ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 Inverse Compton scattering on the optical light dominates the resulting emission at 0.1 TeV. whereas the emission above | TeV is mainly produced via inverse Compton scattering off the IR photons.," Inverse Compton scattering on the optical light dominates the resulting emission at 0.1 TeV, whereas the emission above 1 TeV is mainly produced via inverse Compton scattering off the IR photons."3 On the other hand. we investigate the dynamical and radiative properties of G338.3-0.0 expanding into a medium with a density of 1 .," On the other hand, we investigate the dynamical and radiative properties of G338.3-0.0 expanding into a medium with a density of $1$ $^{-3}$ ."4 With Ej=0810 ere. the radius of the shell of the SNR in radio means that the age of the remnant is SOOO. yr.," With $E_{\rm ej}=0.8\times10^{51}$ erg, the radius of the shell of the SNR in radio means that the age of the remnant is $\sim 8000$ yr."5 With an initial spin-down power of 0.3«10°? ere 1 and other appropriate parameters (see, With an initial spin-down power of $0.3\times10^{39}$ erg $^{-1}$ and other appropriate parameters (see6Beein work on WFEIRST alter JWST is launched This is a reasonable response given NÀSA's fiscal constraints. but in its current form. makes no sense scientifically.,"Begin work on WFIRST after JWST is launched This is a reasonable response given NASA's fiscal constraints, but in its current form, makes no sense scientifically."7 For example. if EUCLID is lanuched in its current mega-nission form. the scientifie drivers for launching WFIRST 5 vears later will be muted at best.," For example, if EUCLID is launched in its current mega-mission form, the scientific drivers for launching WFIRST 5 years later will be muted at best."8 Ancl (his muting. combined with financial pressures and other competing scientific agendas. means thal WEIRST would in [act never be launched.," And this muting, combined with financial pressures and other competing scientific agendas, means that WFIRST would in fact never be launched."9 Nevertheless. (his response does contain (he germ of an approach (hat can vield outstanding science. while taking account of the financial and political realities.," Nevertheless, this response does contain the germ of an approach that can yield outstanding science, while taking account of the financial and political realities."10" As in the NASA response. there would be (wo satellites. but these would be explicitlv designed to be complementary (constituting a single ""mission) and would be linked by binding international agreements."," As in the NASA response, there would be two satellites, but these would be explicitly designed to be complementary (constituting a single “mission”) and would be linked by binding international agreements."11 The first satellite. undertaken by ESA. would be a purely optical wicle-field version of EUCLID. optimized for WL. ie.. with excellent resolution.," The first satellite, undertaken by ESA, would be a purely optical wide-field version of EUCLID, optimized for WL, i.e., with excellent resolution."12 The second satellite. undertaken bv NASA. would be a purely infrared telescope. basically in the model of WFIRST but with weaker constraints on image resolution (and so with larger field of view).," The second satellite, undertaken by NASA, would be a purely infrared telescope, basically in the model of WFIRST but with weaker constraints on image resolution (and so with larger field of view)."13" In the following subsections. P outline how Chis ""dual-satellite mission” is nearly opünmdl [from a science standpoint. while meeting all the financial and political constraints."," In the following subsections, I outline how this “dual-satellite mission” is nearly optimal from a science standpoint, while meeting all the financial and political constraints."14 There are four major scientific objectives. WL. BAO. SN (for DE) and pL (for planets).," There are four major scientific objectives, WL, BAO, SN (for DE) and $\mu$ L (for planets)."15 ] treat these in turn., I treat these in turn.16 WL would benefit (he most., WL would benefit the most.17 If the ESA mission were simplified to a wide-field optical imager. ib would be substantially cheaper aud easier to build and so would be launched sooner.," If the ESA mission were simplified to a wide-field optical imager, it would be substantially cheaper and easier to build and so would be launched sooner."18 Freed from the burden of infrared. (ad other) add-ons. it could probably have a larger [ield or a larger mirror or a longer operational lifetime (or all three). but in any case would be lanuchecl faster and at lower risk.," Freed from the burden of infrared (and other) add-ons, it could probably have a larger field or a larger mirror or a longer operational lifetime (or all three), but in any case would be launched faster and at lower risk."19 There would be immediate science return., There would be immediate science return.20 It is lane (hat the science would not be as strong as it would be with complementary inlrared (IR) photometry for photometric redshifts. but these would eventually be forthcoming when WEIBST was launched.," It is true that the science would not be as strong as it would be with complementary infrared (IR) photometry for photometric redshifts, but these would eventually be forthcoming when WFIRST was launched."21" So there would be initial excellent returns and improvements after several νους,", So there would be initial excellent returns and improvements after several years.22 The second stage of analysis based on It data is not likely to be cuickened (and may be further delaved) by having a single EUCLID mega-mission because it would be more complex. more expensive. and more subject to delavs or even cancellation.," The second stage of analysis based on IR data is not likely to be quickened (and may be further delayed) by having a single EUCLID mega-mission because it would be more complex, more expensive, and more subject to delays or even cancellation."23 Hence. there are substantial gains. reduced risk. and very little lost with this approach.," Hence, there are substantial gains, reduced risk, and very little lost with this approach."24 BAO and SN would also benefit., BAO and SN would also benefit.25 The basic point here is that the WFIRST-like mission in {his scenario is lauu1ched on the same timescale as within the NASA response., The basic point here is that the WFIRST-like mission in this scenario is launched on the same timescale as within the NASA response.26 Its IR capability, Its IR capability27AGN.,AGN.28 For example. water masers have been used to measure the mass of the central supermassive black hole (Miyoshietal.1995:Greenhilletal.2003a:Reid2009:Kuo 2010). the size and structure of the circumnuclear accretion disk (Greenhillal.2008). the shock speeds and densities of radio jets (Pecketal. 2003).. and to determine accurate geometrical distances to the host galaxies (Herrnsteinetal.1999:Braatz2010) — all of which are competitive with the results obtained using other methods.," For example, water masers have been used to measure the mass of the central supermassive black hole \citep{miyoshi95,greenhill03a,reid09,Kuo10}, the size and structure of the circumnuclear accretion disk \citep{greenhill03a,argon07,humphreys08,kondratko08}, the shock speeds and densities of radio jets \citep{peck03}, and to determine accurate geometrical distances to the host galaxies \citep{herrnstein99,braatz10} – all of which are competitive with the results obtained using other methods."29 Surveys for water maser systems at low redshifts ες< 00.06) have tended to focus on Seyfert 2 and LINER galaxies (e.g. Greenhilletal.2003b:Braatz 2003401.," Surveys for water maser systems at low redshifts $z\leq$ 0.06) have tended to focus on Seyfert 2 and LINER galaxies (e.g. \citealt{greenhill03b,braatz04}) )."30 This is because the expected edge-on orientation of the circumnuclear accretion disk for these galaxies should result in a larger path length of velocity-coherent molecular gas which is needed for the maser amplification., This is because the expected edge-on orientation of the circumnuclear accretion disk for these galaxies should result in a larger path length of velocity-coherent molecular gas which is needed for the maser amplification.31 The success rate of thesesurveys is of order —5 per cent. with around ~ 130 water maser galaxies currently known locally (e.g. Braatzetal.2008:: and unpublished).," The success rate of thesesurveys is of order $\sim$ 5 per cent, with around $\sim$ 130 water maser galaxies currently known locally (e.g. \citealt{braatz07}; and unpublished)."32 Surveys have also targeted nearby spiral and starburst galaxies with bright far-infrared (FIR) flux-denisties (Henkeletal.2005:Castangia2008:Surcisetal. 2009).," Surveys have also targeted nearby spiral and starburst galaxies with bright far-infrared (FIR) flux-denisties \citep{henkel05,castangia08,surcis09}."33. These surveys have found [5 water maser galaxies and have a detection rate of 23 per cent (Surcis 2009)., These surveys have found 15 water maser galaxies and have a detection rate of 23 per cent \citep{surcis09}.34 A snapshot survey of 611 luminous active and non-active galaxies in the local Universe (e kKm) was carried out by Braatz&Gugliucei(2008)., A snapshot survey of 611 luminous active and non-active galaxies in the local Universe $v \leq$ km) was carried out by \citet{braatz08}.35 They found only 8 Water maser systems (a detection rate of about ~ | per cent). which demonstrates the rarity of these objects.," They found only 8 water maser systems (a detection rate of about $\sim$ 1 per cent), which demonstrates the rarity of these objects."36 Finding water maser galaxies at higher redshifts has the potential to constrain the nature of dark energy through the accurate measurement of geometrical distances., Finding water maser galaxies at higher redshifts has the potential to constrain the nature of dark energy through the accurate measurement of geometrical distances.37 However. due to the limited sensitivities and the small range of frequencies covered by current radio telescopes. surveys for high redshift water maser galaxies had yielded non-detections (Bennertetal. 2009).. with the notable exception being the detection of the most luminous water maser system known ina type 2 quasar at redshift 0.66 (Barvainis&Antonucci2005).," However, due to the limited sensitivities and the small range of frequencies covered by current radio telescopes, surveys for high redshift water maser galaxies had yielded non-detections \citep{bennert09}, , with the notable exception being the detection of the most luminous water maser system known in a type 2 quasar at redshift 0.66 \citep{barvainis05}."38. To overcome the sensitivity limitations we have started a search for water masers from known gravitationally lensed quasars., To overcome the sensitivity limitations we have started a search for water masers from known gravitationally lensed quasars.39 Observing only gravitationally lensed active galaxies has the advantage of pre-selecting a population of distant AGN (22 L) and of using the magnification provided by the foreground gravitational lensing galaxy to increase the observed flux density of any water masers in the background AGN., Observing only gravitationally lensed active galaxies has the advantage of pre-selecting a population of distant AGN $z > 1$ ) and of using the magnification provided by the foreground gravitational lensing galaxy to increase the observed flux density of any water masers in the background AGN.40 Such lensing magnifications can range up to hundreds. depending on the alignment of the lens and the background source.," Such lensing magnifications can range up to hundreds, depending on the alignment of the lens and the background source."41 The gravitational lensing technique has already been used successfully to study the interstellar medium and molecular gas content of high redshift quasars through the detection of. for example. carbon monoxide (Barvainisetal.1994).," The gravitational lensing technique has already been used successfully to study the interstellar medium and molecular gas content of high redshift quasars through the detection of, for example, carbon monoxide \citep{barvainis94}."42. We found a luminous water maser in the first gravitationally lensed quasar we observed. in the type | quasar MG JO04144-0534. at redshift 2.64 tImpellizzerietal.20085)... which is by far the most distant object to show water vapour emission.," We found a luminous water maser in the first gravitationally lensed quasar we observed, in the type 1 quasar MG J0414+0534 at redshift 2.64 \citep{impellizzeri08}, which is by far the most distant object to show water vapour emission."43 The maser line from MG J041---0534. is broad with a full width at half maximum (FWHM) of ~ 445 kmss+. is blueshifted by 300 ! from the systemic velocity of the quasar and has an cunlensed) apparent isotropic luminosity of ~ 88300 £L...," The maser line from MG J0414+0534 is broad with a full width at half maximum (FWHM) of $\sim$ 45 $^{-1}$, is blueshifted by $-$ 300 $^{-1}$ from the systemic velocity of the quasar and has an (unlensed) apparent isotropic luminosity of $\sim$ 8300 $L_{\odot}$."44 Our initial hypothesis. based on spectra obtained with the Expanded Very Large Array (EVLA) and the Effelsberg radio telescope. was that the maser originates in the relativistic jet of this quasar as it interacts with a molecular cloud lying close to the supermassive black hole.," Our initial hypothesis, based on spectra obtained with the Expanded Very Large Array (EVLA) and the Effelsberg radio telescope, was that the maser originates in the relativistic jet of this quasar as it interacts with a molecular cloud lying close to the supermassive black hole."45 Follow-up VLBI and monitoring observations are currently being analyzed to investigate this hypothesis and will be presented in forthcoming papers., Follow-up VLBI and monitoring observations are currently being analyzed to investigate this hypothesis and will be presented in forthcoming papers.46 We have now begun a survey of other distant gravitationally lensed quasars. and have also extended our search to lensed starbursting galaxies. using the Effelsberg radio telescope and the EVLA.," We have now begun a survey of other distant gravitationally lensed quasars, and have also extended our search to lensed starbursting galaxies, using the Effelsberg radio telescope and the EVLA."47 The aim of this survey is to find additional detections of high redshift water masers. the results of which are presented in this paper.," The aim of this survey is to find additional detections of high redshift water masers, the results of which are presented in this paper."48 In the case of the only currently contirmed high redshift water maser galaxy. MG JO414+0534. we know that the source is highly dust reddened and hosts a powerful AGN that is emitting a relativistic jet.," In the case of the only currently confirmed high redshift water maser galaxy, MG J0414+0534, we know that the source is highly dust reddened and hosts a powerful AGN that is emitting a relativistic jet."49 This suggests that luminous water masers that are associated with AGN are more likely to be found in high redshift galaxies with large amounts of dust and/or powerful radio jets., This suggests that luminous water masers that are associated with AGN are more likely to be found in high redshift galaxies with large amounts of dust and/or powerful radio jets.50 To investigate the first of these possible effects. we have limited the sample studied in this paper to include only radio-quiet quasars ane star-forming galaxies with previously measured strong far-infrarec (FIR) emission or molecular CO line emission.," To investigate the first of these possible effects, we have limited the sample studied in this paper to include only radio-quiet quasars and star-forming galaxies with previously measured strong far-infrared (FIR) emission or molecular CO line emission."51 For all calculations we adopt an. (23= 0.3. Ὃν= 0.7 spatially flat Universe. with a Hubble constan of Ly= 70 + and a solar luminosity of L.—3939. 10° W. We targeted and prioritized those gravitational lens systems with i) background AGN or star-forming galaxies at redshifts between 2.29 and 2.87 (due to the bandpass of the receiver that was used). 1) had high lensing magnitications and 11) showed previous detections of molecular emission (e.g. CO. HCN. ete.)," For all calculations we adopt an $\Omega_{M} =$ 0.3, $\Omega_{\Lambda}=$ 0.7 spatially flat Universe, with a Hubble constant of $H_{0} =$ 70 $^{-1}$ $^{-1}$, and a solar luminosity of $L_{\odot}=$ 3.939 $\times$ $^{26}$ W. We targeted and prioritized those gravitational lens systems with i) background AGN or star-forming galaxies at redshifts between 2.29 and 2.87 (due to the bandpass of the receiver that was used), ii) had high lensing magnifications and iii) showed previous detections of molecular emission (e.g. CO, HCN, etc.)"52 or a large FIR luminosity., or a large FIR luminosity.53 We also limited our sample to those objects that were observable with the Effelsberg radio telescope and the EVLA., We also limited our sample to those objects that were observable with the Effelsberg radio telescope and the EVLA.54 The resulting sample is not statistically complete., The resulting sample is not statistically complete.55 RX 1091120551 is comprised of four images of an X-ray detected quasar at redshift 2.80 (Badeetal.1997:Burud1998)... which is gravitationally lensed by a group of galaxies at redshift 0.769 (Kneibetal.2000).," RX J0911+0551 is comprised of four images of an X-ray detected quasar at redshift 2.80 \citep{bade97,burud98}, which is gravitationally lensed by a group of galaxies at redshift 0.769 \citep{kneib00}."56. The total magnification of the quasar is ~ 222., The total magnification of the quasar is $\sim$ 22.57" CO (3-2) has been tentatively detected by Hainlineetal.(2004) and the large sub-mm luminosity of the lensed quasar suggests a cold dust mass of ~ 10""A7. (Barvainis&Ivison2002)."," CO (3–2) has been tentatively detected by \citet{hainline04} and the large sub-mm luminosity of the lensed quasar suggests a cold dust mass of $\sim$ $^{8}\,M_{\odot}$ \citep{barvainis02}."58. IRAS [021444724 is a highly magnified and dust obscuredgravitationally lensed quasar at redshift 2.285 (Rowan-Robinsonal. 1991)., IRAS 10214+4724 is a highly magnified and dust obscuredgravitationally lensed quasar at redshift 2.285 \citep{rowan-robinson91}.59. This system has four lensed images that are formed by a massive foreground lensing galaxy at redshift 0.893 (Lacey.Rawl-ings&Serjeant 1998)., This system has four lensed images that are formed by a massive foreground lensing galaxy at redshift 0.893 \citep*{lacy98}.60. The magnification of the three merging lensed images is & 1100 (e.g. Broadhurst&Lehar 19955)., The magnification of the three merging lensed images is $\la$ 100 (e.g. \citealt{broadhurst95}) ).61 The lensed source is extremely luminous in the FIR (Rowan-Robinsonetal.1991). and sub-mm (Barvainis&Ivison2002).. and optical polarimetry found evidence of scattered light (Lawrence 1993)).," The lensed source is extremely luminous in the FIR \citep{rowan-robinson91} and sub-mm \citep{barvainis02}, and optical polarimetry found evidence of scattered light \citealt{lawrence93}) )."62 A large quantity of molecular gas from CO (Brown&Van- 1995).. HCN (VandenBoutetal.2004). and. neutral carbon (Weilet has also been reported.," A large quantity of molecular gas from CO \citep{brown91,brown92,solomon92,downes95}, , HCN \citep{vandenbout04} and neutral carbon \citep{weiss05a} has also been reported."63 Finally. near-infrared imaging spectroscopy found evidence for both circumnuclear star formation and an AGN from narrow (butspatially extended) and broad Ha emission (Krokeretal. 1996)..," Finally, near-infrared imaging spectroscopy found evidence for both circumnuclear star formation and an AGN from narrow (butspatially extended) and broad $\alpha$ emission \citep{kroker96}. ."64 All of these data are consistent with the central engine of the AGN being obscured and there being nuclear driven star-formation in a region rich in molecular gas., All of these data are consistent with the central engine of the AGN being obscured and there being nuclear driven star-formation in a region rich in molecular gas.65and second. particularly at late times. by the ionizine cllect of the white ciwarf radiation field.,"and second, particularly at late times, by the ionizing effect of the white dwarf radiation field."66 The observed. photosphere in the isothermal models is initially aclvected outward. with the Yow before. the decreasing density causes the opacity to drop. ancl the photosphere to shrink to zero size., The observed photosphere in the isothermal models is initially advected outward with the flow before the decreasing density causes the opacity to drop and the photosphere to shrink to zero size.67 This gives rise to a liehteurve that rises. and then falls., This gives rise to a lightcurve that rises and then falls.68 Emission lines and edges arise because the photospheric radius is larger at wavelengths. with higher opacity., Emission lines and edges arise because the photospheric radius is larger at wavelengths with higher opacity.69 Surfaces. of constant Doppler shift are perpendicular to the line of sight and hus Gaussian density profiles give rise to Gaussian. velocity »ofiles., Surfaces of constant Doppler shift are perpendicular to the line of sight and thus Gaussian density profiles give rise to Gaussian velocity profiles.70 Direct application of our LPE fitting. method to he uv region would not reproduce the observed. spectra with their variety of ionization states and. semi-forbidden and. permitted. lines., Direct application of our LTE fitting method to the uv region would not reproduce the observed spectra with their variety of ionization states and semi-forbidden and permitted lines.71 We have outlined techniques and »ossible improvements to the model that will increase our understanding further anc enable us to. take the work into the uv., We have outlined techniques and possible improvements to the model that will increase our understanding further and enable us to take the work into the uv.72 More detailed modelling of the ionization and recombination processes is sugeested that also incorporates the heating ellects of both the white dwarl’s ancl the fireballs own photon fields., More detailed modelling of the ionization and recombination processes is suggested that also incorporates the heating effects of both the white dwarf's and the fireballs own photon fields.73 The self-consistent solution to these rate equations will extend the treatment to the non-LPE regimes discussed in Section 5 and Appendix A.., The self-consistent solution to these rate equations will extend the treatment to the non-LTE regimes discussed in Section \ref{sec:discuss} and Appendix \ref{sec:ltevalidity}.74 These modifications may in turn have consequences for the predicted optical spectra., These modifications may in turn have consequences for the predicted optical spectra.75 Improved. modelling of these fireballs ollers us. the chance to probe the chemical composition of the secondary star in NIS. er., Improved modelling of these fireballs offers us the chance to probe the chemical composition of the secondary star in AE Aqr.76 Normally this ds. dillieult because the spectrum of the relatively dim secondary star is contaminated by light from other components in the system., Normally this is difficult because the spectrum of the relatively dim secondary star is contaminated by light from other components in the system.77 We expect that the fireball models will be applicable more generally to the Dickering observed in most CV. systems., We expect that the fireball models will be applicable more generally to the flickering observed in most CV systems.78 We would like to thank Gary Ferlancl ancl Wirk Ixorista for informative discussions regarding atomic ionization aud recombination processes and the routines for calculating them., We would like to thank Gary Ferland and Kirk Korista for informative discussions regarding atomic ionization and recombination processes and the routines for calculating them.79 We are grateful to the anonvmous referee and editor for helpful comments that improved the presentation of this paper., We are grateful to the anonymous referee and editor for helpful comments that improved the presentation of this paper.80binding energies areposilive’.. and where the efficiency of conversion of ionization energv io kinetic energv may differ [rom the elliciency of conversion of other forms of internal energv.,"binding energies are, and where the efficiency of conversion of ionization energy to kinetic energy may differ from the efficiency of conversion of other forms of internal energy."81 As we can see from Eqs. (1)), As we can see from Eqs. \ref{equ1}) )82 ancl (2)). the calenlated values of A depend on the stellar mass. (he core mass. the envelope mass distribution ancl the percentage of internal enerev (hat contributes to envelope ejection.," and \ref{equ2}) ), the calculated values of $\lambda$ depend on the stellar mass, the core mass, the envelope mass distribution and the percentage of internal energy that contributes to envelope ejection."83 Factors like metallicity ancl stellar wind mass loss also affect the stellar evolution and the resulting values of A., Factors like metallicity and stellar wind mass loss also affect the stellar evolution and the resulting values of $\lambda$.84 Additionally. the different definitions of core-envelope boundary lead (o different stellar core mass and therefore allect ihe ealeulated results.," Additionally, the different definitions of core-envelope boundary lead to different stellar core mass and therefore affect the calculated results."85 Our results agree reasonably with that of [anetal.(1994).. Dewi&Tauris(2000) and Dewi&Tauris(2001).. although Chere exist some clifferences.," Our results agree reasonably with that of \citet{han94}, \citet{dew00} and \citet{dew01}, although there exist some differences."86 proposed that A is almost independent of (he stellar chemical composition., \citet{dew00} proposed that $\lambda$ is almost independent of the stellar chemical composition.87 From our calculations we find that given the same mass. Pop.," From our calculations we find that given the same mass, Pop."88 I and Pop., I and Pop.89 IHE stars have different A-values (see Figs., II stars have different $\lambda$ -values (see Figs.90 1. and 2))., \ref{fig1} and \ref{fig2}) ).91" For example. (he mass range in which the values of Aj, can become negative is narrower for Pop."," For example, the mass range in which the values of $\lambda_{\rm b}$ can become negative is narrower for Pop."92 I stars (han lor Pop., II stars than for Pop.93 1 ones., I ones.94 In Fig., In Fig.95 3. we compare the evolution of the two As for a SAL. star with different metallicitv., \ref{fig3} we compare the evolution of the two $\lambda$ s for a $5 M_\odot$ star with different metallicity.96 It is seen that the A evolutions show little difference at the beginning. but later thev diverge [rom each other. and (he values of A lor the Pop.," It is seen that the $\lambda$ evolutions show little difference at the beginning, but later they diverge from each other, and the values of $\lambda$ for the Pop."97 I star always lie above those of the Pop., I star always lie above those of the Pop.98 II star in late evolutionary stages., II star in late evolutionary stages.99 The reason is as follows., The reason is as follows.100 The binding energy. (absolute value) of the Pop., The binding energy (absolute value) of the Pop.101 I star is either comparable with (in early evolutionary stage) or lower than (due to the contribution of ionization energv) that of the Pop., I star is either comparable with (in early evolutionary stage) or lower than (due to the contribution of ionization energy) that of the Pop.102 II star., II star.103 Since the Pop., Since the Pop.104 I star evolves, I star evolves105We have also used the magnitude histograms of the 6601 central region to define the faintest magnitude reachable for cach band.,We have also used the magnitude histograms of the 604 central region to define the faintest magnitude reachable for each band.106 The magnitude limit was set equal to the magnitude of the first bin brighter than the peak of the histogram plot (c.g.. see Figure 6).," The magnitude limit was set equal to the magnitude of the first bin brighter than the peak of the histogram plot (e.g., see Figure 6)."107 Folowing this criterion. the magnitude limits are 21.5Γ mae forJ baud and 20.5 mae for Mand bands.," Following this criterion, the magnitude limits are 21.5 mag for band and 20.5 mag for and bands."108 Figure 6 shows the ustoeram obtained for the heuid where i rant be observed that 20.5 mag biu cotues right before the bin with the highest counts (21.0 mae)., Figure \ref{fig:k_hist} shows the histogram obtained for the band where it can be observed that 20.5 mag bin comes right before the bin with the highest counts (21.0 mag).109 With tjose magnitude limits our sample is reduced to 169 objects., With these magnitude limits our sample is reduced to 1649 objects.110 However. tle photometric uncertainies of the remaining objecs are reduced to 0.01. 0.01. and 0.07 mae.," However, the photometric uncertainties of the remaining objects are reduced to 0.04, 0.04, and 0.07 mag."111 forJ.JL aud Ay. respectively (rearly half of the mean maguitude uncertainties or the whole saniple. see Table 1)).," for, and , respectively (nearly half of the mean magnitude uncertainties for the whole sample, see Table \ref{tab:filters}) )."112 The astrometry was made in two steps. firs an aliguinent iu dare cooinates’ of the finalJ.1. and mages: then a sasequent transformalon from iniage coordinate svstem to World Coordilate System (WCS).," The astrometry was made in two steps, first an alignment in `image coordinates' of the final, and images; then a subsequent transformation from image coordinate system to World Coordinate System (WCS)."113" The traustormation from ""iniage coordinates’ to (6.8) celesial equatorial coordinates was performed using and IRAF tasks. which computed the plate solulon and created the WCS 1uage. respectively."," The transformation from `image coordinates' to $\alpha$ $\delta$ ) celestial equatorial coordinates was performed using and IRAF tasks, which computed the plate solution and created the WCS image, respectively."114 The astronietry was derived using 15 isolated stars. cobunon to our nuages and version 2.3.2 of the Cade Star Catalog (€SC).," The astrometry was derived using 45 isolated stars, common to our images and version 2.3.2 of the Guide Star Catalog (GSC)."115 Special care was taken to select well separated stars in our NIRI field., Special care was taken to select well separated stars in our NIRI field.116 The reference frame of the GSC catalog is the Tuternational Celestial Reference. Frame aid Equinox J2000.0 (2).., The reference frame of the GSC catalog is the International Celestial Reference Frame and Equinox J2000.0 \citep{2008AJ....136..735L}.117 The CCS 2.3.2 typical errors ire 073 aud the trausformatioi RMS are 0.735 aud 0.725 for a aud à. respectively.," The GCS 2.3.2 typical errors are $^{\prime\prime}$ 3 and the transformation RMS are $^{\prime\prime}$ 35 and $^{\prime\prime}$ 25 for $\alpha$ and $\delta$, respectively."118 The imaguitudes obtained iu the individual filters were matched in :| unique list coutainiug 3627 objects in the fek Lan which all threeJ.II. and maenitudes8 were measured.," The magnitudes obtained in the individual filters were matched in a unique list containing 3627 objects in the field in which all three, and magnitudes were measured."119 Table 2 is an excerpt of the photoimetrv table available in its entietv as machine-readable table in the electronic eclition of the Astronomical Journal., Table \ref{tab:latabla} is an excerpt of the photometry table available in its entirety as machine-readable table in the electronic edition of the Astronomical Journal.120 Coluun 1 contains the iuterual object id: Colas 2 and 3 the Celestial Equatorial coordinates (J2000.0) iu degrees: €'oluuimns lL and 5 theJ imnaguitudeOo and dts uuceralutyv. respectively: Columns 6 and 7 theII magnitude and its uncertainty. respectively: Cohuuns Ss aud 9 the uaenitude and its uncertainty. respectively: aud Columns 10 aud 11 list the (€J-ID) aud Ay) colors. respectivelv.," Column 1 contains the internal object id; Columns 2 and 3 the Celestial Equatorial coordinates (J2000.0) in degrees; Columns 4 and 5 the magnitude and its uncertainty, respectively; Columns 6 and 7 the magnitude and its uncertainty, respectively; Columns 8 and 9 the magnitude and its uncertainty, respectively; and Columns 10 and 11 list the ) and ) colors, respectively."121" The otal effective area - takingο into account the ditherimeo pattern - covered by our photometry is of 107""«107 (στους 130 pe? at the distance of As meutioned in Section 1. ortunately, NGC6601 is situated in a privikved. location from an observational poiut of vlew: lu a direction ou of the Galactic plane arxb in the outer part of 333 (a ealaxy which is orientated almost face-ou): hence. it is expectecL that in the 660 field there is little coutanimajon with Oreeroun and backerouud stars."," The total effective area - taking into account the dithering pattern - covered by our photometry is of $^{\prime\prime}\times$ $^{\prime\prime}$ $\sim$ $\times$ 430 $^2$ at the distance of As mentioned in Section 1, fortunately, 604 is situated in a privileged location from an observational point of view: in a direction out of the Galactic plane and in the outer part of 33 (a galaxy which is orientated almost face-on); hence, it is expected that in the 604 field there is little contamination with foreground and background stars."122 We have estimated the contamination bv Galactic objects in our fick using the model of ?jesancon.dfr/. which vields only six objects. LOS of them E. dwart stars belonging to the hick disk.," We have estimated the contamination by Galactic objects in our field using the model of \cite{2003A&A...409..523R}, which yields only six objects, most of them F dwarf stars belonging to the thick disk."123 Our main concern for the study oeseuted. in this paper ds the contaiuination w field objects lie iu the area of he color-color (CC) aud color-maenitude (CAL) diagranx:4. occupied by MYSOs candidates. resulting iu yossible nüsideutificatious and an overestimation of their unnber.," Our main concern for the study presented in this paper is the contamination by field objects lie in the area of the color-color (CC) and color-magnitude (CM) diagrams occupied by MYSOs candidates, resulting in possible misidentifications and an overestimation of their number."124 To examine this effect we eecnerated density plots (two dimensional histogram) or the CC diagram of the 6601. central region and tie. field region surroundius it. as defined in the previous οςction.," To examine this effect we generated density plots (two dimensional histograms) for the CC diagram of the 604 central region and the field region surrounding it, as defined in the previous Section."125 Figure 7 exhibits he CC densiVo lnaps for the central region (top xucel.thefied region scaled to mach the ceutral reglon area {widdle panel) aud ti6 bin to biu subtraction resulting of the ceutra reelon nuinus he field regio1 (bottom panel).," Figure \ref{fig:control} exhibits the CC density maps for the central region (top panel), the field region scaled to match the central region area (middle panel), and the bin to bin subtraction resulting of the central region minus the field region (bottom panel)."126 Eaci bin in these diaegranis is a «0.05 mae., Each bin in these diagrams is a $\times$ 0.05 mag.127 sqare. in (ΠΠ) and (I-A).," square, in ) and )."128 There are three visile differences )etween the oject distribution in fre CC density naps from |voth regions. which become even nore evideut in the botOl Pallet (a) the star opulation around (J-IT)-(0.0 anc Iso. which is preseut in the diagran «4 the ceutral region but uussine 1u he field region.," There are three visible differences between the object distribution in the CC density maps from both regions, which become even more evident in the bottom panel: (a) the star population around $\sim$ 0.0 and $\sim$ 0.0, which is present in the diagram of the central region but missing in the field region."129 These, These1302008)).,).131 Llowever. the photon noise and the pixelation cllect remain to be treated in a more systematic way.," However, the photon noise and the pixelation effect remain to be treated in a more systematic way."132 For example. existing model fitting methods 3ricdlectal.2001:akajima&Bernstein2007:itchingetal. 2008)) use the variance of the noise to weight the pixels in their chi-square ittings.," For example, existing model fitting methods, \citealt{bridle01,bernstein02,kuijken06,miller07,nakajima07,kitching08}) ) use the variance of the noise to weight the pixels in their chi-square fittings."133 Lt is not clear to what level the noise contamination o the shear recovery can be removed in this way. especially or correlated background: noise (sec.c.g... Massey.etal. 20073).," It is not clear to what level the noise contamination to the shear recovery can be removed in this way, especially for correlated background noise (see, \citealt{massey07}) )."134 By integrating the model over the pixels. the mocel itting methods essentially use linear interpolations to treat he pixelation effect.," By integrating the model over the pixels, the model fitting methods essentially use linear interpolations to treat the pixelation effect."135 Phis is found. not accurate as we will discuss in refpixelation.., This is found not accurate as we will discuss in \\ref{pixelation}.136" Indeed. as we will discuss later in the paper. here are two types of noise: the astronomical photon noise and the ""photon counting"" shot noise."," Indeed, as we will discuss later in the paper, there are two types of noise: the astronomical photon noise and the ""photon counting"" shot noise."137 While the second type diminishes when the exposure tine increases. the first type does not.," While the second type diminishes when the exposure time increases, the first type does not."138 We will mainly focus on the astronomical noise in this paper., We will mainly focus on the astronomical noise in this paper.139 For the counting shot noise. we will simply argue in& refsummary that its contamination to the shear recovery can be significantly suppressed. by increasing the exposure time.," For the counting shot noise, we will simply argue in \\ref{summary} that its contamination to the shear recovery can be significantly suppressed by increasing the exposure time."140 Since these issues are not specifically addressed in any previous weak lensing literatures. it. is iniportant to. point them out in this paper.," Since these issues are not specifically addressed in any previous weak lensing literatures, it is important to point them out in this paper."141 1n à recent work bv Zhang (2008) (Z08 hereafter). a new and simple way of measuring the cosmic shear is found.," In a recent work by Zhang (2008) (Z08 hereafter), a new and simple way of measuring the cosmic shear is found."142 Hs main advantages includes: 1., Its main advantages includes: 1.143 it is mathematically simple: 2., it is mathematically simple; 2.144 it is free of assumptions on the morphologies of the galaxies and the PSL 3., it is free of assumptions on the morphologies of the galaxies and the PSF; 3.145 it enables us to probe the shear information rom galaxy substructures. thereby improving the signal-to-noise ratio.," it enables us to probe the shear information from galaxy substructures, thereby improving the signal-to-noise ratio."146 Phese facts encourage us to extend the metho urther by including a treatment of the photon noise am he pixelation effect., These facts encourage us to extend the method further by including a treatment of the photon noise and the pixelation effect.147 Fortunately. we find that these two vpes of systematic errors can be treated in a simple ai mocel-independent wav based on the method. of 205.," Fortunately, we find that these two types of systematic errors can be treated in a simple and model-independent way based on the method of Z08."148 As will become clear later in this paper. the method. we adop o remove the noise contamination can also be considere in other shear measurement methods. and our treatment of he pixelation elfect is Agenerally useful for image Serprocessing5 of all purposes.," As will become clear later in this paper, the method we adopt to remove the noise contamination can also be considered in other shear measurement methods, and our treatment of the pixelation effect is generally useful for image processing of all purposes."149 The paper is organized as follows: in relreview.. we briclly review the shear measurement method of Z0S: in refsystematies.. we show how to treat the photon noise (in refnoise)) and the pixelation οσοι (in refpixelation)) in weak lensing: in refnumerics.. we use computer-generated mock galaxy images to test the performance of our method: finally. we summarize and discuss remaining issues in re Summary.," The paper is organized as follows: in \\ref{review}, we briefly review the shear measurement method of Z08; in \\ref{systematics}, we show how to treat the photon noise (in \\ref{noise}) ) and the pixelation effect (in \\ref{pixelation}) ) in weak lensing; in \\ref{numerics}, we use computer-generated mock galaxy images to test the performance of our method; finally, we summarize and discuss remaining issues in \\ref{summary}."150 ZOS proposes a way of measuring the cosmic shear with the spatial cerivatives of the galaxy surface brightness field., Z08 proposes a way of measuring the cosmic shear with the spatial derivatives of the galaxy surface brightness field.151 To do so. let us define the surface brightness on the image plane as fy(a) and that on the source plane as nom ). where i and 6 are the position angles on the image and source plane respectively.," To do so, let us define the surface brightness on the image plane as $f_I(\vti)$, and that on the source plane as $f_S(\vts)$ , where $\vti$ and $\vts$ are the position angles on the image and source plane respectively."152 These quantities are related in a simple way as: where Aj;=3$;|d; and d;;=0381/087 are the spatial derivatives of the lensing dellection angle.," These quantities are related in a simple way as: where $\mathbf{A}_{ij}=\delta_{ij}+\Phi_{ij}$, and $\Phi_{ij}=\partial\delta\theta^I_i/\partial\theta^S_j$ are the spatial derivatives of the lensing deflection angle."153" Matrix A is often expressed in terms of the convergence s=(0,|,,)/2 and the two shear components 5,=(?,,Dig2lf? and 5»=@,,,.", Matrix $\mathbf{A}$ is often expressed in terms of the convergence $\kappa=(\Phi_{xx}+\Phi_{yy})/2$ and the two shear components $\gamma_1=(\Phi_{xx}-\Phi_{yy})/2$ and $\gamma_2=\Phi_{xy}$.154 Assuming the intrinsic galaxy image fs(067) is statistically isotropic. the shear components can be simply related: to the derivatives of the surface brightness field. as (Seljak&Zaldarriaga1999)): where the averages are taken over the galaxy. Eq.2]]," Assuming the intrinsic galaxy image $f_S(\vts)$ is statistically isotropic, the shear components can be simply related to the derivatives of the surface brightness field as \citealt{seljak99}) ): where the averages are taken over the galaxy. \ref{shear12}] ]"155 is useful only when the angular resolution of the observation is infinitely high., is useful only when the angular resolution of the observation is infinitely high.156 In. practice. the observed galaxy surface brightness distribution fo is always equal to the lensecl galaxy image f; convoluted with the PSE.τι where Wis the PSE.," In practice, the observed galaxy surface brightness distribution $f_O$ is always equal to the lensed galaxy image $f_I$ convoluted with the PSF,: where $W$ is the PSF."157 205 has shown how to modify eq.(2)) when the PSE is an isotropic Gaussian function. which can be written as: where 23 is the scale radius of the Gaussian function.," Z08 has shown how to modify \ref{shear12}) ) when the PSF is an isotropic Gaussian function, which can be written as: where $\beta$ is the scale radius of the Gaussian function."158 The new relation. between the shear components and the derivatives of the surface brightness field is: where For a general PSE. one can transform it into the desired isotropic Gaussian. form through a convolution in Fourier space.," The new relation between the shear components and the derivatives of the surface brightness field is: where For a general PSF, one can transform it into the desired isotropic Gaussian form through a convolution in Fourier space."159 The scale radius 3 of the target. PSE should. be larger than that of the original PSE to avoid singularities in the convolution., The scale radius $\beta$ of the target PSF should be larger than that of the original PSF to avoid singularities in the convolution.160 Furthermore. as shown in 205. the spatial derivatives required by eq.(5)) can also be easily. evaluated in Fourier space.," Furthermore, as shown in Z08, the spatial derivatives required by\ref{shear12PSF}) ) can also be easily evaluated in Fourier space."161" In this section. we introduce the basic ideas for treating the photon noise and the pixelation elfect in and & respectively,"," In this section, we introduce the basic ideas for treating the photon noise and the pixelation effect in \\ref{noise} and \\ref{pixelation} respectively."162 Numerical examples are given in the next section., Numerical examples are given in the next section.163Robust nieasurenients of the cosinic Inicrowave background (0ΠΡ) audsotropy. the source of much of our understanding of the universe's contents. ecometry. and primordial fluctuations. require detailed control over the systematics of the instraueut.,"Robust measurements of the cosmic microwave background (CMB) anisotropy, the source of much of our understanding of the universe's contents, geometry, and primordial fluctuations, require detailed control over the systematics of the instrument."164 The spatial response to a signal on he sky. known as the poiut-spreac-tunction (PSF) or siuplv the telescope beau. is au important svstematic effect because it smooths the anisotropy on the sky. camping hieh spatial frequencies in the angular power spectrun and washing out the encoded cosinological information.," The spatial response to a signal on the sky, known as the point-spread-function (PSF) or simply the telescope beam, is an important systematic effect because it smooths the anisotropy on the sky, damping high spatial frequencies in the angular power spectrum and washing out the encoded cosmological information."165 To recover the power spectrin. we face the challenging task of accurate beam reconstructiou.," To recover the power spectrum, we face the challenging task of accurate beam reconstruction."166 The release of the five-vear results from the (WALAP) lighliehts the issues muaportauce., The release of the five-year results from the (WMAP) highlights the issue's importance.167 substantially refines the model of the instrument beam over the previous version. which is then folded iuto the power spectrum estimate(?7).," substantially refines the model of the instrument beam over the previous version, which is then folded into the power spectrum estimate."168. The result is an increase in the five-vear power spectrum over the thaee-vear of 2 percent at the first acoustic peak aud slightly more at sinaller scales., The result is an increase in the five-year power spectrum over the three-year of 2 percent at the first acoustic peak and slightly more at smaller scales.169 These changes are easily by eve when plotting the three vear aud five vear spectra together. and outside the nominal error bars taken roni the diagonal of the covariance matrix.," These changes are easily by eye when plotting the three year and five year spectra together, and outside the nominal error bars taken from the diagonal of the covariance matrix."170 Because of ie conservative treatment of beam errors in the threc-vear release likelihood method. the cosmological estimates ortuuatelv do not change wachmostly manifesting asa 10 shift iu the preseut-day amplitude of perturbations. σε).," Because of the conservative treatment of beam errors in the three-year release likelihood method, the cosmological estimates fortunately do not change much—mostly manifesting as a $0.7 \sigma$ shift in the present-day amplitude of perturbations, $\sigma_8$."171 Tn this work. we examine the recently launched(??).. the next generation satellite to 1ieasure he CAMB anisotropy.," In this work, we examine the recently launched, the next generation satellite to measure the CMB anisotropy."172 To wring the full cosmological information from the observations. the proper calibration of the beam over a wide rauge of augular scales will prove crucial.," To wring the full cosmological information from the observations, the proper calibration of the beam over a wide range of angular scales will prove crucial."173 Because the seusitivitv of the detectors alone would allow a cosmic variance limited measurement of the temperature power spectrum to high inultipoles. determination of the plysics at high spatial frequency will depend on the removal of systematics. iu particular the quality of the beam reconstruction. especially since errors in the beam imniprint errors on the power spectrum which are strongly correlated between uultipoles.," Because the sensitivity of the detectors alone would allow a cosmic variance limited measurement of the temperature power spectrum to high multipoles, determination of the physics at high spatial frequency will depend on the removal of systematics, in particular the quality of the beam reconstruction, especially since errors in the beam imprint errors on the power spectrum which are strongly correlated between multipoles."174 The beam eror thus will affect a diverse range of science goals for he CC Binaps ane power spectra., The beam error thus will affect a diverse range of science goals for the CMB maps and power spectra.175 These include coustraiuts on the early universe. in the measurement of the primordial spectrmus slope aud running. the CAIB damping tail or any exotic plisics at recombination 2).," These include constraints on the early universe, in the measurement of the primordial spectrum's slope and running, the CMB damping tail, or any exotic physics at recombination ."176 Iu the later universe the 1ο] spectra affects constraints ou the matter distribution from CAIB lensing27777).. lensine-derived limits on neutrino luasses from CAD alone aud in combination with other large scale structure data(77).. information ou cluster physics from the Suuvaev-Zeldovich. (SZ) power spectrmun(??7).. and imodels of correlations ia poiut source populationsο," In the later universe the $l$ spectrum affects constraints on the matter distribution from CMB lensing, lensing-derived limits on neutrino masses from CMB alone and in combination with other large scale structure data, information on cluster physics from the Sunyaev-Zeldovich (SZ) power spectrum, and models of correlations in point source populations."177"ν, Finally. uncertaiutv in the jun shape adds error to cluster SZ and point-source Hux measurements."," Finally, uncertainty in the beam shape adds error to cluster SZ and point-source flux measurements."178 Iu addition. wuderstauding Plenckss jani error ds Huportaut for other experiments when orecasting cosmological performance based ou yprlor parazuneter coustraints.," In addition, understanding s beam error is important for other experiments when forecasting cosmological performance based on prior parameter constraints."179 Ilistoricallv. CMD experiments lave used a colubination of optics calculations ας planet uecasureiieuts to work out the shape of the beamothers).," Historically, CMB experiments have used a combination of optics calculations and planet measurements to work out the shape of the beam."180.. Plauets prove so uscful jcause as bright. compact sources. they rescuible ó-function signal impulses.," Planets prove so useful because as bright, compact sources, they resemble $\delta$ -function signal impulses."181 In Sec. 2.," In Sec. \ref{sec:method},"182 we discuss Planckss. plauct observations during the course of routine operations. our pipeline for simulating planet observations. and two methods for measuring the structure of the iustruinent beam: one in which we use sjenificaut prior information about the beams shape. and another where we use very little.," we discuss s planet observations during the course of routine operations, our pipeline for simulating planet observations, and two methods for measuring the structure of the instrument beam: one in which we use significant prior information about the beam's shape, and another where we use very little."183 We then interpret these beam reconstructionsimm terms of their effect on, We then interpret these beam reconstructionsin terms of their effect on184avoid absorption bands in cool stars) are not.,avoid absorption bands in cool stars) are not.185" So in a first approximation we can assume the “Isophotes” of varying metallicity in a (br) versus (77) two color diagram to be straight lines with a slope of 1. along of which we project the measured colors with (Pry)cL0 onto the mean main sequence track which in the interval 1.0.<(br)xLS is defined by This projection implies that starswith (b.rhea,<LS cannot exist in Fig. 3.."," So in a first approximation we can assume the ”isophotes” of varying metallicity in a $(b-r)$ versus $(r-i)$ two color diagram to be straight lines with a slope of $-1$, along of which we project the measured colors with $(b-r) \geq1861.0$ onto the mean main sequence track which in the interval $1.0 \leq (b-r) \leq 1.8$ is defined by This projection implies that starswith $(b-r)_{{\rm corr}}\ga1871.8$ cannot exist in Fig. \ref{wichtig},"188" which shows the spatial distribution of metallicity corrected (2)7),,44, colors (the limit is indicated by the dashed-dotted line).", which shows the spatial distribution of metallicity corrected $(b-r)_{{\rm corr}}$ colors (the limit is indicated by the dashed–dotted line).189 Both the upper and lower magnitude limits lead to selection effects which have to be taken into account., Both the upper and lower magnitude limits lead to selection effects which have to be taken into account.190 As one can see in Fig., As one can see in Fig.191 3. there is an bimodality of the observed color distribution., \ref{wichtig} there is an bimodality of the observed color distribution.192 The accumulation of red stars 1n the upper left consists mainly of disk stars. and is separated by a void from the blue stars. which predominantly belong to the halo.," The accumulation of red stars in the upper left consists mainly of disk stars, and is separated by a void from the blue stars, which predominantly belong to the halo."193 Thus à erude disk-halo separation can be drawn by a color cut — we take stars with (0rou<0.7 to be halo. stars with (Den:=0.7 to be disk stars.," Thus a crude disk-halo separation can be drawn by a color cut – we take stars with $(b-r)_{{\rm corr}}<0.7$ to be halo, stars with $(b-r)_{{\rm corr}}>0.7$ to be disk stars."194 In the following we will make use of this color cut to derive the distribution of the disk stars separately., In the following we will make use of this color cut to derive the distribution of the disk stars separately.195 In a second step. the distribution as a whole will be analysed.," In a second step, the distribution as a whole will be analysed."196 In Fig., In Fig.197 3. the cut is denoted by a dotted horizontal In the two fields we have analysed so far we find 95 halo and 178 disk stars. that is a factor of two more disk stars than predicted by the standard model (Baheall&Soneira 1980).," \ref{wichtig} the cut is denoted by a dotted horizontal In the two fields we have analysed so far we find 95 halo and 178 disk stars, that is a factor of two more disk stars than predicted by the standard model \cite{BSStandard}. ."198. This surplus of disk stars was already noted by Reid Majewski (1993)., This surplus of disk stars was already noted by Reid Majewski (1993).199 Fig., Fig.200 4. shows the distribution of the ο In the two fields under consideration., \ref{codistr} shows the distribution of the $(b-r)_{{\rm corr}}$ in the two fields under consideration.201 As expected. the detection limit of absolute magnitudes (colors) is distance dependent (see Fig. 3)).4," As expected, the detection limit of absolute magnitudes (colors) is distance dependent (see Fig. \ref{wichtig})"202 We use the two-dimensional distribution of stars in. the (bMeor VS logr diagram (Fig. 3)), We use the two-dimensional distribution of stars in the $(b-r)_{{\rm corr}}$ vs $\log r$ diagram (Fig. \ref{wichtig}) )203 to correct for this incompleteness in the following way: First. we divide the distance in logarithmic bins of 0.2 às indicated in Fig.," to correct for this incompleteness in the following way: First, we divide the distance in logarithmic bins of 0.2 as indicated in Fig."204" 3. and count the stars up to the the upper color limit (b.—3!"" (this is the distance dependent color (luminosity) limit. up to which stars can be detected."," \ref{wichtig} and count the stars up to the the upper color limit $(b-r)^{lim}$ (this is the distance dependent color (luminosity) limit, up to which stars can be detected."205 The metal-poor halo stars are intrinsically fainter (see paragraph 3)). thus the limits are shifted The nearest bins (Sxlogr< 0.2) are assumed to be complete.," The metal-poor halo stars are intrinsically fainter (see paragraph 3)), thus the limits are shifted The nearest bins $-0.8\leq \log r \leq -0.2$ ) are assumed to be complete."206" For the incomplete bins we multiply iteratively with a factor given by the ratio of complete to incomplete number counts in the previous bin. where the limit for the uncorrected counts Is defined by the bin currently under examination (/): where .V; is the number of stars in bin {. Nis the number of stars in the previous bin (j.—1). up to the limit given by the bin j. and AN"" is the number of stars from that limit up to the limit given by bin j—1. see also With the poissonian errors ay=vN. o=VN. and cs=VN"" the error of the corrected number counts becomes: For a detailed deduction see The completeness correction is done for each field With the corrected number counts the density in. the logarithmic spaced volume bins (V;=πω) can than be calculated according to For every logarithmic distance bin we use the mean height : above the Galactic plane <:;>=sm»«r>. where Xlogr-—1logre;PF(ogrjia.loge2=logr;|01><p—1.259w(;. We first study the density distribution of the disk stars by taking only stars in the corresponding color interval into account."," For the incomplete bins we multiply iteratively with a factor given by the ratio of complete to incomplete number counts in the previous bin, where the limit for the uncorrected counts is defined by the bin currently under examination $j$ ): where $N_j$ is the number of stars in bin $j$, $N^{'}$ is the number of stars in the previous bin $j-1$ ), up to the limit given by the bin $j$, and $N^{''}$ is the number of stars from that limit up to the limit given by bin $j-1$, see also With the poissonian errors $\sigma_{N}=\sqrt{N}$, $\sigma_{N^{'}}=\sqrt{N^{'}}$, and $\sigma_{N^{''}}=\sqrt{N^{''}}$ the error of the corrected number counts becomes: For a detailed deduction see The completeness correction is done for each field With the corrected number counts the density in the logarithmic spaced volume bins $V_j=\frac{1}{3}\omega (r_{j+1}^3 - r_j^3)$ ) can than be calculated according to For every logarithmic distance bin we use the mean height $z$ above the Galactic plane $<z_j>=\sin b \cdot <r>$, where $<\log r> = \log207r_j+(\log r_{j+1}-\log r_i)/2=\log r_i+0.1 \Rightarrow <r>=1.259 r_i$ We first study the density distribution of the disk stars by taking only stars in the corresponding color interval into account."208 Although the color-cut at (PMeo=0.7 Is arather crude separation between disk and halo. we gain a clearer insight into the disk distribution since the contamination by halo stars is suppressed As the nearest stars in our fields have still distances of about ppe the normalization at τ=0 has to be established by other means: We take stars from the CNS4 (Jahreif&Wielen 1997b).. which are locatedin a sphere with radius ppc around the sun.," Although the color-cut at $(b-r)_{{\rm corr}}=0.7$ is a rather crude separation between disk and halo, we gain a clearer insight into the disk distribution since the contamination by halo stars is suppressed As the nearest stars in our fields have still distances of about pc the normalization at $z=0$ has to be established by other means: We take stars from the CNS4 \cite{CNS4}, , which are locatedin a sphere with radius pc around the sun."209scaling on the numerically resolved flow.,scaling on the numerically resolved flow.210 A potential problem for the SGS model is the possible lack of isotropy near the flame surface., A potential problem for the SGS model is the possible lack of isotropy near the flame surface.211" However, there is certainly no pronounced anisotropy at the smallest resolved scales if Ao(t)<lx/rt> and statistical isotropy is found for the bulk of turbulent regions."," However, there is certainly no pronounced anisotropy at the smallest resolved scales if $\Delta_{0}(t)<\ell_{\mathrm{K/RT}}$, and statistical isotropy is found for the bulk of turbulent regions."212" Only in the late phase of the explosion, when the transition length becomes smaller than the numerical resolution and the ἐκ/κτresolved small-scale turbulence definitely becomes anisotropic, the notion of SGS turbulence energy cannot be strictly justified."," Only in the late phase of the explosion, when the transition length $\ell_{\mathrm{K/RT}}$ becomes smaller than the numerical resolution and the resolved small-scale turbulence definitely becomes anisotropic, the notion of SGS turbulence energy cannot be strictly justified."213" One should note, however, that this point more or less coincides with the time when a deflagration-to-detonation transition is expected to occur (Gamezoetal.2005;Rópke&Niemeyer2007)."," One should note, however, that this point more or less coincides with the time when a deflagration-to-detonation transition is expected to occur \citep{GamKhok05,RoepNie07}."214". Apart from that, the production of turbulence energy by unresolved buoyancy effects is heuristically included in the SGS model."," Apart from that, the production of turbulence energy by unresolved buoyancy effects is heuristically included in the SGS model."215" In conclusion, the SGS turbulence energy model by Schmidtetal.(2006) to the major part of the explosive burning in the deflagration appliesphase of a Type Ia supernova explosion, but there is no regime for which a pure RT-scaling model holds."," In conclusion, the SGS turbulence energy model by \cite{SchmNie06} applies to the major part of the explosive burning in the deflagration phase of a Type Ia supernova explosion, but there is no regime for which a pure RT-scaling model holds."216" After settling the issue of turbulence scaling in the deflagration phase of a Ia supernova in the present article, we mention that the Typeoccurrence of deflagration-to-detonation transitions can be constrained on the basis of the deflagration model."," After settling the issue of turbulence scaling in the deflagration phase of a Type Ia supernova in the present article, we mention that the occurrence of deflagration-to-detonation transitions can be constrained on the basis of the deflagration model."217" For the DDT mechanism to operate, strong turbulence is necessary in late phases of the burning (e.g.Woosley2007).."," For the DDT mechanism to operate, strong turbulence is necessary in late phases of the burning \citep[e.g.][]{woosley2007a}."218 Rópke(2007) found that this may indeed be realized in deflagration models of SNe Ia with low (but not vanishing) probability., \citet{roepke2007d} found that this may indeed be realized in deflagration models of SNe Ia with low (but not vanishing) probability.219" In order to better quantify this intermittency effect, higher-order structure functions have to be computed."," In order to better quantify this intermittency effect, higher-order structure functions have to be computed."220" Fitting intermittency models to the numerically determined scaling exponents (asproposedbyPanetal. 2008),, the of strong turbulent velocity fluctuations at any instant of probabilitytime can be estimated."," Fitting intermittency models to the numerically determined scaling exponents \citep[as proposed221by][]{Pan08}, the probability of strong turbulent velocity fluctuations at any instant of time can be estimated."222 This analysis will be presented in a future publication., This analysis will be presented in a future publication.223In this paper we have studied the properties of quasi-continuum thermal modes in à prominence thread model which is transversely inhomogencous.,In this paper we have studied the properties of quasi-continuum thermal modes in a prominence thread model which is transversely inhomogeneous.224" We have followed the method of VanderLin- and VanderLinden&Goossens(1991) to investigate, first, the stability of the thermal continuum and, later, the effect of cross-field thermal conduction and magnetic diffusion."," We have followed the method of \citet{vanderlindencont} and \citet{vanderlinden91} to investigate, first, the stability of the thermal continuum and, later, the effect of cross-field thermal conduction and magnetic diffusion."225 We have recovered the general results by VanderLindenetal.(1991) and Goossens(1991) in all cases studied in the present paper., We have recovered the general results by \citet{vanderlindencont} and \citet{vanderlinden91} in all cases studied in the present paper.226" In particular, we have found that the thermal continuum in prominence threads is unstable for PCTR temperatures."," In particular, we have found that the thermal continuum in prominence threads is unstable for PCTR temperatures."227 In agreement with the results of VanderLinden&Goossens(1991).. the effect of cross-field thermal conduction is to replace the thermal continuum by discrete modes which retain the unstable character of the continuum.," In agreement with the results of \citet{vanderlinden91}, the effect of cross-field thermal conduction is to replace the thermal continuum by discrete modes which retain the unstable character of the continuum."228" On the contrary, the role of magnetic diffusion is negligible in our model 1998)."," On the contrary, the role of magnetic diffusion is negligible in our model \citep[see][]{ireland92,ireland98}."229". We have obtained the important result that the instability erowth rate of the most unstable mode is independent of the form of the temperature profile within the PCTR of the thread, and the instability time scale is consistent with the observed lifetime of the threads in Ha observations of solar filaments (c.g.Linetal.2008, 2009).."," We have obtained the important result that the instability growth rate of the most unstable mode is independent of the form of the temperature profile within the PCTR of the thread, and the instability time scale is consistent with the observed lifetime of the threads in $\alpha$ observations of solar filaments \citep[e.g.,][]{lin08,lin09}. ."230" Considering our present results along with those obtained by VanderLinden&Goossens(1991) and VanderLinden(1993)., we conclude that unstable thermal modes may play a relevant role for both the formation of the prominence fine structure and the subsequent instability of the thin prominence threads."," Considering our present results along with those obtained by \citet{vanderlinden91} and \citet{vanderlinden93}, we conclude that unstable thermal modes may play a relevant role for both the formation of the prominence fine structure and the subsequent instability of the thin prominence threads."231 Here we have considered a fully ionized prominence., Here we have considered a fully ionized prominence.232" However, due to the low temperature in the dense core of the threads, the cool prominence plasma is expected to be partially ionized (Gouttebroze&Labrosse2009)."," However, due to the low temperature in the dense core of the threads, the cool prominence plasma is expected to be partially ionized \citep{labrosse}."233". If partial ionization is taken into account, several additional effects have to be considered. namely thermal conduction by neutrals, ambipolar (or Cowling’s) diffusion caused by ion-neutral collisions, and a modified radiation function (see,al.2008;Soleret 2010a)."," If partial ionization is taken into account, several additional effects have to be considered, namely thermal conduction by neutrals, ambipolar (or Cowling's) diffusion caused by ion-neutral collisions, and a modified radiation function \citep[see, e.g.,][]{forteza08,solerhelium}."234. The role of both effects depends strongly on the plasma ionization degree., The role of both effects depends strongly on the plasma ionization degree.235" It requires complicated computations of NLTE radiative transfer and statistical equilibrium of atomic level populations to obtain the precise profile of the ionization degree in prominence threads (seeextensivedetailsin,e.g.,Gouttebroze&Labrosse2009;etal."," It requires complicated computations of NLTE radiative transfer and statistical equilibrium of atomic level populations to obtain the precise profile of the ionization degree in prominence threads \citep[see extensive details in, e.g.,][]{labrosse,labrossereview}."236"2010).. Nevertheless, some relevant conclusions can be obtained if one reasonably assumes that the ionization degree in prominence plasmas is mainly determined by the temperature and. to a lesser extend, by the density."," Nevertheless, some relevant conclusions can be obtained if one reasonably assumes that the ionization degree in prominence plasmas is mainly determined by the temperature and, to a lesser extend, by the density."237" While the ionization degree of the core of the thread is unknown, it is realistic to assume that the condition of full ionization occurs in the PCTR for temperatures higher than à critical temperature, namely 7*."," While the ionization degree of the core of the thread is unknown, it is realistic to assume that the condition of full ionization occurs in the PCTR for temperatures higher than a critical temperature, namely $T^*$."238" The results by Gouttebroze&Labrosse(2009) and Schureetal.(2010) suggest realistic values for the critical ionizing temperature for the hydrogen plasma in prominences of T'z2x10'Kand T«2.5x10! K, respectively, so only the inner part of the transitional layer would be partially ionized."," The results by \citet{labrosse} and \citet{schure} suggest realistic values for the critical ionizing temperature for the hydrogen plasma in prominences of $T^* \approx 2 \times 10^4$ K and $T^* \approx 2.5 \times 10^4$ K, respectively, so only the inner part of the transitional layer would be partially ionized."239" According the instability criterion (Equation (30))), the unstable part of the thermal continuumtakes place at temperatures higher than 8x107 K, which is higher than"," According the instability criterion (Equation \ref{eq:crit1}) )), the unstable part of the thermal continuumtakes place at temperatures higher than $8 \times 10^4$ K, which is higher than"240largee negativee values of «ry are seen only in early-type egalaxies.,large negative values of $x_1$ are seen only in early-type galaxies.241 This causes some clepenclence of ary on colour of host galaxies as seen iu tle ligure., This causes some dependence of $x_1$ on colour of host galaxies as seen in the figure.242" The slope. however. is sinall enough to cause any effect ou brightness of SNe Ia through he correction of Q.12.r,."," The slope, however, is small enough to cause any effect on brightness of SNe Ia through the correction of $0.12x_1$ ."243 Sinilarly Figure 16 shows 1uaxiijun brightuess of SNe Ia ou luminosity of host. galaxies. showing that the former does not depeud on luminosity of the liost.," Similarly Figure \ref{fig:snmhostm} shows maximum brightness of SNe Ia on luminosity of host galaxies, showing that the former does not depend on luminosity of the host."244 Note that we expect tliat average metallicity of the Lost galaxy €langes by 0.59 dex in this luminosity range 2001)., Note that we expect that average metallicity of the host galaxy changes by 0.59 dex in this luminosity range \citep{Tremonti04}.245. The ory parameter changes sligltly as luminosity of host. galaxies chauges., The $x_1$ parameter changes slightly as luminosity of host galaxies changes.246 We also exaimine the dependence of the e parameter and the ory parameter ou the distance [rom the centre of gaaxies as shown in Figure 17.., We also examine the dependence of the $c$ parameter and the $x_1$ parameter on the distance from the centre of galaxies as shown in Figure \ref{fig:dpos_color}.247 The colour excess e=E(B—V) stays nearly at 0.08 with the cisyersion of 0.07 auc does not show a systematic «hange f‘olil spe to 15 kpe., The colour excess $c=E(B-V)$ stays nearly at $\approx$ 0.08 with the dispersion of 0.07 and does not show a systematic change from 1 kpc to 15 kpc.248 In particular. we observe tha the colour excess does not decrease as we go bler «way from the centre of galaxies. wlere star:. txl heice dust. are expected to decrease.," In particular, we observe that the colour excess does not decrease as we go farther away from the centre of galaxies, where stars, and hence dust, are expected to decrease."249 The treud we see here does uot change if we limit t je5 Neto: uw0.2 with which SNe Ia that wou 'ecelve reasonably large extinetion are iucluded iu tl© sallyje., The trend we see here does not change if we limit the SNe to $z<0.2$ with which SNe Ia that would receive reasonably large extinction are included in the sample.250 This fiuding suMODoOests us to Interpr tha the observed reddeuiug is mostly associated with inelvidual superuovae. eiher extinction 1 Supernova itself and/or eireum-supernova dus Or jutrinslc colour variation of tle supernova. rather thau interstellar dust in host galaxies.," This finding suggests us to interpret that the observed reddening is mostly associated with individual supernovae, either extinction from supernova itself and/or circum-supernova dust or intrinsic colour variation of the supernova, rather than interstellar dust in host galaxies."251 This woulc also eive us an upper limi ou the moclel as to the amouu of dust. produced. around SNe Ia. We observe ouly 7 SNe Ia wlich shows colour excess more than E(B—V)>0.3 mag amoug our 137 SN Ia sample., This would also give us an upper limit on the model as to the amount of dust produced around SNe Ia. We observe only 7 SNe Ia which shows colour excess more than $E(B-V)\geq 0.3$ mag among our 137 SN Ia sample.252 In addilOn. we do not observe a systematic depeuclence of the ry parameter ou the distance from the ceire of galaxies: see Figure 17((b).," In addition, we do not observe a systematic dependence of the $x_1$ parameter on the distance from the centre of galaxies; see Figure \ref{fig:dpos_color}( (b)."253 The SNe Ia sample acquired in the SDSS IL containing 137 low redshift SNe from z0.05 to 0.3. indicates that the occurrence of SNe Ia is prinariy proportional to luminosity of galaxies.," The SNe Ia sample acquired in the SDSS II, containing 137 low redshift SNe from $z=0.05$ to 0.3, indicates that the occurrence of SNe Ia is primarily proportional to luminosity of galaxies."254 The LF of SN Ia lost galaxies matches very well wit bthat of field galaxies multiplied by luminosity. aud colour of SN Ia host galaxies does 1ot ciffer from ha of field galaxies.," The LF of SN Ia host galaxies matches very well with that of field galaxies multiplied by luminosity, and colour of SN Ia host galaxies does not differ from that of field galaxies."255 Our low redshift sample does not iudicate au active signature tlal the occurrence of SNe Ia folows star formation activity. except that possible enhancement iu tlje SN [Ia ‘ate ]is noted in late ype galaxies C314 355€)) compared wih the rate in elliptical gaaxles.," Our low redshift sample does not indicate an active signature that the occurrence of SNe Ia follows star formation activity, except that possible enhancement in the SN Ia rate is noted in late type galaxies $\lesssim 31\pm35$ ) compared with the rate in elliptical galaxies."256 Otr low redshift sample is compatible with he two component mocel wiere the effect of the prom(Xt component is mocles (10—30'€ of the total rate). iu tlie cur‘ent models.," Our low redshift sample is compatible with the two component model where the effect of the prompt component is modest $-$ of the total rate), as in the current models."257 We are not adle to differentiate SN rates ar1016 late type galaxies., We are not able to differentiate SN rates among late type galaxies.258 Our iple coutalus ὃ 5e Ta. whose host galaxies were not identified.," Our sample contains 8 SNe Ia, whose host galaxies were not identified."259 It is shown. however. that they COLlslstel with hem occurred iu low lunilous galaxies beyond the survey limit for galaxies.," It is shown, however, that they are consistent with them occurred in low luminous galaxies beyond the survey limit for galaxies."260 ulnosity o (SNe Ia does not appear to ceped upon luminosity or colour of host. galaxies., Luminosity of SNe Ia does not appear to depend upon luminosity or colour of host galaxies.261 'The luiilnosity unuction of SNe La is Gausslan with Aj=—19.12 aud &/=0.2 lmag (FWHM is a factor lL. Li1luminosity). ifthe colour variation is interpreted as recdclening obeying the extinction law of the Milky Way with the standard. να]e liy 3.1.," The luminosity function of SNe Ia is Gaussian with $M_B=-19.42$ and $\sigma=0.24$ mag (FWHM is a factor 1.4 inluminosity), if the colour variation is interpreted as reddening obeying the extinction law of the Milky Way with the standard value $R_V=3.1$ ."262 ThisGaussian distribution is further, ThisGaussian distribution is further263"SINCE Poy=p,.",since $\rho_{\rm ev}=\rho_{\rm c}$.264 However. energy concentrated in the resonant layer of the thread because of resonant wave damping can flow along the field lines and. eventually. supply heating in the evacuated region. where field aligned currents are dominant.," However, energy concentrated in the resonant layer of the thread because of resonant wave damping can flow along the field lines and, eventually, supply heating in the evacuated region, where field aligned currents are dominant."265" Although the energy inflow into the resonance. given by <S, >. and its subsequent divergence along the field lines. given by <S->. are mostly determined by ideal terms in Equations (15)) and (16)). this does not mean that resistivity is not important."," Although the energy inflow into the resonance, given by $<S_r>$ , and its subsequent divergence along the field lines, given by $<S_z>$, are mostly determined by ideal terms in Equations \ref{sr}) ) and \ref{sz}) ), this does not mean that resistivity is not important."266 For instance. the amount of heating. in the form of Ohmic dissipation. will be determined by those currents. resistivity. and their spatial distribution.," For instance, the amount of heating, in the form of Ohmic dissipation, will be determined by those currents, resistivity, and their spatial distribution."267 For this particular case. heating is distributed in a constant manner in the evacuated part of the tube. even if there is no resonant layer in that region (see Fig.," For this particular case, heating is distributed in a constant manner in the evacuated part of the tube, even if there is no resonant layer in that region (see Fig."268 9bb)., \ref{poyntingflux}b b).269 Quiescent filament fine structures are only partially filled with cold and dense absorbing material., Quiescent filament fine structures are only partially filled with cold and dense absorbing material.270 The length of the threads can in principle be measured in events showing transverse oscillations. provided the lifetime of threads is sufficiently large compared to the oscillatory period.," The length of the threads can in principle be measured in events showing transverse oscillations, provided the lifetime of threads is sufficiently large compared to the oscillatory period."271 The length of the supporting magnetic flux tubes are however much larger. and cannot be observed.," The length of the supporting magnetic flux tubes are however much larger, and cannot be observed."272 Density measurements. both in the thread as in the evacuated part of the supporting magnetic tube. are also challenging from the observational point of view.," Density measurements, both in the thread as in the evacuated part of the supporting magnetic tube, are also challenging from the observational point of view."273 It is therefore important to quantify the variations on wave properties due to changes in these equilibrium parameters if we aim to perform an accurate prominence seismology., It is therefore important to quantify the variations on wave properties due to changes in these equilibrium parameters if we aim to perform an accurate prominence seismology.274 It is essential to have computations of periods and damping times for a wide range of thread models to include regimes in which the applicability of simple analytical models could be of limited extent., It is essential to have computations of periods and damping times for a wide range of thread models to include regimes in which the applicability of simple analytical models could be of limited extent.275 For this reason we have computed the oscillatory properties of resonantly damped transverse kink oscillations in rather general two-dimensional fully non-uniform prominence thread models., For this reason we have computed the oscillatory properties of resonantly damped transverse kink oscillations in rather general two-dimensional fully non-uniform prominence thread models.276 This allows for a broad range of prominence threads with very different physical conditions to be modelled and their oscillatory properties characterised., This allows for a broad range of prominence threads with very different physical conditions to be modelled and their oscillatory properties characterised.277 The length of the thread ad the density in the evacuated part of the tube define their longitudinal density structuring., The length of the thread and the density in the evacuated part of the tube define their longitudinal density structuring.278 We find that the length of the thread strongly influences the period and damping time of trasverse kink oscillations. while the damping ratio 1s rather insesitive to this parameter.," We find that the length of the thread strongly influences the period and damping time of transverse kink oscillations, while the damping ratio is rather insensitive to this parameter."279 These results confirm the validity of the analytical approximations made by Soleretal.(2010)., These results confirm the validity of the analytical approximations made by \citet{soler102dthread}.280. In. addition. our modelling has allowed us to identify a new physical parameter with seismological implications. the density in the evacuated part of the thread.," In addition, our modelling has allowed us to identify a new physical parameter with seismological implications, the density in the evacuated part of the thread."281 This quantity also ifluences periods and damping times. and to a lesser extent daniping ratios. and must be taken into account in. the inversion of physical parameters in the context of prominence seismology.," This quantity also influences periods and damping times, and to a lesser extent damping ratios, and must be taken into account in the inversion of physical parameters in the context of prominence seismology."282 Currently available inversion schemes for one-dimensional coronal loops and prominence threads make use of observed periods and damping ratios., Currently available inversion schemes for one-dimensional coronal loops and prominence threads make use of observed periods and damping ratios.283 The first. influence the inferred values for the Alfvénn speed. while the second determine the transverse density structuring.," The first, influence the inferred values for the Alfvénn speed, while the second determine the transverse density structuring."284 Based on our results. we can conclude that ignorance on the length of the thread. the length of the supporting magnetic flux tube. and the density in the evacuated part of the tube will have a significant impact on the inferred values for the Alfvénn speed (hence magnetic field strength) in the thread. depending on whether we use those one-dimensional inversion schemes or the results from two-dimensional models here obtained.," Based on our results, we can conclude that ignorance on the length of the thread, the length of the supporting magnetic flux tube, and the density in the evacuated part of the tube will have a significant impact on the inferred values for the Alfvénn speed (hence magnetic field strength) in the thread, depending on whether we use those one-dimensional inversion schemes or the results from two-dimensional models here obtained."285 On the contrary. because of the smaller sensitivity of the damping ratio to changes in the longitudinal density structuring. seismological estimates of the transverse density structuring will be less affected by our ignorance about the longitudinal density structuring of prominence threads.," On the contrary, because of the smaller sensitivity of the damping ratio to changes in the longitudinal density structuring, seismological estimates of the transverse density structuring will be less affected by our ignorance about the longitudinal density structuring of prominence threads."286 Our study provides additional insight to the physics of resonantly damped kink modes in two-dimensional equilibrium states. by extending previous applications (e.g.Andriesetal..2005:Arreguietal..2005) to more complex non-separable density distributions.," Our study provides additional insight to the physics of resonantly damped kink modes in two-dimensional equilibrium states, by extending previous applications \citep[e.g,][]{Andries05,Arregui05} to more complex non-separable density distributions."287 It also. provides a example of the methods and uses of combining the information from the spatial distribution of eigenfunctions with that obtained from energy arguments., It also provides an example of the methods and uses of combining the information from the spatial distribution of eigenfunctions with that obtained from energy arguments.288 In particular. our energy analysis has allowed us to explain the decrease in damping times for shorter thread lengths found by Soleretal.(2010).," In particular, our energy analysis has allowed us to explain the decrease in damping times for shorter thread lengths found by \cite{soler102dthread}."289. The length of the thread influences the energy of the kink mode. and hence its oscillatory period. but also affects the damping by resonant absorption. through the energy flux into the resonance.," The length of the thread influences the energy of the kink mode, and hence its oscillatory period, but also affects the damping by resonant absorption, through the energy flux into the resonance."290 In an analogous way. the value of the density in the evacuated part of the tube also determines periods and damping times. since both the energy of the kink mode and the energy flux into the resonance vary.," In an analogous way, the value of the density in the evacuated part of the tube also determines periods and damping times, since both the energy of the kink mode and the energy flux into the resonance vary."291 This means that changes m the equilibrium configuration in a non-resonant direction produce variations in the damping properties of kink modes. a result that was qualitatively explained by a detailed examination of the radial and longitudinal profiles of the eigenfunctions.," This means that changes in the equilibrium configuration in a non-resonant direction produce variations in the damping properties of kink modes, a result that was qualitatively explained by a detailed examination of the radial and longitudinal profiles of the eigenfunctions."292 Both the shortening of the length of the thread and the decrease of the density in the evacuated part of the tube produce more marked resonances. with the amplitude of the velocity perturbations at the resonance and the compressibility of the mode in the thread being larger.," Both the shortening of the length of the thread and the decrease of the density in the evacuated part of the tube produce more marked resonances, with the amplitude of the velocity perturbations at the resonance and the compressibility of the mode in the thread being larger."293 Inside the resonant layers shorter transverse spatial scales for the Alfvénnic velocity component are obtained., Inside the resonant layers shorter transverse spatial scales for the Alfvénnic velocity component are obtained.294 In combination with the analysis of the energy of the kink modes and the energy flux into the resonances a quantitative explanation was obtained for both the damping properties obtained in our study and those in Soleretal. (2010)., In combination with the analysis of the energy of the kink modes and the energy flux into the resonances a quantitative explanation was obtained for both the damping properties obtained in our study and those in \citet{soler102dthread}.295 The damping of kink oscillations in two-dimensional fully non-uniform equilibrium configurations can be computed by using energy arguments together with the solution of simpler problems for kink mode and Alfvénn continuum modes., The damping of kink oscillations in two-dimensional fully non-uniform equilibrium configurations can be computed by using energy arguments together with the solution of simpler problems for kink mode and Alfvénn continuum modes.296 This aspect is worth to be considered in future studies of resonant absorption in 2D/3D models of solar atmospheric magnetic structures involving changes of equilibrium parameters that affect the density structuring in a non-resonant direction., This aspect is worth to be considered in future studies of resonant absorption in 2D/3D models of solar atmospheric magnetic structures involving changes of equilibrium parameters that affect the density structuring in a non-resonant direction.297 The use of energy arguments would allow to have a first indication, The use of energy arguments would allow to have a first indication298"|—d will likewise be smaller. rendering it cillicult to ""bury? the problematic &zT scales.","$t=t_i$ will likewise be smaller, rendering it difficult to `bury' the problematic $k \geq T$ scales."299 The issue raised poses a real dilliculty for the theory of inflation., The issue raised poses a real difficulty for the theory of inflation.300 The scenario we have discussed. is only one possibility., The scenario we have discussed is only one possibility.301 Le could. also be that the pre-inllationary. phase is dominated by extremely massive cold. particles. such as monopoles arising [from an carly phase transition. but in such à case it seems likely that there will again be large luctuations on very small scales. unless a supoerlluid. is »ostulated. for the phase (as noted. alter (45)). most Πλας would exhibit a P(A) divergence. towards fo>>1. not just thermal ones).," It could also be that the pre-inflationary phase is dominated by extremely massive cold particles, such as monopoles arising from an early phase transition, but in such a case it seems likely that there will again be large fluctuations on very small scales, unless a superfluid is postulated for the phase (as noted after \ref{DeLargek}) ), most fluids would exhibit a $\cP (k)$ divergence towards $k \gg T$, not just thermal ones)."302 Even if intlation is ‘warm’ it is not obvious how the initial non-perturbative situation of the oe-inllationarv phase may. be averted., Even if inflation is `warm' it is not obvious how the initial non-perturbative situation of the pre-inflationary phase may be averted.303" However. if such a hase is altogether absent. or indeed. consists of a highly. 1i0mogeneous superfluid to start with. then of course that would be the ""ultimate fix’."," However, if such a phase is altogether absent, or indeed consists of a highly homogeneous superfluid to start with, then of course that would be the `ultimate fix'."304 But in this case. inflation cannot ‘laim to haveerpéazned why we fail to find evidence (directly or indirectly) for the relic particles and. the curvature of Space.," But in this case, inflation cannot claim to have why we fail to find evidence (directly or indirectly) for the relic particles and the curvature of space."305 We wish to acknowledge helpful correspondence. with Alassimo Giovannini. Michael Joyce. Joao Alagueijo. lan Aloss and David. Wancds. as well as critical comments [rom Andrei Linde ancl David Lyth.," We wish to acknowledge helpful correspondence with Massimo Giovannini, Michael Joyce, Joao Magueijo, Ian Moss and David Wands, as well as critical comments from Andrei Linde and David Lyth."306in Fig. 3((,in Fig. \ref{fpt}( (307b)( (for µε=10 Ly and (cl) (for py= 10.7).,b) (for $\mu_r=10^{-1}$ ) and (d) (for $\mu_r=10^{-3}$ ).308 Phe volume of the torus is 0.16 (0.04) of the original µε=1 case for pymιο μι 7)., The volume of the torus is 0.16 (0.04) of the original $\mu_r=1$ case for $\mu_r=10^{-1}$ $\mu_r=10^{-3}$ ).309 The magnetic intensity H=D/j in Eq. (, The magnetic intensity $\bf{H}={\bf{B}}/\mu$ in Eq. (3103) is magnified by pp in a superconductor.,3) is magnified by $\mu/\mu_0$ in a superconductor.311 Lhe deformation caused by this field is then calculated by the method given by Mastranoetal.(2011)., The deformation caused by this field is then calculated by the method given by \citet{metal11}.312.. We are unable to derive asimple analytic formula describing the dependence of€ on py=pfpty and A simultaneously. but we find that the general form of eCX) for a given fh is still similar to Eq. (," We are unable to derive asimple analytic formula describing the dependence of $\epsilon$ on $\mu_r=\mu/\mu_0$ and $\Lambda$ simultaneously, but we find that the general form of $\epsilon(\Lambda)$ for a given $\mu_r$ is still similar to Eq. ("313"5). namely The climensionless constants ej,» are quoted in Table 1 for αν=0.5.0.1. 10 107. and 10.7.","5), namely The dimensionless constants $c_{1,2}$ are quoted in Table 1 for $\mu_r=0.5$,$0.1$, $10^{-2}$, $10^{-3}$ , and $10^{-4}$."314 For all jr. the star is oblate for A0.4 and prolate for As0.4.," For all $\mu_r$, the star is oblate for $\Lambda\gtrsim 0.4$ and prolate for $\Lambda\lesssim 0.4$."315" The functional dependence of « on A stavs roughly the same as yt, changes.", The functional dependence of $\epsilon$ on $\Lambda$ stays roughly the same as $\mu_r$ changes.316 The smaller magnetic permeability of the superconducting stellar matter enhances the density perturbation by a factor of ~ji1 as evident from the force balance equation. Eig. (," The smaller magnetic permeability of the superconducting stellar matter enhances the density perturbation by a factor of $\sim \mu_r^{-1}$, as evident from the force balance equation [Eq. ("3173)]. and this is embodied in he approximate: scaling. ejxοlg forsmall yr.,"3)], and this is embodied in the approximate scaling $c_1\propto \mu_r^{-1}$ forsmall $\mu_r$."318 We can now ask how much closer the theoretical curves approach the claa when superconcuctivity is included., We can now ask how much closer the theoretical curves approach the data when superconductivity is included.319 Looking at Fig., Looking at Fig.320 2. for example. the curve for (sav) µε=10. and A=10* (not drawn) is hisher than the lower black dashed curve by three orders of magnitude. which is still below observational limits.," \ref{eps} for example, the curve for (say) $\mu_r=10^{-3}$ and $\Lambda=10^{-3}$ (not drawn) is higher than the lower black dashed curve by three orders of magnitude, which is still below observational limits."321 We conclude that superconducting interiors are easily compatible with current observational upper limits. if the external magneσεο Field is not reduced by accretion in any way.," We conclude that superconducting interiors are easily compatible with current observational upper limits, if the external magnetic field is not reduced by accretion in any way."322" We caution that. in this first pass. we assume yr,=1 throughout the star. instead of only in some region."," We caution that, in this first pass, we assume $\mu_r\neq 1$ throughout the star, instead of only in some region."323 Εις assumption is physically implausible and is only taken to simplifv our (illustrative) caculation., This assumption is physically implausible and is only taken to simplify our (illustrative) calculation.324 :X more thorough caleulation where 65. 15 allowed to vary inside the star is needed before definite conclusions can be drawn regarding the elfects of core superconductivity on cllipticity., A more thorough calculation where $\mu_r$ is allowed to vary inside the star is needed before definite conclusions can be drawn regarding the effects of core superconductivity on ellipticity.325 However. we conjecture that the calculation presented in this section sets the upper limit on the effects of a superconducting interior on stellar deformation: if superconductivity is limited to a smaller region in the star. the changes to € (relative to pe= po) will be less than predicted by Eq. (," However, we conjecture that the calculation presented in this section sets the upper limit on the effects of a superconducting interior on stellar deformation: if superconductivity is limited to a smaller region in the star, the changes to $\epsilon$ (relative to $\mu=\mu_0$ ) will be less than predicted by Eq. ("3267) ancl Table 1.,7) and Table 1.327 Reevcled pulsars are selected as the subjects of this study for two reasons: (1) their spin-down rates are lower than those of other objects. vielding more stringent indirect eravitational-wave limits: and (it) their internal magnetic fields may be much &reater than their surface fields due to diamagnetic screening or burial. vielding larger hycromagnetic deformations than one might otherwise expect.," Recycled pulsars are selected as the subjects of this study for two reasons: (i) their spin-down rates are lower than those of other objects, yielding more stringent indirect gravitational-wave limits; and (ii) their internal magnetic fields may be much greater than their surface fields due to diamagnetic screening or burial, yielding larger hydromagnetic deformations than one might otherwise expect."328 We now examine point (ii)., We now examine point (ii).329 Let us ask what happens if the actual magnetic field strength just. below the surface takes its pre-accretion value (e.g. before diamagnetie screening or burial) Dou=Bootecrved/&. where£ is some dimensionless ‘shielding factor’.," Let us ask what happens if the actual magnetic field strength just below the surface takes its pre-accretion value (e.g., before diamagnetic screening or burial) $B_\mathrm{s, actual}=B_\mathrm{s, observed}/\xi$, where $\xi$ is some dimensionless `shielding factor'."330 We note in passing that this scenario is more realistic than those considered in Sec., We note in passing that this scenario is more realistic than those considered in Sec.331 3.1: it is unlikely that the poloidal feld is reduced to ~10 P in a recycled. pulsa [ike t1e surface field. given the high electrical conductivity expected in the core (Goldreich&Reisenegeer 1992).. except in the special situation where the source currents reside exclusively in the crust.," 3.1: it is unlikely that the poloidal field is reduced to $\sim 10^4$ T in a recycled pulsar like the surface field, given the high electrical conductivity expected in the core \citep{gr92}, , except in the special situation where the source currents reside exclusively in the crust."332 We recalculate € usingEq. (, We recalculate $\epsilon$ usingEq. (333"7). substituting Boch lor D; and keeping j/n=1)"".","7), substituting $B_\mathrm{s, actual}$ for $B_\mathrm{s}$ and keeping $\mu_r=10^{-3}$."334 We plot the results for £=10.1 as the rec curves for A=1 (bottom solid red curve) and A=10* (top solid red curyο) in Fig. 2.., We plot the results for $\xi=10^{-4}$ as the red curves for $\Lambda=1$ (bottom solid red curve) and $\Lambda=10^{-3}$ (top solid red curve) in Fig. \ref{eps}.335 We also plot ο for the case of£=10 ?as the blue curves for A=1 (bottom solid blue curve) and A=10 (top solid. blue curve)., We also plot $\epsilon$ for the case of $\xi=10^{-2}$ as the blue curves for $\Lambda=1$ (bottom solid blue curve) and $\Lambda=10^{-3}$ (top solid blue curve).336 For clarity. the Case µ=fta is not presented: it lies three decades lower than the curves with j=07.," For clarity, the case $\mu=\mu_0$ is not presented; it lies three decades lower than the curves with $\mu_r=10^{-3}$."337 Without a better knowledge of the screening/burial process than is currently at hand. it is best simply to bracket he plausible range 10.1leg<£107? inferred from population svnthesis stuclies (Ixieletal.2008).," Without a better knowledge of the screening/burial process than is currently at hand, it is best simply to bracket the plausible range $10^{-4}\leqslant \xi\leqslant 10^{-2}$ inferred from population synthesis studies \citep{ketal08}."338". Using the potentially stronger shiclelecl pre-accretion fields. the c(D.) curves inl""ug."," Using the potentially stronger shielded pre-accretion fields, the $\epsilon(B_\mathrm{s})$ curves in Fig."339 2. come close to the spin-down limits on c., \ref{eps} come close to the spin-down limits on $\epsilon$.340 Now we can see that. aside from perhaps one object. namely PSR 1599302LA. all the pulsars plotted must have some internal toroidal Lele component. ie. A«1. if one has £x10. 7.," Now we can see that, aside from perhaps one object, namely PSR J1823–3021A, all the pulsars plotted must have some internal toroidal field component, i.e. $\Lambda<1$, if one has $\xi\leqslant 10^{-2}$ ."341 From the red curves and the dots in Fig. 2..," From the red curves and the dots in Fig. \ref{eps},"342 many objects seem to have A« 0.01. even «10.7 in the notable case of PSR J1910.5959€. On the other hand. the heavilv-shieldedcase 4-10 .£—101 (for example) seems to be ruled out by observations (top soid red curve in Fig. 2)).," many objects seem to have $\Lambda<0.01$ , even $\Lambda<10^{-3}$ in the notable case of PSR J1910–5959C. On the other hand, the heavily-shieldedcase $\Lambda=10^{-3}$ , $\xi=10^{-4}$ (for example) seems to be ruled out by observations (top solid red curve in Fig. \ref{eps}) )."343 In general. X cannot," In general, $\Lambda$ cannot"344this surface the radio emission drops to zero abruptly.,this surface the radio emission drops to zero abruptly.345 This causes the discontinuities in the modeled profiles inwards of the peak luminosity towards the cluster center., This causes the discontinuities in the modeled profiles inwards of the peak luminosity towards the cluster center.346" In the GMRT and WSRT images, the relic's surface brightness fades towards the northern and southern ends."," In the GMRT and WSRT images, the relic's surface brightness fades towards the northern and southern ends."347" This could (partly) be explained by the spherical shell model we use for the relic, as the relic's extent into the plane of the sky decreases at the northern and southern ends."," This could (partly) be explained by the spherical shell model we use for the relic, as the relic's extent into the plane of the sky decreases at the northern and southern ends."348 It is also possible that the surface brightness across the shell decreases towards the edges., It is also possible that the surface brightness across the shell decreases towards the edges.349 This effect is not included in our model., This effect is not included in our model.350" However, our goal was not to reproduce the exact profile of the relic, but rather to show that although projection effects can be significant, a clear spectral index gradient can remain."," However, our goal was not to reproduce the exact profile of the relic, but rather to show that although projection effects can be significant, a clear spectral index gradient can remain."351" Even for an opening angle of 40°(a total of into and out of the plane of the sky), a steepening of more than 0.5 units in the spectral index is predicted towards the cluster center."," Even for an opening angle of (a total of into and out of the plane of the sky), a steepening of more than 0.5 units in the spectral index is predicted towards the cluster center."352" Based on this we argue that although relic RE is widened significantly by projection effects, the fact the we see a clear spectral index gradient is not surprising."," Based on this we argue that although relic RE is widened significantly by projection effects, the fact the we see a clear spectral index gradient is not surprising."353 This could also explain the spectral index gradients visible for the relics significantly larger than the maximum intrinsic widths.," This could also explain the spectral index gradients visible for the relics observed by \cite{1997MNRAS.290..577R, 2006AJ....131.2900C, 2007A&A...467..943O, 2008A&A...486..347G, 2009A&A...494..429B}, even though the observed widths are significantly larger than the maximum intrinsic widths."354" Since the width of the relic is larger than the maximum intrinsic width, we estimate the magnetic field at the location of the relics by assuming minimum energy densities in the relics."," Since the width of the relic is larger than the maximum intrinsic width, we estimate the magnetic field at the location of the relics by assuming minimum energy densities in the relics."355" We use the same procedure as described in ? and take k=100, i.e, the ratio of energy in relativistic protons to that in electrons."," We use the same procedure as described in \cite{2009A&A...506.1083V} and take $k=100$, i.e, the ratio of energy in relativistic protons to that in electrons."356" For relic RW, we have a spectral index of -1.49, and a surface brightness of 1.2uJy arcsec?."," For relic RW, we have a spectral index of $1.49$, and a surface brightness of $1.2~\mu$ Jy $^{-2}$."357 We take 290 kpc for the depth (d) along the line of sight., We take 290 kpc for the depth $d$ ) along the line of sight.358 This gives Beg=3.4uGauss., This gives $B_{\mathrm{eq}} = 3.4~\mu$ Gauss.359" For RE we have a spectral index of -1.59, a surface brightness of 0.96 uJy arcsec~*, and we assume d=1 Mpc."," For RE we have a spectral index of $1.59$, a surface brightness of $0.96~\mu$ Jy $^{-2}$, and we assume $d=1$ Mpc."360 This gives Beq=2.5 uGauss., This gives $B_{\mathrm{eq}} = 2.5~\mu$ Gauss.361 The equipartition magnetic field strength scales with (1+k)*/7., The equipartition magnetic field strength scales with $(1+k)^{2/7}$.362" In the above calculation, we used fixed frequency cutoffs (Vmin=10 MHz and v4,=100 GHz), which is not entirely correct (??).. "," In the above calculation, we used fixed frequency cutoffs $\nu_{\mathrm{min}} = 10$ MHz and $\nu_{\mathrm{max}} = 100$ GHz), which is not entirely correct \citep{2005AN....326..414B, 1997A&A...325..898B}."363"With low and high energy cutoffs (Ymins Ymax)> Ymin<Ymax, and fixing Ymin to 100, we find a revised magnetic field strength (B;,) of 7.9 and 6.6 Gauss for RW and RE, respectively."," With low and high energy cutoffs $\gamma_{\mathrm{min}}$, $\gamma_{\mathrm{max}}$ ), $\gamma_{\mathrm{min}} \ll \gamma_{\mathrm{max}}$, and fixing $\gamma_{\mathrm{min}}$ to 100, we find a revised magnetic field strength $B^{\prime}_{\mathrm{eq}}$ ) of 7.9 and 6.6 $\mu$ Gauss for RW and RE, respectively."364" For a lower cutoff of yq,=5000, we get 1.4 and 1.0 wGauss for RW and RE, respectively."," For a lower cutoff of $\gamma_{\mathrm{min}}=5000$, we get 1.4 and 1.0 $\mu$ Gauss for RW and RE, respectively."365" The revised equipartition magnetic field strength (Beg) scales with (14&)!/0-9, for different values of k."," The revised equipartition magnetic field strength $B^{\prime}_{\mathrm{eq}}$ ) scales with $(1+k)^{1/(3-\alpha)}$, for different values of $k$."366" We discovered a double radio relic in the galaxy cluster ZwCl 0008.8+5215, located at z=0.103 (based on a single spectroscopic redshift)."," We discovered a double radio relic in the galaxy cluster ZwCl 0008.8+5215, located at $z=0.103$ (based on a single spectroscopic redshift)."367" A ROSAT X-ray image and galaxy iso-density map show that the cluster is undergoing a binary merger event, with the merger axis oriented roughly east-west."," A ROSAT X-ray image and galaxy iso-density map show that the cluster is undergoing a binary merger event, with the merger axis oriented roughly east-west."368" The two radio relics are located along this merger axis, while their orientation is perpendicular to this axis."," The two radio relics are located along this merger axis, while their orientation is perpendicular to this axis."369" The relics probably trace shocks waves in the ICM, created by the merger event, in which particles are (re)accelerated by the DSA mechanism."," The relics probably trace shocks waves in the ICM, created by the merger event, in which particles are (re)accelerated by the DSA mechanism."370 Integrated radio spectra are consistent with particle acceleration in the shock by DSA and indicate Mach numbers of —2 for the shocks., Integrated radio spectra are consistent with particle acceleration in the shock by DSA and indicate Mach numbers of $\sim 2$ for the shocks.371 The spectral index for both relics shows a steepening towards the cluster center., The spectral index for both relics shows a steepening towards the cluster center.372" Parts of the relics have a polarization fraction in the range of 5—2596, but further observations are needed to better map the polarization properties."," Parts of the relics have a polarization fraction in the range of $5-25\%$, but further observations are needed to better map the polarization properties."373 The relics have an extent of 1.4 Mpc and 290 kpc., The relics have an extent of 1.4 Mpc and 290 kpc.374 This factor of five difference in their linear extent is unlike that of other known double relic systems., This factor of five difference in their linear extent is unlike that of other known double relic systems.375 The size difference could be related to a relatively large mass ratio between the, The size difference could be related to a relatively large mass ratio between the376Iu the Fieure 2dd there is production height distribution (projection of the Figure 2bb on the height8 axis).,In the Figure \ref{fig:all_ph}d d there is production height distribution (projection of the Figure \ref{fig:all_ph}b b on the height axis).377 The bump around 1.3 % 810 loeyy(200 11) is duc to the aceuuulatect signals8 from muons., The bump around 4.3 $\approx$ $_{10}$ (200 m) is due to the accumulated signals from muons.378 From all chareed, From all charged379GeV energies (22222).,GeV energies .380. In this period. this source exhibited variability and alsosome soft y-ray variability).," In this period, this source exhibited X-ray variability and alsosome soft $\gamma$ -ray variability."381. However. found that the EGRET flux was stable during the whole period of observations.," However, found that the EGRET flux was stable during the whole period of observations."382 In 1999 the new X-ray Observatory took images of AA with an unprecedented resolution., In 1999 the new X-ray Observatory took images of A with an unprecedented resolution.383 More than 200 X-ray point sources were identified in those images(?)., More than 200 X-ray point sources were identified in those images.384. AA as a possible source of UHE cosmic rays was early proposed by(?)., A as a possible source of UHE cosmic rays was early proposed by.385 Recently. the Pierre Auger Collaboration reported the existence of anisotropy on the arrival directions of UHE cosmie rays(2).. remarking that at least 2 of this events can be correlated with the AA position (3° circle).," Recently, the Pierre Auger Collaboration reported the existence of anisotropy on the arrival directions of UHE cosmic rays, remarking that at least 2 of this events can be correlated with the A position $3^{\circ}$ circle)."386 Further works have claimed that there are several events that can be associated with AA and its big radio lobes but this correlation is still statistically weak., Further works have claimed that there are several events that can be associated with A and its big radio lobes but this correlation is still statistically weak.387 Finally. Fernti/LAT has detected AA in the first three months of survey with a significance above 10σ(?)..," Finally, /LAT has detected A in the first three months of survey with a significance above $\sigma$."388 The giant radiogalaxy M87 ts located at MMpe within the Virgo cluster(2)., The giant radiogalaxy M87 is located at Mpc within the Virgo cluster.389. It presents a one-sided jet which is inclined with respect to the lineof sight an angle between 20° - 40°(??).., It presents a one-sided jet which is inclined with respect to the lineof sight an angle between $^\circ$ - $^\circ$.390 In addition to its bright and well resolved jet. M87 harbors a very massive black hole (6.0 + 0.5) x 10° M. which is thought to power the relativistic outflow.," In addition to its bright and well resolved jet, M87 harbors a very massive black hole (6.0 $\pm$ 0.5) $\times$ $^9$ $_\odot$ which is thought to power the relativistic outflow."391 Given its proximity. the substructures inside the jet could be resolved in the X-ray. optical. and radio wavebands(2).," Given its proximity, the substructures inside the jet could be resolved in the X-ray, optical, and radio wavebands."392. High frequency VLBI observations have resolved the inner Jet up to about 70 Schwarzschild radii(2)., High frequency VLBI observations have resolved the inner jet up to about 70 Schwarzschild radii.393. Along the jet. nearly stationary components and features moving at superluminal speeds were observed (100pe-scale).," Along the jet, nearly stationary components and features moving at superluminal speeds were observed (100pc-scale)."394 M87 is also a well-known VHE y-rays emitter showing a y-ray flux variability on short time scales with flaing phenomena in VHE. radio. and X-ray wavebands simultaneously(?).," M87 is also a well-known VHE $\gamma$ -rays emitter showing a $\gamma$ -ray flux variability on short time scales with flaring phenomena in VHE, radio, and X-ray wavebands simultaneously."395. Recently. it was detected by Fermni/LAT with à significance greater. than IOc in 10 months of observations(?).," Recently, it was detected by /LAT with a significance greater than $\sigma$ in 10 months of observations."396. Rapid variability constrains the emission region extent to less than =S50Rs. where 6 is the relativistic Doppler factor.," Rapid variability constrains the emission region extent to less than $\approx 5 \delta R_{\rm S}$, where $\delta$ is the relativistic Doppler factor."397 Some suggested explanations for the VHE y-ray emissio were ruled out (e.g. dark matter annihilation (2))) At the same time. various VHE y-ray jet emission models were proposed: leptonic and hadronic ones.," Some suggested explanations for the VHE $\gamma$ -ray emission were ruled out (e.g. dark matter annihilation ) At the same time, various VHE $\gamma$ -ray jet emission models were proposed: leptonic and hadronic ones."398 However. the locatio of the emission region is still unknown.," However, the location of the emission region is still unknown."399 The nucleus. the inner jet(?).. or larger structures in the jet such as the knot HST- have been discussed as possible sites of particle acceleratio(?).," The nucleus, the inner jet, or larger structures in the jet such as the knot HST-1, have been discussed as possible sites of particle acceleration."400. We assume that a population of relativistic particles can be accelerated to very high energies close to the base of the AGN Jet., We assume that a population of relativistic particles can be accelerated to very high energies close to the base of the AGN jet.401 These primary electrons and protons carry a fraction of the total kinetic power of the jet Lihin)™ and as they are dragged along with the jet. they cool giving rise to electromagnetic emission and neutrinos.," These primary electrons and protons carry a fraction of the total kinetic power of the jet $L_{\rm j}^{\rm(kin)}$, and as they are dragged along with the jet, they cool giving rise to electromagnetic emission and neutrinos."402 Assuming that a fraction q; of the Eddington luminosity is carried by the jet and a counter jet. the jet kinetic power Is - This power can be very high if the jet is launched by a dissipationless accretion disk(?).," Assuming that a fraction $q_{\rm j}$ of the Eddington luminosity is carried by the jet and a counter jet, the jet kinetic power is = This power can be very high if the jet is launched by a dissipationless accretion disk."403. Most of the jet content is in the form of a thermal plasma with a constant bulk Lorentz factor Εν., Most of the jet content is in the form of a thermal plasma with a constant bulk Lorentz factor $\Gamma_{\rm b}$.404 This plasma is initially in equipartition with a tangled magnetic field at the Alfvén surface (So=50K. from the central black hole) ?)., This plasma is initially in equipartition with a tangled magnetic field at the $\acute{e}$ n surface $z_0= 50 R_{\rm g}$ from the central black hole) .405. The highly disorganized magnetic field has a root mean square value B(z) at a distance z from the black hole in the observer frame. such that P(z):=(9)+B»(B2).," The highly disorganized magnetic field has a root mean square value $B(z)$ at a distance $z$ from the black hole in the observer frame, such that $B^2(z):= \langle B_x^2 \rangle+ \langle B_y^2 \rangle+ \langle B_z^2 \rangle$."406 The magnetic energy density for z=zo is then οµίςο)=B;(85) with Bo=B(z;) , The magnetic energy density for $z=z_0$ is then $\rho_{\rm m}(z_0)= B_0^2/(8\pi)$ with $B_0=B(z_0)$ .407Equating the magnetic to the kinetic energy density. yields: =——L—.. where zo 1s the distance to the black hole. vy is the jet velocity. and rj(zo)=zotan& is the radius of the jet assuming that it has a conical shape with half-opening angle &.," Equating the magnetic to the kinetic energy density, yields: =, where $z_0$ is the distance to the black hole, $v_{\rm b}$ is the jet velocity, and $r_{\rm j}(z_0)=z_0 \tan{\xi_{\rm j}}$ is the radius of the jet assuming that it has a conical shape with half-opening angle $\xi_{\rm j}$ ."408 A widely accepted view ts that jets are accelerated through the conversion of magnetic energy into kinetic energy?)., A widely accepted view is that jets are accelerated through the conversion of magnetic energy into kinetic energy.409. We adopt a phenomenological dependence on the distance to the black hole for the magnetic field?).. B(z)-(c). Since the density of cold material within the jet decays as οτι using an exponent η€(1.2) in the above expressio=) implies that. as z increases. the magnetic energy decreases more rapidly than the kinetic one.," We adopt a phenomenological dependence on the distance to the black hole for the magnetic field, B(z)=. Since the density of cold material within the jet decays as $z^{-2}$, using an exponent $m\in(1,2)$ in the above expression implies that, as $z$ increases, the magnetic energy decreases more rapidly than the kinetic one."410 The corresponding increase in the bulk Lorentz factor is taken into account as described in Appendix ??.., The corresponding increase in the bulk Lorentz factor is taken into account as described in Appendix \ref{appdx_gam}.411 In the following. we will write simply Dy. but it actually depends on z.," In the following, we will write simply $\Gamma_{\rm b}$, but it actually depends on $z$."412 The particle acceleratior takes place in a compact but inhomogeneous region of size Az<Maccz near the base of the jet. at a distance τς away from the black hole.," The particle acceleration takes place in a compact but inhomogeneous region of size $\Delta z < z_{\rm acc}$ near the base of the jet, at a distance $z_{\rm acc}$ away from the black hole."413 The value of Sace IS fixed by requiring the nagnetic energy density to be in sub-partition with the jet kinetic energy density., The value of $z_{\rm acc}$ is fixed by requiring the magnetic energy density to be in sub-partition with the jet kinetic energy density.414 This condition enables strong shocks to develop(?)., This condition enables strong shocks to develop.415. = Logi, = z_0.416"n(4) For example. if Mp,=10°Mas. i= 1.5. and τρ= 50Ν.. using gm=0.38 yields a distance zy= 132Re."," For example, if $M_{bh}=10^8 M_\odot$, $m=1.5$ , and $z_0=50 R_g$ , using $q_{\rm m}=0.38$ yields a distance $z_{\rm acc}=132 R_{\rm g}$ ."417" The power injected in the form of relativistic particles (L4) is consideredto be a small fraction g,., of the total jet kinetic power and the relation between proton and electron powers ts given by the parameter e such that L,=« ?2)..", The power injected in the form of relativistic particles $L_{\rm rel}$ ) is consideredto be a small fraction $q_{\rm rel}$ of the total jet kinetic power and the relation between proton and electron powers is given by the parameter $a$ such that $L_p= a \ L_e$ .418of the S/Irr population (although the velocity scale of the S/Irr galaxies has not changed).,of the S/Irr population (although the velocity scale of the S/Irr galaxies has not changed).419 Phe values for the biweight velocity centroids and seales of the various samples are given in ‘Table 3.., The values for the biweight velocity centroids and scales of the various samples are given in Table \ref{norsubstr}.420 The offset in the biweight velocity centroid between the 15/80 and S/lrr sample remains albeit slightly lower and is 222 km  with a significance of Sy=Ls., The offset in the biweight velocity centroid between the E/S0 and S/Irr sample remains albeit slightly lower and is 222 km $^{-1}$ with a significance of $S_{\rm V} = 1.8$.421 The velocity centroid of the 1050 galaxies has not changed by extending the sample to a larger radius (see Table 3)). although its biweight. veloci voscale is somewhat smaller: Sp; = 901 km 1 for He za compared to σοι — 0641 km s+ for BR< ," The velocity centroid of the E/S0 galaxies has not changed by extending the sample to a larger radius (see Table \ref{norsubstr}) ), although its biweight velocity scale is somewhat smaller; $S_{\rm BI}$ = 901 km $^{-1}$ for $R < \frac{2}{3} R_A$ , compared to $S_{\rm BI}$ = 964 km $^{-1}$ for $R < \frac{1}{3} R_A$ ."422In terms of their spatial distribution. the elliptical and spiral populations now both reveal significant elongation and have position angles of 102 anc 101. respectively. in the equatorial coordinate frame.," In terms of their spatial distribution, the elliptical and spiral populations now both reveal significant elongation and have position angles of $^{\circ}$ and $^{\circ}$ respectively, in the equatorial coordinate frame."423 This corresponds to »osition angles of 146° and 145° in the Galactic coordinate rane., This corresponds to position angles of $^{\circ}$ and $^{\circ}$ in the Galactic coordinate frame.424 The later are again indicated by arrows in the op-right of t1e middle panels in Fig. S., The latter are again indicated by arrows in the top-right of the middle panels in Fig. \ref{dsall}.425 The à-test. now clearly reveals substructure in the S/LIrr sample (only of he AIC simulations show a larger degree of substructure)., The $\delta$ -test now clearly reveals substructure in the S/Irr sample (only of the MC simulations show a larger degree of substructure).426 Interestingly. the é-test shows that the E/SO population is completely free. of any detectable substructure.," Interestingly, the $\delta$ -test shows that the E/SO population is completely free of any detectable substructure."427" In Fig. δ.,"," In Fig. \ref{dsall},"428" he results from the 2-test of the combined sample out to ""as2Ry is shown in the middle panels. where the upper-DPImicelle panel shows the I2/8O galaxies and the lower-miclelle xuiels clisplavs the S/Irr galaxies."," the results from the $\delta$ -test of the combined sample out to $R < \frac{2}{3} R_A$ is shown in the middle panels, where the upper-middle panel shows the E/SO galaxies and the lower-middle panels displays the S/Irr galaxies."429 Our final sample extends out. to. the full Abell radius of the Norma cluster., Our final sample extends out to the full Abell radius of the Norma cluster.430 Within this region.there are 296 cluster members (Sect.," Within this region,there are 296 cluster members (Sect."431 3.1) of whieh 107 are classifiedE/80, 3.1) of which 107 are classifiedE/S0432"just past f22 0,3Cw reat wluch time it rapidly plinges to very low rotationC rates.","just past $t \approx 4330.3$ Gyr, at which time it rapidly plunges to very low rotation rates."434 We note that the plysical rotation rate of this done restar never actually decreases. aud that the decrease 1ji fy is caused bv the orbital aneular velocity at pc‘Vlastron increasing at a rate faster than the actual rotation rate of the donor star.," We note that the physical rotation rate of this donor star never actually decreases, and that the decrease in $f_1$ is caused by the orbital angular velocity at periastron increasing at a rate faster than the actual rotation rate of the donor star."435 The rapid decrease starts a few «ybis before the svstem transitions from a self-accretion svsteni to a system du which the ejected matter mipacts neither star l nor star 2 during the course of one binary orbit., The rapid decrease starts a few orbits before the system transitions from a self-accretion system to a system in which the ejected matter impacts neither star 1 nor star 2 during the course of one binary orbit.436 Since our code is eureitly not equipped to deal wit1i the evolution of svstenis like this. we teqTuimated the calculation at this point.," Since our code is currently not equipped to deal with the evolution of systems like this, we terminated the calculation at this point."437 We note that at every step of the orbital evolulon calculations. boti the total linear and augular 1ionicuta of the svstei renain constant.," We note that at every step of the orbital evolution calculations, both the total linear and angular momenta of the system remain constant."438 However. nass overflow can exchange niomentum between the spins of the binary conrponents and the orbit.," However, mass overflow can exchange momentum between the spins of the binary components and the orbit."439 We particularly fud that. contrarv fo pas assuniptious (forexame.see?7).. direct iupact ac‘cretion does not necessarily provide a sink of orbital angular 1001101.uu (seealsο2)," We particularly find that, contrary to past assumptions \citep[for example, 440see][]{1988ApJ...332..193V, 2004MNRAS.350..113M}, direct impact accretion does not necessarily provide a sink of orbital angular momentum \citep[see also][]{2007ApJ...670.1314M}."441 dustead. the transferred. natter coiteüus both spiji aud orbital aedlar momentum from the donor star. rt of which cal )o returned o the orbit uyon accretion.," Instead, the transferred matter contains both spin and orbital angular momentum from the donor star, part of which can be returned to the orbit upon accretion."442 This makes it possible for the orbital a18oOular momenta to imerease at t1e expense of thie spin aneulu momentum of the Loux., This makes it possible for the orbital angular momentum to increase at the expense of the spin angular momentum of the donor.443 The mipact of this resIt on the stability of mass OVCLflow iu close biuudes SHCji as double white chwarts will o the subject of a forlhicomiug investigation., The impact of this result on the stability of mass overflow in close binaries such as double white dwarfs will be the subject of a forthcoming investigation.444" Lastly we cau see frou, Figures 9 aud 10. that the tota enerev of the system can chiuge siguificautlv over the course of its lifeiue."," Lastly, we can see from Figures \ref{fig-diev} and \ref{fig-saev} that the total energy of the system can change significantly over the course of its lifetime."445 Due to our assunption that the particle is ejected aud accreted perfectly inclasticallv (Le. with πο loss of Oleiun). ijf is nupossble to conserve the total svseni energy as well.," Due to our assumption that the particle is ejected and accreted perfectly inelastically (i.e., with no loss of momentum), it is impossible to conserve the total system energy as well."446" Even iu the case of selt-accretion where the uias of star d ds uuchauged the ejected particle seIf-accrees with a differcut position aud velocity than tha with w""ich is was ejected.", Even in the case of self-accretion where the mass of star 1 is unchanged the ejected particle self-accretes with a different position and velocity than that with which is was ejected.447 For the systems tested here. he toted euergev can change by as nich as LOW or more. aud ca 1cither increase or decrease depending on the svsen paracters.," For the systems tested here, the total energy can change by as much as $10\%$ or more, and can either increase or decrease depending on the system parameters."448 The energy. added or subtracted from fjo systeni ο] be accompanied by a commensurate cΙΟ 1i the thermal enerev of the stars’ cnvelopes., The energy added or subtracted from the system might be accompanied by a commensurate change in the thermal energy of the stars' envelopes.449 Tn this paper. we extended our previous work on lass transfer du ecceutric binaries by usine ballistic trajectory calculations to determine the parameter space for ciffereit outcomes of mass overflow.," In this paper, we extended our previous work on mass transfer in eccentric binaries by using ballistic trajectory calculations to determine the parameter space for different outcomes of mass overflow."450" Assinine Wass overflow takes place through Roche lobe overflow at the periasron of the binary orbit. we explo(d a broad parameter space to determine the conditions τιider which Wass overflow leads to direct mupact acerction onto he donors companion or fo faIback onto the donor star (""self-accreion)."," Assuming mass overflow takes place through Roche lobe overflow at the periastron of the binary orbit, we explored a broad parameter space to determine the conditions under which mass overflow leads to direct impact accretion onto the donor's companion or to fallback onto the donor star (“self-accretion”)."451 The restIts presented in this paTOS along with those presente¢ i] 777.. provide a Sclf-consistent victure of the orvital evolution of eccenuric effective Πεche lobe overfkiuge binary star svstenis as a fuuction «Xf their iuitial orjtal parameters.," The results presented in this paper along with those presented in \citet{2007ApJ...660.1624S, 4522007ApJ...667.1170S, 2009ApJ...702.1387S}, provide a self-consistent picture of the orbital evolution of eccentric effective Roche lobe overflowing binary star systems as a function of their initial orbital parameters."453" Until τιςny, such a picture has been lacsing both iu binary stelar evolution axd binary populaion svuthesis codes."," Until now, such a picture has been lacking both in binary stellar evolution and binary population synthesis codes."454 We fiud hat svstenis witji a large initial ecceutricitv (6;20. 2) Or a (loon initially rotating near svuchrouiciY (fii 0.5) aro always expected to undergo sef-accretion.," We find that systems with a large initial eccentricity $e_i \gtrsim 4550.2$ ) or a donor initially rotating near synchronicity $f_{1,i} \gtrsim 4560.8$ ) are always expected to undergo self-accretion."457 Direct impact accretion or disk formation is expected to occur mainly for svstems witli low cecentricities. low iuass ratios. and substantially subsvuchrorously rotating donor stars.," Direct impact accretion or disk formation is expected to occur mainly for systems with low eccentricities, low mass ratios, and substantially subsynchronously rotating donor stars."458 Furthermore. seltaccretion is found to always decrease the orbital eccentricity. and can do so over timescales ranging from loss than a Myr to more tiu a Cor depending on the initial binary parameters aud the chosen mass overflow rate.," Furthermore, self-accretion is found to always decrease the orbital eccentricity, and can do so over timescales ranging from less than a Myr to more than a Gyr depending on the initial binary parameters and the chosen mass overflow rate."459" To illustrate the appicability of the presented Ὁπμαντα, we caleulated he orbital evolution due to nass overflow for eccentric binaries consisting of anu evolved iain sequence cOuoL and a zero-age lain sequence accretor as a fiuetion of the initial binary ouwneters"," To illustrate the applicability of the presented formalism, we calculated the orbital evolution due to mass overflow for eccentric binaries consisting of an evolved main sequence donor and a zero-age main sequence accretor as a function of the initial binary parameters."460 For biuaries nudergoing direct impact accretion. niass transfer can increase as well decrease he initial orbital senianajor axis and eccentricity. while or binaries undergoing selfaccretion. lass overflow always decreases both the initial orbital scii-major axis and eccentricity.," For binaries undergoing direct impact accretion, mass transfer can increase as well decrease the initial orbital semi-major axis and eccentricity, while for binaries undergoing self-accretion, mass overflow always decreases both the initial orbital semi-major axis and eccentricity."461 For a iuass overflow rate of My=LOPALxv 1. the time scales of orbital evolution cui range from less than Cr to more than a Dabble time.," For a mass overflow rate of $\dot{M}_0=-10^{-9}\,M_\odot\, 462{\rm yr}^{-1}$ , the time scales of orbital evolution can range from less than Gyr to more than a Hubble time."463 Tn this exploratory study. both the donor star aud," In this exploratory study, both the donor star and"464(KWOS).,(KW08).465 Galaxy pairs only represent mergercandidates. as some apparent pairs will not coalesce even over many Gyr while others rapidly merge. depending on orbital properties. projection effects. and other internal factors.," Galaxy pairs only represent merger, as some apparent pairs will not coalesce even over many Gyr while others rapidly merge, depending on orbital properties, projection effects, and other internal factors."466 Thus. converting from pair fractions to merger rates requires an average timescale. where Ny; is the number of close pairs within a given redshift bin. and Nossa is the number of mergers per unit time within the same galaxy population.," Thus, converting from pair fractions to merger rates requires an average timescale, where $N_{\rm pair}$ is the number of close pairs within a given redshift bin and $\dot{N}_{\rm merger}$ is the number of mergers per unit time within the same galaxy population."467 One theoretical parametrization of this timescale was derived from the Millennium Simulation by KWO8 as a function of galaxy mass. redshift. and maximum projected separation.," One theoretical parametrization of this timescale was derived from the Millennium Simulation by KW08 as a function of galaxy mass, redshift, and maximum projected separation."468 Specifically. their formula for pairs with projected separations Arc30/1 kkpe. ήνΜι>1/4. and photometric redshifts gives a mean merger timescale of In the KWO8 formalism the same mass limit is imposed on both primary and secondary galaxies. while our sample includes lower-mass companions down to the survey limit.," Specifically, their formula for pairs with projected separations $\Delta r< 30 h^{-1}$ kpc, $M_2/M_1>1/4$, and photometric redshifts gives a mean merger timescale of In the KW08 formalism the same mass limit is imposed on both primary and secondary galaxies, while our sample includes lower-mass companions down to the survey limit."469 The resulting KWO8 pair mass ratios depend on primary galaxy mass (1.e.. closer to equal-mass near the survey limit).," The resulting KW08 pair mass ratios depend on primary galaxy mass (i.e., closer to equal-mass near the survey limit)."470 For consistency. we thus repeat the above exercise with their selection method.," For consistency, we thus repeat the above exercise with their selection method."471 Two mass thresholds are applied: log(M./M..)>10.5 as before. and log(M./M..)>10.0 to take full advantage of the survey.," Two mass thresholds are applied: $\log (M_\star/M_\odot)>10.5$ as before, and $\log (M_\star/M_\odot)>10.0$ to take full advantage of the survey."472 These fractions. shown in the top panels of Figure 3.. are somewhat lower than in Figure | where secondary galaxies below the primary mass threshold are included.," These fractions, shown in the top panels of Figure \ref{fig_mergers}, , are somewhat lower than in Figure \ref{fig_pairs} where secondary galaxies below the primary mass threshold are included."473 Figure 3. (bottom) shows the specific merger rates. calculated with eq. (," Figure \ref{fig_mergers} (bottom) shows the specific merger rates, calculated with eq. ("4742). from z20.4—2 at the two mass thresholds.,"2), from $z=0.4-2$ at the two mass thresholds."475 Adopting the KW0O8 timescale. mergers are relatively rare: only about 0.51% of massive quiescent galaxies merge with quiescent companions— each Gyr. and fewer than of galaxies over this redshift range will undergo major mergers.," Adopting the KW08 timescale, mergers are relatively rare: only about $0.5-1$ of massive quiescent galaxies merge with quiescent companions each Gyr, and fewer than of galaxies over this redshift range will undergo major mergers."476 This is due both to the rarity of massive galaxy pairs and to the long effective merger timescales from KWO8. about 2-3.5 Gyr depending on mass and redshift.," This is due both to the rarity of massive galaxy pairs and to the long effective merger timescales from KW08, about 2-3.5 Gyr depending on mass and redshift."477" Although the KWO8 formalism is most applicable to our specific pair selection parameters. a variety of timescales have been employed in the literature: this is discussed further in refsec,imescale.."," Although the KW08 formalism is most applicable to our specific pair selection parameters, a variety of timescales have been employed in the literature; this is discussed further in \\ref{sec_timescale}."478 Close pairs of galaxies are easily identified and detectable to high redshift. and are therefore in principle a robust way of identifying systems that may merge within a relatively short timeseale.," Close pairs of galaxies are easily identified and detectable to high redshift, and are therefore in principle a robust way of identifying systems that may merge within a relatively short timescale."479 But it is less straightforward to determine this timescale and convert the measured pair fractions to merger rates., But it is less straightforward to determine this timescale and convert the measured pair fractions to merger rates.480 Several estimates of the merger timescale have been used in previous work., Several estimates of the merger timescale have been used in previous work.481 Belletal.(2006) assume that galaxy pairs merge within roughly one orbital time. m their case ~O.4 GGyr. while the KWOS8 estimate is nearly an order of magnitude larger for the galaxy masses considered here.," \citet{bell06} assume that galaxy pairs merge within roughly one orbital time, in their case $\sim 0.4$ Gyr, while the KW08 estimate is nearly an order of magnitude larger for the galaxy masses considered here."482 These fundamentally change the interpretation of the measured pair fractions: with the Belletal.(2006) timescale. major mergers play a significant role in the assembly of massive galaxies over time: assuming KWOS. only about of massive quiescent galaxies have undergone major mergers since z=2. with ~2/3 of this occurring at z<I.," These fundamentally change the interpretation of the measured pair fractions: with the \citet{bell06} timescale, major mergers play a significant role in the assembly of massive galaxies over time; assuming KW08, only about of massive quiescent galaxies have undergone major mergers since $z=2$, with $\sim 2/3$ of this occurring at $z<1$."483 This disagreement is largely a result of the KWO8 analysis including physically-associated galaxy pairs which are at relatively large real distances despite having close projected separations. and therefore merge only after a long period (or not at all).," This disagreement is largely a result of the KW08 analysis including physically-associated galaxy pairs which are at relatively large real distances despite having close projected separations, and therefore merge only after a long period (or not at all)."484 Nonetheless. 1f the merger timescale isn't a strong function of redshift. the unchanging pair fraction we measure reflects a similarlyconstant merger rate since z=2.," Nonetheless, if the merger timescale isn't a strong function of redshift, the unchanging pair fraction we measure reflects a similarlyconstant merger rate since $z=2$."485 The pair fractions derived here are in. broad agreement with previous work at zXI., The pair fractions derived here are in broad agreement with previous work at $z\la 1$.486" Bundyetal.(2009). find a low (—ο, non-evolving fraction of massive galaxies in ""major pairs” at z«1.2. with I0!! ΜΜ... galaxies more likely to have close companions than those with 10' ..."," \citet{bundy09} find a low $\sim 4$ ), non-evolving fraction of massive galaxies in “major pairs” at $z<1.2$ , with $10^{11}$ $_\odot$ galaxies more likely to have close companions than those with $10^{10}$ $_\odot$."487 Given their higher mass limit and smaller search radius (20/7! kkpe). our z«| measurements appear to be in agreement a total pair fraction of —4—8%.. depending on whether the primary galaxies are star-forming or quiescent. and no strong evolution in the fraction.," Given their higher mass limit and smaller search radius $20 h^{-1}$ kpc), our $z<1$ measurements appear to be in agreement: a total pair fraction of $\sim 4-8$, depending on whether the primary galaxies are star-forming or quiescent, and no strong evolution in the fraction."488 In addition. we confirm their reported higher incidence of “dry pairs” at lower redshifts. simply due to the coincident increase in the number density of massive quiescent galaxies.," In addition, we confirm their reported higher incidence of “dry pairs” at lower redshifts, simply due to the coincident increase in the number density of massive quiescent galaxies."489" Between 0.4<z«0.8. Belletal.(2006) report a pair fraction of 5+1% for galaxies above 2.5«.10'"" .. also in agreement with our pair fraction measurement MMat the same redshift."," Between $0.4<z<0.8$, \citet{bell06} report a pair fraction of $5\pm 1$ for galaxies above $2.5\times 10^{10}$ $_\odot$, also in agreement with our pair fraction measurement at the same redshift."490 One common theme in these studies is the rarity of dry mergers: even Belletal.(2006).. with their short assumed timescale. find that only ~50% of massive galaxies have undergone major mergers since z=0.8: Bundyetal. estimate at the high-mass (10 MM.) end (butseePadillaetal. 2011).," One common theme in these studies is the rarity of dry mergers: even \citet{bell06}, with their short assumed timescale, find that only $\sim 50$ of massive galaxies have undergone major mergers since $z=0.8$; \citet{bundy09} estimate at the high-mass $10^{11}$ $_\odot$ ) end \citep[but see][]{padilla11}."491.. In their analysis of the environmental dependence of merger rate. Linetal.(2011) estimate a somewhat higher rate inenvironments: 1.240.3 major dry mergers per galaxy since = ]. perhaps not surprising since such environments harbor a larger fraction of quiescent galaxies (e.g.Kauffmannetal.2004;Quadri 2011).," In their analysis of the environmental dependence of merger rate, \citet{lin11} estimate a somewhat higher rate in: $1.2\pm 0.3$ major dry mergers per galaxy since $z=1$ , perhaps not surprising since such environments harbor a larger fraction of quiescent galaxies \citep[e.g.][]{kauffmann04,quadri11}."492. However. these only account for 38+10% of the mass accretion of these galaxies.," However, these only account for $38\pm 10$ of the mass accretion of these galaxies."493 If the merging timescale doesn't vary strongly with redshift. our results suggest that the major merger rate at |<z«2 is comparable to that at z«1: since twice as much cosmic time passes in the latter epoch than the former. this in turn implies that most major mergers occur below z~I.," If the merging timescale doesn't vary strongly with redshift, our results suggest that the major merger rate at $1<z<2$ is comparable to that at $z<1$; since twice as much cosmic time passes in the latter epoch than the former, this in turn implies that most major mergers occur below $z\sim 1$."494 Even when short timescales are assumed. dry mergers from ΞΞ2 to the present occur perhaps once or twice per galaxy at most.," Even when short timescales are assumed, dry mergers from $z=2$ to the present occur perhaps once or twice per galaxy at most."495 Since strong size and mass growth are nonetheless seen over this same redshift interval (e.g.Franxetal.Williamsetal. 2010).. it appears that major mergers are the primary driver behind the observed evolution.," Since strong size and mass growth are nonetheless seen over this same redshift interval \citep[e.g.][]{franx08,williams10}, it appears that major mergers are the primary driver behind the observed evolution."496 Indeed. if major dry mergers were the primary driver behind the smooth evolution in galaxy sizes and surfacedensities seen over 520—2. a much larger number would be required toeliminate the compact quiescent galaxy populationby z=0 (Tayloretal. 2010)..," Indeed, if major dry mergers were the primary driver behind the smooth evolution in galaxy sizes and surfacedensities seen over $z=0-2$, a much larger number would be required toeliminate the compact quiescent galaxy populationby $z=0$ \citep{taylor10}. ."497 Even 1f mergers are more common than the KWOS timescale implies. they still may not account for," Even if mergers are more common than the KW08 timescale implies, they still may not account for"498We estimate the right-haud side of (5.3)).,We estimate the right-hand side of \ref{sec5:eq4}) ).499 Beustein's inequality gives: We let s= nl , Benstein's inequality gives: We let $s=\frac{n+1}{2}$ .500Usiug (5.1)). we compute: Again. using (5.1)). we obtain: which. together with the fact that BS5~ΕΤ. aud estimate (5.1)) of Lemuna 5.1.. vield: The last term of the right-hand side of (5.3)) cau be estimated as follows: We know that Bs.x2C: the homogeneous Hóllder space whose semi-uorm cau be estimated as follows: This.together with (5.7)) vield:," Using \ref{sec5:eq5}) ), we compute: Again, using \ref{sec5:eq5}) ), we obtain: which, together with the fact that $\dot{B}^{s}_{2,2} \simeq501\dot{H}^{s}$, and estimate \ref{sec5:eq2}) ) of Lemma \ref{sec5:lem1}, yield: The last term of the right-hand side of \ref{sec5:eq4}) ) can be estimated as follows: We know that $\dot{B}^{\g}_{\infty,\infty} \simeq \dot{C}^{\g}$; the homogeneous Höllder space whose semi-norm can be estimated as follows: This,together with \ref{sec5:snh}) ) yield:"502with different structural parameters is indeed a crucial requirement for finally uuderstandiug (he formation mechanisms of (hese puzzling stus and their link with the cluster dynamical history.,with different structural parameters is indeed a crucial requirement for finally understanding the formation mechanisms of these puzzling stars and their link with the cluster dynamical history.503 We acknowledge R. Beclin and Y. Momany for useful discussions. and the referee. Robert D. Mathieu. for helpful suggestions in improving the paper.," We acknowledge R. Bedin and Y. Momany for useful discussions, and the referee, Robert D. Mathieu, for helpful suggestions in improving the paper."504 This research. was supported bv the Agenzia Spaziale Italiana (under contract ASLINAF 1/016/07/0). bv the Istituto Nazionale di Astvolisica (INAF. under contract PRIN-INAF2008) and by the Ministero dellIstruzione. cellUniversita e della Ricerea.," This research was supported by the Agenzia Spaziale Italiana (under contract ASI-INAF I/016/07/0), by the Istituto Nazionale di Astrofisica (INAF, under contract PRIN-INAF2008) and by the Ministero dell'Istruzione, dell'Università e della Ricerca."505polarizations both inpal and X4 collisionswell. Fromour analysis. weexpect Chatthe,incident spin-1/2 baryon is mainly composed of a spin-0 valence diquark and a valence506for ὐ= 30° and fy...=ry.,"for $\vartheta\,=$ $^{\circ}$ and $l_{\rm acc}\,=\,r_1$."507" The ambient photon energy density corresponds to a luminosity of if the particle acceleration region is al a distance 307, from the photon source. presumably the accretion disk."," The ambient photon energy density corresponds to a luminosity of if the particle acceleration region is at a distance $30\,r_1$ from the photon source, presumably the accretion disk."508" Comparing (he minimum beam bDuninositv required [or producing the observed. X-rav. and TeV 5-rav. emission ((15)) ) with L,.i we see that ve model requires an accretion flow with a radiative efficiency of or less."," Comparing the minimum beam luminosity required for producing the observed X-ray and TeV $\gamma$ -ray emission \ref{le}) )) with $L_{\rm a}$, we see that the model requires an accretion flow with a radiative efficiency of or less."509 Higher disk --unminosities are possible if (he emission frequency is sullicientiv high (hat Che hieh-energv Tteclrons interact only in the Ixlein-Nishina regime., Higher disk luminosities are possible if the emission frequency is sufficiently high that the high-energy electrons interact only in the Klein-Nishina regime.510 Lovelace(1976) assumed (hat protons may be accelerated. all the wav to ultra. high energies and that quasars thus may be accelerators of ultra high energy cosmic ravs (see Boldt&Ghosh1999:Loewenstein2000:Levinson2000. for similar recent papers).," \citet{Lovelace1976} assumed that protons may be accelerated all the way to ultra high energies and that quasars thus may be accelerators of ultra high energy cosmic rays (see \citealt*{Boldt1999,Boldt2000,Levinson2000} for similar recent papers)."511 Ile stipulated that the flow of high-energy protons may entrain or pick up electrons., He stipulated that the flow of high-energy protons may entrain or pick up electrons.512 Assuming that the high-energv electrons ancl protons would move with identical velocities and Lorentz factors. the acceleration of protons would increase the minimum beam lunmünositv by the proton to electron mass ratio.," Assuming that the high-energy electrons and protons would move with identical velocities and Lorentz factors, the acceleration of protons would increase the minimum beam luminosity by the proton to electron mass ratio."513 For Model 1. the required. luminosity. would exceed the Eddington huninositv by at least (wo orders of magnitude.," For Model 1, the required luminosity would exceed the Eddington luminosity by at least two orders of magnitude."514 For Model 2. accretion with a few times the Eddington rate would be sulficient.," For Model 2, accretion with a few times the Eddington rate would be sufficient."515" With D, 20.17 G. ((39)) predicts proton energies of zz3-10! eV and just the right electron energy of 20 TeV. However. a prohibitively strong toroidal magnetic field exceeding LO° G would be required so that the Povnting flux can power the massive electron-proton beam."," With $B_{\rm p,r_1}\approx$ 0.17 G, \ref{V12}) ) predicts proton energies of $\approx 3\cdot 10^{16}$ eV and just the right electron energy of $\approx$ 20 TeV. However, a prohibitively strong toroidal magnetic field exceeding $10^6$ G would be required so that the Poynting flux can power the massive electron-proton beam."516 The acceleration of electrons or positrons wilh subsequent entrainment of oppositely charged leptons would result in a beam with a much lower power., The acceleration of electrons or positrons with subsequent entrainment of oppositely charged leptons would result in a beam with a much lower power.517" However. ((39)) predicts an adequate voltage drop lor a very weak poloidal magnetic field with D,g& 2.10 1G. Even lor Model 2. the beam power requires again a toroidal magnetic field exceeding 105 C. Stronger {ρε and weaker D,,, would be viable if the leptons are accelerated to energies exceeding 40 TeV. and then entrain both electrons and positrons. slowing them down to a mean enerev of ~20 TeV per lepton."," However, \ref{V12}) ) predicts an adequate voltage drop for a very weak poloidal magnetic field with $B_{\rm p,r_1}\,\approx$ $2 \cdot 10^{-4}$ G. Even for Model 2, the beam power requires again a toroidal magnetic field exceeding $10^6$ G. Stronger $B_{\rm p,r_1}$ and weaker $B_{\rm t,r_1}$ would be viable if the leptons are accelerated to energies exceeding 40 TeV, and then entrain both electrons and positrons, slowing them down to a mean energy of $\sim$ 20 TeV per lepton."518" Acceleration of electrons or positrons with D,~ would produce a few particles with very high energies.", Acceleration of electrons or positrons with $B_{\rm t}\sim B_{\rm p}$ would produce a few particles with very high energies.519 A natural way of transferring the energy [rom a few high-energy particles io many low-energy particles are cascades., A natural way of transferring the energy from a few high-energy particles to many low-energy particles are cascades.520 Electromagnetic cascades in AGN jets have been discussed by (Burns&Lovelace1982:BlandfordLevinson1995: 1995).," Electromagnetic cascades in AGN jets have been discussed by \citep{Burns1982,Blandford1995,Levinson1995}."521. A generic discussion of electromagnetic cascades in the >TeV regime has been given in (Svensson1937)., A generic discussion of electromagnetic cascades in the $>$ TeV regime has been given in \citep{Svensson1987}.522. Unfortunately. cascade models need considerable fine tuning to produce the electron beams with the right. properties.," Unfortunately, cascade models need considerable fine tuning to produce the electron beams with the right properties."523 We briefly go through a specific scenario {ο emphasize some of the relevant difficulties., We briefly go through a specific scenario to emphasize some of the relevant difficulties.524 Electrons are accelerated until (he energy gains in, Electrons are accelerated until the energy gains in525555.,55.526 Table 5 shows the weighted mean values of the (6.2 pam) /(11.8 jm). (7.7 pan) /(11.8 jon). (8.6 pam) /(11.3. jm). (GC. po) /(6.2 pam) aad (8.6 pan) /(6.2 pam) raciance ratios and the standard deviation across (his sample.," Table \ref{ratios55_1} shows the weighted mean values of the (6.2 $\mu$ m)/(11.3 $\mu$ m), (7.7 $\mu$ m)/(11.3 $\mu$ m), (8.6 $\mu$ m)/(11.3 $\mu$ m), (7.7 $\mu$ m)/(6.2 $\mu$ m) and (8.6 $\mu$ m)/(6.2 $\mu$ m) radiance ratios and the standard deviation across this sample."527 Of the 95 ratios given in the table. we identilv 1l ratios that deviatle at more than the 2σ level: no variations at the 36 level or larger are identified.," Of the 95 ratios given in the table, we identify 11 ratios that deviate at more than the $\sigma$ level; no variations at the $\sigma$ level or larger are identified."528 Onulv two regions are identified as having variations in more than one band: 22 has &26 variations in all ratios except (6.2 jm) /(11.3. jm). while 116 has ~2o varialions in the (7.7 jum) /(11.3. jam)m. (8.6 jm)/(11.3. jm). and (7.7 jmm)/(6.2 jm) raciance ratios.," Only two regions are identified as having variations in more than one band: 2 has $\sim$ $\sigma$ variations in all ratios except (6.2 $\mu$ m)/(11.3 $\mu$ m), while 16 has $\sim$ $\sigma$ variations in the (7.7 $\mu$ m)/(11.3 $\mu$ m)m, (8.6 $\mu$ m)/(11.3 $\mu$ m), and (7.7 $\mu$ m)/(6.2 $\mu$ m) radiance ratios."529 We compare the radiances of PAIL features with emission lines and broad-band 224 jam dust continuum in Table 6.., We compare the radiances of PAH features with emission lines and broad-band 24 $\mu$ m dust continuum in Table \ref{ratios55_2}.530 The radiance of the [Ne LI] emission line compared to that of either a neutral (11.3 jm) or an ionized (8.6 san) PAIL emission feature show very few statistically significant variations within 555., The radiance of the [Ne II] emission line compared to that of either a neutral (11.3 $\mu$ m) or an ionized (8.6 $\mu$ m) PAH emission feature show very few statistically significant variations within 55.531 Similarly. the ratio of 24 yan to PAIT racliances is fairly uniform across this galaxy. though the standard deviation is much larger.," Similarly, the ratio of 24 $\mu$ m to PAH radiances is fairly uniform across this galaxy, though the standard deviation is much larger."532 We clraw attention to some interesting complexities in the data in Table 6.., We draw attention to some interesting complexities in the data in Table \ref{ratios55_2}.533 First. 110 is extremely. bright throughout the MIR. and especially so at 24. jn. (compare Figures 19. and 22)): it has the largest (24 pam)/(11.3 jun) ratio (a 236 deviation from the average) and one of the highest ratios of 24 sau radiance to the 8.6 ji PAIL band.," First, 10 is extremely bright throughout the MIR, and especially so at 24 $\mu$ m (compare Figures \ref{contour55} and \ref{region55}) ); it has the largest (24 $\mu$ m)/(11.3 $\mu$ m) ratio (a $>$ $\sigma$ deviation from the average) and one of the highest ratios of 24 $\mu$ m radiance to the 8.6 $\mu$ m PAH band."534 It is Iuninous in the UV and coincident with a high-surface brightness optical cluster., It is luminous in the UV and coincident with a high-surface brightness optical cluster.535 However. its ratio of [Ne IH] to PALL radiance is average.," However, its ratio of [Ne II] to PAH radiance is average."536 116 is also luminous at 24 jm. but contains no UV or optical counterpart and again has average [Ne IH] emission line ratios compared to the PAI features.," 16 is also luminous at 24 $\mu$ m, but contains no UV or optical counterpart and again has average [Ne II] emission line ratios compared to the PAH features."537 Finally. 119 is huninous ad all wavelengths (see Figure 19)) and has the largest [Ne IH] to PAIL ratio in 555.," Finally, 19 is luminous at all wavelengths (see Figure \ref{contour55}) ) and has the largest [Ne II] to PAH ratio in 55."538 Thus. although the conditions within individual regions in 555 are diverse. we lind no statistically significant. variations of PAIL/PAITI racliance ratios in the observed regions of 555.," Thus, although the conditions within individual regions in 55 are diverse, we find no statistically significant variations of PAH/PAH radiance ratios in the observed regions of 55."539 We examine four 52.5 pe radius apertures in 33109., We examine four 52.5 pc radius apertures in 3109.540 The low S/N ratio throughout much of the observed region limits our exploration to onlv those areas (hat are comparatively IR-bright (see Figure 20))., The low S/N ratio throughout much of the observed region limits our exploration to only those areas that are comparatively IR-bright (see Figure \ref{contour3109}) ).541 Negligible scatter is seen in the PDAIL/PAII racdiance ratios for (hese regions. as seen in Table τι only one aperture deviates from the weighted mean at the Yo significance level (the (6.2 j0m)/(11.3 jn) ratio for 44).," Negligible scatter is seen in the PAH/PAH radiance ratios for these regions, as seen in Table \ref{ratios3109_1}; only one aperture deviates from the weighted mean at the $\sigma$ significance level (the (6.2 $\mu$ m)/(11.3 $\mu$ m) ratio for 4)."542 This is in marked contrast to the very significant variations seen in the ratio of 24 jmi to PAIL radiance in these four regions of the galaxy., This is in marked contrast to the very significant variations seen in the ratio of 24 $\mu$ m to PAH radiance in these four regions of the galaxy.543 While (the weighted mean in Table 8 is affected by the difference in (he errorbars over the small number of apertures. it is clear (hat 11 ancl + have very different. properties.," While the weighted mean in Table \ref{ratios3109_2} is affected by the difference in the errorbars over the small number of apertures, it is clear that 1 and 4 have very different properties."544Figure | shows that the distributions of and of young clusters are obviously different from those of old clusters.,Figure \ref{fig1} shows that the distributions of and of young clusters are obviously different from those of old clusters.545 A notable characteristic of the distributions is that there is a gap (RC-MS gap) between the histogram of RC stars and that of MS stars., A notable characteristic of the distributions is that there is a gap (RC-MS gap) between the histogram of RC stars and that of MS stars.546 This gap becomes wider with increasing age., This gap becomes wider with increasing age.547 It nearly disappears when age 7 5 Gyr., It nearly disappears when age $>$ 5 Gyr.548 In spite of there are many stars hotter than the red edge of the instability strip., In spite of there are many stars hotter than the red edge of the instability strip.549 the cluster with age = 0.5 Gyr. the RC-MS gap is located in the range about 6-18 for and about 80-240 forv7.," the cluster with age = 0.5 Gyr, the RC-MS gap is located in the range about 6-18 for and about 80-240 for."5500... For the cluster with age = 1.0 Gyr. Fig.," For the cluster with age = 1.0 Gyr, Fig."551 2. shows that the values of are basically located in the range about 1-10 for RC stars and FGB stars. and in the range about 10-30 for subgiants. while the values are mostly larger than 20 for MS stars.a," \ref{fig2} shows that the values of are basically located in the range about 1-10 for RC stars and FGB stars, and in the range about 10-30 for subgiants, while the values are mostly larger than 20 for MS stars.,"552ddition.. the number of subgiants is very small in this cluster. which leads to the appearance of the gap between about 10 and 20//Hz.," the number of subgiants is very small in this cluster, which leads to the appearance of the gap between about 10 and 20."553. For the cluster with age = 5.0 Gyr. Fig.," For the cluster with age = 5.0 Gyr, Fig."554 3 shows that the values of are almost uniformly distributed in the range about 1-34 for FGB stars and about 34-50 for subgiants. and that the values are mostly larger than 50 for MS stars.," \ref{fig3} shows that the values of are almost uniformly distributed in the range about 1-34 for FGB stars and about 34-50 for subgiants, and that the values are mostly larger than 50 for MS stars."555 However. the of RC stars is almost concentrated in à narrow range of 1-6ΗΣ.," However, the of RC stars is almost concentrated in a narrow range of 1-6."556. Thus. although the number of RC stars is less than that of FGB stars and that of subgiants (see Fig. 43. ," Thus, although the number of RC stars is less than that of FGB stars and that of subgiants (see Fig. \ref{fig4}) ),"557the RC-MS gap still exists in the distribution ofAv., the RC-MS gap still exists in the distribution of.558.. becomes wider with that in the cluster with age = 1.0 Gyr., becomes wider with that in the cluster with age = 1.0 Gyr.559 However. Fig.," However, Fig."560 | shows that the RC-MS gap of clusters with age « 2.0 Gyr can be affected by the stars hotter than the red edge of the instability strip., \ref{fig1} shows that the RC-MS gap of clusters with age $<$ 2.0 Gyr can be affected by the stars hotter than the red edge of the instability strip.561 If these stars are discarded. the RC-MS gap should be enlarged.," If these stars are discarded, the RC-MS gap should be ."562. But for the clusters with age = 2.0 Gyr. the distributions Of and are not affected by these stars.," But for the clusters with age $\gtrsim$ 2.0 Gyr, the distributions of and are not affected by these stars."563 The stars RC-MS gap are primarily subgiants for young clusters. but the stars are mainly subgiants and FGB stars. of RC stars and of MS stars for ‘middle-age’ clusters.," The stars RC-MS gap are primarily subgiants for young clusters, but the stars are mainly subgiants and FGB stars, of RC stars and of MS stars for `middle-age' clusters."564 The number of subgiants and FGB stars increases with age. which results in with increasing age.," The number of subgiants and FGB stars increases with age, which results in with increasing age."565 The mass of stars in old cluster is less than that of the stars in the same evolutionary stage of cluster. which is partly more signiticant mass-loss for old youngerstars. especially for the RC stars.," The mass of stars in old cluster is less than that of the stars in the same evolutionary stage of younger cluster, which is partly more significant mass-loss for old stars, especially for the RC stars."566 In addition. for stars in the same evolutionary phase the lower the mass. the higher the mean density. ie. the larger theAv.," In addition, for stars in the same evolutionary phase the lower the mass, the higher the mean density, i.e. the larger the."567. Thus he values of of subgiants. FGB stars on the bottom of FGB and the MS stars at the end of MS increase with increasing age.," Thus the values of of subgiants, FGB stars on the bottom of FGB and the MS stars at the end of MS increase with increasing age."568 But the values of of RC stars remain in the range about 1-10(;Hz., But the values of of RC stars remain in the range about 1-10.569. Consequently. the RC-MS gap becomes wider with increasing age.," Consequently, the RC-MS gap becomes wider with increasing age."570 Toreover. Fig.," Moreover, Fig."571 3 reveals clearly that another gap (MS gap) exists in the histogram of MS stars., \ref{fig3} reveals clearly that another gap (MS gap) exists in the histogram of MS stars.572 For the cluster with age = 5 Gyr. this gap is located in the range ~ 53-65 for and ~ 920-1200 for£s.," For the cluster with age = 5 Gyr, this gap is located in the range $\sim$ 53-65 for and $\sim$ 920-1200 for."573 Our calculations show that the location of this MS gap changes with age and that it mainly appears in clusters with age = 45 Gyr., Our calculations show that the location of this MS gap changes with age and that it mainly appears in clusters with age $\gtrsim$ 4.5 Gyr.574 This MS gap to the hydrogen-exhausted phase gap. or MS hook (Hurley.Pols&Tout2000).," This MS gap to the hydrogen-exhausted phase gap, or MS hook \citep{hurl00}."575. For stars with AZ— I.I]AL... their core is convective.," For stars with $M >$ 1.1, their core is convective."576 Owing to mixing in the core there is a sudden depletion of fuel over a large region. which leads to a rapid contraction of the core and expansion of radius on a thermal time-scale (Hurley.Pols&Tout2000).," Owing to mixing in the core there is a sudden depletion of fuel over a large region, which leads to a rapid contraction of the core and expansion of radius on a thermal time-scale \citep{hurl00}."577 This causes the appearance of MS hook in the Hertzsprung-Russell diagram and the MS gap in the distributions of andAv., This causes the appearance of MS hook in the Hertzsprung-Russell diagram and the MS gap in the distributions of and.578. There are some sharp edges in the distributions of and in Figs., There are some sharp edges in the distributions of and in Figs.579 | and 2.., \ref{fig1} and \ref{fig2}.580 For example. the edge between the histogram of of MS stars and that of subgiants in the cluster with age = 3.0 Gyr in Fig.," For example, the edge between the histogram of of MS stars and that of subgiants in the cluster with age = 3.0 Gyr in Fig."581 |. and the edge of distribution of RC stars in Fig. 2.., \ref{fig1} and the edge of distribution of RC stars in Fig. \ref{fig2}.582 For the cluster with Z = 0.02 and age = 3.0 Gyr. stars with AJ< 1.45 are MS stars. while most stars with A.o> [1.45 have evolved into MS hook or a later stage.," For the cluster with Z = 0.02 and age = 3.0 Gyr, stars with $M <$ 1.45 are MS stars, while most stars with $M >$ 1.45 have evolved into MS hook or a later stage."583 For stars with AJ 1.45AL... the time spentin MS hook is a thermal time-scale. which is far less than the time of hydrogen burning.," For stars with $M>$ 1.45, the time spentin MS hook is a thermal time-scale, which is far less than the time of hydrogen burning."584 Thus the number of, Thus the number of585referred to as the prior predictive probability for 2. or the elobal likelihood for the entire class of hypotheses.,"referred to as the prior predictive probability for $D$, or the global likelihood for the entire class of hypotheses."586 In the Bavesian framework we start. from. a. prior knowledge we introduce into. the prior. probability distribution. p(44o|/).," In the Bayesian framework we start from a prior knowledge we introduce into the prior probability distribution, $p(H_0|I)$."587 The choice of prior distribution can alfect the posterior distribution. especially if our observed data do not strongly constrain the model parameters.," The choice of prior distribution can affect the posterior distribution, especially if our observed data do not strongly constrain the model parameters."588 Hour prior knowledge is poor. p(Ho|/) can spread over a wide range of possible values for the model parameters.," If our prior knowledge is poor, $p(H_0|I)$ can spread over a wide range of possible values for the model parameters."589 Whenever new cata are available. it is possible to incorporate the new data in our model through the likelihood function. combined with the prior. to obtain a new posterior density probability. p(Ho|D.1). for the parameter.," Whenever new data are available, it is possible to incorporate the new data in our model through the likelihood function, combined with the prior, to obtain a new posterior density probability, $p(H_0|D_1,I)$, for the parameter."590 As soon as we obtain another set of data. Des. we recalculate the posterior density probability in order for it to rellect our new state of knowledge.," As soon as we obtain another set of data, $D_2$, we recalculate the posterior density probability in order for it to reflect our new state of knowledge."591 “Phe possibility to combine new cata sets into the original data we have will allow us to accomplish our goal of detecting transiting exoplanets using scheduled. follow-up observations. Gregory.(2005a)..," The possibility to combine new data sets into the original data we have will allow us to accomplish our goal of detecting transiting exoplanets using scheduled follow-up observations. \cite{2005ApJ...631.1198G},"592 Ford(2006). and others have already shown that Bayesian inference is a useful tool for analyzing precise radial velocity (RV). data of. planet-hosting stars., \cite{2006ApJ...642..505F} and others have already shown that Bayesian inference is a useful tool for analyzing precise radial velocity (RV) data of planet-hosting stars.593 Gregory(2007) used Dayesian inference model selection for the problem of multiple planets. and Gregory(2005a) used it to Construct posterior probability density functions of the lieht-curve parameters.," \cite{2007MNRAS.381.1607G} used Bayesian inference model selection for the problem of multiple planets, and \cite{2005ApJ...631.1198G}594 used it to construct posterior probability density functions of the light-curve parameters."595 Defayctal.(2001). and Xigrain&Favata(2002). demonstrated the use of the Davesian approach to studs planetary transits., \cite{2001A&A...365..330D} and \cite{2002A&A...395..625A} demonstrated the use of the Bayesian approach to study planetary transits.596 Our implementation of Bavesian inference is based on the Metropolis-Hastings (MEI) algorithm. which is à version of the more general Markov-Clhain. Monte Carlo (AICAIC) approach (Ciregory 2005h).. ," Our implementation of Bayesian inference is based on the Metropolis-Hastings (MH) algorithm, which is a version of the more general Markov-Chain Monte Carlo (MCMC) approach \citep{2005blda.book.....G}. ."597Ao Markov chain is calculated. using an initial set of parameter values. Yo. and a transition probability. pCXsa[Xu4). that deseribes the probability of moving from the current state to the next. once.," A Markov chain is calculated using an initial set of parameter values, $\overline{X}_0$, and a transition probability, $p(\overline{X}_{n+1}|\overline{X}_n,I)$, that describes the probability of moving from the current state to the next one."598" The transition probability depends on the acceptance probability. described later in Section ??.. and if properly constructed. then after excluding the so-called. ""burning time”. we can use the chain as à sample from the desired distribution."," The transition probability depends on the acceptance probability, described later in Section \ref{MCMC_approach}, and if properly constructed, then after excluding the so-called “burning time”, we can use the chain as a sample from the desired distribution."599 ALLL Algorithm is an implementation of the MCMCC that is used for obtaining à sequence of random samples from a probability cstribution., MH Algorithm is an implementation of the MCMC that is used for obtaining a sequence of random samples from a probability distribution.600 The ME algorithm does not require good initial guess of the parameters values in order to estimate the posterior distribution., The MH algorithm does not require good initial guess of the parameters values in order to estimate the posterior distribution.601 This is one of the most important advantages of the algorithm., This is one of the most important advantages of the algorithm.602 The algorithm is capable of exploring all regions of the parameter. space having significant probabilities (assuming it meets several basic requirements)., The algorithm is capable of exploring all regions of the parameter space having significant probabilities (assuming it meets several basic requirements).603 The analysis also vields the marginal posterior. probability distribution functions for each of the model parameters. and their uncertainties.," The analysis also yields the marginal posterior probability distribution functions for each of the model parameters, and their uncertainties."604 In Section ?? we describe the follow-up approach we developed. to. detect transiting exoplancts. based on observations [rom low cadence surveys.," In Section \ref{directedFW} we describe the follow-up approach we developed to detect transiting exoplanets, based on observations from low cadence surveys."605 In Section ?? we eive a brief review of Bayesian inference and its applications for our follow-up strategy., In Section \ref{MCMC_approach} we give a brief review of Bayesian inference and its applications for our follow-up strategy.606 Sections ?? and 77. demonstrate the approach by applying it on two stars that are known to harbor hot-Jupiters. LID 200458 and HD. 189733. using the data base.," Sections \ref{HD209458_5pars} and \ref{HD189733_5pars} demonstrate the approach by applying it on two stars that are known to harbor hot-Jupiters, HD 209458 and HD 189733, using the data base."607 Section 77. shows some “sanity checks? we performed on data that do not contain the transit signal at all., Section \ref{sanity_check} shows some “sanity checks“ we performed on data that do not contain the transit signal at all.608 We conclude and describe future applications of the strategv in Section ?7.., We conclude and describe future applications of the strategy in Section \ref{cuncluding}.609 Our ultimate goal is to detect transiting exoplanets using follow-up observations. which will be carefully scheduled to increase the chances to capture transits. should hey exist.," Our ultimate goal is to detect transiting exoplanets using follow-up observations, which will be carefully scheduled to increase the chances to capture transits, should they exist."610 Thus. our approach does not focus on obtaining a detailed. transit. model that best fits the available data. oit. on building a probability distribution function of the xwameters of a simple model. based. on these data.," Thus, our approach does not focus on obtaining a detailed transit model that best fits the available data, but on building a probability distribution function of the parameters of a simple model, based on these data."611 The ransit model that we use in this work is à very simplistic one. based on the BLS philosophy (Ixovácsetal.2002)..," The transit model that we use in this work is a very simplistic one, based on the BLS philosophy \citep{2002A&A...391..369K}."612 Thus. we model a transit’ light. curve as à. box-shaped ransit with two phases. in and out of transit. ancl ignore he duration of the ingress and egress phases. as well as he details of the limb darkening.," Thus, we model a transit light curve as a box-shaped transit with two phases, in and out of transit, and ignore the duration of the ingress and egress phases, as well as the details of the limb darkening."613 These details are less relevant in low-precision. low-cadence surveys. and using ewer parameters makes the model more robust.," These details are less relevant in low-precision, low-cadence surveys, and using fewer parameters makes the model more robust."614 We use the Bavesian MAIL algorithm to obtain a posterior probability distribution for the mocdel parameters. and then use this distribution to prioritize the timing of the observations of he chosen stars for follow-up observations.," We use the Bayesian MH algorithm to obtain a posterior probability distribution for the model parameters, and then use this distribution to prioritize the timing of the observations of the chosen stars for follow-up observations."615 The directed. follow-up approach is not suitable for space missions like or., The directed follow-up approach is not suitable for space missions like or.616fNepler.. Such missions. due o their high cadence. will not benefit from the approach since their phase and. period. coverage are already quite complete.," Such missions, due to their high cadence, will not benefit from the approach since their phase and period coverage are already quite complete."617 Instead. we aim for all-sky surveys like (vanLecuwenοἱal.1997)... or its successor. (Jordi€.etal.2006)... in order to use their extensive low-cadence photometric databases for exoplanets search.," Instead, we aim for all-sky surveys like \citep{1997A&A...323L..61V}, or its successor, \citep{2006MNRAS.367..290J}, in order to use their extensive low-cadence photometric databases for exoplanets search."618 A simplified (BLS-like) transit light curve is pavametrizecl by five quantities. e.g. the period. phase. and width of the transit. and the Lux levels in-transit. ancl ex-transit.," A simplified (BLS-like) transit light curve is parametrized by five quantities, e.g., the period, phase, and width of the transit, and the flux levels in-transit and ex-transit."619 The first step in our proposed. procedure Is to apply the MIT algorithm to the measurements of a target star., The first step in our proposed procedure is to apply the MH algorithm to the measurements of a target star.620 This results in live Markov chains that include the successful iterations for each one of the parameters., This results in five Markov chains that include the successful iterations for each one of the parameters.621 A successful iteration is one that was accepted by the MIL algorithm., A successful iteration is one that was accepted by the MH algorithm.622" After removing the ""burning time”. each. chain represents the stationary distribution of the parameters. which we use as their estimated. Bayesian posterior distributions. for our current state of knowledge (Gregory 2005h).."," After removing the “burning time”, each chain represents the stationary distribution of the parameters, which we use as their estimated Bayesian posterior distributions, for our current state of knowledge \citep{2005blda.book.....G}."623 Unlike the case of precise high-cadence surveys. even if the star does host a transiting exoplanet. due to the low precision ancl low cadence of the observations we clo not expect the distribution to concentrate around a single solution. but rather show cillerent periods that might fit the data. besides the unknown correct one.," Unlike the case of precise high-cadence surveys, even if the star does host a transiting exoplanet, due to the low precision and low cadence of the observations we do not expect the distribution to concentrate around a single solution, but rather show different periods that might fit the data, besides the unknown correct one."624 The next step of our procedure is to assign cach point in time the probability that a transit will occur at that time., The next step of our procedure is to assign each point in time the probability that a transit will occur at that time.625 Calculating this probability is easy. using the posterior distributions we found in the first stage., Calculating this probability is easy using the posterior distributions we found in the first stage.626 Basically for time /. we count the number of MOM successful iterations whose values of P. T; and w predict a transit in time /.," Basically for time $t$, we count the number of MCMC successful iterations whose values of $P$, $T_c$ and $w$ predict a transit in time $t$."627 Normalizing this number by the number of total iterations vield the lnstantaneous Transit Probability (EPI?) for time /., Normalizing this number by the number of total iterations yield the Instantaneous Transit Probability (ITP) for time $t$.628 Lf the ITP has signilicantly high. values for certain times. then a follow-up observation is worthwhile at those preferred times.," If the ITP has significantly high values for certain times, then a follow-up observation is worthwhile at those preferred times."629 When we examine the PPP function. of cillerent observationsand simulations we performed. it is clear that the shape of this function when a transit signal exists contains sharp peaks. where the probability of transit. is," When we examine the ITP function of different observationsand simulations we performed, it is clear that the shape of this function when a transit signal exists contains sharp peaks, where the probability of transit is"630The study of galaxy. evolution is filled with negative results.,The study of galaxy evolution is filled with negative results.631" Despite intensive efforts. only stnall or subtle evolution lias been detected in the uumber density of field galaxies (eg.??22?7?).. iu their masses (eg.22???).. or iu their clustering relative to ""stable clusteriug (eg.2???7?).. from redshift unity to the present day."," Despite intensive efforts, only small or subtle evolution has been detected in the number density of field galaxies \citep[eg,][]{lilly95a,heyl97,hogg98thesis,lin99a,cohen02a}, in their masses \citep[eg,][]{vogt96,vogt97,treu99,brinchmann00,cohen02a}, or in their clustering relative to “stable clustering” \citep[eg,][]{lefevre96a,carlberg97a,small99,hogg00b,carlberg00clust}, from redshift unity to the present day."632 At the same time. the constraints on galaxy. evolution have not been mace stroug enough to allow definitive results from the classical cosmological tests.," At the same time, the constraints on galaxy evolution have not been made strong enough to allow definitive results from the classical cosmological tests."633 Has the ealaxy evolution couumuuity got auything positive to sav?, Has the galaxy evolution community got anything positive to say?634 Indeed it has., Indeed it has.635 For a loug time it has been observed that apparently faint aud (at oue time presumed to be. now largelve known to be) distant egalaxies are. on average.e bluer in color than their local counterparts," For a long time it has been observed that apparently faint and (at one time presumed to be, now largely known to be) distant galaxies are, on average, bluer in color than their local counterparts"636presumably related to their not. having been on the horizontal branch or upper part of the normal branch in the X-ray HID - see section 3.2 below.,presumably related to their not having been on the horizontal branch or upper part of the normal branch in the X-ray HID - see section 3.2 below.637 The upper limits to the linear polarisation of the radio emission [rom GX 17|2 and GX 5-1. of order (Lable 1) clo not seriously constrain he emissive mechanism or optical depth of the ejecta.," The upper limits to the linear polarisation of the radio emission from GX 17+2 and GX 5-1, of order (Table 1) do not seriously constrain the emissive mechanism or optical depth of the ejecta."638 While he brightest racio transients may show very high linear polarisation (e.g. Fender ct al., While the brightest radio transients may show very high linear polarisation (e.g. Fender et al.639 19992: ILjellming et al., 1999a; Hjellming et al.640 1999). the total radio Hux (ie. not resolving individual components) from. X-ray. binaries is generally polarised at less than the level.," 1999), the total radio flux (i.e. not resolving individual components) from X-ray binaries is generally polarised at less than the level."641 Exclucling 28 0021-630 (see discussion. below). we have observed several atol-ivpe X-ray binaries ancl not detected any ο ‘them.," Excluding 2S 0921-630 (see discussion below), we have observed several atoll-type X-ray binaries and not detected any of them."642 Significantly. this eroup includes GX 9|1 and GX 9|9. amongst the most luminous (in X-rays) of the class.," Significantly, this group includes GX 9+1 and GX 9+9, amongst the most luminous (in X-rays) of the class."643 In particular we do detect the atoll source iU 1820-30 in he globular cluster NGC 6624., In particular we do detect the atoll source 4U 1820-30 in the globular cluster NGC 6624.644" All the limits presented. in Ἱ""ble 1 are either improveniens on previous limits (typicalv by a factor of ~ 2) or the first. limits presented. for he sources.", All the limits presented in Table 1 are either improvemens on previous limits (typically by a factor of $\sim 2$ ) or the first limits presented for the sources.645 These upper limits are further evidence that he atoll sources as a class are significantly fainter at racic» (as well as N-rav) wavelengths than the Z πουσος. à poln explored further in Fender Llencry (2000).," These upper limits are further evidence that the atoll sources as a class are significantly fainter at radio (as well as X-ray) wavelengths than the Z sources, a point explored further in Fender Hendry (2000)."646 While originally classified as an atoll source. the bright," While originally classified as an atoll source, the bright"647 also show indications for late time activity of such an engine 2006a)., also show indications for late time activity of such an engine \citep{Soderberg06}.648. Due to their low luninositv. are detected only from low redshilts (2S0.1).," Due to their low luminosity, are detected only from low redshifts $z\lesssim0.1$ )."649 With these redshifts. the four observed imply an event rate of 230.Guetla155 vr (Soderbergetal.2006a.seealsoCowardnan2005:Cobb2006:Liang2007: 2011).. about. 100-1000. (mes higher than the rate of LGRDs pointing toward earth mum2005:Liangetal.2007:Guetta&DellaValleWanclermanPira," With these redshifts, the four observed imply an event rate of $230^{+490}_{-190}$ $^{-3}$ $^{-1}$ \citep[][see also Coward 2005; Cobb 2006; Liang650et al. 2007; Guetta \& Della Valle 2007; Fan et al.6512011]{Soderberg06}, about 100-1000 times higher than the rate of LGRBs pointing toward earth \citep[][]{Coward05,Liang07,Guetta07,Wanderman10}."652n 2010).. Soderbergetal.(20064) -- ihe rate of broad line Ibe SNe to be of the same order as (he rate of/-GRBs. implving that cannot be significantly beamed and that they could very. well be isotropic.," \citet{Soderberg06} estimated the rate of broad line Ibc SNe to be of the same order as the rate of, implying that cannot be significantly beamed and that they could very well be isotropic."653 Using (he overall ratio of the rates of broad line (wpe Ib.c SNe and we find that the beaming factor of is <10. corresponding to opening angles =30.," Using the overall ratio of the rates of broad line type Ib,c SNe and we find that the beaming factor of is $\lesssim 10$, corresponding to opening angles $\gtrsim30^\circ$."654 The lack of bright. late time. radio emission from stronely constrain the total energv of any relativistic outflow volved in these events (Waxman2004:Soderberg.Frail.&Wieringa2004:Soderbergetal. 2006b).," The lack of bright, late time, radio emission from strongly constrain the total energy of any relativistic outflow involved in these events \citep{Waxman04,Soderberg04a,Soderberg06b}."655". Additionally. statistical arguments rule out the possibility ""i are regular LGRBDs viewed at a large angle (e.g. 2007).", Additionally statistical arguments rule out the possibility that are regular LGRBs viewed at a large angle \citep[e.g.][]{Daigne07}.656. -Thus. if are generated by relativistic jets these jets must be weak and have a opening angle.," Thus, if are generated by relativistic jets these jets must be weak and have a large opening angle."657 We review. briefly. the essential features of jet propagation in a stellar envelope (B11).," We review, briefly, the essential features of jet propagation in a stellar envelope (B11)."658" Consider a cold relativistic jet with a power £; and an initial opening angle 0, that is injected into à star.", Consider a cold relativistic jet with a power $L_j$ and an initial opening angle $\theta_0$ that is injected into a star.659 As the jet propagates it pushes the stellar material in [ront of it. leading to the formation of a double shock structure al the jets front. the jets head.," As the jet propagates it pushes the stellar material in front of it, leading to the formation of a double shock structure at the jet's front, the jet's head."660 The pressure of the shocked material is much higher than (he pressure of the surrounding medium. (hus matter that enters (he head is heated ancl pushed sidewavs forming a pressurized cocoon surrounding the jet.," The pressure of the shocked material is much higher than the pressure of the surrounding medium, thus matter that enters the head is heated and pushed sideways forming a pressurized cocoon surrounding the jet."661 The cocoon. in turus. applies a pressure on (he jet and if the jet power is not too large it collimates the jet into a cvlindrical shape.," The cocoon, in turns, applies a pressure on the jet and if the jet power is not too large it collimates the jet into a cylindrical shape."662" The material in the collimated jet remains relativistic and its Lorentz factor is D;&6,!.", The material in the collimated jet remains relativistic and its Lorentz factor is $\Gamma_j\simeq\theta_0^{-1}$.663 The jets head propagates. however. at a nich lower velocity and it effectively clissipates all the jets energy into the cocoon.," The jet's head propagates, however, at a much lower velocity and it effectively dissipates all the jet's energy into the cocoon."664 Thus. in order for the jet to breakout. (he engine must operate continually and supply power to (he jet practically until the jets head reaches the surface. al which stage the dissipation stops.," Thus, in order for the jet to breakout, the engine must operate continually and supply power to the jet practically until the jet's head reaches the surface, at which stage the dissipation stops."665 Thejet propagation depends on the stellar density profile., The jet propagation depends on the stellar density profile.666 Above the stellar core through a considerable fraction of the envelope. where the jet spends most of its propagation time.," Above the stellar core through a considerable fraction of the envelope, where the jet spends most of its propagation time,"667disc.,disc.668 The minimum particle size is 1 km and there are no planetesimals at {=0., The minimum particle size is 1 km and there are no planetesimals at $t=0$.669 Binary parameters are those of the yCephei system., Binary parameters are those of the $\gamceph$ system.670" Figure 2. shows the resulting (a,e) distribution after 1000 binary orbits (~90000 yr)."," Figure \ref{figae} shows the resulting $(a,e)$ distribution after 1000 binary orbits $\sim 90000$ yr)."671" It is clear that significant accretion has taken place in the inner parts of the disc, where particles have grown from 1 km to 50 km."," It is clear that significant accretion has taken place in the inner parts of the disc, where particles have grown from 1 km to 50 km."672" Outside a/ay=0.1, which corresponds to 2 AU, perturbations due to the binary are too strong for accretion to occur."," Outside $a/\ab=0.1$, which corresponds to $2$ AU, perturbations due to the binary are too strong for accretion to occur."673" Inside 2 AU, planetesimals grow to sizes up to 70 km."," Inside $2$ AU, planetesimals grow to sizes up to $70$ km."674 The top panel of Fig., The top panel of Fig.675 3 shows the time evolution of the maximum and mean planetesimal size over the whole disc., \ref{figsizemass} shows the time evolution of the maximum and mean planetesimal size over the whole disc.676" 'The mean size is dominated by the large number of small planetesimals in the outer disc, and stays between 10 and 20 km."," The mean size is dominated by the large number of small planetesimals in the outer disc, and stays between $10$ and $20$ km."677 The maximum particle size goes up rapidly by dust accretion (smooth parts of the curve) and by accreting collisions (jumps)., The maximum particle size goes up rapidly by dust accretion (smooth parts of the curve) and by accreting collisions (jumps).678" The largest planetesimal is destroyed a few times as well, but in general the trend is to grow to larger sizes."," The largest planetesimal is destroyed a few times as well, but in general the trend is to grow to larger sizes."679" After ~600 binary orbits, there is no source of small dust remaining from which to create new planetesimals, and the collision time scale goes up."," After $\sim 600$ binary orbits, there is no source of small dust remaining from which to create new planetesimals, and the collision time scale goes up."680 The maximum size that can be reached in this scenario is limited by the total amount of solid material present in the disc., The maximum size that can be reached in this scenario is limited by the total amount of solid material present in the disc.681" In the 2D approximation, this mass has to be artificially low in order to end up with a realistic collision time scale."," In the 2D approximation, this mass has to be artificially low in order to end up with a realistic collision time scale."682" For a disc that has twice the solid material, keeping all other parameters the same, growth up to 100 km was observed."," For a disc that has twice the solid material, keeping all other parameters the same, growth up to 100 km was observed."683" Increasing N by a factor of 2 while decreasing Ring by the same factor, which amounts to increasing the resolution of the simulation, showed growth up to 150 km."," Increasing $N$ by a factor of 2 while decreasing $\rinf$ by the same factor, which amounts to increasing the resolution of the simulation, showed growth up to 150 km."684 In these higher resolution runs more accreting collisions are observed than depicted in Fig. 3.., In these higher resolution runs more accreting collisions are observed than depicted in Fig. \ref{figsizemass}.685" This is because the system goes through phases of low particle number density, for which the collision statistics in the lower resolution runs are not optimal."," This is because the system goes through phases of low particle number density, for which the collision statistics in the lower resolution runs are not optimal."686" In this sense, Fig."," In this sense, Fig."687" 3 represents a worst-case scenario, where many low-velocity collisions are missed, and adding particles, while keeping the collision time scale the same, will only favour planetesimal growth more."," \ref{figsizemass} represents a worst-case scenario, where many low-velocity collisions are missed, and adding particles, while keeping the collision time scale the same, will only favour planetesimal growth more."688 Gas drag appears to play only a minor role in determining the qualitative outcome of the model., Gas drag appears to play only a minor role in determining the qualitative outcome of the model.689 This is mainly due to the fact that the planetesimals are weak enough so that any small eccentricity difference of ~0.01 will lead to destructive collisions., This is mainly due to the fact that the planetesimals are weak enough so that any small eccentricity difference of $\sim 0.01$ will lead to destructive collisions.690 Whether this difference is due to differential orbital phasing (under the influence of gas drag) or simply orbital crossing (in the absence of gas) does not matter., Whether this difference is due to differential orbital phasing (under the influence of gas drag) or simply orbital crossing (in the absence of gas) does not matter.691" A simulation without any gas showed the same trend as in Fig. 3,,"," A simulation without any gas showed the same trend as in Fig. \ref{figsizemass},"692 with growth up to 80 km., with growth up to 80 km.693 Crucial parameters are the efficiency of planetesimal formation and dust accretion., Crucial parameters are the efficiency of planetesimal formation and dust accretion.694 Accretion as shown in Fig., Accretion as shown in Fig.695 3 can only occur if the small debris created in catastrophic collisions is swept up by larger bodies rather than forming new planetesimals., \ref{figsizemass} can only occur if the small debris created in catastrophic collisions is swept up by larger bodies rather than forming new planetesimals.696" Increasing ερ by a factor of 10 still results in accretion up to 80 km, but for a factor 100 growth stalls at 10 km."," Increasing $\epsilon_\mathrm{p}$ by a factor of 10 still results in accretion up to 80 km, but for a factor 100 growth stalls at 10 km."697" Similarly, the resultdepicted in Fig."," Similarly, the resultdepicted in Fig."698 3 is robust to changes in ea up to a factor of 10., \ref{figsizemass} is robust to changes in $\epsilon_\mathrm{d}$ up to a factor of 10.699 We consider the aCentauri system in Figs., We consider the $\alphacen$ system in Figs.700 4 and 5.., \ref{figaealphacen} and \ref{figsizemassalphacen}. .701rolational velocities comparable to. DIID stars. when allowance is made for their different radii?,"rotational velocities comparable to BHB stars, when allowance is made for their different radii?"702 C2003 suggested four follow-up studies., C2003 suggested four follow-up studies.703 First. expancl the sample (ο ascertain il our original saniple was unusual in some manner.," First, expand the sample to ascertain if our original sample was unusual in some manner."704" This paper reports on the results of a ""hast study of £5 additional metal-poor fieldstars?."," This paper reports on the results of a “hasty"" study of 45 additional metal-poor field."705.. We also sought to obtain line broadening measures for giants in globular clusters. and initial data for [four clusters are in hand.," We also sought to obtain line broadening measures for giants in globular clusters, and initial data for four clusters are in hand."706 Results will be published later., Results will be published later.707 Η planets do exist around metal-poor stars. (his would suggest that disk instability is a viable mechanism lor planet formation. ancl a high-precision racial velocily survey of roughly (wo hundred metal-poor field stars was begun (see Sozzelli οἱ 22006).," If planets do exist around metal-poor stars, this would suggest that disk instability is a viable mechanism for planet formation, and a high-precision radial velocity survey of roughly two hundred metal-poor field stars was begun (see Sozzetti et 2006)."708 Finally. we have (ο ask if the line broadening is due (ο rotation or to macroturbulence.," Finally, we have to ask if the line broadening is due to rotation or to macroturbulence."709 This can be determined with verv high-resolution. hieh-S/N spectra that. enable Fourier transform studies of line profiles. Following Grav (1982: 1984) and Gray Toner (1956: 1987).," This can be determined with very high-resolution, high-S/N spectra that enable Fourier transform studies of line profiles, following Gray (1982; 1984) and Gray Toner (1986; 1987)."710 We have completed. acquisition aud analvsis of such spectra and will report on the results in a future paper (Carnev οἱ 22007)., We have completed acquisition and analysis of such spectra and will report on the results in a future paper (Carney et 2007).711 C2003 discussed in detail the criteria by which thev assembled their list of RGB and RIB field stars lor study., C2003 discussed in detail the criteria by which they assembled their list of RGB and RHB field stars for study.712 These followed mostly the earlier study of Carney Latham (1986). which in turn relied on the kinematically-ubiased metal-poor (ΤΟ € -1.5)samplesidenti fiedbyBond( 19}," These followed mostly the earlier study of Carney Latham (1986), which in turn relied on the kinematically-unbiased metal-poor ([Fe/H] $\leq$ $-1.5$ ) samples identified by Bond (1980)."713 linedslarsusingobjectiveprismspectroscopyas wellsfollow—upuvbyphotometry., He identifield weak-lined stars using objective prism spectroscopy as well as follow-up $uvby$ photometry.714 Norris. Bessell. &Pichl photometry.," Norris, Bessell, Pickles (1985) undertook a large program of additional DDO and $RI_{\rm C}$ photometry."715 Most of the stars in (he C2003 sample had been classified as metal-poor red eiants by Anthony-Twarog Twarog (1994: herealter ATT). who supplied estimated reddening and Ady values.," Most of the stars in the C2003 sample had been classified as metal-poor red giants by Anthony-Twarog Twarog (1994; hereafter ATT), who supplied estimated reddening and $M_{\rm V}$ values."716 All but one of the 45 stars cliseussecl in this paper come from ATT. and most of them were identified originally by Bond (1980).," All but one of the 45 stars discussed in this paper come from ATT, and most of them were identified originally by Bond (1980)."717 The stars in (his study. represent most of the stars in ATT with [Fe/H] <—1.5 not studied by C2003 but within reach of the telescopes emploved in that program., The stars in this study represent most of the stars in ATT with [Fe/H] $\leq\ -1.5$ not studied by C2003 but within reach of the telescopes employed in that program.718 To make certain that the stars are all luminous. we have emploved the Hipparcos database. which includes 44 of (he stars (only BD+25 3410 lacks a measured (rigonomelric parallax).," To make certain that the stars are all luminous, we have employed the Hipparcos database, which includes 44 of the stars (only BD+25 3410 lacks a measured trigonometric parallax)."719" The parallaxes are all very. small ancl consistent with high luminosities,", The parallaxes are all very small and consistent with high luminosities.720 , 721using well established galaxy evolution mocels (Moetcalfe οἱ al..,"using well established galaxy evolution models (Metcalfe et al.,"722 2001. 2006).," 2001, 2006)."723 We find that of these four galaxy populations. the ο νο galaxies eross-correlate most strongly with ssources and also have the brightest stacked. sub-nim flux.," We find that of these four galaxy populations, the $z$ red galaxies cross-correlate most strongly with sources and also have the brightest stacked sub-mm flux."724 The left panel of reff-optstacks— shows that the high-z red galaxies are strongly detected at aat 5.9o significance with 5u;y= 142-2440., The left panel of \\ref{f-optstacks} shows that the $z$ red galaxies are strongly detected at at $5.9\sigma$ significance with $S_{870}=142\pm24 \umu$ Jy.725 The Low-z red ealaxies also show a significant detection at 4.7e significance while the blue galaxy. populations are more weakly detected al ~2c significance (see Table 2))., The $z$ red galaxies also show a significant detection at $4.7\sigma$ significance while the blue galaxy populations are more weakly detected at $\sim2\sigma$ significance (see Table \ref{t-seds}) ).726 reff-optstacks. and Table 2. also show the stacking analyses al aand24pum., \\ref{f-optstacks} and Table \ref{t-seds} also show the stacking analyses at and.727. In. both of these wavebands. the high-z: red ealaxies are the only sources to show any significant. lux.," In both of these wavebands, the $z$ red galaxies are the only sources to show any significant flux."728 ‘Taken together. the analyses shown in reff-optstacks| suggest that the zz0.5 red. galaxies emit most strongly at long Li wavelengths.," Taken together, the analyses shown in \\ref{f-optstacks} suggest that the $z>0.5$ red galaxies emit most strongly at long IR wavelengths."729 Almaini ct ((2005) suggested that the explanation for a cross-correlation between bright sub-nim sources and Hox 23.2%0.5 galaxies could be gravitational lensing.," Almaini et (2005) suggested that the explanation for a cross-correlation between bright sub-mm sources and $R<23$, $z\approx0.5$ galaxies could be gravitational lensing."730 ut where Almaini et [found a cross-correlation between Sus210 mdyw sub-mnm sources and galaxies and Little cross-correlation for fainter sub-mun sources. we find that the cross-correlation with high-z red. galaxies is dominatecl bv Suo10 mJy sub-nin sources. so any such lensing signature is absent [rom our data.," But where Almaini et found a cross-correlation between $S_{850}>10$ mJy sub-mm sources and galaxies and little cross-correlation for fainter sub-mm sources, we find that the cross-correlation with $z$ red galaxies is dominated by $S_{850}<10$ mJy sub-mm sources, so any such lensing signature is absent from our data."731 We have compared the optical (and NIR see Lill. 2010) colours of the high-z red galaxies to those of the X-ray sources and find that the faint X-ray sources and the high-z red galaxies tend to occupy the same colour space. while the locus of the bright X-ray sources is elsewhere reff-brixrav)).," We have compared the optical (and NIR – see Hill, 2010) colours of the $z$ red galaxies to those of the X-ray sources and find that the faint X-ray sources and the $z$ red galaxies tend to occupy the same colour space, while the locus of the bright X-ray sources is elsewhere \\ref{f-brixray}) )."732 This raises the possibility that the high-z red ealaxics and the faint X-ray sources. the two object classes seen to exhibit the strongest correlation. with sub-nun sources. both sample the same population.," This raises the possibility that the $z$ red galaxies and the faint X-ray sources, the two object classes seen to exhibit the strongest correlation with sub-mm sources, both sample the same population."733 The faint A-ray sources are expected to include the more absorbed AGN: the question. then. is whether the high-z red galaxies also host. absorbed AGN. specifically Conipton-thick ACN which are X-rav-uncdetectecd.," The faint X-ray sources are expected to include the more absorbed AGN; the question, then, is whether the $z$ red galaxies also host absorbed AGN, specifically Compton-thick AGN which are X-ray-undetected."734 We have attempted to estimate the AGN fraction within the high-z τοῦ population using the LRAC 3.6]4.5]:5.8]8.0] colour-colour plot. which is known to be a robust wav of distinguishing AXGN from other sources (Stern et al..," We have attempted to estimate the AGN fraction within the $z$ red population using the IRAC [3.6]--[4.5]:[5.8]--[8.0] colour-colour plot, which is known to be a robust way of distinguishing AGN from other sources (Stern et al.,"735 2005)., 2005).736 However. LRAC observations are available only for the small. central field (SWIRL survey: Lonsdale et al..," However, IRAC observations are available only for the small, central field (SWIRE survey; Lonsdale et al.,"737 2004) and the rresolution is relatively poor. so only 1% of our X-ray and optical populations are detected in all four URAC bands.," 2004) and the resolution is relatively poor, so only $\sim1$ of our X-ray and optical populations are detected in all four IRAC bands."738 of the faint X-ray sources and of the high-z red galaxies which are detected in all four LRAC channels Lie within or very close to the AGN wedge. but with such small saniples itis hard to draw any firm conclusions here.," of the faint X-ray sources and of the $z$ red galaxies which are detected in all four IRAC channels lie within or very close to the AGN wedge, but with such small samples it is hard to draw any firm conclusions here."739 Looking at longer. mid/far-11t wavelengths. we can use the stacked.244un.. aand fluxes given in Table 2. to characterise the average far-L1 SEDs of the high-z red galaxies and the X-ray sources reff-secds)).," Looking at longer, mid/far-IR wavelengths, we can use the stacked, and fluxes given in Table \ref{t-seds} to characterise the average far-IR SEDs of the $z$ red galaxies and the X-ray sources \\ref{f-seds}) )."740 We find that two temperature components are needed to fit the data lor the X-ray. sources and the high-z red galxics galaxies. with both populations well matched » a model with 71=27 lx. 15=120 lx and a ratio tween the components of {ον=0.3. although this is » no means a unique solution.," We find that two temperature components are needed to fit the data for the X-ray sources and the $z$ red galxies galaxies, with both populations well matched by a model with $T_1=27$ K, $T_2=120$ K and a ratio between the components of $L_2/L_1=0.3$, although this is by no means a unique solution."741 With only three datapoints constraining the fits the derived: temperatures should. not oe taken too seriously. particularly as the 24pun.. aand fluxes παν vet arise. from dillerent populations.," With only three datapoints constraining the fits the derived temperatures should not be taken too seriously, particularly as the , and fluxes may yet arise from different populations."742 evertheless. does serve to illustrate that the. FIIV colours (and. implied dust temperatures) of red galaxies ancl X-ray sources appear broadly consistent.," Nevertheless, \\ref{f-seds} does serve to illustrate that the FIR colours (and implied dust temperatures) of red galaxies and X-ray sources appear broadly consistent."743 We conclude that there is no inconsistency in the optical/NIRΓΗ colours of z>0.5 red. galaxies ancl faint X-ray sources., We conclude that there is no inconsistency in the optical/NIR/FIR colours of $z>0.5$ red galaxies and faint X-ray sources.744 Phese galaxies could therefore host Compton-thick AGN., These galaxies could therefore host Compton-thick AGN.745 We shall see later that the sub-mm background associated with these red galaxies is also consistent. with the predicted contribution from such AGN., We shall see later that the sub-mm background associated with these red galaxies is also consistent with the predicted contribution from such AGN.746 Using the sub-nim stacking analyses presented. for X-rav sources and high-z red galaxies. we can quantify the contribution made by these populations to the sub-nin background. by multiplying by the stacked average Lux by the total number of sources and dividing bv 0.25 deg?.," Using the sub-mm stacking analyses presented for X-ray sources and $z$ red galaxies, we can quantify the contribution made by these populations to the sub-mm background, by multiplying by the stacked average flux by the total number of sources and dividing by 0.25 $^2$."747 For the high-z red galaxies this contribution is 4.84:1.2 Jv 5 while for the X-ray sources it is 1.50.1 Jv ," For the $z$ red galaxies this contribution is $4.8 \pm 1.2$ Jy $^{-2}$, while for the X-ray sources it is $1.5 \pm 0.1$ Jy $^{-2}$."748The total sub-mm background. based on measurements. is estimated to be ~45 Jv = (Puge et aL.," The total sub-mm background, based on measurements, is estimated to be $\sim45$ Jy $^{-2}$ (Puget et al.,"749 1996: Fixsen ct al.," 1996; Fixsen et al.,"750 1998)., 1998).751 Therefore. the value of 1.530.1 Jv 7 from X-ray sources represents only a very small fraction. of the total extragalactic background. ligh (EBL).," Therefore, the value of $1.5 \pm 0.1$ Jy $^{-2}$ from X-ray sources represents only a very small fraction of the total extragalactic background light (EBL)."752 However. the total contribution that could be mace by active galactic nuclei; à significant proportion of which could be missing from the X-ray sample we have used. coul vet be fargreater.," However, the total contribution that could be made by active galactic nuclei, a significant proportion of which could be missing from the X-ray sample we have used, could yet be fargreater."753 In order to quantify this. we make use of a mocdel of obscured AGN.," In order to quantify this, we make use of a model of obscured AGN."754 We have detected. significant sub-mim [lux from X-ray sources which are likely to bea population of obscured AGN., We have detected significant sub-mm flux from X-ray sources which are likely to be a population of obscured AGN.755few molecules remain in the gas.,few molecules remain in the gas.756 Those that are present show a considerable increase in deuteration ratios compared to the input values., Those that are present show a considerable increase in deuteration ratios compared to the input values.757 The effect is most marked [or multiply deuterated molecules., The effect is most marked for multiply deuterated molecules.758" Thus D, /IL, increases from 3.1 x * in the molecular cloud to 25 in the midplane of Models B aud C. due to the high molecular depletions in the disk."," Thus $_3^+$ $_3^+$ increases from 3.1 $\times$ $^{-3}$ in the molecular cloud to 25 in the midplane of Models B and C, due to the high molecular depletions in the disk."759" This compares to ID /IL, in the same models. which increases from an input. value of 0.16 to ~ 0.95."," This compares to $_2$ $^+$ $_3^+$ in the same models, which increases from an input value of 0.16 to $\sim$ 0.95."760" The increase in the midplane D, /IL, ratio is even higher in Models A and D where the depletion is also higher due to the absence of CRII.", The increase in the midplane $_3^+$ $_3^+$ ratio is even higher in Models A and D where the depletion is also higher due to the absence of CRH.761 The radial variation of the ratios of the cobumn densities of deuterated to non.deuterated molecules in Model B (Figure 16)) shows a similar distribution to those of 2.., The radial variation of the ratios of the column densities of deuterated to non–deuterated molecules in Model B (Figure \ref{fig:comp_dh}) ) shows a similar distribution to those of \citet{ah01}.762" There is a decrease in D/II with decreasing radius for molecules such as δω. ΠΟ and whose deuteration depends on the isotopomers of IL,."," There is a decrease in D/H with decreasing radius for molecules such as $_3$, $_2$ O and $^+$ whose deuteration depends on the isotopomers of $_3^+$."763" The deuteration of IL, decreases towards the star because of the increase in temperature.", The deuteration of $_3^+$ decreases towards the star because of the increase in temperature.764" Molecules such as IeCO and CHI, whose deuteration depends on CIL, are less affected by the temperature increase. since the barrier to the reverse reaction of al 310 IN is higher than that for sD + Ib — IL, + HD (220 IN)."," Molecules such as $_2$ CO and $_4$ whose deuteration depends on $_3^+$ are less affected by the temperature increase, since the barrier to the reverse reaction of at 370 K is higher than that for $_2$ $^+$ + $_2$ $\rightleftharpoons$ $_3^+$ + HD (220 K)."765" Theratio (ο) /CIL, decreases towards (he center of the disk. because there is a [all in ID in the molecular laver. caused partlv by the increase in photodissociation as the UV field increases and partly by (he efficient incorporation of deuterium into water and ammonia ices."," The ratio $_2$ $^+$ $_3^+$ decreases towards the center of the disk, because there is a fall in HD in the molecular layer, caused partly by the increase in photodissociation as the UV field increases and partly by the efficient incorporation of deuterium into water and ammonia ices."766 With the inclusion of photodesorption (he picture changes., With the inclusion of photodesorption the picture changes.767 This mechanism can alter the molecular Ο/Η ratios bv injecting molecules formed on (he erains into the eas ancl therefore the gaseous D/II ratios partly depend on the ratios in the ice mantles., This mechanism can alter the molecular D/H ratios by injecting molecules formed on the grains into the gas and therefore the gaseous D/H ratios partly depend on the ratios in the ice mantles.768" Some molecules do nol show much variation in the D/IL ratio with photodesorption HIIDCO/IISCO and are roughly the same in both Models Band €. In Models C and D the high abundance of molecules in the molecular laver keep the abundances of IL, and D, in this region low.", Some molecules do not show much variation in the D/H ratio with photodesorption $_2$ CO and $^+$ $^+$ are roughly the same in both Models B and C. In Models C and D the high abundance of molecules in the molecular layer keep the abundances of $_3^+$ and $_3^+$ in this region low.769" Ilence most of the IL, and D, are in the midplane where the temperatures are low and (he deuteration is high leading to high D. /IL, ralios for these models.", Hence most of the $_3^+$ and $_3^+$ are in the midplane where the temperatures are low and the deuteration is high leading to high $_3^+$ $_3^+$ ratios for these models.770" In Models A and B the higher depletion al z > 50 AU means (hat the destruction rate of IL, is lower and we see a secondary peak in abundance for this molecule ab z o 10 - 120 AU (Figure 11)).", In Models A and B the higher depletion at $z$ $>$ 50 AU means that the destruction rate of $_3^+$ is lower and we see a secondary peak in abundance for this molecule at $z$ $\sim$ 70 - 120 AU (Figure \ref{fig:frac_250_a}) ).771" Consequently NULL, ) is higher andN(D,4 )/N(IT, ) is lower in these models.", Consequently $_3^+$ ) is higher and$_3^+$ $_3^+$ ) is lower in these models.772 In Model D we find that DCN/IIC'N decreases slehtly with decreasing radius in agreement with ?.. although our ratios are slightlv higher than theirs.," In Model B we find that DCN/HCN decreases slightly with decreasing radius in agreement with \citet{ah01}, although our ratios are slightly higher than theirs."773 At large radii DC'N. forms from which is produced by the reaction of either or D. with INC., At large radii DCN forms from $^+$ which is produced by the reaction of either $^+$ or $_3^+$ with HNC.774" Hence DCN/IICN is dependent on the deuteration of IL,.", Hence DCN/HCN is dependent on the deuteration of $_3^+$ .775 At R= 50 AU. neutral.neutral reactions," At $R$ = 50 AU, neutral–neutral reactions"776pattern.,pattern.777 The current. observational boundaries (Table 1)) are MM hieh A forPuls Pal 200. VY Sel 1200 τςPuls) z;2500. --Low state ye2500," The current observational boundaries (Table \ref{tab1}) ) are: Stable high $\dot{M}$ for $P_{orb}$ (s) $\la 1200$ , VY Scl 1200 $\la P_{orb}$ (s) $\la 2500$, Low state $P_{orb}$ (s) $\ga 2500$."778 As more AM CVn stars are it will be interesting to see whether the boundaries are sharp or if there is some fuzziness to them.," As more AM CVn stars are discovered, it will be interesting to see whether the boundaries are sharp or if there is some fuzziness to them."779 PAW is supported by funds made available from the National Research Foundation and by strategic funds mace available to DW from the University of Cape Town., PAW is supported by funds made available from the National Research Foundation and by strategic funds made available to BW from the University of Cape Town.780 DW's research is supported by the University., BW's research is supported by the University.781 We thank Retha Pretorius for kindlytaking snapshots of -2O003aw., We thank Retha Pretorius for kindlytaking snapshots of `2003aw'.782"anti-parallel magnetic fields coexist in a narrow and compressed region, does not form between the filaments.","anti-parallel magnetic fields coexist in a narrow and compressed region, does not form between the filaments."783" Instead, magnetic fields vanish at the middle point between the filaments."," Instead, magnetic fields vanish at the middle point between the filaments."784" From the analysis of results of the simulation, it is found that there is a quasi-equilibrium after the formation of the filamentary structure in the saturation phase of the Weibel instability."," From the analysis of results of the simulation, it is found that there is a quasi-equilibrium after the formation of the filamentary structure in the saturation phase of the Weibel instability."785" In this section, we consider whether the equilibrium is expressed by an analytical solution of the Vlasov-Maxwell system."," In this section, we consider whether the equilibrium is expressed by an analytical solution of the Vlasov-Maxwell system."786" Suzuki&Shigeyama(2008) present a method to construct stationary solutions of the Vlasov-Maxwell system and derive a new two-dimensional equilibrium configuration of collisionless plasmas, whose velocity distribution really resembles the quasi-equilibrium shown above."," \cite{ss08} present a method to construct stationary solutions of the Vlasov-Maxwell system and derive a new two-dimensional equilibrium configuration of collisionless plasmas, whose velocity distribution really resembles the quasi-equilibrium shown above."787" Slightly modifying the result of Suzuki&Shigeyama (2008), we construct the following equilibrium, with the electromagnetic potentials expressed as which lead to One can easily check that these expressions satisfy the Vlasov-Maxwell system exactly."," Slightly modifying the result of \cite{ss08}, we construct the following equilibrium, with the electromagnetic potentials expressed as which lead to One can easily check that these expressions satisfy the Vlasov-Maxwell system exactly."788 In the following we considerthe case of m;=me because ions are not thermalized in the case of m;— 16me., In the following we considerthe case of $m_\mathrm{i}=m_\mathrm{e}$ because ions are not thermalized in the case of $m_\mathrm{i}=16m_\mathrm{e}$ .789" If the stationary solution described above actually corresponds to the quasi-equilibrium, free parameters in the solution must be determined from results of the simulation."," If the stationary solution described above actually corresponds to the quasi-equilibrium, free parameters in the solution must be determined from results of the simulation."790" The value of is derived by fitting a gaussian to the x,y,v4,,v,-integratedυο{ο) velocity distribution of electrons in the quasi-equilibrium phase."," The value of $v_\mathrm{e}(=v_\mathrm{i})$ is derived by fitting a gaussian to the $x,y,v_x,v_z$ -integrated velocity distribution of electrons in the quasi-equilibrium phase."791" Then, we obtain ve=0.14."," Then, we obtain $v_\mathrm{e}=0.14$."792" The value of Bo is derived from the magnetic energy Ej using Equations (21)) and (28)) as On the other hand, Ej,~0.2 from the result."," The value of $B_0$ is derived from the magnetic energy $E_\mathrm{m}$ using Equations \ref{Em}) ) and \ref{Beq}) ) as On the other hand, $E_\mathrm{m}\simeq 0.2$ from the result."793 So we obtain Bo~v/0.2/ 0.14., So we obtain $B_0\simeq \sqrt{0.2}/\pi\simeq 0.14$ .794" Figure 9 shows the density distribution and the configuration of the magnetic field B, and B, of", Figure \ref{figure9} shows the density distribution and the configuration of the magnetic field $B_x$ and $B_y$ of795"are much brighter than AKp=2 would most likely have been detected spectroscopically (seeBatalha2011), so we consider those to be ruled out as well.","are much brighter than $\Delta K\!p = 2$ would most likely have been detected spectroscopically \citep[see][]{Batalha:11}, so we consider those to be ruled out as well."796 We indicate this with the green hatched region in the lower right-hand side of the figure., We indicate this with the green hatched region in the lower right-hand side of the figure.797" Finally, the colors of the background/foreground configurations simulated with pprovide a further constraint which is represented by the blue hatched area on the lower left of the figure."," Finally, the colors of the background/foreground configurations simulated with provide a further constraint which is represented by the blue hatched area on the lower left of the figure."798" This swath of parameter space is excluded because the blends are significantly redder than the color index measured for Kepler-10 (Kp—K,=1.465+ 0.029), by more than three times the uncertainty in the observed index."," This swath of parameter space is excluded because the blends are significantly redder than the color index measured for Kepler-10 $K\!p-K_s = 1.465 \pm 0.029$ ), by more than three times the uncertainty in the observed index."799" As a result of these complementary constraints, the only section of parameter space remaining for viable blends involving star+planet pairs is the area under the 3c contour and limited from below and on the left by the hatched areas (color and brightness conditions) and shaded area cconstraint), respectively."," As a result of these complementary constraints, the only section of parameter space remaining for viable blends involving star+planet pairs is the area under the $\sigma$ contour and limited from below and on the left by the hatched areas (color and brightness conditions) and shaded area constraint), respectively."800 All of these blends have the eclipsing pair behind the target (foreground scenarios are all ruled , All of these blends have the eclipsing pair behind the target (foreground scenarios are all ruled out).801"We note out).that in this star+planet blend scenario white dwarfs can also act as tertiaries, as long as they are cooler than the secondaries so that they do not lead to deep occultation events that would have been seen in the light curve ofKOI-072."," We note that in this star+planet blend scenario white dwarfs can also act as tertiaries, as long as they are cooler than the secondaries so that they do not lead to deep occultation events that would have been seen in the light curve of."802"02.. The above range of tertiary radii Rjup to R35) excludes essentially all cool carbon-oxygen and oxygen-neon white dwarfs more massive than about 0.4Mo, as these are smaller than the lower limit set byBLENDER,, which corresponds to Re (see,e.g.,Paneietal.2000)."," The above range of tertiary radii $R_{\rm Jup}$ to $R_{\rm Jup}$ ) excludes essentially all cool carbon-oxygen and oxygen-neon white dwarfs more massive than about $M_{\sun}$, as these are smaller than the lower limit set by, which corresponds to $R_{\earth}$ \citep[see,803e.g.,][]{Panei:00}."804". Low-mass helium-core or oxygen-core white dwarfs that are the product of common-envelope evolution in binary stars can be considerably larger in size, although they appear to be very rare."," Low-mass helium-core or oxygen-core white dwarfs that are the product of common-envelope evolution in binary stars can be considerably larger in size, although they appear to be very rare."805 The MMission itself has uncovered only three examples to date (Roweetal.2010;Carter2011).," The Mission itself has uncovered only three examples to date \citep{Rowe:10, Carter:11}."806". However, all of them are very hot (T.g>10,000 KK), and produce deep and unmistakable flat-bottomed occultation signals."," However, all of them are very hot $T_{\rm eff} > 10,000$ K), and produce deep and unmistakable flat-bottomed occultation signals."807" Model calculations such as those of Paneietal.(2007) show that as these helium-core white dwarfs cool, their radii quickly become Earth-size or smaller."," Model calculations such as those of \cite{Panei:07} show that as these helium-core white dwarfs cool, their radii quickly become Earth-size or smaller."808" Therefore, we do not consider white dwarfs to be a significant source of blends forKOI-072."," Therefore, we do not consider white dwarfs to be a significant source of blends for."809"02.. Eclipsing binaries composed of two stars physically associated with the target are clearly ruled out byBLENDER,, as they produce very poor fits to the llight curves."," Eclipsing binaries composed of two stars physically associated with the target are clearly ruled out by, as they produce very poor fits to the light curves."810" For cases in which the tertiaries are planets, viable scenarios identified by sspan a range of secondary masses and tertiary radii within the 3e contour shown in Figure 5.."," For cases in which the tertiaries are planets, viable scenarios identified by span a range of secondary masses and tertiary radii within the $\sigma$ contour shown in Figure \ref{fig:htp_r3}."811" Most of these configurations turn out to involve eccentric orbits, with transit durations longer than those corresponding to circular orbits along with secondary stars that are smaller than the primary (see Figure 6))."," Most of these configurations turn out to involve eccentric orbits, with transit durations longer than those corresponding to circular orbits along with secondary stars that are smaller than the primary (see Figure \ref{fig:htp_dur}) )."812" Once again other observational constraints are very complementary, and in this case they are sufficient to exclude all of these blends."," Once again other observational constraints are very complementary, and in this case they are sufficient to exclude all of these blends."813" For example, the shaded area of parameter space to the left of Mo is eliminated by the oobservations, as described earlier."," For example, the shaded area of parameter space to the left of $M_{\sun}$ is eliminated by the observations, as described earlier."814" The constraint on the Kp—K, color (hatched area on the left) is partly redundant with the NIR observations, but extends to slightly larger secondary masses."," The constraint on the $K\!p - K_s$ color (hatched area on the left) is partly redundant with the NIR observations, but extends to slightly larger secondary masses."815" And finally, the spectroscopic constraint removes the remaining scenarios corresponding to higher-mass (brighter) secondaries."," And finally, the spectroscopic constraint removes the remaining scenarios corresponding to higher-mass (brighter) secondaries."816" We conclude that of all the hierarchical triple blend scenarios that are capable of precisely reproducing the detailed shape of the ttransit light curve, would have escaped detection by one or more of our follow-up efforts, including NIR oobservations, high-resolution spectroscopy, or absolute"," We conclude that of all the hierarchical triple blend scenarios that are capable of precisely reproducing the detailed shape of the transit light curve, would have escaped detection by one or more of our follow-up efforts, including NIR observations, high-resolution spectroscopy, or absolute"817choices. we followed the recipe described by Fan&Piran(2008).,"choices, we followed the recipe described by \citet{FaPi08}."818" The SSC process essentially generates a new spectral component superposed {ο the underlying synchrotron spectrum. with the same global shape up to a cut-off frequency: where [is the fireball bulk motion Lorentz factor. 71, the electron mass. c the speed of light and /: the Planck constant."," The SSC process essentially generates a new spectral component superposed to the underlying synchrotron spectrum, with the same global shape up to a cut-off frequency: where $\Gamma$ is the fireball bulk motion Lorentz factor, $m_{\rm e}$ the electron mass, $c$ the speed of light and $h$ the Planck constant."819 Above this frequency the SSC emission is no more in the Thomson regime and becomes much weaker (Klein-Nishina regime)., Above this frequency the SSC emission is no more in the Thomson regime and becomes much weaker (Klein-Nishina regime).820 For typical bulk motion Lorentz factors (T~200attheafterglowonset.Molinarietal.2007) the SSC emission of the afterglow is in the Thomson regime.," For typical bulk motion Lorentz factors \citep[$\Gamma \sim 200$ at the afterglow onset,][]{Mol07} the SSC emission of the afterglow is in the Thomson regime."821" Assuming we are in a constant density circumburst environment. the predictec SSC spectrum ts characterized by two typical frequencies (Fan&Piran2008) às the synchrotron afterglow spectrum refsec:aft)): where Yssc=US,/Uy, is the rest frame synchrotron to magnetic field energy density ratio."," Assuming we are in a constant density circumburst environment, the predicted SSC spectrum is characterized by two typical frequencies \citep{FaPi08} as the synchrotron afterglow spectrum \\ref{sec:aft}) ): where $Y_{\rm SSC} = U'_{\rm syn}/U'_{\rm B}$ is the rest frame synchrotron to magnetic field energy density ratio."822" Defining £.=(vy/v,)7*7. it can be shown (Sari&Esin2001) ναί The synchrotron injection to cooling synchrotron frequency ratio for the slow-cooling case is: The numerical factor2C; in front ofS refeq:nmner is not exactly the one derived from refeq:v and 2. since. as already mentioned in refsec:aft.. IC cooling. also affects the location of the synchrotron cooling frequency making. the. problem numerically difficult to solve."," Defining $\xi_{\rm c} = (\nu_{\rm m}/\nu_{\rm c})^{(p-2)/2}$, it can be shown \citep{SaEs01} : The synchrotron injection to cooling synchrotron frequency ratio for the slow-cooling case is: The numerical factor in front of \\ref{eq:nmncr} is not exactly the one derived from \\ref{eq:v} and \ref{eq:c} since, as already mentioned in \\ref{sec:aft}, IC cooling also affects the location of the synchrotron cooling frequency making the problem numerically difficult to solve."823 We now apply an approximate solution fully adequate for our goals (seeFan&Piran2008.forafull discussion).., We now apply an approximate solution fully adequate for our goals \citep[see][for a full discussion]{FaPi08}.824 With our parameters ((8)) becomes v/v.=0.00025 and Yssc.= 2.1. ((5)), With our parameters \ref{eq:nmncr}) ) becomes $ \nu_{m} / \nu_{c} \simeq 0.00025 $ and $Y_{\rm SSC} \simeq 2.1$ . \ref{eq:mssc}) )825 and (6)) become vinsse505 HHz LEX(=5 kKkeV) and visse=7.4%107? HHz (zx310 MMeV).," and \ref{eq:vssc}) ) become $\nu_{\rm m,SSC} \approx 1.1 \times 10^{18}$ Hz $\simeq 5$ keV) and $\nu_{\rm c,SSC} \approx 7.4 \times 10^{22}$ Hz $\simeq 310$ MeV)."826 The cooling SSC frequency is at much lower energy than the band covered by the MAGIC observations (EuacGic~ 90GGeV)., The cooling SSC frequency is at much lower energy than the band covered by the MAGIC observations $_{\rm MAGIC} \sim 90$ GeV).827" We are therefore in the spectral range where the SSC spectrum ts softer. following a power-law behaviour, 7777=y—1.05"," We are therefore in the spectral range where the SSC spectrum is softer, following a power-law behaviour, $\nu^{-p/2} \simeq \nu^{-1.05}$."828 In order to derive the expected flux density at the MAGIC energy we have to compute the flux density at the typical SSC frequency (Fan&Piran2008) at the epoch of the MAGIC observation: where Dj is the luminosity distance of the source. 4.8 GGpe (~1.5x1075 cem).," In order to derive the expected flux density at the MAGIC energy we have to compute the flux density at the typical SSC frequency \citep{FaPi08} at the epoch of the MAGIC observation: where $D_{\rm L}$ is the luminosity distance of the source, $D_{\rm L} \sim 4.8$ Gpc $\sim 1.5 \times 10^{28}$ cm)."829 With our parameters. Fyssc=5.2«1078 cem ss?! MMeV7! which is much lower than the synchrotron flux at the same frequency. vjssc. well within the Swift--XRT energy range with these parameters.," With our parameters, $F_{\nu_{\rm m},{\rm SSC}} \simeq 5.2 \times 10^{-13}$ $^{-2}$ $^{-1}$ $^{-1}$ which is much lower than the synchrotron flux at the same frequency, $\nu_{\rm m,SSC}$, well within the -XRT energy range with these parameters."830 Then. finally. from the peak energy to the MAGIC band we have to extrapolate the SSC spectrum as: and. again with our parameters. Foogey~2.9x107 cem ssMMeV!.," Then, finally, from the peak energy to the MAGIC band we have to extrapolate the SSC spectrum as: and, again with our parameters, $F_{90\,{\rm GeV}} \sim 2.9 \times 10^{-18}$ $^{-2}$ $^{-1}$ $^{-1}$."831" The flux integrated in the MAGIC band. the parameter to be compared to the reported upper limits. can be well approximated by vF, at about GGeV. and we have Fyacic~2.6x107 eerecem ss7! at the epoch of the MAGIC observation. f~8 kks from the burst."," The flux integrated in the MAGIC band, the parameter to be compared to the reported upper limits, can be well approximated by $\nu F_\nu$ at about GeV, and we have $F_{\rm MAGIC} \sim 2.6 \times 10^{-13}$ $^{-2}$ $^{-1}$ at the epoch of the MAGIC observation, $t \sim 8$ ks from the burst."832 Any uncertainty in the underlying afterglow parameters affects of course the VHE predictions., Any uncertainty in the underlying afterglow parameters affects of course the VHE predictions.833 Some of these uncertainties have. however. a rather limited (considering the present observational limits) impact because one of the relevant factors. the ratio between the injection and cooling synchrotron frequency. is constrained by the afterglow SED and uncertainties for micro-physical parameters should. still keep the ratio close to the observed value.," Some of these uncertainties have, however, a rather limited (considering the present observational limits) impact because one of the relevant factors, the ratio between the injection and cooling synchrotron frequency, is constrained by the afterglow SED and uncertainties for micro-physical parameters should still keep the ratio close to the observed value."834" The v/v, ratio drives the importance of the IC component and the position of the cooling SSC frequency. 1.8. where the VHE flux begins to decrease steeply moving toward higher energies."," The $\nu_{\rm m}/\nu_{\rm c}$ ratio drives the importance of the IC component and the position of the cooling SSC frequency, i.e. where the VHE flux begins to decrease steeply moving toward higher energies."835 The total energy on the contrary is estimated assuming ai efficiency for the GRB prompt emission process., The total energy on the contrary is estimated assuming an efficiency for the GRB prompt emission process.836 This is a weakly known factor given that at present no satisfactory description of the GRB prompt emission process exists (Lyutikov2009)., This is a weakly known factor given that at present no satisfactory description of the GRB prompt emission process exists \citep{Lyut09}.837 It is therefore possible (Zhang2007) that the efficiency is substantially higher. modifying the total energy and therefore the expected flux.," It is therefore possible \citep{Zha07} that the efficiency is substantially higher, modifying the total energy and therefore the expected flux."838 Circumburst matter density has an important effect on the expectedSSC flux., Circumburst matter density has an important effect on the expectedSSC flux.839 With the present afterglow data it can essentially only be estimated coupled to the micro-physical parameters., With the present afterglow data it can essentially only be estimated coupled to the micro-physical parameters.840 A higher density would make the SSC component more important and possibly detectable at lower energies (seee.g.Harrisonetal.2001)., A higher density would make the SSC component more important and possibly detectable at lower energies \citep[see e.g.][]{Harr01}.841. However. the value of the circumburst density derived for afterglows with data allowing a detailed modeling is consistent with the value we report for 0080430 (Panaitescu2005).," However, the value of the circumburst density derived for afterglows with data allowing a detailed modeling is consistent with the value we report for 080430 \citep{Pan05}."842 A milder than expected temporal decay in the X-rays band together with the consistency of the observed SED with the reference afterglow model prediction. raises some concern about thereliability of the adopted theoretical scenario.," A milder than expected temporal decay in the X-rays band together with the consistency of the observed SED with the reference afterglow model prediction, raises some concern about thereliability of the adopted theoretical scenario."843 A shallower afterglow decay showing a synchrotron spectrum can be explained with late-time energy injection in the outflow (Panaitescu2006;Zhangetal. 2006).," A shallower afterglow decay showing a synchrotron spectrum can be explained with late-time energy injection in the outflow \citep{Pan06,Zha06}. ."844. In this case the VHE SSC flux temporal decay could be slowed in a way related to the time evolution of the energy injection (Wel&Fan2007: 20058).," In this case the VHE SSC flux temporal decay could be slowed in a way related to the time evolution of the energy injection \citep{WeFa07,GoMe07,GaPi07,FaPi08}. ."845ionized mecdium (IMAL).,ionized medium (HIM).846 Neutral atomic hydrogen (HI) is main observed constituent of the ISM: estimates of its filling [actor ranges from to (Burton 1988)., Neutral atomic hydrogen (HI) is the main observed constituent of the ISM; estimates of its filling factor ranges from to (Burton 1988).847" The WNAM is the ISAL component which occupies most of the space (Ixulkarni IHeiles 1988) in an ""intercloud medium"" or surrounding the cold clouds.", The WNM is the ISM component which occupies most of the space (Kulkarni Heiles 1988) in an “intercloud medium” or surrounding the cold clouds.848 Since information about the WNM mainly comes from 21 em line emission data. direct measurements of temperatures (estimated to be in the range 5000 Ix to 8000 IX) or densities (estimated to be between 0.1 and 1 cm.?) ave difficult to obtain.," Since information about the WNM mainly comes from 21 cm line emission data, direct measurements of temperatures (estimated to be in the range 5000 K to 8000 K) or densities (estimated to be between 0.1 and 1 $cm^{-3}$ ) are difficult to obtain."849 Therefore. the WNM is the least well understood component of the ISM.," Therefore, the WNM is the least well understood component of the ISM."850 A typical density of 0.4 cm.7 characterizes much of the ISM (Burton 1983)., A typical density of 0.4 $cm^{-3}$ characterizes much of the ISM (Burton 1988).851 We assume (hat the ISM has the characteristics of the WAAL because it is the main observed constituent of the ISM., We assume that the ISM has the characteristics of the WNM because it is the main observed constituent of the ISM.852 We have considered that the ISAT pressure is simply the standard gas kinetic pressure. 2?=nv (where n is (he number of particles per unit volume. Jv is the Boltzmans constant and Tis the gas temperature).," We have considered that the ISM pressure is simply the standard gas kinetic pressure, $P = nKT$ (where $n$ is the number of particles per unit volume, $K$ is the Boltzman's constant and $T$ is the gas temperature)."853 Spitzer (LOTS) demonstrated that the contribution of cosmic ravs. and magnetic pressures can be comparable (ο the eas kinetic pressure.," Spitzer (1978) demonstrated that the contribution of cosmic rays, and magnetic pressures can be comparable to the gas kinetic pressure."854 As we have not included. magnetic fields or turbulent motions in our simulations. we therefore. use large densiües (1 cimoE instead of the twpical 0.4 en7)a to compensate [or the cosmic ravs. magnetic (D7/8x) and turbulent pressure components (assuming all of them contribute (o an isotropic pressure).," As we have not included magnetic fields or turbulent motions in our simulations, we therefore, use large densities (1 $cm^{-3}$ instead of the typical 0.4 $cm^{-3}$ ) to compensate for the cosmic rays, magnetic $B^2/8\pi$ ) and turbulent pressure components (assuming all of them contribute to an isotropic pressure)."855 In our models. at time zero the grid is filled homogeneously with neutral atomic hydrogen with a density of Lem7. a temperature of 6000 IX. and zero macroscopic velocity.," In our models, at time zero the grid is filled homogeneously with neutral atomic hydrogen with a density of 1 ${\rm cm^{-3}}$, a temperature of 6000 K, and zero macroscopic velocity."856 Hence. we have to consider wind propagation in an homogeneous (cloudless). stationary (ο=0) and warm (7=6000 K) ISM (the isothermal sound speed is e;=7 ))," Hence, we have to consider wind propagation in an homogeneous (cloudless), stationary $v = 0$ ) and warm $T=6000~{\rm K}$ ) ISM (the isothermal sound speed is $c_{\rm s} = 7~$ )."857 The interaction of the stellar wind with the ISM influences the subsequent evolution of the flow., The interaction of the stellar wind with the ISM influences the subsequent evolution of the flow.858 The adopted ISM. pressure of 6000. IXcm7 lies in the range of the ISM pressure which allows pressure equilibrium between the different ISAT components., The adopted ISM pressure of 6000 ${\rm K~cm^{-3}}$ lies in the range of the ISM pressure which allows pressure equilibrium between the different ISM components.859 In. particular. thermal equilibrium between the WNM and the CNM is only possible over a range of pressure between ~2500 and ~6500 Gn units of the the Doltzman's constant).," In particular, thermal equilibrium between the WNM and the CNM is only possible over a range of pressure between $\sim$ 2500 and $\sim$ 6500 (in units of the the Boltzman's constant)."860 The gas evolution considering an ISM with pressure ten times lower is also presented in 14., The gas evolution considering an ISM with pressure ten times lower is also presented in $\S4$.861 In order to keep (he ISM at its original temperature. no source of gas heating is included.," In order to keep the ISM at its original temperature, no source of gas heating is included."862 The temperature of the ISM decreases due to radiative cooling during the evolution a negligible amount (~ 500 IX)., The temperature of the ISM decreases due to radiative cooling during the evolution a negligible amount $\sim$ 500 K).863 The beginning of the evolution is the same for all the stellar models., The beginning of the evolution is the same for all the stellar models.864 A wind with a constant mass-loss rate (10 !1)) and velocity (£3 )) starts to interact, A wind with a constant mass-loss rate $\sim 10^{-8}$ ) and velocity $\sim$ 3 ) starts to interact865PWNe powered by the Crab and Vela pulsars in the level of structure.,PWNe powered by the Crab and Vela pulsars in the level of structure.866 We first note that there is a clear axis of symmetry associated with the system., We first note that there is a clear axis of symmetry associated with the system.867 This is manifested not only in the overall elongation of the nebula. but also in the orientation of features C. D and E in Figure 2.. and features 1. 2. 4 and 5 in Figure 5..," This is manifested not only in the overall elongation of the nebula, but also in the orientation of features C, D and E in Figure \ref{fig_g320_acisi}, and features 1, 2, 4 and 5 in Figure \ref{fig_g320_center}."868" This axial structure is manifested on an extremely wide range of scales. from 10""20.2 pe up to 10215 pe (where here and in further discussion we adopt a distance to the system of 5 kpe)."," This axial structure is manifested on an extremely wide range of scales, from $10'' = 0.2$ pc up to $10' = 15$ pc (where here and in further discussion we adopt a distance to the system of 5 kpc)."869 It is hard to see how this orientation could be due to interaction with a pre-existing structure in the ISM citecas79)). which should generally only influence the overall morphology of the system but not its small-scale structure.," It is hard to see how this orientation could be due to interaction with a pre-existing structure in the ISM \\cite{cas79}) ), which should generally only influence the overall morphology of the system but not its small-scale structure."870 We therefore believe that this overall symmetry reflects the properties of the central pulsar., We therefore believe that this overall symmetry reflects the properties of the central pulsar.871 The only well-defined axes associated with a neutron star are its spin axis and its velocity vector., The only well-defined axes associated with a neutron star are its spin axis and its velocity vector.872 PWNe whose morphologies are dominated by the pulsar’s motion are axially symmetric but inevitably show a clear cometary morphology with the pulsar at one end (Frailetal.1996:; Olbertetal. 2001:: 2001))., PWNe whose morphologies are dominated by the pulsar's motion are axially symmetric but inevitably show a clear cometary morphology with the pulsar at one end \cite{fggd96}; \cite{ocw+01}; \cite{kggl01}) ).873 In contrast. the PWN surrounding PSR eextends to many parsees on both sides of the pulsar.," In contrast, the PWN surrounding PSR extends to many parsecs on both sides of the pulsar."874 It is thus clear that any motion of the pulsar cannot explain the overall PWN morphology (see additional discussion in refdisesowshock)). andthatthemainaxisofthenebulamustcorrespondtdi," It is thus clear that any motion of the pulsar cannot explain the overall PWN morphology (see additional discussion in \\ref{disc_bowshock}) ), and that the main axis of the nebula must correspond to the pulsar spin axis."875dteyfusenossiumaei Pula (B forth ," Such a correspondence was previously proposed for this pulsar by Brazier Becker \nocite{bb97}) ), has also been argued for the Crab and Vela pulsars and their nebulae \cite{hss+95}; \cite{psg+00}; \cite{hgh01}) ), and is seen in some models for pulsar magnetospheres \cite{sl90}) )."876The region of the PWN where the symmetry of the system is most clearly broken is in the vicinity of feature F. the promontory extending several aremin west of the main axis.," The region of the PWN where the symmetry of the system is most clearly broken is in the vicinity of feature F, the promontory extending several arcmin west of the main axis."877" The sharp edges seen for this feature are reminiscent of the ""bay"" observed along the western edge of the Crab Nebula (Weisskopfetal. 2000a)).", The sharp edges seen for this feature are reminiscent of the “bay” observed along the western edge of the Crab Nebula \cite{wht+00}) ).878 Such morphologies are presumably produced by strong confining pressure in these regions. possibly the result of pre-existing magnetic fields or circumstellar material.," Such morphologies are presumably produced by strong confining pressure in these regions, possibly the result of pre-existing magnetic fields or circumstellar material."879 The spectral fits in Table 2. show that X-ray emission from the diffuse component of the PWN can be well fit by a power- with photon index I22.05+0.04 and absorbing column (9.5+0.3)«107! em., The spectral fits in Table \ref{tab_spec} show that X-ray emission from the diffuse component of the PWN can be well fit by a power-law with photon index $\Gamma = 2.05\pm0.04$ and absorbing column $(9.5\pm0.3)\times10^{21}$ $^{-2}$.880 This measurement is in good agreement with previous spectral measurements for this source Okayasu.&Sekimoto 1993:; Trussonietal. 1996:: 1996:; Marsdenetal. 1997:: Mineoetal.2001 ))., This measurement is in good agreement with previous spectral measurements for this source \cite{kos93}; \cite{tmc+96}; ; \cite{tkyb96}; \cite{mbg+97}; \cite{mcm+01}) ).881 Du Plessis ((1995)) have argued from a collection of archival X-ray data sets that the photon index for this PWN steepens from [«1.6 below 2—6 keV to P=2.15 above this energy.," Du Plessis \nocite{ddb+95}) ) have argued from a collection of archival X-ray data sets that the photon index for this PWN steepens from $\Gamma \approx 1.6$ below $\varepsilon \sim 6$ keV to $\Gamma \approx8822.15$ above this energy."883 We see no evidence for this spectral curvature in our data set: a fit to the data using à broken power-law model rules out any spectral break across the bbandpass larger than AT>0.2.," We see no evidence for this spectral curvature in our data set: a fit to the data using a broken power-law model rules out any spectral break across the bandpass larger than $\Delta884\Gamma \ga 0.2$."885 It is more likely that the effect claimed by du Plessis ((1995)) was a result of the large uncertainties in and calibration differences between their lower sensitivity data sets., It is more likely that the effect claimed by du Plessis \nocite{ddb+95}) ) was a result of the large uncertainties in and calibration differences between their lower sensitivity data sets.886 The standard observational picture for a PWN is that its spectrum is comparatively flat at radio wavelengths. |< 1.3. but is steeper at X-ray energies. typically with Dz2 (see Gaensler200]. and references therein).," The standard observational picture for a PWN is that its spectrum is comparatively flat at radio wavelengths, $1 \la \Gamma \la 1.3$ , but is steeper at X-ray energies, typically with $\Gamma \approx 2$ (see \cite{gae01b}887 and references therein)."888 It is usually assumed that synchrotron losses are at least partly responsible for this steepening. generatinga spectral break at energy ον across which we expect a change in photon index AT20.5 (Kardashev 1962:; Pacint&Salvati 1973)).," It is usually assumed that synchrotron losses are at least partly responsible for this steepening, generatinga spectral break at energy $\varepsilon_s$, across which we expect a change in photon index $\Delta\Gamma = 0.5$ \cite{kar62}; \cite{ps73}) )."889 The frequency of this break. along with an estimate of the age of the PWN. can be used to infer the nebular magnetic field citefged96)).," The frequency of this break, along with an estimate of the age of the PWN, can be used to infer the nebular magnetic field \\cite{fggd96}) )."890" In the case of the PWN around PSRB1509-58.. the observed photon index [=2.05 implies 7,<1 keV. If the system is f vr old. then the implied nebular magnetic field is so that we can infer B,=8 μα. There are a number of other. independent. estimates of the nebular magnetic field: the assumption of equipartition etal.1984)). a simple model of PWN evolution (Sewardetal.1984)) and a detection of inverse-Compton emission from the PWN at TeV energies (Sakoetal.2000)) all result in estimates in the range B,25—8 j/G. while Chevalier (2000)) points out that the comparatively low efficiency. with which this pulsar converts its spin-down energy into X-ray synchrotron emission also implies 7,<4 keV and hence B,z μα. We thus conclude from a variety of arguments that the synchrotron break in this PWN is just below the X-ray band. and that the magnetic field is correspondingly low. B,z 8 μα B,=100 μα for the Crab)."," In the case of the PWN around PSR, the observed photon index $\Gamma = 2.05$ implies $\varepsilon_s \la 1$ keV. If the system is $t$ yr old, then the implied nebular magnetic field is so that we can infer $B_n \ga 8$ $\mu$ G. There are a number of other, independent, estimates of the nebular magnetic field: the assumption of equipartition \cite{shss84}) ), a simple model of PWN evolution \cite{shss84}) ) and a detection of inverse-Compton emission from the PWN at TeV energies \cite{smm+00b}) ) all result in estimates in the range $B_n = 5-8$ $\mu$ G, while Chevalier \nocite{che00}) ) points out that the comparatively low efficiency with which this pulsar converts its spin-down energy into X-ray synchrotron emission also implies $\varepsilon_s \la 4$ keV and hence $B_n \ga 5$ $\mu$ G. We thus conclude from a variety of arguments that the synchrotron break in this PWN is just below the X-ray band, and that the magnetic field is correspondingly low, $B_n \approx$ 8 $\mu$ G $B_n \ga 100$ $\mu$ G for the Crab)."891 This low magnetic field presumably results from the low confining pressure into which this PWN expands (Sewardetal...1984))., This low magnetic field presumably results from the low confining pressure into which this PWN expands \cite{shss84}) ).892" It is worth noting s age spondzokym Raposedmani 1988:; Gvaramadze2001). we infer fromhmndford/&p Equation (1)) that B,~1.5 μα. which ts at odds with all the other estimates deseribed above."," It is worth noting that if we assume a pulsar age $t\sim20$ kyr \cite{br88}; \cite{gva01b}) ), we infer from Equation \ref{eqn_b}) ) that $B_n \sim 1.5$ $\mu$ G, which is at odds with all the other estimates described above."893 The above arguments imply that the synchrotron lifetime of electrons emitting in the bband is similar to the age of the pulsar and nebula., The above arguments imply that the synchrotron lifetime of electrons emitting in the band is similar to the age of the pulsar and nebula.894 Thus while synchrotron-emitting electrons in the outer diffuse nebula have now radiated most of their energy. we expect that regions closer to the pulsar will have a flow time less than the synchrotron lifetime and thus should have a spectrum flatter by AT~0.5. Indeed the various X-ray features seen close to the pulsar (to be discussed in more detail in subsequent sections) all have photon indices D—1.5.," Thus while synchrotron-emitting electrons in the outer diffuse nebula have now radiated most of their energy, we expect that regions closer to the pulsar will have a flow time less than the synchrotron lifetime and thus should have a spectrum flatter by $\Delta\Gamma 895\sim 0.5$ Indeed the various X-ray features seen close to the pulsar (to be discussed in more detail in subsequent sections) all have photon indices $\Gamma \sim 1.5$."896" Figure 3 demonstrates a good match between the perimeter of the diffuse X-ray PWN and the ""tongue"" of emission apparent in radio observations.", Figure \ref{fig_g320_tongue} demonstrates a good match between the perimeter of the diffuse X-ray PWN and the “tongue” of emission apparent in radio observations.897 This correspondence was not apparent in earlier observations. which lacked the sensitivity to delineate the full extent of the X-ray PWN in this region FFig.," This correspondence was not apparent in earlier observations, which lacked the sensitivity to delineate the full extent of the X-ray PWN in this region Fig."898 [4 of G99)., 4 of G99).899" Given that the two regions of emission occupy similar extents and have similar shapes. we propose that the ""tongue"" simply corresponds to emission fromthe pulsar nebula. and ts the long-sought radio PWN in this system."," Given that the two regions of emission occupy similar extents and have similar shapes, we propose that the “tongue” simply corresponds to emission fromthe pulsar nebula, and is the long-sought radio PWN in this system."900 This region has previously been shown to have a very high fractional polarization and a well-ordered magnetic field (Milne.Caswell.&Haynes 1993:: G99). properties typical of a radio PWN.," This region has previously been shown to have a very high fractional polarization and a well-ordered magnetic field \cite{mch93}; G99), properties typical of a radio PWN."901 The radio PWN ts comparatively faint. sits on a complicated background. and tis in close proximity to the very bright radio emission from89.," The radio PWN is comparatively faint, sits on a complicated background, and is in close proximity to the very bright radio emission from."902. The parameters of the radio PWN are thus poorlyconstrained. but using the data of Whiteoak Green (1996)) and of G99. we estimate an approximate flux density of 2+1 Jy at both 0.8 GHz and 1.4 GHz.," The parameters of the radio PWN are thus poorlyconstrained, but using the data of Whiteoak Green \nocite{wg96}) ) and of G99, we estimate an approximate flux density of $2\pm1$ Jy at both 0.8 GHz and 1.4 GHz."903" If we assume that this represents ~50% of the total radio flux density of the PWN (the rest being hidden by 89)). we find fiGu,~(30.5)&107 ere s! en keV7!."," If we assume that this represents $\sim$ of the total radio flux density of the PWN (the rest being hidden by ), we find $f_{1~{\rm GHz}} \sim904(1\pm0.5)\times 10^{-5}$ erg $^{-1}$ $^{-2}$ $^{-1}$ ."905 By comparing this to the flux density for the diffuse X-ray nebula inferred from Table 2.. fywy=19ς107! erg em Καντ. we can infer that," By comparing this to the flux density for the diffuse X-ray nebula inferred from Table \ref{tab_spec}, , $f_{1~{\rm keV}} = 1.9\times10^{-11}$ erg $^{-2}$ $^{-1}$ , we can infer that"906HD 181420 (cf BO9) and HD 49933 (cf A08).,HD 181420 (cf B09) and HD 49933 (cf A08).907" In this section, we present the specific choices of the priors, then the astrophysical results."," In this section, we present the specific choices of the priors, then the astrophysical results."908" According to the excess power observed in both stars at low frequency, the splitting prior was set at the maximum value of the peaks."," According to the excess power observed in both stars at low frequency, the splitting prior was set at the maximum value of the peaks."909" For HD 181420, a quite large excess power appears in the frequency range [3.2,6] uHz (cf B09), with a maximum at 4.2 uHz; the prior splitting frequency was fixed at 4.2+0.5 wHz."," For HD 181420, a quite large excess power appears in the frequency range $[3.2, 6]\ \mu$ Hz (cf B09), with a maximum at 4.2 $\mu$ Hz; the prior splitting frequency was fixed at $4.2 \pm 0.5\ \mu$ Hz."910" For HD 49933, the excess power appears as a peak at almost 3.3 uHz (cf A08); the prior splitting frequency was fixed at 3.3+0.4 wHz."," For HD 49933, the excess power appears as a peak at almost $3.3\ \mu$ Hz (cf A08); the prior splitting frequency was fixed at $3.3 \pm 0.4\ \mu$ Hz."911" Secondly, the low frequency stellar noise is well fitted by a sum of 3 background components, i.e. 3 Lorentzian functions of the frequency, centered on 0 and with 3 different amplitudes and widths (?).."," Secondly, the low frequency stellar noise is well fitted by a sum of 3 background components, i.e. 3 Lorentzian functions of the frequency, centered on 0 and with 3 different amplitudes and widths \citep{Harvey_85}."912" In our case, we do not consider the too low frequency range (v<600 4Hz), hence one background component is sufficient."," In our case, we do not consider the too low frequency range $\nu < 600\ \mu$ Hz), hence one background component is sufficient."913 The power term p (see Eq. 3)), The power term $p$ (see Eq. \ref{Eq_harvey}) )914 is set to 2., is set to 2.915" Despite the simple analytic expression, the background is the most sensitive contribution of the global fitting to be determined."," Despite the simple analytic expression, the background is the most sensitive contribution of the global fitting to be determined."916" Indeed, the consequences of an incorrect fitting are damaging for the oscillation parameter estimates."," Indeed, the consequences of an incorrect fitting are damaging for the oscillation parameter estimates."917" As an example, overestimating the background leads to an underestimate of the mode amplitudes and may strongly change the output values for the height even more for the width."," As an example, overestimating the background leads to an underestimate of the mode amplitudes and may strongly change the output values for the height even more for the width."918" Therefore, to diminish the risk of incorrectly estimating the background, we fitted it alone over a wide frequency range [600,5000] “Hz, where the oscillation mode energy is negligible and the interval large enough to constrain the background parameters."," Therefore, to diminish the risk of incorrectly estimating the background, we fitted it alone over a wide frequency range $[600,5000]\ \mu$ Hz, where the oscillation mode energy is negligible and the interval large enough to constrain the background parameters."919" Then, when fitting the oscillation spectrum, these parameters are set as free with the previous estimates as priors."," Then, when fitting the oscillation spectrum, these parameters are set as free with the previous estimates as priors."920" The tolerance is set to of the prior values, to permit small fluctuations due to the influence of the oscillation power when fitting the background profile alone."," The tolerance is set to of the prior values, to permit small fluctuations due to the influence of the oscillation power when fitting the background profile alone."921" For the GHA, the prior value of the width of the Gaussian function σῃ is set to the approximate width at half maximum estimated on a smoothed power spectrum over 25 uHz."," For the GHA, the prior value of the width of the Gaussian function $\sigma\ind{H}^2$ is set to the approximate width at half maximum estimated on a smoothed power spectrum over 25 $\mu$ Hz."922 For, For923differences.,differences.924 First. rings are edge-brightened when seen in projection. yet most of the linear features in Fig.," First, rings are edge-brightened when seen in projection, yet most of the linear features in Fig."925 | do not get brighter towards their ends., \ref{nutfig} do not get brighter towards their ends.926 Second. axisymmetric rings project to straight lines that pass through the systemic velocity of the galaxy at its centre. whereas many of the linear features in Fig.," Second, axisymmetric rings project to straight lines that pass through the systemic velocity of the galaxy at its centre, whereas many of the linear features in Fig."927 | are not quite straight. and have non-zero velocities at the centres of their galaxies.," \ref{nutfig} are not quite straight, and have non-zero velocities at the centres of their galaxies."928 Such features cannot occur in an axisymmetric potential. so we are once again forced to conclude that these galaxies are barred.," Such features cannot occur in an axisymmetric potential, so we are once again forced to conclude that these galaxies are barred."929 A third possibility. finally. is that we are seeing variations in ionization structure of the gas. and not inhomogeneities in the distribution of the gas itself.," A third possibility, finally, is that we are seeing variations in ionization structure of the gas, and not inhomogeneities in the distribution of the gas itself."930 In fact. in barred galaxies we also expect such variations. induced by the shocks: and they are indeed observed as systematic variations 1n N[II]/Ho ratio over the (2.>ο) diagrams (see also Bureau 1998)).," In fact, in barred galaxies we also expect such variations, induced by the shocks; and they are indeed observed as systematic variations in $\alpha$ ratio over the $(R,v)$ diagrams (see also Bureau \cite{b98}) )."931 A detailed analysis of the line ratio lies beyond the scope of this paper., A detailed analysis of the line ratio lies beyond the scope of this paper.932 However. as far as can be ascertained. the Ho emission line shows identical structure to the [NII] line. providing further evidence that the structure cannot be attributed to the details of the gas” ronization state.," However, as far as can be ascertained, the $\alpha$ emission line shows identical structure to the [NII] line, providing further evidence that the structure cannot be attributed to the details of the gas' ionization state."933 In summary. this spectral study of edge-on galaxies quite firmly establishes the link between boxy bulges and galactic bars.," In summary, this spectral study of edge-on galaxies quite firmly establishes the link between boxy bulges and galactic bars."934 However. we have only just begun to tap into the wealth of information that the spectra provide.," However, we have only just begun to tap into the wealth of information that the spectra provide."935 Modelling the full complexity of spectral data such as those shown in Fig., Modelling the full complexity of spectral data such as those shown in Fig.936 | should yield a wealth of information about barred galaxies. allowing us to map out their complete three-dimensional structure for the first time.," \ref{nutfig} should yield a wealth of information about barred galaxies, allowing us to map out their complete three-dimensional structure for the first time."937 The Galactic halo provides important clues for understanding the evolution and structure of the Galaxy.,	 The Galactic halo provides important clues for understanding the evolution and structure of the Galaxy.938 In the past few decades. considerable observational and theoretical efforts have been made to investigate its chemical evolution. details. of its structure. and its kinematical characteristics.," In the past few decades, considerable observational and theoretical efforts have been made to investigate its chemical evolution, details of its structure, and its kinematical characteristics."939 Very poor stars in the halo. those with metallicity < —2.0. are regarded as fossils of the earliest generations of stars.," Very metal-poor stars in the halo, those with metallicity $\lesssim940-2.0$ , are regarded as fossils of the earliest generations of stars."941 They preserve the chemical information created by their stellar progenitors. providing fundamental insights regarding the properties of the very first generation of stars. the chemical history of our Galaxy (and other large spirals like it). the modes of star formation in the proto-Ailky Way. the formation of the Galactic halo. and physical nechanisms such as feedback processes in the early stages of galaxy evolution.," They preserve the chemical information created by their stellar progenitors, providing fundamental insights regarding the properties of the very first generation of stars, the chemical history of our Galaxy (and other large spirals like it), the modes of star formation in the proto-Milky Way, the formation of the Galactic halo, and physical mechanisms such as feedback processes in the early stages of galaxy evolution."942 Although we are gaining a deeper understanding over time. much remains to be explored.," Although we are gaining a deeper understanding over time, much remains to be explored."943 It is particularly. revealing that. after many decades of assuming that the Galactic halo comprises a single stellar population. recent work (??) has provided additional support to suspicions that emerged from previous efforts that the halo is indeed divisible into two structural components. with notably different spatial density profiles. stellar orbits. and stellar metallicities.," It is particularly revealing that, after many decades of assuming that the Galactic halo comprises a single stellar population, recent work \citep{Carollo2007Nature,Carollo2010ApJ}944 has provided additional support to suspicions that emerged from previous efforts that the halo is indeed divisible into two structural components, with notably different spatial density profiles, stellar orbits, and stellar metallicities."945 Recently. new theoretical. models (e.g.. ???)) and observational constraints (e.g. 2222??22)) have greatly enhanced our understanding of the nature of the halo components of our Galaxy. Those are enabling the development of plausible assembly histories based on the degree of detectable spatial and phase-space substructures.," Recently, new theoretical models (e.g., \citealt{Helmi2008AARv,Prantzos2008AA, Salvadori2010MNRAS}) ) and observational constraints (e.g., \citealt{Carollo2007Nature,Carollo2010ApJ,Bell2008ApJ,Ivezic2008ApJ,946Juric2008ApJ,Bond2009astroph,deJong2010ApJ}) ) have greatly enhanced our understanding of the nature of the halo components of our Galaxy, Those are enabling the development of plausible assembly histories based on the degree of detectable spatial and phase-space substructures."947 The possible association of at least some presently observed dwarf galaxies with the formation of the halo populations. as invoked by ? to account for their dual halo structure. has received additional support based on high-resolution spectroscopic analysis of individual stars in. ultra-faint and dwarf spheroidal galaxies(e.g.. 222??))," The possible association of at least some presently observed dwarf galaxies with the formation of the halo populations, as invoked by \citet{Carollo2007Nature} to account for their dual halo structure, has received additional support based on high-resolution spectroscopic analysis of individual stars in ultra-faint and dwarf spheroidal galaxies(e.g., \citealt{Munoz2006ApJ,Kirby2008ApJ,Geha2009ApJ,948Frebel2010Nature,Norris2010ApJ}) )."949 Finally. the identification and detailed analysis of the elemental abundance patterns for the most chemically primitive stars. e.g.. the ultra—4.0:: 2)) and hyper-5.0:; ??2)) metal-poor stars allow one to trace back close to the very beginning of star formation in the Galaxy.," Finally, the identification and detailed analysis of the elemental abundance patterns for the most chemically primitive stars, e.g., the ultra; \citealt{Norris2007ApJ}) ) and hyper; \citealt{Christlieb2002Nature,Frebel2005Nature,Aoki2006ApJ}) ) metal-poor stars allow one to trace back close to the very beginning of star formation in the Galaxy."950 The observed metallicity distribution function (MDF) of halo stars provides strong constraints on models for the formation and chemical evolution of the Galaxy., The observed metallicity distribution function (MDF) of halo stars provides strong constraints on models for the formation and chemical evolution of the Galaxy.951 Any accepted model must be able to predict the relative numbers of halo stars as a function of their metallicity (22).. and in the case of a dual- model. as a function of location and kinematics.," Any accepted model must be able to predict the relative numbers of halo stars as a function of their metallicity \citep{Beers2005ARAA,Helmi2008AARv}, and in the case of a dual-halo model, as a function of location and kinematics."952 Early investigations on the shape of the halo MDF were hampered by the small numbers of very metal-poor stars known at the time (222?)..," Early investigations on the shape of the halo MDF were hampered by the small numbers of very metal-poor stars known at the time \citep{Hartwick1976ApJ,Bond1981ApJ, Ryan1991AJ,Carney1996AJ}."953 Other attempts (e.g.. 22)). based on samples of metal-poor stars from the HK survey of Beers and colleagues (??).. suffer from poorly constrained selection eriteria. except perhaps at the lowest metallicities.," Other attempts (e.g., \citealt{Bonifacio2000AJ,Schuster2004AA}) ), based on samples of metal-poor stars from the HK survey of Beers and colleagues \citep{Beers1985AJ,Beers1992AJ}, suffer from poorly constrained selection criteria, except perhaps at the lowest metallicities."954" More recent efforts have made use of statistically well-understood selection. criteria to identify large numbers of metal-poor candidates from objective-prism surveys. such as the Hamburg/ESO survey (HES—?).. as reported in a series of papers (?2???).. ὧν,"," More recent efforts have made use of statistically well-understood selection criteria to identify large numbers of metal-poor candidates from objective-prism surveys, such as the Hamburg/ESO survey \citep[HES --][]{Wisotzki1996AA}, as reported in a series of papers \citep{Barklem2005AA,HESstellarIV,HESstellarV,Placco2010AJ}. \citet{HESstellarV},"955 for example. used a sample of 1638 metal-poor giants to study the shape of the low-metallicity tail of the halo MDF. and made detailed comparisons with MDFs of Galactic globular clusters and satellite galaxies. as well as with theoretical models.," for example, used a sample of 1638 metal-poor giants to study the shape of the low-metallicity tail of the halo MDF, and made detailed comparisons with MDFs of Galactic globular clusters and satellite galaxies, as well as with theoretical models."956 Main-sequence turnoff (MSTO) stars have long been used to explore Galactic structure. including the recognition of stellar substructures in the Galactic halo (??).. searches for kinematic streams (e.g.. 2)). and statistical analyses of the amount of cold halo substructure in the Milky Way (e.g.. ?)).," Main-sequence turnoff (MSTO) stars have long been used to explore Galactic structure, including the recognition of stellar substructures in the Galactic halo \citep{Majewski2004ApJ,An2009ApJ}, searches for kinematic streams (e.g., \citealt{Klement2009ApJ}) ), and statistical analyses of the amount of cold halo substructure in the Milky Way (e.g., \citealt{Schlaufman2009ApJ}) )."957 In addition. MSTO stars have also been proven important to the field of Galactic chemical evolution. through the analysis of high-resolution. high signal-to-noise spectroscopic observations to derive elemental abundances for metal-poor dwarf stars (2).. chemically interesting metal-poor turnoff stars (?).. and investigations of the so-called Spite Plateau (?) through Li abundance measurements for metal-poor turnoff stars(22)..," In addition, MSTO stars have also been proven important to the field of Galactic chemical evolution, through the analysis of high-resolution, high signal-to-noise spectroscopic observations to derive elemental abundances for metal-poor dwarf stars \citep{Cohen2004ApJ}, chemically interesting metal-poor turnoff stars \citep{Aoki2008ApJ}, and investigations of the so-called Spite Plateau \citep{spite1982AA}958 through Li abundance measurements for metal-poor turnoff \citep{Aoki2009ApJ,Sbordone2010astroph}."959 In this paper we construct the MDF of Galactic halo MSTO stars based on follow-up moderate-resolution (R~ 2000) spectroscopic observations of candidate metal- turnoff stars from the HES., In this paper we construct the MDF of Galactic halo MSTO stars based on follow-up moderate-resolution $R \sim 2000$ ) spectroscopic observations of candidate metal-poor turnoff stars from the HES.960 We also compare our results, We also compare our results961"Lyman-a Nyy=2<10°? high-z z Wolfeetal.2005)). 7;~5000—8000 7;~40—200 Lyman-a 7, Kanekar&Briggs High-z 7, (@700 Wolfe&Davis1979;Kanekar&Chengalur2003)) (7.=300 Braun high-z (Carillietal.1996;","$\alpha$ $\nhi \ge 2 \times 10^{20}$ $z$ $z$ \citealp{wolfe05}) $T_k \sim 5000 - 8000$ \citealt{wolfire95}) $T_k \sim 40 -200$ $\alpha$ $\ts$ \citealp{kanekar04}) $z$ $\ts$ $\gtrsim 700$ \citealt{wolfe79,carilli96,kanekar03}) $\ts \lesssim 300$ \citealt{braun92}) $z$ \citep{carilli96,chengalur00}."962" Wolfeetal.(2003) high-z indicates sizeable CNM fractions (~50%)) in half the DLA population,", \citet{wolfe03b} $z$ indicates sizeable CNM fractions $\sim 50$ ) in half the DLA population.963 A plausible cause for the putative higher WNM fractions in high-z DLAs ts their typically-Iow metallicity |Z/H]. implying fewer routes for gas cooling.," A plausible cause for the putative higher WNM fractions in $z$ DLAs is their typically-low metallicity [Z/H], implying fewer routes for gas cooling."964" If high 7; values in DLAs arise due to low absorber metallicities. one would expect an anti-correlation between 7; and [Z/H]. with low-7, DLAs having high metallicities. and vice-versa (Kanekar&Chengalur 2001).."," If high $\ts$ values in DLAs arise due to low absorber metallicities, one would expect an anti-correlation between $\ts$ and [Z/H], with $\ts$ DLAs having high metallicities, and vice-versa \citep{kanekar01a}."965 We report here the detection of the predicted anti-correlation between 7; and [Z/H]. supporting the conclusion that tin high-z DLAs is predominantly in the warm. neutral medium.," We report here the detection of the predicted anti-correlation between $\ts$ and [Z/H], supporting the conclusion that in $z$ DLAs is predominantly in the warm neutral medium."966 Over the last decade. we have carried out 21em absorption studies of DLAs towards compact. radio-loud quasars to measure their spin temperatures (e.g. Kanekar 2007). and have also measured the DLA covering factors through low-frequency very long baseline interferometry (VLBI) studies (Kanekaretal.2009a).," Over the last decade, we have carried out 21cm absorption studies of DLAs towards compact, radio-loud quasars to measure their spin temperatures (e.g. \citealt{kanekar03,kanekar06,kanekar07,york07}) ), and have also measured the DLA covering factors through low-frequency very long baseline interferometry (VLBI) studies \citep{kanekar09a}."967". The VLBI images yleld the fraction of compact radio emission, and thus a lower limit to the DLA covering factor,"," The VLBI images yield the fraction of compact radio emission, and thus a lower limit to the DLA covering factor."968" However. no additional radio emission is detected up to scales of ~1"". indicating that the remaining emission arises from much larger scales (z10 kpe). and is unlikely to be covered by the foreground DLA."," However, no additional radio emission is detected up to scales of $\sim 1''$, indicating that the remaining emission arises from much larger scales $\gtrsim 10$ kpc), and is unlikely to be covered by the foreground DLA."969 The radio core fraction thus provides a good estimate of the DLA covering factor., The radio core fraction thus provides a good estimate of the DLA covering factor.970 We have also obtained metallicity estimates for most of the DLAs with 21em studies from our own observations or the literature., We have also obtained metallicity estimates for most of the DLAs with 21cm studies from our own observations or the literature.971" There are 26 DLAs. at 0.09©2=3.45. with estimates of both 7, and [Z/H]. ofwhich20havecovering factor estimates from low-frequencyVLBIstudies.and2]have estimates of dust depletion. [Z/Fe]."," There are 26 DLAs, at $0.09 \lesssim z \lesssim 3.45$, with estimates of both $\ts$ and [Z/H], ofwhich20havecovering factor estimates from low-frequencyVLBIstudies,and21have estimates of dust depletion, [Z/Fe]."972 The 26 absorbers of the sample are listed in Table 1.. whose columns contain (1) the quasar name. (2) the DLA redshift. (3) the," The 26 absorbers of the sample are listed in Table \ref{tab:met}, , whose columns contain (1) the quasar name, (2) the DLA redshift, (3) the"973variable event is that observed: around: 200000 s which is shown in Figure 3 with a smaller bin size of 200 s. The 9 data points to the left of the variability event (only two data points are shown in Fie.,variable event is that observed around 200000 s which is shown in Figure 3 with a smaller bin size of 200 s. The 9 data points to the left of the variability event (only two data points are shown in Fig.974 3) have a mean value of (2.21+1.26)107counts1. we treat this average count rate as the rate in the quiescent state from which the Dare arises.," 3) have a mean value of $(2.21\pm1.26)\times10^{-2}{\rm~count~s^{-1}}$, we treat this average count rate as the rate in the quiescent state from which the flare arises."975 After 200200 s from the beginning of the observation. he PSPC count rate began to increase reaching a maximum. of (8.9042.25)0?counts+ at 200600 s. Thus. an increase in the count rate by a factor of ~4 in <4004141s (<3804134s in the rest frame) is detected.," After 200200 s from the beginning of the observation, the PSPC count rate began to increase reaching a maximum of $(8.90\pm2.25)\times10^{-2}{\rm~count~s^{-1}}$ at 200600 s. Thus, an increase in the count rate by a factor of $\sim4$ in $\le 400\pm141{\rm~s}$ $\le 380\pm134{\rm~s}$ in the rest frame) is detected."976 Thus. for he variability event shown in Fig.," Thus, for the variability event shown in Fig."977 3. we find a change in the PSPC count rate of (6.60+2.58) 107>counts1," 3, we find a change in the PSPC count rate of $6.69\pm2.58$ $\times10^{-2}{\rm~count~s^{-1}}$."978 This corresponds a change in the intrinsic luminosity of (7.43+2.86).10!ergs in the energy band of 0.1.2.0 keV in the rest frame time interval of «380+134s., This corresponds a change in the intrinsic luminosity of $(7.43\pm2.86)\times10^{44}{\rm~erg~s^{-1}}$ in the energy band of 0.1–2.0 keV in the rest frame time interval of $<380\pm134{\rm~s}$.979" In order to investigate the long term (time scale of a [ew vears) variation in the X-ray intensity of RN 1334.2|3759. we have converted the WRI count rate. obtained. from the observation of 1997 June 4. into equivalent PSPC count rate using the best-fit model parameters (Ng.=3.310""em7. Py— 3x8) obtained from the joint-fit to the PSPC spectra (see $3.2)."," In order to investigate the long term (time scale of a few years) variation in the X-ray intensity of RX J1334.2+3759, we have converted the HRI count rate, obtained from the observation of 1997 June 4, into equivalent PSPC count rate using the best-fit model parameters $\NH=3.3\times10^{20}{\rm~cm^{-2}}$, $\Gamma_{X}=3.8$ ) obtained from the joint-fit to the PSPC spectra (see $\S 3.2$ )."980 The equivalent PSPC count rate is estimated to be (2.22+0.11)10counts+. which is similar to the PSPC count rates obtained during the observations of 1991 and 1993.," The equivalent PSPC count rate is estimated to be $(2.22\pm0.11)\times10^{-2}{\rm~count~s^{-1}}$, which is similar to the PSPC count rates obtained during the observations of 1991 and 1993."981 Photon energy spectra of RA J1334.213150. were accumulated from their PSPC observations shown in Table 2., Photon energy spectra of RX J1334.2+3759 were accumulated from their PSPC observations shown in Table 2.982 The same regions for the source anc the background. as stated. above. were used.," The same regions for the source and the background, as stated above, were used."983 The AOSAT PSPC. pulse height data obtained in 256 pulse height channels were appropriately re-grouped. to improve the statistics., The $ROSAT$ PSPC pulse height data obtained in 256 pulse height channels were appropriately re-grouped to improve the statistics.984 X-ray spectra of IN. 1334.213759. from the two observations thus obtained are shown in Figure 4., X-ray spectra of RX J1334.2+3759 from the two observations thus obtained are shown in Figure 4.985 We used the ASPEC (Version 11.0) spectral. analysis, We used the XSPEC (Version 11.0) spectral analysis986being “source free’. show a higher soft X-ray flux on average compared to the flux from the same regions at the re-projected blank sky observation.,"being `source free', show a higher soft X-ray flux on average compared to the flux from the same regions at the re-projected blank sky observation."987 At the same time the hard X-ray count rate measured for the same regions shows no offset. confirming that the diffuse source emission has fully covered the chip.," At the same time the hard X-ray count rate measured for the same regions shows no offset, confirming that the diffuse source emission has fully covered the chip."988 The background estimated from surrounding chips cannot be directly used since they are all front-illuminated., The background estimated from surrounding chips cannot be directly used since they are all front-illuminated.989 Thus a background image from blank sky observations with energies from 0.3-2.0 keV was used for the spatial analysis., Thus a background image from blank sky observations with energies from 0.3-2.0 keV was used for the spatial analysis.990 Point sources were detected using the CIAO wavelet detection algorithm and replaced by their local surrounding mean counts before any spatial analysis was carried out., Point sources were detected using the CIAO wavelet detection algorithm and replaced by their local surrounding mean counts before any spatial analysis was carried out.991 Vignetting and other sensitivity variations were corrected for. using an exposure map appropriate for the energy band.," Vignetting and other sensitivity variations were corrected for, using an exposure map appropriate for the energy band."992 We extracted ACIS spectra in successive circular annuli. in the energy range limited to 0.5-2.0 keV. excluding point sources.," We extracted ACIS spectra in successive circular annuli, in the energy range limited to 0.5-2.0 keV, excluding point sources."993 The background was chosen for each annulus separately. to account for vignetting. and from the same region on the re-projected blank sky observations to account for local variations within the chip.," The background was chosen for each annulus separately, to account for vignetting, and from the same region on the re-projected blank sky observations to account for local variations within the chip."994 The spectra were fit with absorbed hot plasma models and the APEC (Smithetal.2001). model was found to give the best fit., The spectra were fit with absorbed hot plasma models and the APEC \citep{smith01} model was found to give the best fit.995 A fixed hydrogen column density of μι=0.051077em7 was included in the model to account for Galactic absorption.," A fixed hydrogen column density of $N_{H, gal}=0.08\times 10^{22}\, {\rm cm}^{-2}$ was included in the model to account for Galactic absorption."996 An integrated spectrum was first extracted from a large region. of 3 aremin radius.," An integrated spectrum was first extracted from a large region, of 3 arcmin radius."997 A mean temperature of 0.66 keV and an unabsorbed flux of 0.1751077 erg  7 (05-2 keV) was found., A mean temperature of 0.66 keV and an unabsorbed flux of $0.175\times10^{-12}$ erg $^{-1}$ $^{-2}$ (0.5-2 keV) was found.998 However. the fit was poor with a reduced of about 1.5.," However, the fit was poor with a reduced of about 1.5."999 This high from a single temperature fit is not surprising. since we will see below that a strong temperature gradient is present in the X-ray emitting plasma.," This high from a single temperature fit is not surprising, since we will see below that a strong temperature gradient is present in the X-ray emitting plasma."1000 Absorption intrinsic to the source was also required by the fit. with a value of 0.1H(-E0.02)107?em7," Absorption intrinsic to the source was also required by the fit, with a value of $0.14(\pm0.02)\times 10^{22}\, {\rm1001cm}^{-2}$."1002 The diffuse 0.3-2 keV X-ray emission. (see Fig 29) shows a rather relaxed X-ray morphology in the core of the group.," The diffuse 0.3-2 keV X-ray emission, (see Fig \ref{fig2}) ), shows a rather relaxed X-ray morphology in the core of the group."1003 For the production of this image. point sources were removed. the background subtracted. and the 5S binned image adaptively smoothed and corrected using the exposure map.," For the production of this image, point sources were removed, the background subtracted, and the $8\times8$ binned image adaptively smoothed and corrected using the exposure map."1004 The X-ray contours in the central region are circularly symmetric but the outer region shows a moderate NE-SW elongation., The X-ray contours in the central region are circularly symmetric but the outer region shows a moderate NE-SW elongation.1005 To study any possible substructure and its signiticance. a hardness ratio map was obtained by dividing the adaptively smoothed diffuse emission image in a hard band by that in a soft band.," To study any possible substructure and its significance, a hardness ratio map was obtained by dividing the adaptively smoothed diffuse emission image in a hard band by that in a soft band."1006 The soft image contains photons with energies from 400 eV to 900 eV. The hard image consisted of photons with energv ranging from 900 eV to 2 keV. We replaced the point sources by their surrounding counts before smoothing the hard band S binned image using the esmoolh task in., The soft image contains photons with energies from 400 eV to 900 eV. The hard image consisted of photons with energy ranging from 900 eV to 2 keV. We replaced the point sources by their surrounding counts before smoothing the hard band $8\times8$ binned image using the $csmooth$ task in.1007. The same scale map is then used to smooth the soft band image as well as the associated background and exposure map images (i.e. all these are smoothed in an identical way)., The same scale map is then used to smooth the soft band image as well as the associated background and exposure map images (i.e. all these are smoothed in an identical way).1008 Before deriving the hardness ratio map. the corresponding backgrounds were subtracted from the soft and hard images.," Before deriving the hardness ratio map, the corresponding backgrounds were subtracted from the soft and hard images."1009 The hardness ratio map is limited to the central 3 aremin radius. Fig. 3..," The hardness ratio map is limited to the central 3 arcmin radius, Fig. \ref{fig3},"1010 due to the limited S/N. The relatively uniform and circularly symmetric distribution of the hardness ratio suggests that. the system is relatively relaxed and has not experienced a violent or strong perturbation recently. as caused by mergers and starforming activity.," due to the limited S/N. The relatively uniform and circularly symmetric distribution of the hardness ratio suggests that the system is relatively relaxed and has not experienced a violent or strong perturbation recently, as caused by mergers and starforming activity."1011 Small scale features at the centre of the map are found to be at the noise level and most likely due to the unresolved point sources or inaccurate repacement by local diffuse emission after point source removal., Small scale features at the centre of the map are found to be at the noise level and most likely due to the unresolved point sources or inaccurate replacement by local diffuse emission after point source removal.1012 We therefore assume that the system is spherically symmetric anc is in hydrostatic equilibrium., We therefore assume that the system is spherically symmetric and is in hydrostatic equilibrium.1013raction of the explosion euecrev than in GRDs/XREs. naking them detectable in the radio oulv for very nearby events.,"fraction of the explosion energy than in GRBs/XRFs, making them detectable in the radio only for very nearby events."1014 Current observational constraints (6.8.7) on hese radio sources could be satisfied by low-energv wildly relativistic jets and/or relatively low values for external density or shock nmücrophnysies parameters., Current observational constraints \citep[e.g.][]{soderberg2006apj} on these radio sources could be satisfied by low-energy mildly relativistic jets and/or relatively low values for external density or shock microphysics parameters.1015 It ws also been proposed (2?) that some Ib/e SNe are xoducimg relativistic GBRD/XRE jets that point away roni our line of sight aud are thus not detected at carly times in optical and N-ravs. but could be detected in radio bands after a few months to wears (orphan afterelows).," It has also been proposed \citep[][]{paczynski2001aca,granot2003apj} that some Ib/c SNe are producing relativistic GRB/XRF jets that point away from our line of sight and are thus not detected at early times in optical and X-rays, but could be detected in radio bands after a few months to years (orphan afterglows)."1016 After initially unsuccessful efforts to find this latter type of afterelow (e.g.72).. there is currently evidence that two of these eveuts 22007er (7) and SN22009b0b. (7) ave been observed. although their radio enuission was already detected within davs of the initial explosions.," After initially unsuccessful efforts to find this latter type of afterglow \citep[e.g.][]{soderberg2006apj}, there is currently evidence that two of these events, 2007gr \citep[][]{paragi2010nature} and 2009bb \citep[][]{soderberg2010nature} have been observed, although their radio emission was already detected within days of the initial explosions."1017 Iu. the case of the very. nearby 22007er (dzz11 Alpe). mildly-relativistic expansion of the radio source was iueasured using αν Long Daseline Tuterferometry (VLDBI) observatious. while for the radio-Inninous 22009bb (at dzzLO Mpc) a similar expansion speed of the ejecta was inferred from modeling the broadband radio data.," In the case of the very nearby 2007gr $d \approx 11\;$ Mpc), mildly-relativistic expansion of the radio source was measured using Very Long Baseline Interferometry (VLBI) observations, while for the radio-luminous 2009bb (at $d \approx 40\;$ Mpc) a similar expansion speed of the ejecta was inferred from modeling the broadband radio data."1018 We note that the interred energv in (mnildlv) relativistic ejecta was siguificantlv lavecr in SN22000bb than in 22007er. ~107? and ERU org. respectively.," We note that the inferred energy in (mildly) relativistic ejecta was significantly larger in 2009bb than in 2007gr, $\sim 10^{49}$ and $\sim 10^{46}$ erg, respectively."1019 The discovery of a bright N-ray transicut in the nearby (dz27 AIpe) galaxy 22770 (2)... which was later identified as a type Ib SN. has provided au unprecedented opportunity to study the sequence: GRBs NRFs normal core-collapse SNe.," The discovery of a bright X-ray transient in the nearby $d \approx 27\;$ Mpc) galaxy 2770 \citep[][]{berger2008atel}, which was later identified as a type Ib SN, has provided an unprecedented opportunity to study the sequence: GRBs – XRFs – normal core-collapse SNe."1020 The very carly discovery aud close proxiuitv of this source has enabled a multitude of observations across the whole electromagnetic spectrum., The very early discovery and close proximity of this source has enabled a multitude of observations across the whole electromagnetic spectrum.1021 Follow-up observations of 00850109. detected the ransicut at optical wavelengths (2? ).. and also in racio wave bands with the Verv Large Array (VLA:7) and Westerbork Svuthesis Radio Telescope (WSRT:?)..," Follow-up observations of 080109 detected the transient at optical wavelengths \citep[][]{deng2008gcn7160,thone2008gcn7161}, , and also in radio wave bands with the Very Large Array \citep[VLA;][]{soderberg2008gcn7178} and Westerbork Synthesis Radio Telescope \citep[WSRT;][]{vanderhorst2008gcn7190}."1022 As he optical counterpart rapidly brightened. spectroscopic observations revealed broad features possibly related to an emereiue supernova. 22008D. and the source was first classified as a type Ic SN. but later re-classified to ype Ib based on the emerging presence of helium in he spectra (22??3..," As the optical counterpart rapidly brightened, spectroscopic observations revealed broad features possibly related to an emerging supernova, 2008D, and the source was first classified as a type Ic SN, but later re-classified to type Ib based on the emerging presence of helium in the spectra \citep[][]{soderberg2008nature,mazzali2008science,malesani2009apj,modjaz2008arxiv}."1023 All the SNe associated with CRBs. aud the two SNe that have shown wildly relativistic expansion. are type Ic SNe. but some type Ib SNe are also expected to produce (at least mildly) relativistic outflows (?) although this has not been observed so far.," All the SNe associated with GRBs, and the two SNe that have shown mildly relativistic expansion, are type Ic SNe, but some type Ib SNe are also expected to produce (at least mildly) relativistic outflows \citep[][]{macfadyen1999apj} although this has not been observed so far."1024 The uature of the N-ray outburst observed iu SN22008D has been extensively discussed. in the literature., The nature of the X-ray outburst observed in 2008D has been extensively discussed in the literature.1025 It has been claimed that the outburst is a weak NRF caused by a inddlv relativistic outflow (?27)..," It has been claimed that the outburst is a weak XRF caused by a mildly relativistic outflow \citep{li2008mnras,mazzali2008science}."1026 However. there are counter-clainis that we have wituessed the X-ray cussion from a supernova shock breakout (?7?77).. caused bw the transition from a radiation donunated to a collisionless shock.," However, there are counter-claims that we have witnessed the X-ray emission from a supernova shock breakout \citep{soderberg2008nature,chevalier2008apj,wang2008aipc}, caused by the transition from a radiation dominated to a collisionless shock."1027 ? lave collected 2 mouths of Chigh-frequeney) radio observations from this source aud couchided that the outflow is freely expanding at non-relativistic velocities., \citet{soderberg2008nature} have collected 2 months of (high-frequency) radio observations from this source and concluded that the outflow is freely expanding at non-relativistic velocities.1028 They also argued that the flux of the N-rav outburst im combination with the nou-detectious at UV/optical wavelengths is inconsistent with a (iuildlv) relativistic outflow., They also argued that the flux of the X-ray outburst in combination with the non-detections at UV/optical wavelengths is inconsistent with a (mildly) relativistic outflow.1029 Ou the other haud. the measured variable optical polarization sugeests al axisvnunetre aspherical expansion with variable eccentricity (?).. as expected in the collapsar model (?)..," On the other hand, the measured variable optical polarization suggests an axisymmetric aspherical expansion with variable eccentricity \citep{gorosabel2008arxiv}, as expected in the collapsar model \citep{woosley1993apj}."1030 This asphericitv has also been inferred. frou optical spectra of 22008D (?).., This asphericity has also been inferred from optical spectra of 2008D \citep{modjaz2008arxiv}.1031 We note. however. that the optical aud radio enüssioü are not unanbiguouslv conineg from the sanie emission region.," We note, however, that the optical and radio emission are not unambiguously coming from the same emission region."1032 Iu this paper we present the results from our extensive radio follow-up campaigus of 22008D with the WSRT and the (απ Aletrewave Badio Telescope (GAIRT)., In this paper we present the results from our extensive radio follow-up campaigns of 2008D with the WSRT and the Giant Metrewave Radio Telescope (GMRT).1033 Combined with VLA and CARMA data from ?.. we study. these well-sampled elt curves up to 17 mouths after the initial explosion. across a broad frequency range. frou 325 MIIz to 95 GIIz.," Combined with VLA and CARMA data from \citet{soderberg2008nature}, we study these well-sampled light curves up to 17 months after the initial explosion, across a broad frequency range, from 325 MHz to 95 GHz."1034 We also discuss here the Type Ib SN22007uv. (77) that went off in the sale Galaxy teu days before the discovery of 22008D. Besides radio photometry measurements. we present the results of two epochs of global VLBI observations of 22008D and discuss their implications for the nature of the source.," We also discuss here the Type Ib 2007uy \citep[][]{nakano2008iauc,blondin2008cbet} that went off in the same galaxy ten days before the discovery of 2008D. Besides radio photometry measurements, we present the results of two epochs of global VLBI observations of 2008D and discuss their implications for the nature of the source."1035 Finally. CO observations with the Arizona Radio Observatory 12au telescope and lower resolution WSRT ineasurements provide the opportunity to study 22770. the host galaxy of both SNe. in detail.," Finally, CO observations with the Arizona Radio Observatory 12-m telescope and lower resolution WSRT measurements provide the opportunity to study 2770, the host galaxy of both SNe, in detail."1036" Tn particular. with two Type Ib SNe occiumiung in the sale galaxy within 10 days. aud three of those SNe within 10 vears (thethirdouebeiugSN1999eh:2?).. we discuss the properties of the molecular gas in 22770 compared to other galaxies. which make this galaxy a οποιο ""SN factorv (asalsodiscussedim?).."," In particular, with two Type Ib SNe occurring in the same galaxy within 10 days, and three of those SNe within 10 years \citep[the third one being SN\,1999eh;][]{hurst1999iauc,jha1999iauc}, we discuss the properties of the molecular gas in 2770 compared to other galaxies, which make this galaxy a possible “SN factory” \citep[as also discussed in][]{thone2009apj}. ."1037 All the measureients of the SNe and their lost ealaxy are prescuted in Section 77.. the modeling aud interpretation of the SNe data in Section ??.. aud of the jost galaxy data in Section ??..," All the measurements of the SNe and their host galaxy are presented in Section \ref{sec:obs}, the modeling and interpretation of the SNe data in Section \ref{sec:natsne}, and of the host galaxy data in Section \ref{sec:host}."1038 In Section 2°? we discuss he tuplications for future observations of such SNe and their host ealaxics with the Low Frequency Array (LOFAR). the first new ecueration meter waveleneth clescope.," In Section \ref{sec:lofar} we discuss the implications for future observations of such SNe and their host galaxies with the Low Frequency Array (LOFAR), the first new generation meter wavelength telescope."1039 We cud with our conclusions in Section ??.., We end with our conclusions in Section \ref{sec:conclusions}.1040 We have performed observations of 22008D. and 220070 with the WSRT at 1.1.2.8. L8 aud 8.1 GIIz.," We have performed observations of 2008D and 2007uy with the WSRT at 1.4, 2.3, 4.8 and 8.4 GHz."1041 We used the Multi Frequency Frout Ends (7) in coubination with the IVC|DZB back in coutimmun mode. with a bandwidth of 8x20 MIIz at all observiug frequencics.," We used the Multi Frequency Front Ends \citep{tan1991} in combination with the IVC+DZB back in continuum mode, with a bandwidth of 8x20 MHz at all observing frequencies."1042 Cain aud phase calibrations were performed with the calibrator 3€ 286 for το observations. although for a few epochs 3C 18 was used.," Gain and phase calibrations were performed with the calibrator 3C 286 for most observations, although for a few epochs 3C 48 was used."1043 The observations were analyzed using the Multichaunel huage Reconstruction nage Analysis and Display (MIRIAD:7) software package. except for the WSRT data hat were obtained duriug VLBI observations. whichwere analyzed with the Astronomical Bhuaee Processing System (AIPS:?)..," The observations were analyzed using the Multichannel Image Reconstruction Image Analysis and Display \citep[MIRIAD;][]{sault1995} software package, except for the WSRT data that were obtained during VLBI observations, whichwere analyzed with the Astronomical Image Processing System \citep[AIPS;][]{wells1985}."1044 All the results of our observations. for )otli 22008D and 220070. are detailed iu Table 1:: he resulting light curves are shown im Figures 1 aud 2..," All the results of our observations, for both 2008D and 2007uy, are detailed in Table \ref{table:wsrtgmrtdata}; the resulting light curves are shown in Figures \ref{fig:SN2008Dlcs} and \ref{fig:SN2007uylcs}."1045 The first observation at Ls CGIIz. at ~6.5 davs after he Nav detection. was reported asa detection of 22008D (?7).. but the reported flux was siguificautlv üsgher than theone in Table Ἐν," The first observation at 4.8 GHz, at $\sim$ 6.5 days after the X-ray detection, was reported asa detection of 2008D \citep{vanderhorst2008gcn7190}, , but the reported flux was significantly higher than theone in Table \ref{table:wsrtgmrtdata}."1046",This discrepancy was caused by a contribution of the host galaxy. which is"," This discrepancy was caused by a contribution of the host galaxy, which is"1047"adopted a teiiplate of the form. eiven in equation (8) with my=2.0. v2=1.2 aud 4,=1500 to model the IR. clustered power when extracting cosmolosical information from the small-scale SPT 150 Giz CAIB power spectrum.","adopted a template of the form given in equation (8) with $n_1=2.0$ , $n_2=1.2$ and $\ell_{\textrm b}=1500$ to model the IR clustered power when extracting cosmological information from the small-scale SPT 150 GHz CMB power spectrum."1048" We fd that this form is uot a good fit to the data. with \?/d.o.f.=22/7 when we fit to the 217 GIIz spectra with n. ny and (, fixed to the above values and using τήο priors on the clusterimg amplitude aud shot-oise."," We find that this form is not a good fit to the data, with $\chi^2/\textrm{d.o.f.}=22/7$ when we fit to the 217 GHz spectrum with $n_1$, $n_2$ and $\ell_{\textrm b}$ fixed to the above values and using uniform priors on the clustering amplitude and shot-noise."1049 II12 found that the BLAST and BLAST « ACT data are well-fit by assuming the IR ealaxy power spectrum in equation (5) is eiven by where b is the linear bias factor aud Pp is the near dark matter power spoectrun. and using the Bll predictions for the redshift-distzibution of the fiux. d/d:.," H12 found that the BLAST and BLAST $\times$ ACT data are well-fit by assuming the IR galaxy power spectrum in equation (5) is given by where $b$ is the linear bias factor and $P_{\textrm{DM}}$ is the linear dark matter power spectrum, and using the B11 predictions for the redshift-distribution of the flux, $\textrm{d}S/\textrm{d}z$."1050 PLL also found that the data from cach baud are wellit by this model (also using the Bil d5/d:) if uo prior is euforced on the shot-noise levels., P11 also found that the data from each band are well-fit by this model (also using the B11 $\textrm{d}S/\textrm{d}z$ ) if no prior is enforced on the shot-noise levels.1051 The bias levels are uot consisteut: when fitting for a suele. redshift-indepeudent value of 5. IT12 found 5.040. Lwhereas P11 found 5=2.18+0.11 for the 515 CGIIz spectiua.," The bias levels are not consistent; when fitting for a single, redshift-independent value of $b$, H12 found $b=5.0\pm0.4$ whereas P11 found $b=2.18\pm0.11$ for the 545 GHz spectrum."1052 The linear bias model is strouslv rejected by the combined aud BLAST data even without anv VArot-noise prior., The linear bias model is strongly rejected by the combined and BLAST data even without any shot-noise prior.1053 Fitting a linear bias model with sinele-value bias to the 857 CIIz aud BLAST 350 jam spectra together. usine the Bll d5/d: aud iultiplviug 1ο BLAST data points by a color correction factor of 1.07 (see Section [). vields \7/d.o.f.=39/16.," Fitting a linear bias model with single-value bias to the 857 GHz and BLAST 350 $\mu$ m spectra together, using the B11 $\textrm{d}S/\textrm{d}z$ and multiplying the BLAST data points by a color correction factor of 1.07 (see Section 4), yields $\chi^2/\textrm{d.o.f.}=39/16$."1054 For a joint fit to the 515 GIIz and BLAST 500 jan spectra we fud \?/d.o.f.=13/16., For a joint fit to the 545 GHz and BLAST 500 $\mu$ m spectra we find $\chi^2/\textrm{d.o.f.}=43/16$.1055 This result is driven x the shape of the linear matter power spectrum rather ian the choice of d5/dz: we repeated the fitting using jo predictions of rather thau Blt aud found \?/d.o.f. of 36/16 aud 12/16 for the sh? Gz 350 gan and 515 Cz 500 jun spectra respectively., This result is driven by the shape of the linear matter power spectrum rather than the choice of $\textrm{d}S/\textrm{d}z$: we repeated the fitting using the predictions of rather than B11 and found $\chi^2/\textrm{d.o.f.}$ of 36/16 and 42/16 for the 857 GHz / 350 $\mu$ m and 545 GHz / 500 $\mu$ m spectra respectively.1056/ Neither the nor BLAST/ data alone were able to rule out the Lucar bias model without a shot- prior because of the liiüted angular scales probed., Neither the nor BLAST data alone were able to rule out the linear bias model without a shot-noise prior because of the limited angular scales probed.1057 similarly fouud that the linear bias model could not be ruled out using only SPT data from (z2000., similarly found that the linear bias model could not be ruled out using only SPT data from $\ell\gtrsim2000$.1058 Table 5 shows the mareinalized deeree of clustering correlation. fear. for cach cross-spectimm included iu our fitting. along with the separation Av letween the two bands.," Table 5 shows the marginalized degree of clustering correlation, $f_{\textrm{corr}}$, for each cross-spectrum included in our fitting, along with the separation $\Delta\nu$ between the two bands."1059 The BLAST « BLAST cross-spectra are consistent with correlation., The BLAST $\times$ BLAST cross-spectra are consistent with correlation.1060 Our couchisious reearding the deerees of correlation between the ACT and BLAST hands are limited bv the data quality., Our conclusions regarding the degrees of correlation between the ACT and BLAST bands are limited by the data quality.1061 A decrease in correlation with increasing band separation would be consistent with the sources h|iug at a ranec of redshifts. with the ligher-redshift sources being of ercater relative importance at the longer waveleneths2000).. however the current data are not of sufficient quality to confirm this.," A decrease in correlation with increasing band separation would be consistent with the sources lying at a range of redshifts, with the higher-redshift sources being of greater relative importance at the longer wavelengths, however the current data are not of sufficient quality to confirm this."1062 Two of the mareimalized mean feoy values lie more than lo above unity. while noue lie more than lo below.," Two of the marginalized mean $f_{\rm corr}$ values lie more than $\sigma$ above unity, while none lie more than $\sigma$ below."1063" Aleasuriuge fea,2lo omav indicate that the angular scale dependence of the cross-spectra clusteringpowcr is not described by the same smgledudex power law as the auto-spectra clusterius. since there is no plysical explanation for a correlation iu excess ofLOO%."," Measuring $f_{\rm corr}>1$ may indicate that the angular scale dependence of the cross-spectra clusteringpower is not described by the same single-index power law as the auto-spectra clustering, since there is no physical explanation for a correlation in excess of."1064. This could be the case even with no worsening of the 47. due to the limited aneular scales probed by the BLAST. & ACT data.," This could be the case even with no worsening of the $\chi^2$ , due to the limited angular scales probed by the BLAST $\times$ ACT data."1065 To test that this is uot having a siguificaut effect on our clustering template. we repeated the MICAIC fitting described iu Section 3.2 using ouly the auto-spectrum data.," To test that this is not having a significant effect on our clustering template, we repeated the MCMC fitting described in Section 3.2 using only the auto-spectrum data."1066 We found that the values of η. οἱ and fy change by <0.50 compared to the ft with the cross-spectra included.," We found that the values of $n$, $\beta$ and $I_0$ change by $<0.5\sigma$ compared to the fit with the cross-spectra included."1067 We couclude that. while the data may o hinting that the cross-spectrum clustering power has a differeut shape to the auto-spectium clustering. this is rot significantly biasing our template.," We conclude that, while the data may be hinting that the cross-spectrum clustering power has a different shape to the auto-spectrum clustering, this is not significantly biasing our template."1068 Further subnuucunnn cross-correlation studies (e.e.. « ACT / SPT) axe clearly required to provide nore insight iuto the distribution of the sources with redshift. to coustrain the angular scale cependeuce of the cross-spectrum clustering. aud to investigate row the clustering shape changes with increasing baud separation.," Further submm-mm cross-correlation studies (e.g., $\times$ ACT / SPT) are clearly required to provide more insight into the distribution of the sources with redshift, to constrain the angular scale dependence of the cross-spectrum clustering, and to investigate how the clustering shape changes with increasing band separation."1069 We lave found that a power-law model for the IR point source clustering is adequate to simultaucously fitPlanck. BLAST aud cross-corrclated BLAST / ACT power spectrum data over a broad range of frequency (150«ν1200 CIIz) aud augular scale (muultipole moment 100«(6 9000).," We have found that a power-law model for the IR point source clustering is adequate to simultaneously fit, BLAST and cross-correlated BLAST / ACT power spectrum data over a broad range of frequency $150<\nu<1200$ GHz) and angular scale (multipole moment $100<\ell<9000$ )."1070" We find the clustering power varies with angular scale as ("" with υ=1.2540.06 and that the SED of the clustering can be described as a iocified blackbody with enuüssvitv index ο)=2.2040.07 and effective temperature Tig=9.7 K. Our work does not rely on iu assuniptious regarding the phlivsical properties of the IR sources (host halos. redshift distribution. ete.)."," We find the clustering power varies with angular scale as $\ell^{-n}$ with $n=1.25\pm0.06$ and that the SED of the clustering can be described as a modified blackbody with emissivity index $\beta=2.20\pm0.07$ and effective temperature $T_{\rm eff}=9.7$ K. Our work does not rely on any assumptions regarding the physical properties of the IR sources (host halos, redshift distribution, etc.)."1071 As well as providing a simple template for use in CXMB foreground. subtraction. we have established that the and BLAST / BLAST « ACT data sets appearcompatible when bandpass filters. flux cut aud calibration are accounted for. as described m Sections 2 and 3.," As well as providing a simple template for use in CMB foreground subtraction, we have established that the and BLAST / BLAST $\times$ ACT data sets appearcompatible when bandpass filters, flux cut and calibration are accounted for, as described in Sections 2 and 3."1072 We imake predictions for the IR clustering power for the Planck. ACT aud SPT CAB bands: our predictions for the ACT and SPT bands at around 150 aud220 νε are fully consistent with existing lmuecasurements 2011)..," We make predictions for the IR clustering power for the , ACT and SPT CMB bands; our predictions for the ACT and SPT bands at around 150 and220 GHz are fully consistent with existing measurements ."1073Models of atmospheres of protoplanets that do not fill in the Roche lobe have been considered by different authors.,Models of atmospheres of protoplanets that do not fill in the Roche lobe have been considered by different authors.1074 Papaloizou Nelson (2005) have constructed models assuming that the only source of energv for the atmosphere is that produced by (he gravitational contraction of the eas., Papaloizou Nelson (2005) have constructed models assuming that the only source of energy for the atmosphere is that produced by the gravitational contraction of the gas.1075 This corresponds to (he runaway gas accretion phase. ie. alter the core has reached (he critical mass. as the Iuminosity due to the accretion of planetesimadls is (hen negligible.," This corresponds to the runaway gas accretion phase, i.e. after the core has reached the critical mass, as the luminosity due to the accretion of planetesimals is then negligible."1076 The radius of the protoplanet atinosphere mav become smaller (han the Roche lobe radius in (his phase if the protostellar disc cannot supply gas to the atmosphere rapidly enough., The radius of the protoplanet atmosphere may become smaller than the Roche lobe radius in this phase if the protostellar disc cannot supply gas to the atmosphere rapidly enough.1077 Such models have also been considered by Mordasini et al. (, Such models have also been considered by Mordasini et al. (10782011).,2011).1079 Lissaner et al. (, Lissauer et al. (10802009. see also Movshowvitz et al.,"2009, see also Movshovitz et al."1081 2010) have caleulated models of aàmospheres prior to the runaway gas accretion pliase by assuming (hat the surface radius is either (he Bondi radius or a fraction of the Roche lobe radius depending on whether the Bondi radius is s1iadler or larger than the Roche lobe radius., 2010) have calculated models of atmospheres prior to the runaway gas accretion phase by assuming that the surface radius is either the Bondi radius or a fraction of the Roche lobe radius depending on whether the Bondi radius is smaller or larger than the Roche lobe radius.1082 This is based on the results of three dimensional numerical simulations indicating that only gas within about one quarter of the Roche lobe remains bound to the protoplanet., This is based on the results of three dimensional numerical simulations indicating that only gas within about one quarter of the Roche lobe remains bound to the protoplanet.1083 The gas accretion rate onto the atmosphere is then (taken (o be the value required (o match the outer radius to the desired value., The gas accretion rate onto the atmosphere is then taken to be the value required to match the outer radius to the desired value.1084 To calculate the volume of gas which is bound to the protoplanet from the numerical simulations. they identifv the trajectories of (racer particles that are trapped. inside (he gravitational potential of the protoplanet.," To calculate the volume of gas which is bound to the protoplanet from the numerical simulations, they identify the trajectories of tracer particles that are trapped inside the gravitational potential of the protoplanet."1085 We comment (hat (his approach may severely uuclerestimate the amount of gas that can be accretecl as in reality particles can collide with the protoplanet ancl become bound if energy is dissipated into shocks. as in Che process of star formation.," We comment that this approach may severely underestimate the amount of gas that can be accreted as in reality particles can collide with the protoplanet and become bound if energy is dissipated into shocks, as in the process of star formation."1086 Hf (he protoplanet fills in a signilicant part of its Roche lobe. collisions may become frequent.," If the protoplanet fills in a significant part of its Roche lobe, collisions may become frequent."1087 Three dimensional numerical simulations of protoplanets of a few earth masses in protoplanetary disces usually start with a core surrounded by a static atmosphere (hat extends at most up to the Bondi radius., Three dimensional numerical simulations of protoplanets of a few earth masses in protoplanetary discs usually start with a core surrounded by a static atmosphere that extends at most up to the Bondi radius.1088 The Keplerian flow located within ancl around the Roche lobe. and in which the protoplanet is embedded. is (hen perturbed and material [rom (he disc rains down on the protoplanet [rom above and below. because the gravitational potential from the protoplanet is not balanced (DAngelo et al.," The Keplerian flow located within and around the Roche lobe, and in which the protoplanet is embedded, is then perturbed and material from the disc rains down on the protoplanet from above and below, because the gravitational potential from the protoplanet is not balanced (D'Angelo et al."1089 2003. Paarclekooper Mellema 2008. Machida et al.," 2003, Paardekooper Mellema 2008, Machida et al."1090 2010)., 2010).1091 This results in the protoplanet being surrounded by a cold isothermal accretion flow., This results in the protoplanet being surrounded by a cold isothermal accretion flow.1092 In the present paper. we show that. if a core below the critical mass is surrounded. by an atmosphere that does not fill in entirely the Roche lobe aud is embedded in an accretion flow. it is going to evolve into a static model expanded to the Roche lobe radius on a very short timescale. even when the Dondi radius is smaller (han the Roche lobe radius.," In the present paper, we show that, if a core below the critical mass is surrounded by an atmosphere that does not fill in entirely the Roche lobe and is embedded in an accretion flow, it is going to evolve into a static model expanded to the Roche lobe radius on a very short timescale, even when the Bondi radius is smaller than the Roche lobe radius."1093 In (he calculations presented here. we include the accretion of planetesimals onto (he core.," In the calculations presented here, we include the accretion of planetesimals onto the core,"1094Soon after the launch of the N-vav Timing Explorer TE) in late 1995. observations with it of neutron star low-mass N-rav binaries revealed kilohertz quasi-periodic briehtuess oscillations (QPOs) in the accretion-powered. emission roni niv of these sources (sec. e.g. van der Rlis 2000 for a review).,"Soon after the launch of the X-ray Timing Explorer ) in late 1995, observations with it of neutron star low-mass X-ray binaries revealed kilohertz quasi-periodic brightness oscillations (QPOs) in the accretion-powered emission from many of these sources (see, e.g., van der Klis 2000 for a review)."1095 These oscillations have high frequencies (up to ~1300 Tz [vai Straaten et al., These oscillations have high frequencies (up to $\sim$ 1300 Hz [van Straaten et al.1096 2000]. the ehest-frequency astrophysical oscillations ever observed). laugh amzplitucles (up to ruis in the 2-60 keV band of the Proportional Counter Array on RANTE). aud veh coherences (with quality factors Q=vf/FWIIM>100 in many cases). and often appear as two (but no more) high-frequency oscillations iu a single power density spectrum.," 2000], the highest-frequency astrophysical oscillations ever observed), high amplitudes (up to rms in the 2-60 keV band of the Proportional Counter Array on ), and high coherences (with quality factors $Q\equiv \nu/{\rm FWHM}>100$ in many cases), and often appear as two (but no more) high-frequency oscillations in a single power density spectrum."1097 Beat-frequency models (BFAIs) were quickly proposed for this phenomenon (Strolunaver 11996: Miller. Lamb. Psaltis 1998).," Beat-frequency models (BFMs) were quickly proposed for this phenomenon (Strohmayer 1996; Miller, Lamb, Psaltis 1998)."1098 In these models. the higher-trequency of the wo oscillations is attributed to the orbital frequeney of gas at some special radius rear the star. aud the lower-frequency oscillation is a beat between this orbital yequency and the stellar spin frequency.," In these models, the higher-frequency of the two oscillations is attributed to the orbital frequency of gas at some special radius near the star, and the lower-frequency oscillation is a beat between this orbital frequency and the stellar spin frequency."1099 These models are consistent with many of the trends evident in the carly data. including the approximate coustaucy of he frequency difference between the two simultaucous kilolertz oscillations aud he close match in four sources of this frequency difference with the stellar spin yequency mferred from brightuess oscillations durius thermonuclear ταν bursts.," These models are consistent with many of the trends evident in the early data, including the approximate constancy of the frequency difference between the two simultaneous kilohertz oscillations and the close match in four sources of this frequency difference with the stellar spin frequency inferred from brightness oscillations during thermonuclear X-ray bursts."1100 It has. however. been established receutly that in several sources the frequency clifference isnot coustaut. aud deed cau vary by more than 50 Iz.," It has, however, been established recently that in several sources the frequency difference is constant, and indeed can vary by more than 50 Hz."1101 Moreover. tliis variation is systematic: the higher the lower peak frequency. the lower the frequency difference.," Moreover, this variation is systematic: the higher the lower peak frequency, the lower the frequency difference."1102 The explicability of this beliavior in the beat-frequeney picture has direct bearing on some of the most important infereuces drawn from the kilohertz QPOs., The explicability of this behavior in the beat-frequency picture has direct bearing on some of the most important inferences drawn from the kilohertz QPOs.1103 For example. ouly in BFAIs is the leveline-off of the frequency of both kilohertz," For example, only in BFMs is the leveling-off of the frequency of both kilohertz"1104We have fitted the continuum of all the five source spectra with two mmodels.,We have fitted the continuum of all the five source spectra with two models.1105 The first one represents the dominating TC component of the spectrum. while the second one provides deseription of the strong ~ | keV BB-like feature and. when observed. the hard X-ray emission above 30 keV. In our scenario (see also F07 and F08). the thermal ccomponent is an emission originated in the relatively cold (~3 keV) and optically thick (7~ 5) outer transition layer due to the Comptonization of cold disk seed photons (~ kkeV).," The first one represents the dominating TC component of the spectrum, while the second one provides description of the strong $\sim$ 1 keV BB-like feature and, when observed, the hard X-ray emission above 30 keV. In our scenario (see also F07 and F08), the thermal component is an emission originated in the relatively cold $(\sim$ keV) and optically thick $\tau\sim$ 5) outer transition layer due to the Comptonization of cold disk seed photons $\sim$ keV)."1106 This component ts rather stable as can be seen in Tab. 1.., This component is rather stable as can be seen in Tab. \ref{tab_fit}.1107 The second ccomponent originates in the innermost region of the system and is subjected to major changes (Tab. 1)., The second component originates in the innermost region of the system and is subjected to major changes (Tab. \ref{tab_fit}) ).1108 Hot BB photons (1 kkeV) are emitted by the neutron star surface and within the transition layer itself., Hot BB photons $\sim$ keV) are emitted by the neutron star surface and within the transition layer itself.1109 The spectral shape variability is due to the changing Comptonization of this seed photon population. mainly led by the accretion rate. as discussed Section 4.2..," The spectral shape variability is due to the changing Comptonization of this seed photon population, mainly led by the accretion rate, as discussed Section \ref{evolution}. ."1110" The difference in the measured values of &T, (a factor about three) strongly points in favour of two distinct populations of seed photons. which in turn translates in the existence of two physically separated regions."," The difference in the measured values of $\kts$ (a factor about three) strongly points in favour of two distinct populations of seed photons, which in turn translates in the existence of two physically separated regions."1111 Looking at the best-fit parameters in Tab. |..," Looking at the best-fit parameters in Tab. \ref{tab_fit},"1112 the source spectral stability of the persistent continuum is almost evident. in generalagreement with high-luminosity LMXBs," the source spectral stability of the persistent continuum is almost evident, in generalagreement with high-luminosity LMXBs"1113In summary the presence of CO» in ice mixtures with CO does not strongly affect the reactivity of CO with H-atoms.,In summary the presence of $_2$ in ice mixtures with CO does not strongly affect the reactivity of CO with H-atoms.1114" Figure 5. shows the differencespectrum for vs(C=O) at -1710 em""! of pure HCOOH ice bombarded with H-atoms as well as control experiments with bombardment of H> molecules at 12 K (foranoverviewofallinfraredfeaturesofHCOOHseeCyriac&Pradeep. 2005).."," Figure \ref{hcooh} shows the differencespectrum for $\nu_{\rm1115 S}$ (C=O) at $\sim$ 1710 $^{-1}$ of pure HCOOH ice bombarded with H-atoms as well as control experiments with bombardment of $_2$ molecules at 12 K \citep[for an overview of all infrared features of1116HCOOH see][]{cyriac2005}. ."1117 The growth of an infrared feature around ~ 1050 em! indicative for CH:(OH)- formation has not been observed (foranoverviewofthe 1994).," The growth of an infrared feature around $\sim$ 1050 $^{-1}$ indicative for $_2$ $_2$ formation has not been observed \citep[for an overview of the infrared features of1118CH$_2$(OH)$_2$ see][]{lugez1994}."1119 A decrease on the blue side of the vs(C=O) mode at 1710 em?! of HCOOH is seen at 1750 em! as well as an increase at 1730 em!'. which means that the overall HCOOH band profile changes slightly.," A decrease on the blue side of the $\nu_{\rm S}$ (C=O) mode at 1710 $^{-1}$ of HCOOH is seen at 1750 $^{-1}$ as well as an increase at 1730 $^{-1}$ , which means that the overall HCOOH band profile changes slightly."1120 At 1730 em'!. the C=O stretch for H:CO is located. but other features of H»CO. such as the 1500 em! band. are missing.," At 1730 $^{-1}$, the C=O stretch for $_2$ CO is located, but other features of $_2$ CO, such as the 1500 $^{-1}$ band, are missing."1121 The decrease corresponds to «0.1 ML derived from our caleulated RAIR band strength., The decrease corresponds to $<$ 0.1 ML derived from our calculated RAIR band strength.1122 These features are present in difference spectra for HCOOH ice bombarded with H-atoms at 12 and 40 K. A similar shift is seen for transmission infrared experiments with pure HCOOH ice that is heated to ~60 K (Bisschopetal..2007a)., These features are present in difference spectra for HCOOH ice bombarded with H-atoms at 12 and 40 K. A similar shift is seen for transmission infrared experiments with pure HCOOH ice that is heated to $\sim$ 60 K \citep{bisschop2007b}.1123. At the same time the vs(CH) and vs(OH) vibrational modes increase due to conversion of HCOOH in dimeric form to HCOOH organized in chains., At the same time the $\nu_{\rm S}$ (CH) and $\nu_{\rm S}$ (OH) vibrational modes increase due to conversion of HCOOH in dimeric form to HCOOH organized in chains.1124 In the RAIRS spectra these bands are also seen to increase., In the RAIRS spectra these bands are also seen to increase.1125 Furthermore. the same change in RAIR profile is found for HCOOH ices of 40 K. where H-atoms cannot stick any longer onto the surface. but can only collide.," Furthermore, the same change in RAIR profile is found for HCOOH ices of 40 K, where H-atoms cannot stick any longer onto the surface, but can only collide."1126 Since the ice has a temperature of 40 K. the reorganization of the ice is less and consequently the signal of the difference spectrum is smaller.," Since the ice has a temperature of 40 K, the reorganization of the ice is less and consequently the signal of the difference spectrum is smaller."1127 In conclusion. the RAIR data do suggest that some restructuring takes place in the surface but no reaction.," In conclusion, the RAIR data do suggest that some restructuring takes place in the surface but no reaction."1128 With TPD the masses of 48 (CH2(OH)>). 46 (HCOOH). 45 (HCOO). 44 (CO). 32/31 (CH30H). 30/29 (H2CO). and 28 amu (CO) have been monitored during warm-up.," With TPD the masses of 48 $_2$ $_2$ ), 46 (HCOOH), 45 (HCOO), 44 $_2$ ), 32/31 $_3$ OH), 30/29 $_2$ CO), and 28 amu (CO) have been monitored during warm-up."1129 No products are detected at 48. 32. 31. or 30 amu to upper limits of «0.01 ML. indicating that HCOOH is neither hydrogenated nor dissociated.," No products are detected at 48, 32, 31, or 30 amu to upper limits of $<$ 0.01 ML, indicating that HCOOH is neither hydrogenated nor dissociated."1130 Thus. consistent with the lack of a 1500 em! H;CO absorption feature in the RAIRS data. no evidence for H»CO formation is observed in the TPD experiment.," Thus, consistent with the lack of a 1500 $^{-1}$ $_2$ CO absorption feature in the RAIRS data, no evidence for $_2$ CO formation is observed in the TPD experiment."1131 The detected masses 45. 44. and 29 amu are assigned to HCOOH dissociating in the mass spectrometer. because the same relative mass ratios are seen for a TPD spectrum of pure HCOOH ice that is not bombarded by H-atoms.," The detected masses 45, 44, and 29 amu are assigned to HCOOH dissociating in the mass spectrometer, because the same relative mass ratios are seen for a TPD spectrum of pure HCOOH ice that is not bombarded by H-atoms."1132 We conclude that within the limits of our experimental set-up the reaction of HCOOH with H-atoms is not efficient at 12 K. Since no unambiguous evidence for HCOOH destruction in the ice is found. it is only possible to derive an upper limit on its reaction rate. presented in Table 3..," We conclude that within the limits of our experimental set-up the reaction of HCOOH with H-atoms is not efficient at 12 K. Since no unambiguous evidence for HCOOH destruction in the ice is found, it is only possible to derive an upper limit on its reaction rate, presented in Table \ref{dest}. ."1133 It is clear thatthe HCOOH destruction rates are below 2.3x107' em?s! as derived from the limit on the column density after | min of H-atom bombardment (see 3.2))., It is clear thatthe HCOOH destruction rates are below $\times$ $^{-17}$ $^2$$^{-1}$ as derived from the limit on the column density after 1 min of H-atom bombardment (see \ref{rate_sec}) ).1134 As for CO» these reaction rates are very low., As for $_2$ these reaction rates are very low.1135By fitting the data with an absorbed power-law (fixing E to 1.8 for two of them). these sources appear to be Compton-Thin.,"By fitting the data with an absorbed power-law (fixing $\Gamma$ to 1.8 for two of them), these sources appear to be Compton-Thin."1136 All the detected sources have relatively bright observed 2-10 keV fluxes (>3x1077 1j) and rest-frame 2-10 keV luminosities >107 erg sv! (prior to absorption corrections). placing them among luminous obscured QSOs.," All the detected sources have relatively bright observed 2-10 keV fluxes $>3\times 10^{-14}$ ) and rest-frame 2-10 keV luminosities $\gtrsim 10^{44}$ erg $^{-1}$ (prior to absorption corrections), placing them among luminous obscured QSOs."1137 As for the archival undetected object. SDSSJO85600. by assuming a 10 keV banddetection limit of 4 photons and a pure reflection spectrum in XSPEC). the corresponding upper limits to the observed 2-10 keV flux and rest frame luminosity. are 2x1077 erg E «!EN and 3»10? erg EM|. respectively.," As for the archival undetected object, SDSSJ085600, by assuming a 2-10 keV band detection limit of 4 photons and a pure reflection spectrum in XSPEC), the corresponding upper limits to the observed 2-10 keV flux and rest frame luminosity, are $2\times10^{-14}$ erg $^{-2}$ $^{-1}$ and $3\times10^{43}$ erg $^{-1}$, respectively."1138 A summary of the results obtained from the X-ray analysis of these 9 objects 1s given in Table 3.., A summary of the results obtained from the X-ray analysis of these 9 objects is given in Table \ref{xspec}.1139 Optical and X-ray spectra are shown in Fig., Optical and X-ray spectra are shown in Fig.1140 A+ We placed the new 9 SDSS obscured QSOs at z~] on the X/NeVdiagram (see Fig. 2)).," \ref{oxspec}1141 We placed the new 9 SDSS obscured QSOs at $z\sim 1$ on the X/NeV diagram (see Fig. \ref{xnevsdss}) )."1142 The archival undetected X-ray object SDSSJOS5600 shows X/NeV«15. which is strongly suggestive of CT absorption.," The archival undetected X-ray object SDSSJ085600 shows $<15$, which is strongly suggestive of CT absorption."1143 Indeed. most SDSS QSOs classified as CT by VIO do show X/NeV ratio below this value. while all those classified as Compton-Thin by V10 lie at X/NeV>100.," Indeed, most SDSS QSOs classified as CT by V10 do show X/NeV ratio below this value, while all those classified as Compton-Thin by V10 lie at $>100$."1144 The faint (18 X-ray counts) object SDSSJOSOS has X/NeV=45., The faint (18 X-ray counts) object SDSSJ0808 has $=45$.1145 This value has been observed in both CT and Compton-Thin AGN (see Fig. 1)).," This value has been observed in both CT and Compton-Thin AGN (see Fig. \ref{xnev_local}) ),"1146 but. based on the X-raye spectral analysis. SDSSJOSOS appears to have all the signatures of CT absorption.," but, based on the X-ray spectral analysis, SDSSJ0808 appears to have all the signatures of CT absorption."1147 We then classify SDSSJO85600 and SDSSJO808 as likely CT candidates., We then classify SDSSJ085600 and SDSSJ0808 as likely CT candidates.1148 The four objects with highest photon statistics have also the lowest absorption and the highest X/NeV ratios (X/NeV 150). and occupy a region of the diagram which is populated by the Compton-Thin QSOs in VIO and by the obscured tail of the broad line SDSS QSOs studied by?.," The four objects with highest photon statistics have also the lowest absorption and the highest X/NeV ratios $>150$ ), and occupy a region of the diagram which is populated by the Compton-Thin QSOs in V10 and by the obscured tail of the broad line SDSS QSOs studied by."1149. Finally. three objects do show X/NeV ratios in the range 20-60. which are typical of heavily obscured AGN. but are still consistent with Compton-Thin absorption.," Finally, three objects do show X/NeV ratios in the range 20-60, which are typical of heavily obscured AGN, but are still consistent with Compton-Thin absorption."1150 Both the X-ray and optical spectra of these sources would favor mild absorption. hence a Compton-Thin interpretation. although no clear cut classification can be made with the current data.," Both the X-ray and optical spectra of these sources would favor mild absorption, hence a Compton-Thin interpretation, although no clear cut classification can be made with the current data."1151 We note that the actual degree of obscuration of our two CT candidates SDSSJO8S5600 and SDSSJOSOS8 can only be confirmed by obtaining good-quality X-ray spectra through deeper observations., We note that the actual degree of obscuration of our two CT candidates SDSSJ085600 and SDSSJ0808 can only be confirmed by obtaining good-quality X-ray spectra through deeper observations.1152 Indeed. any selection method which simply relies on X-ray hardness ratios. or on the comparison between the measured. obscured X-ray emission (f any) with some other indicator of the intrinsic nuclear power (e.g. dust-reprocessed [R-emission. or high-tonization. narrow optical emission lines) can only provide an indirect way to select CT AGN.," Indeed, any selection method which simply relies on X-ray hardness ratios, or on the comparison between the measured, obscured X-ray emission (if any) with some other indicator of the intrinsic nuclear power (e.g. dust-reprocessed IR-emission, or high-ionization, narrow optical emission lines) can only provide an indirect way to select CT AGN."1153 Different caveats and limitations affect these different selection methods. like e.g. reddening in line-selected sources and contamination from star formation in IR-selected sources.," Different caveats and limitations affect these different selection methods, like e.g. reddening in line-selected sources and contamination from star formation in IR-selected sources."1154 The X/NeV vs X-ray absorption diagnostic diagram presented in Section 3 has been derived directly from observed values. without applying any corrections to either [Ne V] or X-ray emission.," The X/NeV vs X-ray absorption diagnostic diagram presented in Section 3 has been derived directly from observed values, without applying any corrections to either [Ne V] or X-ray emission."1155 The rationale behind is that X-rays come from the innermost nuclear regions and can be depressed by small scale («1 pe) absorption. while the [Ne V] is instead a good indicator of the intrinsic nuclear luminosity being 1) emitted on larger (kpe) scales. free from nuclear obscuration and 11) being Isotropic.," The rationale behind is that X-rays come from the innermost nuclear regions and can be depressed by small scale $<$ 1 pc) absorption, while the [Ne V] is instead a good indicator of the intrinsic nuclear luminosity being i) emitted on larger (kpc) scales, free from nuclear obscuration and ii) being isotropic."1156 We now discuss whether this two hypotheses are satisfied and what happens if they are not., We now discuss whether this two hypotheses are satisfied and what happens if they are not.1157 As for the first point. we note that significant extinetion towards the NLR is commonly observed in local AGN?).," As for the first point, we note that significant extinction towards the NLR is commonly observed in local AGN."1158. Under the simplest version of the AGN unification schemes. one would expect to observe similar properties in the NLR of both type-1 and type-2 AGN.," Under the simplest version of the AGN unification schemes, one would expect to observe similar properties in the NLR of both type-1 and type-2 AGN."1159 Therefore. if the dust content in the NLR of both AGN types ts the same. the same extinction correction should be applied to the measured values of the [Ne V] lines. which would simply shift towards lower X/NeV values all the datapoints in Fig.," Therefore, if the dust content in the NLR of both AGN types is the same, the same extinction correction should be applied to the measured values of the [Ne V] lines, which would simply shift towards lower X/NeV values all the datapoints in Fig."1160 | and2.. without altering the X/NeV vs Nj; trend.," \ref{xnev_local}1161 and \ref{xnevsdss}, without altering the X/NeV vs $N_H$ trend."1162" However. based on the results by?.. there is some evidence that the extinction to the NLR in local Seyfert 2s is somewhat higher than in Seyfert Is. the following relation holding for the median values: ae(e.g.~Ay""L5-1.020.5."," However, based on the results by, there is some evidence that the extinction to the NLR in local Seyfert 2s is somewhat higher than in Seyfert 1s, the following relation holding for the median values: $A_V^{Sy2}-A_V^{Sy1}\sim1.5-1.0=0.5$."1163 Using standard extinction curves?).. this translates into an average correction in the [Ne V] flux a factor of ~2.3 larger in Seyfert 2s than in Seyfert Is.," Using standard extinction curves, this translates into an average correction in the [Ne V] flux a factor of $\sim 2.3$ larger in Seyfert 2s than in Seyfert 1s."1164 Applying some extinetion correction to both the observed data and model curves would then produce a stronger shift towards lower X/NeV ratios in Seyfert 2s than in Seyfert Is. making the anti-correlation between X/NeV ratio and obscuration even more evident.," Applying some extinction correction to both the observed data and model curves would then produce a stronger shift towards lower X/NeV ratios in Seyfert 2s than in Seyfert 1s, making the anti-correlation between X/NeV ratio and obscuration even more evident."1165 At any rate. our main interest is on the possibility of applying the X/NeV diagnostic ratio to zo1 objects. for which it is often impossible to measure the extinetion to the NLR because of the lack of strong Balmer lines in the optical spectrum.," At any rate, our main interest is on the possibility of applying the X/NeV diagnostic ratio to $z\sim 1$ objects, for which it is often impossible to measure the extinction to the NLR because of the lack of strong Balmer lines in the optical spectrum."1166 Therefore. we simply do not apply any reddening correction to the measured [Ne V] fluxes.," Therefore, we simply do not apply any reddening correction to the measured [Ne V] fluxes."1167images.,images.1168 In Fig. 10..," In Fig. \ref{bp},"1169 we plot the distribution of €' for the red and blue passive galaxies respectively., we plot the distribution of $C$ for the red and blue passive galaxies respectively.1170 The blue passive galaxies clearly show late-type morphologies (C'«2.6:Stratevaetal.2001)... while their red counterparts mostly have values of concentration typical of spheroidals.," The blue passive galaxies clearly show late-type morphologies \citep[$C\! <\! 2.6$;][]{strateva}, while their red counterparts mostly have values of concentration typical of spheroidals."1171 The galaxies in the former category are similar to the passive blue spirals that have been studied by other authors (e.g.Gotoetal.2003a:Wild 2008).," The galaxies in the former category are similar to the passive blue spirals that have been studied by other authors \citep[\eg][]{goto03,wild}."1172. By analysing the 5yatially resolved long-slit spectroscopy of a few passive spiral galixies. selected from the SDSS. Ishigakietal.(2007). find that the Hs absorption lines are more prominent," By analysing the spatially resolved long-slit spectroscopy of a few passive spiral galaxies selected from the SDSS, \citet{ishigaki07} find that the $_\delta$ absorption lines are more prominent"1173sample should be very similar to that viewed by CDM onFermi. and this should enable a test of assumption (1).,"sample should be very similar to that viewed by GBM on, and this should enable a test of assumption (1)."1174 We use the 145 bursts that were observed bySif? BAT between January 2005 and June 2009 and have known redshift., We use the 145 bursts that were observed by BAT between January 2005 and June 2009 and have known redshift.1175 Figure | shows these bursts plotted as a function of redshift and BAT luence., Figure \ref{fig:batfluxdist} shows these bursts plotted as a function of redshift and BAT fluence.1176 A considerable amount of Llucnee rom these bursts arises largely from a few bright events: he brightest 10 per cent of bursts in the sample accounts or approximately 55 per cent of the Huence., A considerable amount of fluence from these bursts arises largely from a few bright events; the brightest 10 per cent of bursts in the sample accounts for approximately 55 per cent of the fluence.1177 While high energv Dux has only been seen from a handful of bright GRBs using EGRET. the fact that these bursts account or a large fraction of Uuence seen at lower energies means hat our assumption (2) should. be reasonable even if the »oportionalitv does not hold for faint bursts.," While high energy flux has only been seen from a handful of bright GRBs using EGRET, the fact that these bursts account for a large fraction of fluence seen at lower energies means that our assumption (2) should be reasonable even if the proportionality does not hold for faint bursts."1178 We have not included: LAT bursts such as GIUDOSOO916C in this analysis. although we do show where this event would have en in Fig.," We have not included LAT bursts such as GRB080916C in this analysis, although we do show where this event would have been in Fig."1179 P. based on its CDM Iuence., \ref{fig:batfluxdist} based on its GBM fluence.1180 We will discuss his event in the context of our emission. model in Section 1.1l.., We will discuss this event in the context of our emission model in Section \ref{080916C}.1181 Two recent papers have related the keV/MeV flux from GRBs to high. energy. emission. anc have estimated. the ratio of Hluences in these regimes.," Two recent papers have related the keV/MeV flux from GRBs to high energy emission, and have estimated the ratio of fluences in these regimes."1182 Le&Dermer(2009) estimated the count rate for GeV photons in the LAL based on the bursts seen by the EGRET spark chamber., \citet{le&dermer09} estimated the count rate for GeV photons in the LAT based on the bursts seen by the EGRET spark chamber.1183 The ινων Huence ratio inferred there varied from 0.05 to over 0.3., The $F_{LAT}/F_{BATSE}$ fluence ratio inferred there varied from 0.05 to over 0.3.1184 Based on the deadtime factors allecting some EGRET GRB observations these authors argue that a ratio of greater than 30. per cent between. BATSE and EGRET is reasonable., Based on the deadtime factors affecting some EGRET GRB observations these authors argue that a ratio of greater than 30 per cent between BATSE and EGRET is reasonable.1185 Ando.Nakar.&Sari(2008). made the assumption that there. is à. log-normal clistribution Of heavyflay8Beener{Pearse in the roughly 100 BATSE bursts that were in the field. of view of EGRET., \citet*{ando08} made the assumption that there is a log-normal distribution of $F_{GeV}/F_{MeV} \approx F_{EGRET}/F_{BATSE}$ in the roughly 100 BATSE bursts that were in the field of view of EGRET.1186 A maximum likelihood fit to the available data suggested a ratio of 0.003xPeayοι0.06., A maximum likelihood fit to the available data suggested a ratio of $0.003 \leq F_{GeV}/F_{MeV} \leq 0.06$.1187 Shehthy cüllerent assumptions about the high-energy spectrum. and energy range of the GeV emission. were used in each case., Slightly different assumptions about the high-energy spectrum and energy range of the GeV emission were used in each case.1188 Le Dermer used an index of -2. similar to the best [it to the EGRET bursts of -10.95 from Dingus(1995).. while the Anco et al.," Le Dermer used an index of -2, similar to the best fit to the EGRET bursts of -1.95 from \citet{dingus95}, while the Ando et al."1189 work assumed an spectral index of -2.4 in the EGRET enerev range., work assumed an spectral index of -2.4 in the EGRET energy range.1190 The latter noted that hardening the spectral index could have increased thei values for the [lux ratio cocllicient significantLv., The latter noted that hardening the spectral index could have increased their values for the flux ratio coefficient significantly.1191 In this paper we use the energy. fluences seen bySuciffBAT to predict GeV. emission using the ratio where {οΠΕ and fpyrs\e refer to the Huence (ine-integrated [lux over the duration of the event) across the EGRET spark chamber and BATS energy ranges. taken to be 100 AleV 5 GeV. and 20 keV 2 MeV. respectively.," In this paper we use the energy fluences seen byBAT to predict GeV emission using the ratio where $F_{\mbox{\scriptsize EGRET}}$ and $F_{\mbox{\scriptsize BATSE}}$ refer to the fluence (time-integrated flux over the duration of the event) across the EGRET spark chamber and BATSE energy ranges, taken to be 100 MeV – 5 GeV and 20 keV – 2 MeV respectively."1192 A constant spectral index is assumed to be valid from the EGRET energy range up to 100 GeV. We take an approach similar to Le Dermer and use a value of p0.1 for prompt phase emission. which does a reasonable job matching the high energy uence seen for the recent. GRB 0SQO16C. sce Section 4.1..," A constant spectral index is assumed to be valid from the EGRET energy range up to $>100$ GeV. We take an approach similar to Le Dermer and use a value of $\rho=0.1$ for prompt phase emission, which does a reasonable job matching the high energy fluence seen for the recent GRB 080916C, see Section \ref{080916C}."1193 This is higher than the range of values for p proposed by Ando et al.," This is higher than the range of values for $\rho$ proposed by Ando et al.,"1194 but that paper also assumed a softer high energy spectrum in deriving results., but that paper also assumed a softer high energy spectrum in deriving results.1195 Afterglows are. also a possible source. of. high-energv emission. though one that is even more poorly constrained than the prompt phase.," Afterglows are also a possible source of high-energy emission, though one that is even more poorly constrained than the prompt phase."1196 There. are. various mechanisms that have been hypothesized as possible sources of GeV οποίος., There are various mechanisms that have been hypothesized as possible sources of GeV photons.1197 | popular assumption invokes inverse-C'ompton upscattering of svnchrotron photons in the GRB outflow (SSC mechanism). although a variety of other sources are »ossible. such as SSC emission from the internal x-ray Lares seen in afterelows or Compton upscattering of these photons »w electrons accelerated in the external shock (lanetal.2008).," A popular assumption invokes inverse-Compton upscattering of synchrotron photons in the GRB outflow (SSC mechanism), although a variety of other sources are possible, such as SSC emission from the internal x-ray flares seen in afterglows or Compton upscattering of these photons by electrons accelerated in the external shock \citep{fan08}."1198. Limits from EGRET observations suggest a typical luence ratio of 0.01 to 0.1. and a spectral index of -1.5 to -2 (Andoetal.2008).," Limits from EGRET observations suggest a typical fluence ratio of 0.01 to 0.1, and a spectral index of -1.5 to -2 \citep{ando08}."1199. For afterglows. we assume a ratio of p—0.01 in this work.," For afterglows, we assume a ratio of $\rho=0.01$ in this work."1200" 1n converting the uence seen byοι BAT (15 150 keV) to BALTSE tus. we assume a common Bane (Bandetal.1993). functional form over the BATBATSE energy range: llere à, and o» are the low and high energy. indices. {ει is the break energy. ancl © is the Leaviside step function."," In converting the fluence seen by BAT (15 – 150 keV) to BATSE flux, we assume a common Band \citep{band93} functional form over the BAT–BATSE energy range: Here $\alpha_1$ and $\alpha_2$ are the low and high energy indices, $E_{br}$ is the break energy, and $\Theta$ is the Heaviside step function."

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