shock and at a smaller radius in the deceleration region., These waves are coupled via two coupling processes taking place at the shock and at a smaller radius in the deceleration region.1200 The acoustic wave reaching the shock makes it oscillate and. creates the entropy-vorticity wave., The acoustic wave reaching the shock makes it oscillate and creates the entropy-vorticity wave.

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