ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2. For practical purposes. we take (he sum of the spectra of two appropriate stars for comparison with the observed spectra al large frequencies.," For practical purposes, we take the sum of the spectra of two appropriate stars for comparison with the observed spectra at large frequencies."3 This process results in (wo speciral (vpes for both stars., This process results in two spectral types for both stars.4 Ireland&Kraus(2008) propose the stars have M1.5c0.5 spectral (vpes. consistent with an age of 4 Myrs.," \citet{Ireland} propose the stars have $M1.5\pm 0.5$ spectral types, consistent with an age of $4 \ Myrs$ ."5 Here. we use the evolutionarv tracks from Siessοἱal.(2000) given in 0.L.M.. steps.," Here, we use the evolutionary tracks from \citet{Siess} given in $0.1M_\odot$ steps."6 Taking (his restriction.," Taking this restriction,"7with the position angle of the isophotes at large scales (NGC 821. NGC 3610. NGC 3377. NGC 4621. and NGC 7619).,"with the position angle of the isophotes at large scales (NGC 821, NGC 3610, NGC 3377, NGC 4621, and NGC 7619)."8 While some of these ellipticals are known to be diskv (e.g.. NGC 3377. 4621). ellipticals are known to be triaxial and do not all rotate about their apparent minor axis (Ryden1992).," While some of these ellipticals are known to be disky (e.g., NGC 3377, 4621), ellipticals are known to be triaxial and do not all rotate about their apparent minor axis \citep{ryden92}."9. We return to this point in Section 5.2.., We return to this point in Section \ref{sec:ndisks}.10 In contrast to the early-(wpe. inactive sample. we do not detect many nuclear stellar disks in the early-iwpe active sample or in either of the late-tvpe subsamples.," In contrast to the early-type, inactive sample, we do not detect many nuclear stellar disks in the early-type active sample or in either of the late-type subsamples."11 The absence of nuclear disks in late-tvpe galaxies has already been reported in the work of (2002).. who studied WWFPC? F606W images of a sample of 38 spiral galaxies and did not find nuclear disks in barred galaxies or galaxies of Hubble (wpe later than Sb.," The absence of nuclear disks in late-type galaxies has already been reported in the work of \citet{pizzella02}, who studied WFPC2 F606W images of a sample of 38 spiral galaxies and did not find nuclear disks in barred galaxies or galaxies of Hubble type later than Sb."12 These dillerences [rom our inactive sample are most likely due to (he substantial cireumnuclear dust in the active sample and the Iate-tvpe galaxies.," These differences from our early-type, inactive sample are most likely due to the substantial circumnuclear dust in the early-type active sample and the late-type galaxies."13 In almost all cases. (he dust (and emission-line reeions) present in these galaxies make it difficult to identify nuclear stellar disks. particularly if (he dust and stellar disks have the same position angle.," In almost all cases, the dust (and emission-line regions) present in these galaxies make it difficult to identify nuclear stellar disks, particularly if the dust and stellar disks have the same position angle."14 Some notable exceptions are NGC 4111. which has a prominent stellar disk almost perpendicular to a smaller dust disk. and NGC 4026. which has a nuclear stellar disk and weak dust features.," Some notable exceptions are NGC 4111, which has a prominent stellar disk almost perpendicular to a smaller dust disk, and NGC 4026, which has a nuclear stellar disk and weak dust features."15 If nuclear stellar disks are also present in the dusty. active earlv-tvpes. the fraction may be higher in a study of near-inlrared observations of early-twpe galaxies. allhough as nole the nuclear stellar disk fraction may also depend on the central surface briehtness profile.," If nuclear stellar disks are also present in the dusty, active early-types, the fraction may be higher in a study of near-infrared observations of early-type galaxies, although as \citet{ravindranath01} note the nuclear stellar disk fraction may also depend on the central surface brightness profile."16 A histogram5 of the presence of stellar disks as a function of Hubble type in our matched and extended. samples is shown in Figure 3.., A histogram of the presence of stellar disks as a function of Hubble type in our matched and extended samples is shown in Figure \ref{fig-hdisk}.17 The presence of nuclear stellar disks in our saniple is identified in column 12 of Tables 1. and 2. and provides the radial extent of the structure in kiloparsecs., The presence of nuclear stellar disks in our sample is identified in column 12 of Tables \ref{tab-active} and \ref{tab-control} and provides the radial extent of the structure in kiloparsecs.18 In Figure 4 we present a histogram showing the radial extent in kiloparsecs of the nuclear disks. which have a mean value of 0.2420.31 kkpe ancl never extend bevond kkpe.," In Figure \ref{fig-hdisksize} we present a histogram showing the radial extent in kiloparsecs of the nuclear disks, which have a mean value of $\pm$ kpc and never extend beyond kpc."19 Therefore the bright stellar disks in the inactive galaxies are nuclear structures. similar or somewhat smaller in average spatial extent to the dust structures observed in the active galaxies.," Therefore the bright stellar disks in the inactive galaxies are nuclear structures, similar or somewhat smaller in average spatial extent to the dust structures observed in the active galaxies."20 In the previous section we identified (wo significant dillerences between active and inactive earlv-0(vpe galaxies., In the previous section we identified two significant differences between active and inactive early-type galaxies.21 First. all active earlv-tvpe galaxies have circummnuclear dust. while dust is only present in of inactive earlv-tvpe galaxies.," First, all active early-type galaxies have circumnuclear dust, while dust is only present in of inactive early-type galaxies."22 Second. of the inactive early-(vpe galaxies have nuclear stellar disks. while they are detectable in almost none of the active. earlv-tvpe galaxies.," Second, of the inactive early-type galaxies have nuclear stellar disks, while they are detectable in almost none of the active, early-type galaxies."23 In the subsections below we discuss these results in the content, In the subsections below we discuss these results in the context24The mass loss in late evolutionary stages of low- and intermediate-mass stars (1.€M/Ma 8). particularly in the red giant branch (RGB) and asymptotic. giant branch (AGB). is Important not only for the evolution of the star itself but also for the chemical enrichment of the interstellar medium.,"The mass loss in late evolutionary stages of low- and intermediate-mass stars $1 \la M/\MSOL \la 8$ ), particularly in the red giant branch (RGB) and asymptotic giant branch (AGB), is important not only for the evolution of the star itself but also for the chemical enrichment of the interstellar medium."25 The mass loss in Mira-type stars in the AGB phase. which are characterized by the large-amplitude pulsation with a variability amplitudeof AV=6 mag and a period of -ἰ year. has been observationally and theoretically intensively studied.," The mass loss in Mira-type stars in the AGB phase, which are characterized by the large-amplitude pulsation with a variability amplitudeof $\Delta V 26\approx 6$ mag and a period of $\sim$ 1 year, has been observationally and theoretically intensively studied."27 However. normal (1.e.. non-Mira-type) M giants show variability amplitudes of AV = 1-2 mag. much smaller than those of Mira stars. without clear periodic variations.," However, normal (i.e., non-Mira-type) M giants show variability amplitudes of $\Delta V$ = 1–2 mag, much smaller than those of Mira stars, without clear periodic variations."28 They are classified as semi-regular or irregular variables., They are classified as semi-regular or irregular variables.29 Among the M stars listed in the General Catalog of Variable Stars (GCVS) 4.2 (Samus et al. 2009)).," Among the M stars listed in the General Catalog of Variable Stars (GCVS) 4.2 (Samus et al. \cite{samus09}) ),"30 there are 1987 stars classified as semi- or irregular giants (denoted as SRA. SRB. SRD. LB. and L in the catalog). in contrast to 1524 stars classified as Mira variables (see. however. Lebzelter et al.," there are 1987 stars classified as semi-regular or irregular giants (denoted as SRA, SRB, SRD, LB, and L in the catalog), in contrast to 1524 stars classified as Mira variables (see, however, Lebzelter et al."31 1995. for the uncertainty in the classificatio in the GCVS)., \cite{lebzelter95} for the uncertainty in the classification in the GCVS).32 There are also many semi-regular or irregular variables whose spectral type is unknown in the catalog., There are also many semi-regular or irregular variables whose spectral type is unknown in the catalog.33 Therefore. normal M giants with small variability amplitudes may outnumber Mira stars considerably.," Therefore, normal M giants with small variability amplitudes may outnumber Mira stars considerably."34 Despite their small variability amplitudes and the absence of clear periodicities. normal A giants are experiencing mass loss with mass-loss rates comparable to those in (optically bright) Mira stars.," Despite their small variability amplitudes and the absence of clear periodicities, normal M giants are experiencing mass loss with mass-loss rates comparable to those in (optically bright) Mira stars."35 Normal M giants in the AGB phase show mass-loss rates of 107—109 aand expansion velocities of 5-15 ((e.g.. Gonzalez Delgado et al. 20091:," Normal M giants in the AGB phase show mass-loss rates of $10^{-7}$ $10^{-6}$ and expansion velocities of 5–15 (e.g., Gonzalez Delgado et al. \cite{gonzalez_delgado03};"36 Winters et al. 2003::, Winters et al. \cite{winters03};37 De Beck et al. 2010))., De Beck et al. \cite{debeck10}) ).38 There is also evidence of significant mass loss in RGB stars with mass-loss rates up to 10 nnear the tip of the RGB (Origlia et al. 2007.. 20101:," There is also evidence of significant mass loss in RGB stars with mass-loss rates up to $10^{-6}$ near the tip of the RGB (Origlia et al. \cite{origlia07}, \cite{origlia10};"39 Ita et al. 2007)., Ita et al. \cite{ita07}) ).40 The mass-loss rate (M)) and expansion velocity (Gp of our target. BK Vir. are estimated to be (1.5—4)1077 and 4—7.5s!.. respectively (Gonzalez Delgado et al. 2003:;," The mass-loss rate ) and expansion velocity ) of our target, BK Vir, are estimated to be $(1.5-4)\times10^{-7}$ and 4–7.5, respectively (Gonzalez Delgado et al. \cite{gonzalez_delgado03};"41 Winters et al. 2003)).," Winters et al. \cite{winters03}) ),"42 which are approximately the same as in the prototypical Mira o Cet with ==(1-8)x107 and == 1.5-7 kms-'(Winters et al. 2003))., which are approximately the same as in the prototypical Mira $o$ Cet with = $(1-8)\times 10^{-7}$ and = 1.5–7 (Winters et al. \cite{winters03}) ).43 This illustrates the importance of the contribution of normal M giants to the chemical enrichment of the interstellar medium., This illustrates the importance of the contribution of normal M giants to the chemical enrichment of the interstellar medium.44 For understanding the mass-loss mechanism in red giant stars. it is indispensable to probe the physical properties of the region between the upper photosphere and the innermost part of the circumstellar envelope. where the energy and momentum are expected to be deposited for the wind acceleration.," For understanding the mass-loss mechanism in red giant stars, it is indispensable to probe the physical properties of the region between the upper photosphere and the innermost part of the circumstellar envelope, where the energy and momentum are expected to be deposited for the wind acceleration."45 It is now known from spectroscopy and high-spatial resolution observations 1n the IR that there is a dense. warm (~ 1000-2000 K) molecular outer atmosphere. the called MOLsphere. extending to a few nnot only in Mira stars but also in normal K-M giants. with small variability amplitudes (e.g.. Tsuji et al. [ου," It is now known from spectroscopy and high-spatial resolution observations in the IR that there is a dense, warm $\sim$ 1000–2000 K) molecular outer atmosphere, the so-called MOLsphere, extending to a few not only in Mira stars but also in normal K–M giants with small variability amplitudes (e.g., Tsuji et al. \cite{tsuji97};"46 Tsuji 20011: Perrin et al. 2004:;, Tsuji \cite{tsuji01}; Perrin et al. \cite{perrin04};47 Ohnaka 2004:; Ohnaka et al. 2005::, Ohnaka \cite{ohnaka04}; Ohnaka et al. \cite{ohnaka05};48 Takami et al. 2009))., Takami et al. \cite{takami09}) ).49 This MOLsphere is considered to play an important role in driving mass outflows., This MOLsphere is considered to play an important role in driving mass outflows.50 In the case of Mira stars. the large-amplitude pulsation can fairly explain the presence of the extended atmosphere (e.g.. Ohnaka et al. 2006:," In the case of Mira stars, the large-amplitude pulsation can fairly explain the presence of the extended atmosphere (e.g., Ohnaka et al. \cite{ohnaka06};"51 Wittkowski et al. 2007.. 2008::," Wittkowski et al. \cite{wittkowski07}, \cite{wittkowski08};"52 Woodruff et al. 2009))., Woodruff et al. \cite{woodruff09}) ).53 However. the amplitudes of the stellar pulsation are much smaller in normal K-M giants. and therefore. the origin of the MOLsphere m these stars is by no means clear.," However, the amplitudes of the stellar pulsation are much smaller in normal K–M giants, and therefore, the origin of the MOLsphere in these stars is by no means clear."54 In additior to the MOLsphere. the detection of UV emission lines as well as the Ha line suggests the presence of a chromosphere (e.g.. Eaton 1995; Isabel Pérrez Martinez et al.," In addition to the MOLsphere, the detection of UV emission lines as well as the $\alpha$ line suggests the presence of a chromosphere (e.g., Eaton \cite{eaton95}; ; Isabel Pérrez Martínnez et al."55 2011]. anc references therein)., \cite{isabel_perez_martinez11} and references therein).56 This means that there are both hot and warm components in the outer atmosphere of normal K-M giants. and the spectral analyses of the CO lines in the UV anc," This means that there are both hot and warm components in the outer atmosphere of normal K–M giants, and the spectral analyses of the CO lines in the UV and"57sensitive to extinction by dust.,sensitive to extinction by dust.58 From the well known luminosity- relationship (foreg.Pilyugin2001:Leeetal.2006).. one would expect that the typical metallicity in our sample galaxies is substantially less than solar.," From the well known luminosity-metallicity relationship \citep[for eg.][]{pilyugin01, lee06}, one would expect that the typical metallicity in our sample galaxies is substantially less than solar."59 We do not know the detailed star ormation history for many of our sample galaxies. nor do we dave information on what their IMF may be.," We do not know the detailed star formation history for many of our sample galaxies, nor do we have information on what their IMF may be."60 We discuss these issues in some more detail in Section 4.. but this uncertainty in he application of this calibration to our sample galaxies should be borne in mind.," We discuss these issues in some more detail in Section \ref{sec:dis}, but this uncertainty in the application of this calibration to our sample galaxies should be borne in mind."61 As discussed above. a uniform optical diameter is not available or all the galaxies in our sample.," As discussed above, a uniform optical diameter is not available for all the galaxies in our sample."62 For those galaxies for which he Holmberg diameter is available. the average UV flux at this diameter corresponds to a star formation rate of ~ 1.85.10 I. | kpe7.," For those galaxies for which the Holmberg diameter is available, the average UV flux at this diameter corresponds to a star formation rate of $\sim$ $\times$ $^{-4}$ $_\odot$ $^{-1}$ $^{-2}$."63 The globally averaged star formation rate for all galaxies was hence computed to be the average star formation rate inside the region where the SFR falls to 185-10 4M. yr.! > (where the average SFR was calculated along ellipses in the smoothed FUV image having ellipticity b/a as given in Table |)., The globally averaged star formation rate for all galaxies was hence computed to be the average star formation rate inside the region where the SFR falls to $\times$ $^{-4}$ $_\odot$ $^{-1}$ $^{-2}$ (where the average SFR was calculated along ellipses in the smoothed FUV image having ellipticity b/a as given in Table \ref{tab:samp}) ).64 Both and were averaged over identical regions for each galaxy., Both and were averaged over identical regions for each galaxy.65 Table 2. summarizes the parameters derived during the analysis., Table \ref{tab:beam} summarizes the parameters derived during the analysis.66" The columns are as follows: CI) the galaxy name. (2) the synthesized beam sizes of the HI maps (which are also the smoothed resolution of the FUV images) in areseconds. (3) the synthesised beam sizes in parsecs. (4) the rms noise per channel of the cleaned HI data cubes. (5) the major axis in areseconds of the ellipse. defining the ""star forming region"". as described above."," The columns are as follows: (1) the galaxy name, (2) the synthesized beam sizes of the HI maps (which are also the smoothed resolution of the FUV images) in arcseconds, (3) the synthesised beam sizes in parsecs, (4) the rms noise per channel of the cleaned HI data cubes, (5) the major axis in arcseconds of the ellipse, defining the ”star forming region”, as described above."67 For 10 galaxies in the sample. it was possible to make maps at a linear resolution of ~200 pc. and so a pixel by pixel comparison of FUV and HI data was carried out for these galaxies at this higher resolution.," For 10 galaxies in the sample, it was possible to make maps at a linear resolution of $\sim$ 200 pc, and so a pixel by pixel comparison of FUV and HI data was carried out for these galaxies at this higher resolution."68 All other processing was done in a manner analogous to the procedure discussed above., All other processing was done in a manner analogous to the procedure discussed above.69 The relevant details are listed in. 3.. whose columns are as follows: (1) the galaxy name. (2) the synthesized beam sizes for the HI maps (which are also the smoothed resolution of the FUV images) in areseconds. (3) the synthesised beam sizes in parsees. (+) the rms noise per channel of the cleaned HI data cubes.," The relevant details are listed in \ref{tab:beam1}, whose columns are as follows: (1) the galaxy name, (2) the synthesized beam sizes for the HI maps (which are also the smoothed resolution of the FUV images) in arcseconds, (3) the synthesised beam sizes in parsecs, (4) the rms noise per channel of the cleaned HI data cubes."70 The HI contours overlaved on the FUV grey seales are shown in Figure 2.., The HI contours overlayed on the FUV grey scales are shown in Figure \ref{fig:olay1}.71 Figure 3. shows the relationship between the disk-averaged SFR and HI gas surface density for the galaxies in our sample. along with those from Kennicutt(1998).," Figure \ref{fig:tot} shows the relationship between the disk-averaged SFR and HI gas surface density for the galaxies in our sample, along with those from \citet{ken98}."72. The scatter in the disk-averaged data from the galaxies in our sample is too large. and the range of star formation rates and gas densities covered is too small. to provide any meaningful constraints on the power law slope.," The scatter in the disk-averaged data from the galaxies in our sample is too large, and the range of star formation rates and gas densities covered is too small, to provide any meaningful constraints on the power law slope."73 The fit shown in Figure 3. is instead taken from Kennicutt(1998)., The fit shown in Figure \ref{fig:tot} is instead taken from \cite{ken98}.74" The shaded area covers various estimates of the ""threshold density"" as tabulated in Kennicutt(1989):Martin&(2001)."," The shaded area covers various estimates of the “threshold density” as tabulated in \cite{ken89,mar01}."75".. As can be seen our sample galaxies (1) have gas densities that are senerally around or below the expected ""threshold density"" and (2) have star formation rates below that predicted by the Kennicutt(1998) relation.", As can be seen our sample galaxies (1) have gas densities that are generally around or below the expected ”threshold density” and (2) have star formation rates below that predicted by the \cite{ken98} relation.76 These two facts are probably not independent., These two facts are probably not independent.77 Two galaxies. viz.," Two galaxies, viz."78 E321-014 and UGC 6541 lie significantly above the Kennicutt(1998) relation inspite of being sufficiently below the threshold gas density region. and we discuss them in more detail below.," E321-014 and UGC 6541 lie significantly above the \cite{ken98} relation inspite of being sufficiently below the threshold gas density region, and we discuss them in more detail below."79 The solid line in Fig., The solid line in Fig.80 4. shows the Kennicutt(1998) fit to the data for spiral galaxies alone., \ref{fig:spc} shows the \citet{ken98} fit to the data for spiral galaxies alone.81 As can be seen. the dwarf galaxy data is in reasonable agreement with this fit.," As can be seen, the dwarf galaxy data is in reasonable agreement with this fit."82 As such. it appears that dwarf and spiral galaxies have a steeper dependence of SFR on the gus density than predicted by the Kennicutt-Schmidt law.," As such, it appears that dwarf and spiral galaxies have a steeper dependence of SFR on the gas density than predicted by the Kennicutt-Schmidt law."83 The “pixel by pixel” (see Section 2.1. for the detinition of a pixel) comparison of and Που. individual galaxies at 400 pe resolution are shown in Figure Al.., The “pixel by pixel” (see Section \ref{ssec:hidata} for the definition of a pixel) comparison of and for individual galaxies at 400 pc resolution are shown in Figure \ref{fig:plots}.84 Plots for the representative galaxies UGC 4459 and E321-014. along with the combined scatter plot for all the galaxies in the sample are shown in Figure 5..," Plots for the representative galaxies UGC 4459 and E321-014, along with the combined scatter plot for all the galaxies in the sample are shown in Figure \ref{fig:tp}."85 The left panel for each galaxy shows the seatter plot. each point denoting a pixel.," The left panel for each galaxy shows the scatter plot, each point denoting a pixel."86 The right panel shows the binned data (0.05 dex sized bins were used along the x-axis). with the scatter (10) in each bin denoted bv errorbars.," The right panel shows the binned data (0.05 dex sized bins were used along the x-axis), with the scatter $\sigma$ ) in each bin denoted by errorbars."87albeit highlv anisotropic.,albeit highly anisotropic.88 The structure and location of the molecular torus and the BLR. as well as the location of the emitting plasma in the jet. determine the angular dependence of external emission.," The structure and location of the molecular torus and the BLR, as well as the location of the emitting plasma in the jet, determine the angular dependence of external emission."89 This anisotropy is further amplifiecl by relativistic aberration and Doppler boosting or de-boosting in the rest frame of the plasma., This anisotropy is further amplified by relativistic aberration and Doppler boosting or de-boosting in the rest frame of the plasma.90 For simplicity we assume that external radiation is static during the flare., For simplicity we assume that external radiation is static during the flare.91 The properties of the putative dusty torus in (he nucleus of an active galaxy. are poorly known., The properties of the putative dusty torus in the nucleus of an active galaxy are poorly known.92 In particular. despite what its name implies. the geometrical shape ancl size of this structure in a quasar or BL Lac object are poorly constrained.," In particular, despite what its name implies, the geometrical shape and size of this structure in a quasar or BL Lac object are poorly constrained."93 According to one model 2000).. the obscuration is provided by a conical outflow of the material [rom (he accretion disk.," According to one model \citep{elv00}, the obscuration is provided by a conical outflow of the material from the accretion disk."94 Recent. interferometric observations of the nucleus of NGC LOGS (Jaffeetal.2004) reveal (he presence of warm dust at temperature 7300Ix in a structure ~2.1pc in size. surrounding a smaller. warmer (7>800 1x) structure of size ~0.7pe.," Recent interferometric observations of the nucleus of NGC 1068 \citep{jaf04} reveal the presence of warm dust at temperature $T\sim300\,\mbox{K}$ in a structure $\sim2.1\,\mbox{pc}$ in size, surrounding a smaller, warmer $T>800\,\mbox{K}$ ) structure of size $\sim0.7\,\mbox{pc}$."95 The mass of the black hole in NGC 1068 is 1.4x10*M. according to VLBI measurements of water maser emission (Greenhill&Gwinn1997)., The mass of the black hole in NGC 1068 is $1.4\times 10^7\mbox{M}_\sun$ according to VLBI measurements of water maser emission \citep{gri97}.96. For quasars and BL Lac objects harboring more massive black holes. one should expect the size of the torus to scale accordingly.," For quasars and BL Lac objects harboring more massive black holes, one should expect the size of the torus to scale accordingly."97 Fie., Fig.98 l illustrates the geometry and size of (he molecular torus in relation to the position of the emitting blob of plasma in (hie jet [or a representative set of assumptions about the external sources of seed emission and the location of the radiating plasma., \ref{ext} illustrates the geometry and size of the molecular torus in relation to the position of the emitting blob of plasma in the jet for a representative set of assumptions about the external sources of seed emission and the location of the radiating plasma.99" The (torus is characterized by semi-opening angle 6,, and radius Πο.", The torus is characterized by semi-opening angle $\theta_{{{}} op}$ and radius $r_{{{}} tor}$.100" We assume that the emission from the Corus is dominated by dust Chat racdiates as a black body at temperature 7. so that the intensity of emission from (he torus is 5,4,(L)isthePlanckfunction."," We assume that the emission from the torus is dominated by dust that radiates as a black body at temperature $T$, so that the intensity of emission from the torus is where ${\cal %script 101B}_{ (T) is the Planck function."102where Here and below. the primed quantities associated with (he emitting plasma are given in the rest Irae of the host galaxy. whereas unprimed quantities are reserved [or use in the plasma rest fune (this convention follows the one adopted in Paper I).," Here and below, the primed quantities associated with the emitting plasma are given in the rest frame of the host galaxy, whereas unprimed quantities are reserved for use in the plasma rest frame (this convention follows the one adopted in Paper I)."103 We only take into account the emission from the portion of the torus that faces (he central continuum source., We only take into account the emission from the portion of the torus that faces the central continuum source.104 However. the details of (iis approach are not crucial to the final results.," However, the details of this approach are not crucial to the final results."105 The only essential parameters are (he angle 8/mine Which determines the maximum Doppler boosting. ancl the dust temperature 7.," The only essential parameters are the angle $\theta'_{{{}} min}$, which determines the maximum Doppler boosting, and the dust temperature $T$ ."106 The DLhR. can be represented by a uniform spherical source of emission at a fiducial frequency ή.," The BLR can be represented by a uniform spherical source of emission at a fiducial frequency $\nu'_{{{}}107blr}$."108" The integrated intensity of the incident emission from the BLR is given by where ry, is the radius of the BLR.Ar(@) is the geometric thickness of the BLR in"," The integrated intensity of the incident emission from the BLR is given by where $r_{{{}} blr}$ is the radius of the BLR,$\Delta{r}(\theta')$ is the geometric thickness of the BLR in"109"higher mass outllow rates are needed. as the observations anc modelling of optically thick sources require a absorbing column depth. and hence in a good fraction of sources N,N. may actually be much higher than 1 g 2","higher mass outflow rates are needed, as the observations and modelling of optically thick sources require a absorbing column depth, and hence in a good fraction of sources $N_{\rm a} \Sigma_{\rm c}$ may actually be much higher than 1 g $^{-2}$."110 Considering a specilic case of the local obscured GN studied by (2)... we note that he bolometric luminosities of these objects in the infrared. X-ray ancl optical bands are in the range Lb—107(ον©107. org/sec. which implies SMLDLIL accretion rates of “only” ~0.01M. vear+ for the stancard radiative clliciency.," Considering a specific case of the local obscured AGN studied by \citep{Guainazzi05}, we note that the bolometric luminosities of these objects in the infrared, X-ray and optical bands are in the range $L \sim 10^{43} -111\hbox {few} \times 10^{44}$ erg/sec, which implies SMBH accretion rates of “only” $ \sim 0.01 \msun$ $^{-1}$ for the standard radiative efficiency."112 Hence if the obscuration of the optically thick objects in that sample were provided by the winds. we would conclude that the SALBLE accretion process must be very wasteful. with 1000.000 times more mass Ilowing out of the inner parsec than acereting on the SALBLI.," Hence if the obscuration of the optically thick objects in that sample were provided by the winds, we would conclude that the SMBH accretion process must be very wasteful, with $\sim 100-10,000$ times more mass flowing out of the inner parsec than accreting on the SMBH."113 (Notethatthesemoderatelybright.AGNarenotlikelytooutllowsareinfactexpected.e.g...2). I would also require avery high mass influx into the inner parsec to sustain such winds.," \citep[Note that these114moderately bright AGN are not likely to be in the non-radiative accretion flow115regime when vigorous outflows are in fact expected, e.g., ][]{Blandford99} It would also require a very high mass influx into the inner parsec to sustain such winds."116 Given the cdillieulty of delivering enough fuel to the SAIBUs even in the earlier gas-rich epochs (2). it is hard to see how such high mass inlluxes could be maintained.," Given the difficulty of delivering enough fuel to the SMBHs even in the earlier gas-rich epochs \citep{Thompson05}, it is hard to see how such high mass influxes could be maintained."117 Accreting black holes can drive strong outllows via X-ray heating. line or continuous radiation pressure. and hvdromagnetic forces (e.g...2?22?)..," Accreting black holes can drive strong outflows via X-ray heating, line or continuous radiation pressure, and hydromagnetic forces \citep[e.g.,][]{Begelman83,Konigl94,Kartje99,Proga03c}."118 Star formation on the outskirts of cool massive ο cdises will also result. in winds driven by outflows from the voung massive stars (7).., Star formation on the outskirts of cool massive accretion discs will also result in winds driven by outflows from the young massive stars \citep{Cuadra05}.119 Llere we studied the obseuration. properties of the ACN outllows of the last type., Here we studied the obscuration properties of the AGN outflows of the last type.120 We found that these outflows are quite clumpy. with the lines of sight passing through the clumps becoming moderately optically thick to Thomson scattering for high enough wind mass loss rates.," We found that these outflows are quite clumpy, with the lines of sight passing through the clumps becoming moderately optically thick to Thomson scattering for high enough wind mass loss rates."121 Such outllows are bound to play a role in the obscuration of AGN., Such outflows are bound to play a role in the obscuration of AGN.122 We however feel that this role cannot be dominant. as the sky averaged column depth is significantly smaller than, We however feel that this role cannot be dominant as the sky averaged column depth is significantly smaller than123stars and their metallicities were documented. for example. by Caputo (1997).,"stars and their metallicities were documented, for example, by Caputo (1997)."124 There is also observational evidence for a non-linear dependence of Ady on Ve/L]., There is also observational evidence for a non-linear dependence of $M_V$ on [Fe/H].125 For example. Caputo ct al. (," For example, Caputo et al. ("1262000). analvzing WR Lyr variables from Galactic eglobular clusters. obtained a=O17+0.04 for ο]«1.5 and à=0.27£0.06 for Fe/L]1.5.,"2000), analyzing RR Lyr variables from Galactic globular clusters, obtained $\alpha=0.17\pm0.04$ for ${\rm [Fe/H]}<-1.5$ and $\alpha=0.27\pm0.06$ for ${\rm [Fe/H]}>-1.5$."127 The systematic errors caused by using the simple linear relation (1) may be minimized by excluding highly evolved RR. Lyraes from the analvzed samples., The systematic errors caused by using the simple linear relation (1) may be minimized by excluding highly evolved RR Lyraes from the analyzed samples.128 In particular. the formula should be used with caution when analyzing variables belonging to globular clusters showing “blue” horizontal branches.," In particular, the formula should be used with caution when analyzing variables belonging to globular clusters showing ""blue"" horizontal branches."129" Recent investigation of the Ad,Fefll] relation in αἱ Cen was done by Rev et al. (", Recent investigation of the $M_V - {\rm [Fe/H]}$ relation in $\omega$ Cen was done by Rey et al. (1302000).,2000).131 They used οΗ] metallicities derived from the h& index of the Caby photometric system., They used [Fe/H] metallicities derived from the $hk$ index of the $by$ photometric system.132 Their intensity averaged: magnitudes of RR Lyr stars were taken from photographic photometry from Butler et al. (, Their intensity averaged magnitudes of RR Lyr stars were taken from photographic photometry from Butler et al. (1331978). and. the CCD photometry of Ixaluznuv ct al. (,1978) and the CCD photometry of Kaluzny et al. (1341997).,1997).135 The new V-band light curves of Ri Lyr variables in o Con obtained by the CASE (Ixaluzny et al., The new $V$ -band light curves of RR Lyr variables in $\omega$ Cen obtained by the CASE (Kaluzny et al.136 2003) have about three times as many observed points as the photometry of Ixaluzny et al (1997)., 2003) have about three times as many observed points as the photometry of Kaluzny et al (1997).137 Ehis permits a more precise estimation of mean «V magnitudes and allows the elimination of objects with unstable light curves from the calibration sample., This permits a more precise estimation of mean $<V>$ magnitudes and allows the elimination of objects with unstable light curves from the calibration sample.138 Adopting an apparent distance modulus. of the cluster as determined by (Ixaluzny et al., Adopting an apparent distance modulus of the cluster as determined by (Kaluzny et al.139 2002) we can obtain a new AdyFe/H] relation for RR. Lyr stars in ew Con.," 2002) we can obtain a new $M_V - {\rm140[Fe/H]}$ relation for RR Lyr stars in $\omega$ Cen."141 The sample used for the calibration includes 122 stars with stable light curves of &ood. quality ancl metallicities determined. by Rev et al. (, The sample used for the calibration includes 122 stars with stable light curves of good quality and metallicities determined by Rey et al. (1422000).,2000).143 We decided. to remove the RRab variable V52 from the further analysis., We decided to remove the RRab variable V52 from the further analysis.144 Its high luminosity suggests that it isa foreground star., Its high luminosity suggests that it is a foreground star.145 Accditionallv. van Leeuwen et al. (," Additionally, van Leeuwen et al. ("1462000) gives only a membership probability for this variable. based on a proper motion study.,"2000) gives only a membership probability for this variable, based on a proper motion study."147 The AdyFe/M] dependence for remaining 121 variables is shown in the upper panel of Fig., The $M_V - {\rm [Fe/H]}$ dependence for remaining 121 variables is shown in the upper panel of Fig.148 1., 1.149 The middle panel shows the same relation but [or à sub-sample consisting of 67 RRab variables., The middle panel shows the same relation but for a sub-sample consisting of 67 $ab$ variables.150 Phe errors in. Ο/Η) were taken directly from Table 5 of Rev et al. (, The errors in [Fe/H] were taken directly from Table 5 of Rey et al. (1512000).,2000).152 For objects with unknown error of. Fe/LI] we assumed it to be equal to 0.2 dex as suggested by Rey et al. (, For objects with unknown error of [Fe/H] we assumed it to be equal to 0.2 dex as suggested by Rey et al. (1532000).,2000).154 The intensity. averaged magnitudes «V obtainec from light curves containing 500-800. points have interna errors smaller than 0.001 mag., The intensity averaged magnitudes $<V>$ obtained from light curves containing 500-800 points have internal errors smaller than 0.001 mag.155 We estimate that the externa error of the zero point of the photometry is about 0.02 mag., We estimate that the external error of the zero point of the photometry is about 0.02 mag.156 ‘Transformation from «V to My. ds. performec using the distance modulus whose error is 0.04 mae., Transformation from $<V>$ to $M_V$ is performed using the distance modulus whose error is 0.04 mag.157 Thus. combining this value with 0.02 mae error resulting [roni computing the mean niagnitudes «V.2. we assumed tha individual absolute magnitudes of RR Lyr stars from our sample have uncertainties of 0.05 mag.," Thus, combining this value with 0.02 mag error resulting from computing the mean magnitudes $<V>$, we assumed that individual absolute magnitudes of RR Lyr stars from our sample have uncertainties of 0.05 mag."158 Knowing the errors we were able to fit straight lines to the graphs shown in Fig., Knowing the errors we were able to fit straight lines to the graphs shown in Fig.159 1., 1.160 For the sample of 121 RR bvr stars we obtain the following relation: and for 67 Ritab variables is in the form: The values of a derived above may. be biased. by the presence of extremely evolved. objects in our sample., For the sample of 121 RR Lyr stars we obtain the following relation: and for 67 $ab$ variables is in the form: The values of $\alpha$ derived above may be biased by the presence of extremely evolved objects in our sample.161 The evolutionary models of Lee (1990) suggest. that the IH Lyrac stars in clusters having a very blue LB and with a metallicity in the range of ο&—1.6 are highly evolved stars., The evolutionary models of Lee (1990) suggest that the RR Lyrae stars in clusters having a very blue HB and with a metallicity in the range of $-2.0 < {\rm [Fe/H]} < -1.6$ are highly evolved stars.162 They have signilicantly. brighter magnitudes and longer periods than those near the zero age horizontal branch (ZALIB)., They have significantly brighter magnitudes and longer periods than those near the zero age horizontal branch (ZAHB).163 ltey et al. (, Rey et al. (1642000) investigated this problem in c Cen.,2000) investigated this problem in $\omega$ Cen.165 Analvzing the period-amplitucde relations for different ranges of metallicities they discovered that a significant sample of evolved RR Lyr stars exists only for variables with 19<οLI]< 5., Analyzing the period-amplitude relations for different ranges of metallicities they discovered that a significant sample of evolved RR Lyr stars exists only for variables with $-1.9 \leq {\rm [Fe/H]} < -1.5$ .166 Looking at their Fig., Looking at their Fig.167 9b one can clearly see that these evolved objects are also characterized by periods longer than 0.7 davs., 9b one can clearly see that these evolved objects are also characterized by periods longer than 0.7 days.168 Thus we decided to exclude from our sample all RIvib variables having metallicities in a range of, Thus we decided to exclude from our sample all $ab$ variables having metallicities in a range of169In addition to a variable warm absorber. à constant warm absorber. which is primarily seen in the active state. has been introduced.,"In addition to a variable warm absorber, a constant warm absorber, which is primarily seen in the active state, has been introduced."170 All the parameters of the two. warm-absorbers are shared. by the two datasets apart [rom the ionization parameter of the variable. absorber (g-.)., All the parameters of the two warm-absorbers are shared by the two datasets apart from the ionization parameter of the variable absorber $\xi_{\rm v}$ ).171" The temperatures of the constant ancl variable absorbers are assuemcd to be 110 WK and 110"" Ix. respectively."," The temperatures of the constant and variable absorbers are assuemd to be $1\times 10^5$ K and $1\times 10^6$ K, respectively."172 This model gives a good fit to the data despite the only one parameter (£.) in the two warnme-absorboers dillering between the two datasets (Table 3. Fig.," This model gives a good fit to the data despite the only one parameter $\xi_{\rm v}$ ) in the two warm-absorbers differing between the two datasets (Table 3, Fig."173 9. Fig.," 9, Fig."174 10)., 10).175 The ionization parameter (ἐςc30 13) and column density CN=21072em. 7)) of the constant absorber are found to be similar to those seen in higher luminosity Seyfert 1: nuclei (Revnolds 1997: George οἱ al 1998)., The ionization parameter $\xi_{\rm c}\simeq 30$ ) and column density $N_{\rm W}^{\rm c}\simeq 2\times 10^{22}$ ) of the constant absorber are found to be similar to those seen in higher luminosity Seyfert 1 nuclei (Reynolds 1997; George et al 1998).176 The variable absorber has a higher and variable, The variable absorber has a higher and variable177" ???????????)) ???)). ?77)) ?7)). (?), has recently produced a successful detection of CO at K-band wavelengths (?).."," \citealt{charbonneau2005p523, Deming2005p740, Harrington2007p691, charbonneau2008p1341, swain2008p482, knutson2009p822, Pont2008p109, Pont2009pL6, grillmair2008p767, fressin2010p374, gillon2010p3}) \citealt{charbonneau2002p377, snellen2009p543, Snellen2010p76}) \citealt{burrows2007pL171, Burrows2008p1436, fortney2008p1419, showman2009p564, fortney2010p1396}; \citealt{seager2010p631} \citealt{Gillon2009p359, alonso2010p1481, deMooij2009pL35}) \citealt{redfield2008pL87, Snellen2008p357}) \citep{Hook2009p225}, \citealt{Wiedemann2001p1068, Brown2002p826, Richardson2003p581, Richardson2003p1053, Barnes2007p473, Barnes2008p1258,Barnes2010p445}) has recently produced a successful detection of CO at K-band wavelengths \citep{Snellen2010p1049}."178" Additionally, a startling ground-based detection of molecular spectral features from the atmosphere of the exoplanet 1189733b at moderate spectral resolution was announced by Swain et al. ("," Additionally, a startling ground-based detection of molecular spectral features from the atmosphere of the exoplanet 189733b at moderate spectral resolution was announced by Swain et al. ("1792010; hereafter S10).,2010; hereafter S10).180" S10 analyzed observations of 1189733 taken with the SpeX instrument on the Infrared Telescope Facility, focusing on two wavelength regions: 2.0—2.4 uum (K band) and 3.1—4.1 um (L band)."," S10 analyzed observations of 189733 taken with the SpeX instrument on the Infrared Telescope Facility, focusing on two wavelength regions: $2.0 - 2.4\,\mu$ m (K band) and $3.1 - 4.1\,\mu$ m (L band)."181" Their results show a rising spectral slope at um, which they attribute to CO» absorption, and a bright peak at wm. The shape of the spectrum at K-band wavelengths matched previous results at these wavelengths taken with the Hubble Space Telescope which the authors construed as support for their ground-based(?),, detection in the L band."," Their results show a rising spectral slope at $\mu$ m, which they attribute to $_2$ absorption, and a bright peak at $\mu$ m. The shape of the spectrum at K-band wavelengths matched previous results at these wavelengths taken with the Hubble Space Telescope \citep{swain2009pL114}, which the authors construed as support for their ground-based detection in the L band."182 $10 interpret the flux peak in the L-band portion of their data as evidence of bright non-LTE (NLTE) emission from radiatively-excited methane., S10 interpret the flux peak in the L-band portion of their data as evidence of bright non-LTE (NLTE) emission from radiatively-excited methane.183" The wavelength region for the claimed emission roughly corresponds with the R-branch region of CH, (3.1—3.34 um), but no corresponding feature is seen in the region of the methane P-branch (3.35—3.5 $10 note this disparity but offer no theoretical explanationum); for it."," The wavelength region for the claimed emission roughly corresponds with the R-branch region of $_4$ $3.1-3.34\,\mu$ m), but no corresponding feature is seen in the region of the methane P-branch $3.35-3.5\,\mu$ m); S10 note this disparity but offer no theoretical explanation for it."184 The L-band emission claimed by S10 is also remarkable for having high intensity at low spectral resolution., The L-band emission claimed by S10 is also remarkable for having high intensity at low spectral resolution.185" The wavelength extent of each spectral bin analyzed by S10 is um, (λ/δλ~ 30), and they find a planet-to-star contrast ratio of approximately in their brightest bin at um. This result is even more surprising considering that NLTE emission lines are unlikely to be collisionally broadened, since the pressures needed to achieve significant line broadening would produce high collisionratesthat would drive"," The wavelength extent of each spectral bin analyzed by S10 is $\mu$ m, $\lambda / \delta\lambda \sim 30$ ), and they find a planet-to-star contrast ratio of approximately in their brightest bin at $\mu$ m. This result is even more surprising considering that NLTE emission lines are unlikely to be collisionally broadened, since the pressures needed to achieve significant line broadening would produce high collisionratesthat would drive"186value of Vis only slightly larger than that listed in Table 1. so 980425 remains to have the largest scatter in Fig. 5..,"value of $V$ is only slightly larger than that listed in Table 1, so 980425 remains to have the largest scatter in Fig. \ref{fig5}."187 With our definition of variability. GIU 980425 is much less olfset from the variabilitv-Iuminosity correlation.," With our definition of variability, GRB 980425 is much less offset from the variability-luminosity correlation."188" Figure 5. also shows that the values of the variability with our definition are systematically smaller than that with LOLs definition. as indicated by the straight line defined by ""our variability = 1019 variability. divided by 157."," Figure \ref{fig5} also shows that the values of the variability with our definition are systematically smaller than that with R01's definition, as indicated by the straight line defined by “our variability $=$ R01's variability divided by 15”."189 This is mainly caused by the cilferent normalization in the two definitions., This is mainly caused by the different normalization in the two definitions.190 In addition. our variability generally has a larger error bar than that of ROL. caused by the fact. that we normalize our variability by peak count. which sullers larger photon noise than total counts: and we have considered. the error arising [rom changing που by £1.," In addition, our variability generally has a larger error bar than that of R01, caused by the fact that we normalize our variability by peak count, which suffers larger photon noise than total counts; and we have considered the error arising from changing $N_{\rm iter}$ by $\pm 1$."191 Now. we make a linear fit to logLV. for the overlapping sample. with our definition of variability.," Now, we make a linear fit to $\log L-\log V$ for the overlapping sample, with our definition of variability."192 For the reasons explained in Section 4. we exclude GRBs 980425 and 030329 during the fit.," For the reasons explained in Section 4, we exclude GRBs 980425 and 030329 during the fit."193 Phus. the total number ofGRBs to fit is N=17.," Thus, the total number ofGRBs to fit is $N=17$."194 We obtain log£L=3.101ogV.|59.13. and v2fdof=2.43 where dof=15.," We obtain $\log L = 3.10 \log V + 59.13$, and $\chi^2/\dof 195= 2.43$ where $\dof = 15$."196 The weighted average of the individual data variance is (07)= 0.02715., The weighted average of the individual data variance is $\langle \sigma_i^2\rangle = 0.02715$ .197" Vhus. for the fit with our definition of variability. we have s,=0.257. 5,—0,0829. and s=0.079."," Thus, for the fit with our definition of variability, we have $s_y = 0.257$, $s_x = 0.0829$, and $s =0.079$."198 Next. we make a linear fit to log£L.V. for the overlapping sample. with ROL’s definition of variability.," Next, we make a linear fit to $\log L-\log V$ for the overlapping sample, with R01's definition of variability."199 Again. GRBs 980425 and 030329 are excluded. to make the GRB members in the sample remain the same.," Again, GRBs 980425 and 030329 are excluded, to make the GRB members in the sample remain the same."200 We obtain logL=L77logV.|54.06. and νο=19.89.," We obtain $\log L = 2011.77\log V + 54.06$, and $\chi^2/\dof = 19.89$."202 The weighted. average ofthe individual data variance is 0675[=0.009240., The weighted average ofthe individual data variance is $\langle \sigma_i^2\rangle = 0.009240$.203" Thus. for the fit with ROI's definition of variability. we have s,=0.429. s,=0.2417. and s=0.211."," Thus, for the fit with R01's definition of variability, we have $s_y = 0.429$, $s_x = 0.2417$, and $s =0.211$."204 The scatter of data points around. the fitted. logLlogV. is measured. by os., The scatter of data points around the fitted $\log L-\log V$ is measured by $s$.205 Thus. according to the above numbers. for the overlapping sample the scatter in our logL—V. relation is smaller than that of €05 hy a [actor of 0.211/0.079=2.7. although the weighted average of the individual data error (=fo2X2i ) with our definition of variability is larger than that with BROLI's definition of variability by a factor of (0.02715/0.000240)7=1.7.," Thus, according to the above numbers, for the overlapping sample the scatter in our $\log L - \log V$ relation is smaller than that of G05 by a factor of $0.211/0.079 = 2.7$, although the weighted average of the individual data error $\equiv \langle\sigma_i^2\rangle^{1/2}$ ) with our definition of variability is larger than that with R01's definition of variability by a factor of $(0.02715/0.009240)^{1/2} = 1.7$."206 The above results. for the overlapping sample excluding 980425 ancl 030329. are summarized in Fig. 6..," The above results, for the overlapping sample excluding 980425 and 030329, are summarized in Fig. \ref{fig6},"207 where filled circles and thick lines represent the results with our data. open circles and thin lines represent the results with 65's data.," where filled circles and thick lines represent the results with our data, open circles and thin lines represent the results with G05's data."208 Phe luminosities in the two set of data are the same. so each GARB is represented. by a horizontal line in Fig. 6..," The luminosities in the two set of data are the same, so each GRB is represented by a horizontal line in Fig. \ref{fig6}."209 Llowever. the variabilities in the two set of data cüller. so we see the olfset along the horizontal direction of each pair of filled. civele-open circle.," However, the variabilities in the two set of data differ, so we see the offset along the horizontal direction of each pair of filled circle-open circle."210" Phe solid straight lines are least-\2 fits to each data set. and the dashed straight lines mark the “leo width"" of the fits."," The solid straight lines are $\chi^2$ fits to each data set, and the dashed straight lines mark the $\sigma$ width” of the fits."211 We have set the unitary length of the horizontal ancl vertical axes to be the same. so that for each data set 5 is just the half-width of the region bounded by the two dashed lines that vou see.," We have set the unitary length of the horizontal and vertical axes to be the same, so that for each data set $s$ is just the half-width of the region bounded by the two dashed lines that you see."212 The figure clearly shows that the scatter of our data set is much narrower than that ol ος data set., The figure clearly shows that the scatter of our data set is much narrower than that of G05's data set.213 With the above defined scatter parameter 5. we can also check how sensitive our results are to the parameters in our definition of variabilitv.," With the above defined scatter parameter $s$, we can also check how sensitive our results are to the parameters in our definition of variability."214 As mentioned in Section 2. we have chosen Jy= duüz;. and Ni to bethe integer closest to the ratio Zoo Z7.," As mentioned in Section 2, we have chosen $T_f = T_{0.5}$ , and $N_{\rm iter}$ to bethe integer closest to the ratio $T_{90}/T_f$ ."215 With this choice. for the GIU sample described in Section7 3 (excluding 980425. 030329. ancl 030528) we obtained ν΄/dof=1.93 and s=0.078 ," With this choice, for the GRB sample described in Section 3 (excluding 980425, 030329, and 030528), we obtained $\chi^2/\dof = 1.93$ and $s = 0.078$ "216the nucleus.,the nucleus.217 We extrapolated the gas density model derived from the surface brightness to the radii of the radio lobes and combined the derived pressure there with the expected cavity volumes created through adiabatic bubble expansion to estimate the required energy of ~5«1055 ergs., We extrapolated the gas density model derived from the surface brightness to the radii of the radio lobes and combined the derived pressure there with the expected cavity volumes created through adiabatic bubble expansion to estimate the required energy of $\sim5\times10^{58}$ ergs.218 The mass accreted onto the SMBH to produce this energy would be ~2«I0?...., The mass accreted onto the SMBH to produce this energy would be $\sim2\times10^{5}$.219 Centaurus A (NGC 5128) and NGC 1316 (Pornax A) are both nearby elliptical. galaxies. which have recently undergone a merger event that has produced strong nuclear activity.," Centaurus A (NGC 5128) and NGC 1316 (Fornax A) are both nearby elliptical galaxies, which have recently undergone a merger event that has produced strong nuclear activity."220 They each host a 10° M.. black hole that are low lummosity AGNs and have dust lanes roughly perpendicular to their radio lobes (Marconietal.2006.Nowak 2008).," They each host a $^{8}$ $_{\odot}$ black hole that are low luminosity AGNs and have dust lanes roughly perpendicular to their radio lobes \citep{mar06, now08}."221. However. they differ significantly with regards to the observational characteristies of their Jets.," However, they differ significantly with regards to the observational characteristics of their jets."222 At a distance of 3.7 Mpe (Ferrareseetal.2007).. Cen A has a 1.5 kpe radio and X- emitting jet that extends from the nucleus to the northeast radio lobe (Feigelsonetal.1981.Kraft2002.Hardeastleetal. 2003).," At a distance of 3.7 Mpc \citep{fer07}, Cen A has a 1.5 kpc radio and X-ray emitting jet that extends from the nucleus to the northeast radio lobe \citep{fei81, kra02, har03}."223. In contrast. NGC 1316s ~3 kpe radio jet does not extend to the large radio lobes and coincides not with an X-ray jet. but with a soft X-ray cavity.," In contrast, NGC 1316's $\sim$ 3 kpc radio jet does not extend to the large radio lobes and coincides not with an X-ray jet, but with a soft X-ray cavity."224 This phenomenological contrast suggests that. while the jet of Cen A is dissipative (Nulsenetal.2010).. NGC 1316's jet is not or that the NGC 1316 jet has shut off on larger scales.," This phenomenological contrast suggests that, while the jet of Cen A is dissipative \citep{nul10}, NGC 1316's jet is not or that the NGC 1316 jet has shut off on larger scales."225 We detected considerably more dust emission. for NGC 1316 than expected in an early type galaxy with its K-band lummosity and observed evidence of recent AGN outbursts in the form of X-ray cavities and radio lobes., We detected considerably more dust emission for NGC 1316 than expected in an early type galaxy with its K-band luminosity and observed evidence of recent AGN outbursts in the form of X-ray cavities and radio lobes.226 We presented images of the infrared dust emission. including the first image of dust emission atjjm.," We presented images of the infrared dust emission, including the first image of dust emission at."227. We determined that the dust has a clumpy morphology. mostly confined to two regions. one 28788 (3.1 kpe) southeast of the nucleus and a 43/99 (4.8 kpe) protrusion ending in an are northwest of the nucleus.," We determined that the dust has a clumpy morphology, mostly confined to two regions, one 8 (3.1 kpc) southeast of the nucleus and a 9 (4.8 kpc) protrusion ending in an arc northwest of the nucleus."228 Molecular emission is detected from these regions., Molecular emission is detected from these regions.229 The resolved radio jet is not detected bySpirzer.. and it does hot coincide with regions of dust emission.," The resolved radio jet is not detected by, and it does not coincide with regions of dust emission."230 Since the dust nust be almost entirely external in origin. we use the dust Nass to constrain the type and mass of the merger galaxy.," Since the dust must be almost entirely external in origin, we use the dust mass to constrain the type and mass of the merger galaxy."231 We calculated a present dust mass of 2.4370.9ς107 based on the integratedMIPS fluxes. which agrees with the dust nass of 3.2«10’ predicted by the model of Draineet (2007).," We calculated a present dust mass of $2.4\pm 0.9\times 10^{7}$ based on the integratedMIPS fluxes, which agrees with the dust mass of $3.2 \times 10^{7}$ predicted by the model of \citet{dra07}."232. We estimate the merger galaxy was a late type galaxy with a stellar mass of 1—6«10 and a gas mass of 2—4«10°... some of which was likely ionized or used to form stars in the merger event.," We estimate the merger galaxy was a late type galaxy with a stellar mass of $1-6\times10^{10}$ and a gas mass of $2-4\times10^{9}$, some of which was likely ionized or used to form stars in the merger event."233 We constrained the age and energy of the merger and outburst events based on the X-ray and radio emission., We constrained the age and energy of the merger and outburst events based on the X-ray and radio emission.234 The image shows a pair of X-ray cavities to the west and southeast of the nucleus. likely created by the expansion of radio plasma. which are closer to the nucleus than the 1.4 GHz radio lobes.," The image shows a pair of X-ray cavities to the west and southeast of the nucleus, likely created by the expansion of radio plasma, which are closer to the nucleus than the 1.4 GHz radio lobes."235 The relative locations of these cavities and the radio lobes suggests that either the 1.4 GHz radio emission. does not show the full extent of the radio emission from the outburst or that there have been at least two AGN outbursts., The relative locations of these cavities and the radio lobes suggests that either the 1.4 GHz radio emission does not show the full extent of the radio emission from the outburst or that there have been at least two AGN outbursts.236" We calculate buoyant rise times for the X-ray cavities of 0.1 Gyr and for the radio lobes of 0.4 Gyr.assuming expansion in the plane of the sky at 0.6c,.. which agrees with the synchrotron age estimated by (1983)."," We calculate buoyant rise times for the X-ray cavities of 0.1 Gyr and for the radio lobes of 0.4 Gyr,assuming expansion in the plane of the sky at 0.6, which agrees with the synchrotron age estimated by \citet{eke83}."237. Since the age of the radio lobes is close to the 0.5 Gyr age estimated by Mackie&Fabbiano(1998) for the merger. the outburst was likely triggered by the accretion of material onto the SMBH from this merger.," Since the age of the radio lobes is close to the 0.5 Gyr age estimated by \citet{mac98} for the merger, the outburst was likely triggered by the accretion of material onto the SMBH from this merger."238 Finally. we constrained the kinetic energy of the outbursts based on the energy required to create the cavities and the radio lobes bubbles.," Finally, we constrained the kinetic energy of the outbursts based on the energy required to create the cavities and the radio lobes bubbles."239 We estimate the outburst that created the X-ray cavities had à kinetic energy of 1075 eres and that the creation of the radio lobes required ~5 times more power., We estimate the outburst that created the X-ray cavities had a kinetic energy of $10^{58}$ ergs and that the creation of the radio lobes required $\sim5$ times more power.240 We are grateful to C. Horellou and J. H. Black for providing us with CO(2-1) intensities and for their comments regarding the molecular gas kinematics. to Zhtyuan Li for his useful discussion on calculating masses. to Dharam Lal for his assistance in creating the 4.89 GHz map. and to Ramesh Narayan for his comments.," We are grateful to C. Horellou and J. H. Black for providing us with CO(2-1) intensities and for their comments regarding the molecular gas kinematics, to Zhiyuan Li for his useful discussion on calculating masses, to Dharam Lal for his assistance in creating the 4.89 GHz map, and to Ramesh Narayan for his comments."241 We thank the anonymous referee for the many comments that improved this work., We thank the anonymous referee for the many comments that improved this work.242 This work was based on archival data obtained from the Spitzer Science. Archive. the Chandra Data Archive. and the XMM-Newton Science Data Archive.," This work was based on archival data obtained from the Spitzer Science Archive, the Chandra Data Archive, and the XMM-Newton Science Data Archive."243 Archived images were also obtained from the Hubble Legacy Archive. the NASA/IPAC Extragalactic Database. and the National Radio Astronomy Observatory Archive.," Archived images were also obtained from the Hubble Legacy Archive, the NASA/IPAC Extragalactic Database, and the National Radio Astronomy Observatory Archive."244 We thank Z. Levay of the Space Telescope Science Institute for his assistance in obtaining the Hubble ACS image., We thank Z. Levay of the Space Telescope Science Institute for his assistance in obtaining the Hubble ACS image.245 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. X," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"246 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XM," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"247 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"248 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"249 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM..," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"250 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.. ," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"251 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.. V," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"252 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.. VL," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"253 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.. VLA," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"254 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.. VLA.," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"255 This work was supported by the Smithsonian Institution. the Chandra X-ray Center. and NASA contract NNX07AQISG. Spitzer.. CXO.. XMM.. VLA..," This work was supported by the Smithsonian Institution, the Chandra X-ray Center, and NASA contract NNX07AQ18G. , , ,"256A total of 26 individual CORALIE spectra of WASP-23 were co-added to produce a single spectrum with a typical S/N of around 50:1.,A total of 26 individual CORALIE spectra of WASP-23 were co-added to produce a single spectrum with a typical S/N of around 50:1.257 The standard pipeline reduction products were used in the analysis., The standard pipeline reduction products were used in the analysis.258 The analysis was performed using the methods given in 2.., The analysis was performed using the methods given in \citet{Gillon:2009p3869}.259 The lline was used to determine the effective temperature (T4). while the Na D and Me b lines were used às surface gravity g)) diagnostics.," The line was used to determine the effective temperature ), while the Na D and Mg b lines were used as surface gravity ) diagnostics."260 The parameters obtained from the analysis are listed in Table 1.., The parameters obtained from the analysis are listed in Table \ref{wasp23-params}.261 The elemental abundances were determined from equivalent width measurements of several clean and unblended lines., The elemental abundances were determined from equivalent width measurements of several clean and unblended lines.262 A value for microturbulence (&)) of 0.8 kmss! wwas determined from Fe using ?. method., A value for microturbulence ) of 0.8 $^{-1}$ was determined from Fe using \citet{Magain:1984p4690} method.263 The quoted error estimates include that given by the uncertainties inTey. aand&.. as well as the scatter due to measurement and atomic data uncertainties.," The quoted error estimates include that given by the uncertainties in, and, as well as the scatter due to measurement and atomic data uncertainties."264 The projected stellar rotation velocity was determined by fitting the profiles of several unblended Fe lines., The projected stellar rotation velocity was determined by fitting the profiles of several unblended Fe lines.265 Because the value of was paramount to the model fitting. we used the combined HARPS spectra.," Because the value of was paramount to the model fitting, we used the combined HARPS spectra."266 A value for macroturbulence (44:0) of 0.8€0.3 kmss! was assumed. based on work by ?. and an instrumental FWHM of0.0604... determined from the telluric lines around6300A.," A value for macroturbulence ) of $0.8\pm0.3$ $^{-1}$ was assumed, based on work by \citet{Bruntt:2010p8010} and an instrumental FWHM of, determined from the telluric lines around."267. A best fitting value of -z-22 z0.3kmss! wwas obtained., A best fitting value of = 2.2 $\pm$ 0.3 $^{-1}$ was obtained.268 Using a macroturbulence based on the tabulation by ? of 1.2 kmss! wwe obtain the same result for showing tts robustness., Using a macroturbulence based on the tabulation by \citet{Gray:2008p4677} of 1.2 $^{-1}$ we obtain the same result for showing its robustness.269 The HARPS spectra show that there is weak emission in the cores of the Calctum H K lines., The HARPS spectra show that there is weak emission in the cores of the Calcium H K lines.270"Activity levels on the star are estimated by means of the logAj, (22?) and obtained using a B-Vz0.88+0.05 estimated from the effective temperature.","Activity levels on the star are estimated by means of the $\log\,R'_\mathrm{HK}$ \citep{Noyes:1984p6855,Santos:2000p6686,Boisse:2009p1077} and obtained using a $B-V = 0.88 \pm 0.05$ estimated from the effective temperature."271 The Mount Wilson index. Syyy is also given.," The Mount Wilson index, $S_\mathrm{MW}$ is also given."272 The spectroscopic. data were reduced using the online Data Reduction Software (DRS) for the HARPS instrument., The spectroscopic data were reduced using the online Data Reduction Software (DRS) for the HARPS instrument.273 The radial velocity information was obtained by removing the instrumental blaze function and cross-correlating each spectrum with a KS mask., The radial velocity information was obtained by removing the instrumental blaze function and cross-correlating each spectrum with a K5 mask.274 This correlation is compared with the Th-Ar spectrum acting as a reference: see ? ? for details., This correlation is compared with the Th-Ar spectrum acting as a reference; see \citet{Baranne:1996p1069} \citet{Pepe:2002p1068} for details.275 Recently the DRS was shown to achieve remarkable precision (2) thanks to a revision of the reference lines for Thorium and Argon by ?.., Recently the DRS was shown to achieve remarkable precision \citep{Mayor:2009p1088} thanks to a revision of the reference lines for Thorium and Argon by \citet{Lovis:2007p1122}.276 A similar software package is used for CORALIE data., A similar software package is used for CORALIE data.277 A resolving power R=110000 for HARPS yields a cross-correlation function (CCF) binned in 0.25 ss! increments. while for CORALIE. with a lower resolution of 0000. we used 0.5 kkmss7!.," A resolving power $R=110\,000$ for HARPS yields a cross-correlation function (CCF) binned in $0.25$ $^{-1}$ increments, while for CORALIE, with a lower resolution of 000, we used $0.5$ $^{-1}$."278 The CCF window was adapted to be three times the size of the full width at half maximum (FWHM) of the CCF., The CCF window was adapted to be three times the size of the full width at half maximum (FWHM) of the CCF.279 ]c error bars on individual data points were estimated from photon noise alone.," $1\,\sigma$ error bars on individual data points were estimated from photon noise alone."280 HARPS is stable on the long term within mmss! and CORALIE to better than ss!., HARPS is stable on the long term within $^{-1}$ and CORALIE to better than $^{-1}$.281 These are smaller than our individual error bars and thus have not been taken into account., These are smaller than our individual error bars and thus have not been taken into account.282 The absence of a variation 1n bisector span correlated with the phase. or of a variation of the FWHM. indicate that the photometric and spectroscopic signals are indeed those of a planet.," The absence of a variation in bisector span correlated with the phase, or of a variation of the FWHM, indicate that the photometric and spectroscopic signals are indeed those of a planet."283 For comparison we invite the reader to look at ?.. on 441004 for which it has been proven a blend by a star and its brown dwarf companion was causing a spectroscopic Doppler shift similar to that of a planet on a foreground object.," For comparison we invite the reader to look at \citet{Santos:2002p2845}, on 41004 for which it has been proven a blend by a star and its brown dwarf companion was causing a spectroscopic Doppler shift similar to that of a planet on a foreground object."284 The data was fitted using a Markov Chain Monte-Carlo (MCMC) method., The data was fitted using a Markov Chain Monte-Carlo (MCMC) method.285 A code allowing to combine both photometry and spectroscopy has been developed., A code allowing to combine both photometry and spectroscopy has been developed.286 It has been used in several occasions (??) and is deseribed in length in 2..," It has been used in several occasions \citep{Bouchy:2008p229, Gillon:2008p767} and is described in length in \citet{Triaud:2009p7520}."287 It is similar to those presented in ?.., It is similar to those presented in \citet{Cameron:2007p2879}.288 The code uses a common set of free parameters from which physical parameters can be derived to construct models for the photometric and the spectroscopic signal., The code uses a common set of free parameters from which physical parameters can be derived to construct models for the photometric and the spectroscopic signal.289 Free parameters used are: P the period of the object. Το the mid-transit time. D the depth of the transit. W its width. > the Impact parameter. K the semi-amplitude of the Doppler reflex motion by the star.," Free parameters used are: $P$ the period of the object, $T_0$ the mid-transit time, $D$ the depth of the transit, $W$ its width, $b$ the impact parameter, $K$ the semi-amplitude of the Doppler reflex motion by the star."290 To fit the Rossiter-MeLaughlin. we use VVsin/cosB and VVsin7B where Vsin/ ts the projected stellar rotation and f the projected spin-orbit angle.," To fit the Rossiter-McLaughlin, we use $\sqrt{V\,\sin\,I}\,\cos\,\beta$ and $\sqrt{V\,\sin\,I}\,\sin\,\beta$ where $V\,\sin\,I$ is the projected stellar rotation and $\beta$ the projected spin-orbit angle."291 Trying to estimate if the orbit is eccentric we used Vecosw and vesinw where e is the eccentricity and c is the argument of the periastron.," Trying to estimate if the orbit is eccentric we used $\sqrt{e}\,\cos\,\omega$ and $\sqrt{e}\,\sin\,\omega$ where $e$ is the eccentricity and $\omega$ is the argument of the periastron."292 In addition we at times added y. a radial-velocity drift with time. in order to assess the presence of an additional body in the system.," In addition we at times added $\,\dot{\gamma}\,$, a radial-velocity drift with time, in order to assess the presence of an additional body in the system."293 In addition to these we need to fit one normalisation factor for each photometric dataset (5 1n our case) and 2 y velocities for the radial velocities. one for each set.," In addition to these we need to fit one normalisation factor for each photometric dataset (5 in our case) and 2 $\gamma$ velocities for the radial velocities, one for each set."294 We used Gaussian priors to draw randomly each parameter., We used Gaussian priors to draw randomly each parameter.295 We decided to use Vecosw vVesinw as free parameters instead of the more traditional ecosc) esinc because this would amount to impose a prior proportional to e as noted in 2.," We decided to use $\sqrt{e}\,\cos\,\omega$ $\sqrt{e}\,\sin\,\omega$ as free parameters instead of the more traditional $e\,\cos\,\omega$ $e\,\sin\,\omega$ because this would amount to impose a prior proportional to $e^2$ as noted in \citet{Ford:2006p7023}."296.Figure 5. shows the difference between both runs., Figure \ref{fig:pdfecc} shows the difference between both runs.297 Having Vecosw Vesinw makes the eccentricity less biased towards high values.," Having $\sqrt{e}\,\cos\,\omega$ $\sqrt{e}\,\sin\,\omega$ makes the eccentricity less biased towards high values."298 We therefore made a similar change to another pair of variables. making VVsin/cosB and VVsin7B. as free parameters rather than using Vsin/cosB and Vsin/ f. Checks for these were also conducted and validated our choice of jump parameters.," We therefore made a similar change to another pair of variables, making $\sqrt{V\,\sin\,I}\,\cos\,\beta$ and $\sqrt{V\,\sin\,I}\,\sin\,\beta$ as free parameters rather than using $V\,\sin\,I\,\cos\,\beta$ and $V\,\sin\,I\,\sin\,\beta$ Checks for these were also conducted and validated our choice of jump parameters."299bulge and the disks line of nodes of 15. and a bulge ratio Y& = 6.5.,"bulge and the disk's line of nodes of $^{\circ}$, and a bulge mass-to-light ratio $\Upsilon_B$ = 6.5."300" This in turns implies a bulge OLH Mis, = M)23 107""IML...", This in turns implies a bulge of $M_{\rm bulge}$ = 2.3 $\times$ $^{10}$.301" rpThe diskH mass within. the racius of the bulge (3.5 kpe) can be estimated by assuming reasonable values for the BD-band. exponential disk. central surface brightness (fg=21.6 mag arcsec Walterbos Ixennicutt. 1988). which when corrected for. absorption and inclination using the values of Berman (2001) gives a central surface density of X, —113 L. pc5 disk scale length (Ay=5.8 kpe. Walterbos Ixennicutt. 1988). and a conservative disk mass-to-light ratio of 4."," The disk mass within the radius of the bulge (3.5 kpc) can be estimated by assuming reasonable values for the $B$ -band exponential disk central surface brightness $I_0 = 21.6$ mag $^{-1}$, Walterbos Kennicutt 1988) which when corrected for absorption and inclination using the values of Berman (2001) gives a central surface density of $\Sigma_d = $ 113 $_{\odot}$ $^{-3}$, disk scale length $R_d = 5.8$ kpc, Walterbos Kennicutt 1988), and a conservative disk mass-to-light ratio of 4."302 Γης gives a disk mass of Mai = 1.2 10/7NL..., This gives a disk mass of $M_{\rm disk}$ = 1.2 $\times$ $^{10}$.303 The total (bulge | disk | halo) mass inside the same radius can be estimated from the rotation curve as Miu = 3.7 OfAL. leaving little room for a dark halo.," The total (bulge + disk + halo) mass inside the same radius can be estimated from the rotation curve as $M_{\rm total}$ = 3.7 $\times$ $^{10}$, leaving little room for a dark halo."304 The value of R=1.2 found by Berman (2001) implies that the bulge of M31 is a fast rotator., The value of $\mathcal R = 1.2$ found by Berman (2001) implies that the bulge of M31 is a fast rotator.305 Interestingly. of the few galaxies for which ὃν has been measured. all have been found to be fast rotators.," Interestingly, of the few galaxies for which $\mathcal R$ has been measured, all have been found to be fast rotators."306 Debattista Selbwood (2000) showed that such fast. bars cannot co-exist with massive dark halos because cvnamical friction would slow them down rapidly., Debattista Sellwood (2000) showed that such fast bars cannot co-exist with massive dark halos because dynamical friction would slow them down rapidly.307 Phe value of R=1.2 is further evidence that N31 must have a minimal halo. and. consequently. a maximum disk.," The value of $\mathcal R = 1.2$ is further evidence that M31 must have a minimal halo, and, consequently, a maximum disk."308 This is notably at odds with the contentions of Cold Dark Matter (CDM) models. which predict. massive halos and minimal clisks NNavarro ct 11997).," This is notably at odds with the contentions of Cold Dark Matter (CDM) models, which predict massive halos and minimal disks Navarro et 1997)."309 As mentioned earlier. Berman (2001) pointed. out that a much stronger case could be mace for his best. ft model if it were found to be consistent with observations of CO velocities away [rom the line of nodes of the disk.," As mentioned earlier, Berman (2001) pointed out that a much stronger case could be made for his best fit model if it were found to be consistent with observations of CO velocities away from the line of nodes of the disk."310 Moreover. it is important to utilize these olfaxis velocities. since dvnanmical models of the bulge of M31 have never before been constrained by real 3D. CX.Y.V) observations.," Moreover, it is important to utilize these off–axis velocities since dynamical models of the bulge of M31 have never before been constrained by real 3D $(X,Y,V)$ observations."311 ‘To compare accurately the output of the model with the data. we performed a 2D convolution of the model with a realistic description of the PCRAO bean. consisting of the sum of a sinc function and two Caussians.," To compare accurately the output of the model with the data, we performed a 2D convolution of the model with a realistic description of the FCRAO beam, consisting of the sum of a sinc function and two Gaussians."312 The sidelobes of the sine function. provide a σου (very conservative) description of any coma ellects. while the two Gaussians describe the response of the FORAO antenna to emission coming from angles far away from the nominal pointing position. due to the imperfections of the dish (μάς Lever 1996).," The sidelobes of the sinc function provide a good (very conservative) description of any coma effects, while the two Gaussians describe the response of the FCRAO antenna to emission coming from angles far away from the nominal pointing position, due to the imperfections of the dish (Ladd Heyer 1996)."313 The sinc function has a full width at. half maximum (EWII of 45 arcsec and the two Gaussians have, The sinc function has a full width at half maximum (FWHM) of 45 arcsec and the two Gaussians have314"In the unstratified shearing sheet. (he solution for the incompressive shwave is given bv: Ov, = ypον -.. ΓονSa, and lll δει (027 10) where A?=kh?+ he (Ay. 00,5) are the values of (A000) at /=0 and an overdot denotes a time The kinetic enereve. [or a singleex incompressive shwave can be defined as ο2 οδις NDSΕντ an expression which varies with Gime aid peaks al f=0.","In the unstratified shearing sheet, the solution for the incompressive shwave is given by: v_x = v_y = v_x and = = + 2(q - 1) ), where $k^2 = k_x^2 + k_y^2$ , $k_0,\delta v_{x0}$ ) are the values of $k, \delta v_x$ ) at $t=0$ and an overdot denotes a time The kinetic energy for a single incompressive shwave can be defined as E_k _0 v_x^2 + v_y^2) = _0 an expression which varies with time and peaks at $k_x = 0$."315" If one defines an amplification factor for an individual shwave.A= m0) σα ια. ibis apparent that an arbitrary amount of (ransient amplification in the kinetic energy ol an individual shwave can be obtained as one increases (he amount of swing for a leading shwave (hyo< hy)Vy,."," If one defines an amplification factor for an individual shwave, = 1 + it is apparent that an arbitrary amount of transient amplification in the kinetic energy of an individual shwave can be obtained as one increases the amount of swing for a leading shwave $k_{x0} \ll -k_y$ )."316" This transient. amplification of local nonaxisvnuuelric disturbances is reminiscent of the ""swing amplification. mechanism that occurs in disks (hat are marginallv-stable tothe axisvmmeltrie gravitational instability", This transient amplification of local nonaxisymmetric disturbances is reminiscent of the “swing amplification” mechanism that occurs in disks that are marginally-stable tothe axisymmetric gravitational instability317" This transient. amplification of local nonaxisvnuuelric disturbances is reminiscent of the ""swing amplification. mechanism that occurs in disks (hat are marginallv-stable tothe axisvmmeltrie gravitational instability.", This transient amplification of local nonaxisymmetric disturbances is reminiscent of the “swing amplification” mechanism that occurs in disks that are marginally-stable tothe axisymmetric gravitational instability318NGC 2316 is a auch studied bipolar plauctary nebula at a distance DS00 pe (Acker et al.,NGC 2346 is a much studied bipolar planetary nebula at a distance $D\sim 800$ pc (Acker et al.319"199231... At its center lies a binary svstem. formed by a main-sequence star of spectral type ADV. with mass ~La AL... teniperature ~SOOO TS and Iuniuositv ~1E (Ménuudez and Nicuicla 1981: Walsh 1983) and hot star. not detected in the visual. with T,~130000 I& (Ιόνιο 1978). which excites the nebula."," At its center lies a binary system, formed by a main-sequence star of spectral type A5V, with mass $\sim 1.8$ , temperature $\sim 8000$ K and luminosity $\sim 14$ (Ménndez and Niemela 1981; Walsh 1983) and hot star, not detected in the visual, with $\sim 130000$ K (Ménndez 1978), which excites the nebula."320 Its blunünositv is verv uncertain. as we will discuss iu &L.," Its luminosity is very uncertain, as we will discuss in 4."321" Estimates in the literature eive L,~17.90 L.(Ménudez 1978: Calvet aud Peiubert 1983).", Estimates in the literature give $\sim 17-90$ (Ménndez 1978; Calvet and Peimbert 1983).322 Iu the optical. the nebula as a butterfly shape (Balic- 1987: Walsh et al.," In the optical, the nebula has a butterfly shape (Balick 1987; Walsh et al."323 1991). with well developed. bipolar lobes aud a bright torus which surrounds the central star.," 1991), with well developed bipolar lobes and a bright torus which surrounds the central star."324" The temperature aud density of the ionized gas in the torus have been estimate to be T~12000 IK aud D,FOO7. respectivelve (Liu e al."," The temperature and density of the ionized gas in the torus have been estimated to be $T\sim 12000$ K and $n_e\simless 700$, respectively (Liu et al."325 1995: MeIxenua I&eenan 1996)., 1995; McKenna Keenan 1996).326 The nebula contains a large amount of material i the form of molecular gas. as revealed by the CO observations of Knapp (1986). IIugeius Iealy (1986). Tlealyv ILIugeius (1988).," The nebula contains a large amount of material in the form of molecular gas, as revealed by the CO observations of Knapp (1986), Huggins Healy (1986), Healy Huggins (1988)."327 Bachiller ot al. (, Bachiller et al. (328"1989) have mapped the entire nebula iu the two CO lines J=1-0 and J=2-1: the morphology of the molecular gas follows very closely that of the ionized gas. showing a παν. inhomogeucous torus. tilted with respect to the line of sieht by about 56"". which is expaudiug outward.","1989) have mapped the entire nebula in the two CO lines J=1-0 and J=2-1; the morphology of the molecular gas follows very closely that of the ionized gas, showing a clumpy, inhomogeneous torus, tilted with respect to the line of sight by about $^o$, which is expanding outward."329 Scaled to our adopted distance D=800 pc. the torus has a radius of ~0.05 pe. and mass ~ 0.26AD... much larger than the lnass of Jonized eas (~0.01AD... Walsh 1983).," Scaled to our adopted distance $D=800$ pc, the torus has a radius of $\sim$ 0.05 pc, and mass $\sim$ 0.26, much larger than the mass of ionized gas $\sim$ 0.01, Walsh 1983)."330 The radial velocity of the mest intense CO condeusatious is of the order of 15351. which results in a dynamical age of about 2500 vr (but see also Walsh et al.," The radial velocity of the most intense CO condensations is of the order of 15–35, which results in a dynamical age of about 2500 yr (but see also Walsh et al."331 1991)., 1991).332 NGC 2316 is à Type I nebula. originated by a massive progenitor (Calvet and Peimbert 1983).," NGC 2346 is a Type I nebula, originated by a massive progenitor (Calvet and Peimbert 1983)."333 As many PNe of siuuilar type. NCC 2316 is detected in the vibrationally excited lines of IT;(Webster et al.," As many PNe of similar type, NGC 2346 is detected in the vibrationally excited lines of (Webster et al."334 1988)., 1988).335 Zuckerman Gatley (1988) have mapped the nebula in the 1-0S(1) ime at 2.12 ming a ποιοσα 120 arcsec spectrometer with resolution ~ 200., Zuckerman Gatley (1988) have mapped the nebula in the 1-0S(1) line at 2.12 using a single-beam 12 arcsec spectrometer with resolution $\sim$ 200.336 The morphology of the uebula iu this lue is again very simular to the morphology observed iu the optical lines and in CO., The morphology of the nebula in this line is again very similar to the morphology observed in the optical lines and in CO.337 This result wasconfirmed more recently by the nuages obtained iu the same line, This result wasconfirmed more recently by the images obtained in the same line338volume deusitv within a factor of two.,volume density within a factor of two.339 Uneertainties in the estimate of the central deusitv iu addition to variations in the shape of the profile due to the varvine amount of magnetic support can arise through the followingo effects., Uncertainties in the estimate of the central density in addition to variations in the shape of the profile due to the varying amount of magnetic support can arise through the following effects.340 Depending on the detailed conditions iu the parent coud and the details of the formation mechamnisi. the structure of the core at radii significantly larger than the ceutral flat part of the cobuun density profile cau differ appreciably from the profiles preseuted here.," Depending on the detailed conditions in the parent coud and the details of the formation mechannism, the structure of the core at radii significantly larger than the central flat part of the column density profile can differ appreciably from the profiles presented here."341 However. ie deusitv profile iu the most central parts of the core iat we consider here. aud at central deusities m excess of ~LOcm7. has lost. to a laree extent. nienioryv of the details of the core formation mechanism) aud of je pliavsical couitiois at the enveloye.," However, the density profile in the most central parts of the core that we consider here, and at central densities in excess of $\sim 10^5 {\rm \, cm^{-3}}$, has lost, to a large extent, memory of the details of the core formation mechanism and of the physical conditions at the envelope."342 This has been shown bv. e.g. Eug (2002) for the xresence of MIID waves in the pareie oud. Tassis \ouschovias (2007) Or varviug amounts of initial maeuetic support. aud Cong Ostriker (2011) in the case of cores created by colliding turbulent flows.," This has been shown by, e.g., Eng (2002) for the presence of MHD waves in the parent cloud, Tassis Mouschovias (2007) for varying amounts of initial magnetic support, and Gong Ostriker (2011) in the case of cores created by colliding turbulent flows."343 The latter. once collapse starts. exhibit density profiles compatible with models of pure walrodvuanucal collapse (Larson-Peustou solutious).," The latter, once collapse starts, exhibit density profiles compatible with models of pure hydrodynamical collapse (Larson-Penston solutions)."344 For lis reason. and as loug as the analysis is performed very close to the central part of the core. our results are relatively robust to the nature of the processes which are operating at larger scales aud which are responsible for the formation of the core.," For this reason, and as long as the analysis is performed very close to the central part of the core, our results are relatively robust to the nature of the processes which are operating at larger scales and which are responsible for the formation of the core."345 Although more detailed modeling can vield more accurate estimates of the volume deusitv aud au array of other paraineters. the remarkable simplicity of our proposed technique aud its insensitivity to the details of the adopted dynamical model render it a fast and useful tool.," Although more detailed modeling can yield more accurate estimates of the volume density and an array of other parameters, the remarkable simplicity of our proposed technique and its insensitivity to the details of the adopted dynamical model render it a fast and useful tool."346 We have verified that the results obtained iu wellstudied molecular cloud cores LI5tL and D68. are consistent with other. more detailed estimates available in the literature.," We have verified that the results obtained in well-studied molecular cloud cores L1544 and B68 are consistent with other, more detailed estimates available in the literature."347"sPhe formation⋅. and evolution. of. carly-type galaxics.. cDaractertzec|Utwed d» Alla dominantionim central stellar""| bulneeal component.| Cjis still |a matter: of (deldebate.","The formation and evolution of early-type galaxies, characterized by a dominant central stellar bulge component, is still a matter of debate."348", m1ο semina—. paper by.?.. postulated that stars are formed) in. a single burst of⋅ star formation⋅. from⋅ gas FallingDa towards the center. and the evolution. is. passive. thereafter."," The seminal paper by\citet[][]{Eggen62}, postulated that stars are formed in a single burst of star formation from gas falling towards the center, and the evolution is passive thereafter."349⋅ Such. a simple. scenario might be dillicult to reconcile with the standard cosmological paradigm of structure formation. in. which dark matter halos grow hierarchically: through merging.," Such a simple scenario might be difficult to reconcile with the standard cosmological paradigm of structure formation, in which dark matter halos grow hierarchically through merging."350": Phe most advanced. and up-to-date Semi-analytical models still clo not completely agree on the type of evolution MM¢ byJU stfalas¢3h N h H . the fraction Ol4 stellarar massass pussylornmnecformed in the initial.(etial. igas-rich@as-Lic.) ""bur MUSÉοἱ star formation. and on the role plavecl by the late evolution driven by major and minor mergers."," The most advanced and up-to-date Semi-analytical models (SAMs) of galaxy formation \citep[e.g.,][]{Cole00,Benson03,Granato04,Granato06,Menci04,Ciras05,Khochfar05,Vittorini05,Bower06,Cattaneo06,Croton06,DeLucia06,Hopkins06,Lapi06,Shankar06,Monaco07,Somerville08SAM,Cook09,Fontanot09} still do not completely agree on the type of evolution undergone by massivegalaxies (see also \citealt{Dekel09}) ), on the fraction of stellar mass formed in the initial, gas-rich burst of star formation, and on the role played by the late evolution driven by major and minor mergers."351 Nevertheless. all models agree that. galaxies.. must have been much more compact at the epoch of ⋅⋅formation.⊀ owing⊀ to a denser universe.. larecr σας fractions⋅⋠ in⋠ the progenitors.⋠ and. more dissipation.," Nevertheless, all models agree that galaxies must have been much more compact at the epoch of formation, owing to a denser universe, larger gas fractions in the progenitors, and more dissipation."352VEN ..“Phe latter prediction⊀⊀ has been confirmed:⋅ by a number of⋅ deep observations⊀ of ⋅⊀high τοςκ. ealaxies. (eg.T?TT). which have independently found carly-tvpe hieh-redshilt. massive ealaxies to bea [ποσο of a few more compactompac thanE localval counterpartscoparts ofο the| same stellarstell: massass.," The latter prediction has been confirmed by a number of deep observations of high redshift galaxies \citep[e.g.,][]{Trujillo06,Vandokkum08,Cimatti08,Saracco08}, which have independently found early-type, high-redshift, massive galaxies to be a factor of a few more compact than local counterparts of the same stellar mass."353 Note. however. that several observational biases nieht limit the quality. and reliability⊀↔ of⋅ some of⋅ these measurements (og.?2?7)..," Note, however, that several observational biases might limit the quality and reliability of some of these measurements \citep[e.g.,][]{HopkinsRez,Mancini09,vandokkum09ALL}."354 lt is. still. debated. how these compact galaxies.H have evolved from⋅ hisgh-redshifts. ⋠⋅⋠increasing⋠ their ↶↶sizes in⋠ à wav to fall on the size-mass relation we observe today., It is still debated how these compact galaxies have evolved from high-redshifts increasing their sizes in a way to fall on the size-mass relation we observe today.355 As already extensively discussed by. e.g. 2.. 2.. (2009a).. the scatter around the local. median," As already extensively discussed by, e.g., \citet{Shen03}, \citet{ShankarBernardi}, , , the scatter around the local, median"356The fundamental role Type Ia supernovae (SNe la) play in determining the expansion history of the universe (Riessetal.2004.andreferencestherein) has added new urgency to understanding the progenitors of these events.,The fundamental role Type Ia supernovae (SNe Ia) play in determining the expansion history of the universe \citep[][and references therein]{rie04} has added new urgency to understanding the progenitors of these events.357 Theoretical work on modeling the flame dynamics of exploding white dwarfs (WDs) has demonstrated that the composition and energy of ejecta depend sensitively on the competition between flame propagation. instabilities driven by turbulence. and expansion of the WD (e.g..Hillebrandt&Niemeyer 2000).," Theoretical work on modeling the flame dynamics of exploding white dwarfs (WDs) has demonstrated that the composition and energy of ejecta depend sensitively on the competition between flame propagation, instabilities driven by turbulence, and expansion of the WD \citep[e.g.,][]{hn00}."358.. This has sparked new interest in understanding the evolution that occurs up until the onset of explosive burning., This has sparked new interest in understanding the evolution that occurs up until the onset of explosive burning.359 A critical aspect is the shear profile present within the WD., A critical aspect is the shear profile present within the WD.360 SNe [a are expected to occur exclusively in accreting binary systems., SNe Ia are expected to occur exclusively in accreting binary systems.361 Therefore the role of angular momentum may be unique or at least especially important for these SNe., Therefore the role of angular momentum may be unique or at least especially important for these SNe.362 If aceretion drives shear in the WD. it may be an important source of viscous heating and material mixing.," If accretion drives shear in the WD, it may be an important source of viscous heating and material mixing."363 Yoon&Langer(2004) considered the accretion and evolution of a WD up until the point of ignition. including the effects of rotation.," \citet{yl04} considered the accretion and evolution of a WD up until the point of ignition, including the effects of rotation."364 In their study angular momentum was transported primarily via the Kelvin-Helmholtz instability. which lead to significant shear throughout the WD.," In their study angular momentum was transported primarily via the Kelvin-Helmholtz instability, which lead to significant shear throughout the WD."365 Saio&Nomoto(2004) focused on whether rotational effects can prevent accretion induced collapse (ALC) at high accretion rates (M=3«10-°—10?M..yr! )., \citet{sn04} focused on whether rotational effects can prevent accretion induced collapse (AIC) at high accretion rates $\dot{M}=3\times10^{-6}-10^{-5}\ M_\odot\ {\rm yr^{-1}}$ ).366 They found that an AIC was not averted. but perhaps more interestingly. they also found the WD should be rotating nearly uniformly when the baroclinic instability is included. a process that was neglected by Yoon&Langer (2004).," They found that an AIC was not averted, but perhaps more interestingly, they also found the WD should be rotating nearly uniformly when the baroclinic instability is included, a process that was neglected by \citet{yl04}."367 There is further opportunity for shear to develop during the carbon simmering phase., There is further opportunity for shear to develop during the carbon simmering phase.368 At carbon ignition. a convective region grows at the WD center.," At carbon ignition, a convective region grows at the WD center."369 This envelops ~1M over a timescale of ~1000yrs. until a burning wave commences...," This envelops $\sim1\ M_\odot$ over a timescale of $\sim1000\ {\rm yrs}$, until a burning wave commences."370 This phase has received increasing attention in recent years due tothe realization that it sets the initial temperature and density for the explosive burning (Lesaffreetal.2006).. as well as the distribution of ignition points (Woosleyetal.2004:Wunsch&WoosleyKuhlenetal. 2006).," This phase has received increasing attention in recent years due tothe realization that it sets the initial temperature and density for the explosive burning \citep{les06}, as well as the distribution of ignition points \citep{woo04,ww04,kuh06}."371. Nuclear reactions during this time may change the neutron excess of the WD core. which is crucial for determining the ratio of radioactive to non-radioactive nickel produced in the Type Ia explosion (Piro&Bildsten2007b:Chamulaketal. 2007).," Nuclear reactions during this time may change the neutron excess of the WD core, which is crucial for determining the ratio of radioactive to non-radioactive nickel produced in the Type Ia explosion \citep{pb07b,cha07}. ."372 In the present study we consider the opportunity. for developing shear during both of these stages., In the present study we consider the opportunity for developing shear during both of these stages.373 In 82. we revisit estimates for the shear of accreting WDs and confirm that the baroclinic instability limits the shear before the Kelvin-Helmholtz instability initiates., In \ref{sec:accretion} we revisit estimates for the shear of accreting WDs and confirm that the baroclinic instability limits the shear before the Kelvin-Helmholtz instability initiates.374 In addition we explore whether magnetohydrodynamie effects could reduce the shear even further., In addition we explore whether magnetohydrodynamic effects could reduce the shear even further.375 We conclude that the WD is nearly uniformly rotating at the onset of carbon ignition. and that viscous heating is negligible.," We conclude that the WD is nearly uniformly rotating at the onset of carbon ignition, and that viscous heating is negligible."376 In $3. we explore the properties of convection present prior to explosive burning., In \ref{sec:convection} we explore the properties of convection present prior to explosive burning.377 The growth of this convective region allows further opportunity for shearing., The growth of this convective region allows further opportunity for shearing.378 The three-dimensional nature of this problem makes 1t hard to definitively determine what occurs during this stage., The three-dimensional nature of this problem makes it hard to definitively determine what occurs during this stage.379 We illustrate some general features expected for the interaction of spin and convection by appealing to observations and numerical experiments., We illustrate some general features expected for the interaction of spin and convection by appealing to observations and numerical experiments.380 In $4 we conclude with a summary of our results and a discussion of future research., In \ref{sec:theend} we conclude with a summary of our results and a discussion of future research.381 Aaterial accreted at a rate M reaches the WD surface with a nearly Keplerian spin frequency of Oy2(GM/R?)7., Material accreted at a rate $\dot{M}$ reaches the WD surface with a nearly Keplerian spin frequency of $\wk=(GM/R^3)^{1/2}$.382 The majority of the kinetic energy associated with this flow is dissipated in a boundary layer of thickness Πμι«R (asstudiedbyPiro&Bildsten2004) and never reaches far into the surface., The majority of the kinetic energy associated with this flow is dissipated in a boundary layer of thickness $H_{\rm BL}\ll R$ \citep[as studied by][]{pb04} and never reaches far into the surface.383 Nevertheless. angular momentum is added at arate MR Οκ. so a torque of this magnitude should be communicated into the WD.," Nevertheless, angular momentum is added at arate $\dot{M}R^2\wk$ , so a torque of this magnitude should be communicated into the WD."384 It remains an open question, It remains an open question385With this anecdote in munud I recall some of the early developments iu our soundings of the Sum.,With this anecdote in mind I recall some of the early developments in our soundings of the Sun.386 LEiuust attempt to confine attention to our inerences: from low-deerce modes alone. appreciating the «iffieulv of attempting to ignore the bias that was inevitably exerted bv what we had learned from modes of high degree. kuowledee of a kind that will rot be obtainable from other stars.," I must attempt to confine attention to our inferences from low-degree modes alone, appreciating the difficulty of attempting to ignore the bias that was inevitably exerted by what we had learned from modes of high degree, knowledge of a kind that will not be obtainable from other stars."387 Mau of our early iiferencees were derived from calibrations of solar ποσο]. a necessary procedure because localized information (ciher iu coufiguraion space or in the space of more geueral euquirv) genuinely uncontamineatccL by elobal properties is almost impossible to obtain Yolu modes of only low degree.," Many of our early inferences were derived from calibrations of solar models, a necessary procedure because localized information (either in configuration space or in the space of more general enquiry) genuinely uncontaminated by global properties is almost impossible to obtain from modes of only low degree."388 Moreover. accurate nonseisuiic information of the kind we have about the Sun. such as lnass and radius. will uot always be available.," Moreover, accurate nonseismic information of the kind we have about the Sun, such as mass and radius, will not always be available."389" We must bear in nuid tha seine data alone can at best provide informatioi about ouly he functional forms of sou speed aud density with respect to fractional radius. which we nieht use. wit ithe hel» of some idea about the value of the first acjabatic exponeit 24 (which for stars compose of ""ordinary matter is close to 5/3 evervwhere excep in the outer lavers where the effects of ionization of the abundant elements are naxortant). to create a scismuically calibrated model of the star."," We must bear in mind that seismic data alone can at best provide information about only the functional forms of sound speed and density with respect to fractional radius, which we might use, with the help of some idea about the value of the first adiabatic exponent $\gamma _1$ (which for stars composed of `ordinary' matter is close to 5/3 everywhere except in the outer layers where the effects of ionization of the abundant elements are important), to create a seismically calibrated model of the star."390 Nouseisiuic infoxiuation. be i in the form: of observation or prejudice (otherwise known as prior information). is essential to ‘complete’ the picture.," Nonseismic information, be it in the form of observation or prejudice (otherwise known as prior information), is essential to `complete' the picture."391 The first solar calibration from low-deeree imiodes was carried out by Jorgen Christenscn-Dalseaard and ie for the purpose of estimating t1ο helimm abundance Y. a quantity of extreme importance at that time to the solar neutrik» problein.," The first solar calibration from low-degree modes was carried out by rgen Christensen-Dalsgaard and me for the purpose of estimating the helium abundance $Y$, a quantity of extreme importance at that time to the solar neutrino problem."392 Two differcut fits to the dai were better fian the rest: they were about equalv good (although not as good as we had hoped). so the calibration was indeterminate.," Two different fits to the data were better than the rest; they were about equally good (although not as good as we had hoped), so the calibration was indeterminate."393 One had Y above our expecation. the other beknw.," One had $Y$ above our expectation, the other below."394 The reasol hat ueither of TOS was as good as we had hope was (partlv) fji the surface lavers iad been inadequately inodelled. a natter which was not of priucipal interest after all. those lavers are irelevaut to neutrino production — but which impeded our diagnostic endeavours.," The reason that neither of the fits was as good as we had hoped was (partly) that the surface layers had been inadequately modelled, a matter which was not of principal interest – after all, those layers are irrelevant to neutrino production – but which impeded our diagnostic endeavours."395 Actually. we knew which o| the two fits to choose. from an earlier calibration of the convection zone with high-deerce data. but that was outside the information space within which asteroseisiiologists can work. so I cannot discuss it here.," Actually, we knew which of the two fits to choose, from an earlier calibration of the convection zone with high-degree data, but that was outside the information space within which asteroseismologists can work, so I cannot discuss it here."396 Nevertheless. I zu renduded of some of the Bayesian discussions at this workshop.," Nevertheless, I am reminded of some of the Bayesian discussions at this workshop."397 Wad we used Bayesian statisics With à notrino-flux-sed prior. we would have chose- he wrong fit.," Had we used Bayesian statistics with a neutrino-flux-based prior, we would have chosen the wrong fit."398 Tha firs calibration was extremely naive: it sought to fid the inodeI that vest repronuces all he frequency data iuciscrininaely. baselona philosophy which. broadly speaking. is still in use today to provide “inversions” by (reguarized) data fitting by Cast SqlOYCS. a Procelure COMMLOilv ktown as regularized least squares’.," That first calibration was extremely naive: it sought to find the model that best reproduces all the frequency data indiscriminately, based on a philosophy which, broadly speaking, is still in use today to provide `inversions' by (regularized) data fitting by least squares, a procedure commonly known as `regularized least squares'."399 Τιday we seek ryequency combinations that are sjeuaros of SOLO secifüie featíre Q| the structure of the sar tha we nüeght wish to iwestigate., Today we seek frequency combinations that are signatures of some specific feature of the structure of the star that we might wish to investigate.400 An example is to trauslate the 1iner pliase of oscillation of some stellar mode into à ueastrable frequeicy signature. in the nianner proftered by Eui Roxburg rand Sergei Vorontsov.," An example is to translate the inner phase of oscillation of some stellar model into a measurable frequency signature, in the manner proffered by Ian Roxburgh and Sergei Vorontsov."401 Chiistensen-Dalseaare has done tlat at this workshop for a Procvon-likο Wael to reproduc f1ο behaviour of the siiall frequeacy separations repored w Tin Bedding., Christensen-Dalsgaard has done that at this workshop for a Procyon-like model to reproduce the behaviour of the small frequency separations reported by Tim Bedding.402 He is no doubt οOrrec in clainuug flat the convective core provides the explavation of the curvative of Beddine’s echelle plot. but I here play devil's acvocate to warn that such behaviour cotld also he reproducec by appropriate sole woul rightvy say contrived aspherical structure in the outer avers of the star.," He is no doubt correct in claiming that the convective core provides the explanation of the curvature of Bedding's echelle plot, but I here play devil's advocate to warn that such behaviour could also be reproduced by appropriate – some would rightly say contrived – aspherical structure in the outer layers of the star."403 We must invoke what one nuelt misnarne a xior (actually a posterior coustraint) to reject the latter., We must invoke what one might misname a prior (actually a posterior constraint) to reject the latter.404 The same commen applies to the lneasures of he different integralC» properties of the core (which do ixot require a specific stellar model for their cdeteruinuatioi) that were discussed lere by Cotter IIloudek iu the coutext of the Sun., The same comment applies to the measures of the different integral properties of the core (which do not require a specific stellar model for their determination) that were discussed here by Günnter Houdek in the context of the Sun.405 I nüsht poiut out that in the early days onlookers regarded our eudeavours to be impossible., I might point out that in the early days onlookers regarded our endeavours to be impossible.406 I recall. for exinuple. eiviug a lecture explaining how we expecte to determine the solar Y directly by measuring the effect of helm ionization on 54.," I recall, for example, giving a lecture explaining how we expected to determine the solar $Y$ directly by measuring the effect of helium ionization on $\gamma _1$."407 Iu the audieuce was Donald Lyudeu-Bell. who cousidered it impossible: he wagered that even after ten vears’ work we will not have suceceeced.," In the audience was Donald Lynden-Bell, who considered it impossible; he wagered that even after ten years' work we will not have succeeded."408 Ile was right., He was right.409 But not for the right reason., But not for the right reason.410 We id actually achieved the impossible: iu fact. we had meastped," We had actually achieved the impossible: in fact, we had measured"411The NICMOS UDF Treasury observations were designed to complement and enhance the ACS optical UDF observations.,The NICMOS UDF Treasury observations were designed to complement and enhance the ACS optical UDF observations.412 They provide an extension in wavelength (o 1.6jun and provide two additional bands which extend the rest band energv and morphology measurements (o longer wavelengths., They provide an extension in wavelength to $1.6 \micron$ and provide two additional bands which extend the rest band energy and morphology measurements to longer wavelengths.413 They also provide the potential lor viewing objects ab redshilts bevond 7.5 where Lyman line anc continuum. absorption quench the flux in ihe ACS bands., They also provide the potential for viewing objects at redshifts beyond 7.5 where Lyman line and continuum absorption quench the flux in the ACS bands.414 The additional wawelength coverage also helps distinguish between the influences of age. metallicity ancl extinction.," The additional wavelength coverage also helps distinguish between the influences of age, metallicity and extinction."415 As with previous deep field catalog publications (Williamsοἱal. (1996).. Thompsonetal.(1999) and Beckwithetal. (2004))) this paper is intended primarily as a description of the observations. data analysis aid source photometric properties rather (han a scientific evaluation of the implications of the observations.," As with previous deep field catalog publications \citet{wil96}, , \citet{thm99} and \citet{beck04}) ) this paper is intended primarily as a description of the observations, data analysis and source photometric properties rather than a scientific evaluation of the implications of the observations."416" Due to the relatively small field of the NICMOS camera 3 used in this program 51). only a subsection (144""x 1447) of the optical UDF was covered."," Due to the relatively small field of the NICMOS camera 3 used in this program $51\arcsec 417\times 51\arcsec$ ), only a subsection $144\arcsec \times 144\arcsec$ ) of the optical UDF was covered."418 This was done with a 3x3 Gling of the NICAIOS images., This was done with a 3x3 tiling of the NICMOS images.419 The NICMOS images extend a few arc seconds bevond {his subsection but al a significantly decreased signal to noise., The NICMOS images extend a few arc seconds beyond this subsection but at a significantly decreased signal to noise.420 All of the individual processed NICMOS images are available in the HIST treasury archive (MASTH)) as are the raw images in the main IIST archive., All of the individual processed NICMOS images are available in the HST treasury archive ) as are the raw images in the main HST archive.421 The primary purpose of (his paper is (o provide a very detailed account. of the data reduction steps used to produce the treasury. images and catalog stored in NLAST so that users are aware of the pedigree of (he data and can reproduce (he analvsis if they wish., The primary purpose of this paper is to provide a very detailed account of the data reduction steps used to produce the treasury images and catalog stored in MAST so that users are aware of the pedigree of the data and can reproduce the analysis if they wish.422 Other users may wish to alter the reduction steps if (hev prefer other choices than theones made, Other users may wish to alter the reduction steps if they prefer other choices than theones made423ull σοιtribution aud the quadratic term of the equation. respectively.,"full contribution and the quadratic term of the equation, respectively."424 The middle panel displays he residuals from the linear aud quadratic terms. aud the bottom panel the CCD residuals row he full ephemeris.," The middle panel displays the residuals from the linear and quadratic terms, and the bottom panel the CCD residuals from the full ephemeris."425 These appear as O-C' ui in the fourth column of Table 3., These appear as $O$ $C_{\rm full}$ in the fourth column of Table 3.426 As indicated by the igure. he quadraticplus LIT ephemeris eives a satisfactory representation o“the ensemlheo “the 'esicduas.," As indicated by the figure, the quadratic LTT ephemeris gives a satisfactory representation of the ensemble of the residuals."427 In addition. another moculatoÜiowith a period o. about five vrs aid a semii-aripliucle of abot| 0.000E clneg exist in the O-C' ul residuals.," In addition, another modulaton with a period of about five yrs and a semi-amplitude of about 0.0004 d exist in the $O$ $C_{\rm full}$ residuals."428 Ao laree number of futIre accurate timings is require before his can be tested at :ui acceptable leve., A large number of future accurate timings is required before this can be tested at an acceptable level.429 I£ the orbit of the outer coliipolet is coplanar with tlat of tlie close eclipsilg pair (/4288.2). its thass is abott Aly=0.83 l. aud COLTeESPOles to a spectral type of IXK1-IN2 [or a normal mait-sequeuce star.," If the orbit of the outer component is coplanar with that of the close eclipsing pair $i_3$ $^\circ$ .2), its mass is about $M_3$ =0.83 $_\odot$ and corresponds to a spectral type of K1–K2 for a normal main-sequence star."430 Tus would coiriute about to the otal light of the triple system. so it will e difficult to detect such a third ligh source [rom light-cuve analysis.," This would contribute about to the total light of the triple system, so it will be difficult to detect such a third light source from light-curve analysis."431 Tle ¢uacratic term (A) of the ephemeris indicates a continuous period iicrease at a rate of --(1.12:0.2 x10 '« +t., The quadratic term (A) of the ephemeris indicates a continuous period increase at a rate of $\pm$ $\times$ $^{-7}$ d $^{-1}$.432 Because CL Aur is a semi-detached system with the less massive auc coo secondary component filling its inner Roche lobe (cL., Because CL Aur is a semi-detached system with the less massive and cool secondary component filling its inner Roche lobe (cf.433 Section [). its Roche-geouetry configuration permits mass transer from the secondary to the more massive primary star. aldit is conventiona to ascribe such an inerease to conservative mass transfer between the stars iu the system.," Section 4), its Roche-geometry configuration permits mass transfer from the secondary to the more massive primary star, and it is conventional to ascribe such an increase to conservative mass transfer between the stars in the system."434 The calculated mass transfer rate is of the order of 1.3x10 “ONL. voL amone tie largest rates [or Aleoltype systems.," The calculated mass transfer rate is of the order of $\times$ $^{-7}$ $_\odot$ $^{-1}$, among the largest rates for Algol-type systems."435 Alternatively. the 21.6 vr oscillation in the O-C' residuals could be cause by magnetice mocdulatic)li due to au activity cycle in the convective envelope of the late-type star (Applegate 1992. Lanza et al.," Alternatively, the 21.6 yr oscillation in the $O$ $C$ residuals could be caused by magnetic modulation due to an activity cycle in the convective envelope of the late-type star (Applegate 1992, Lanza et al."436 1998)., 1998).437 The hot prijury component of CL Aur likely bas a radiative envelope as surmised from its spectral type while the less massive aud cool secoudary should have as allow convective shell aud at most weak magletic activity., The hot primary component of CL Aur likely has a radiative envelope as surmised from its spectral type while the less massive and cool secondary should have a shallow convective shell and at most weak magnetic activity.438 We applied the period (23) and amplitude (A to Applegate's formulae and obtained model parameters for possible magnetic activity., We applied the period $P_3$ ) and amplitude $K$ ) to Applegate's formulae and obtained model parameters for possible magnetic activity.439 Te results are listed in Table 5. where the boOnmetric magiitude cliffereice (Nina) relative to the meat light level was obtained with equation (1) in the payer of Ixim οἱ al. (," The results are listed in Table 5, where the bolometric magnitude difference $\Delta m_{\rm rms}$ ) relative to the mean light level was obtained with equation (4) in the paper of Kim et al. ("4401997).,1997).441" lost of the abtlatect values are close to those derived for several other bitaries that apyear LO supsort Appleg:1tes theory.""N. btt the light variation predicted [ron aL active CἘν Aur seconeary is at tLe upper litii ο ‘thetheoretical value CLLa~ 0.1) proposed by him."," Most of the tabulated values are close to those derived for several other binaries that appear to support Applegate's theory, but the light variation predicted from an active CL Aur secondary is at the upper limit of the theoretical value $\Delta L/L_2 \sim 0.1$ ) proposed by him."442" 5ince our obse'""vatious represent the first conipele light curve. we canuot test whether such a variaion has occurred in the yast."," Since our observations represent the first complete light curve, we cannot test whether such a variation has occurred in the past."443 However. because the secondary star is expected not to be strongly magnetically aclive. we tlink the mos probable explauation of the periodic oscillation to be the LTT effect due ο a low-luiHnosity Ix-tvpe tertiary Companion.," However, because the secondary star is expected not to be strongly magnetically active, we think the most probable explanation of the periodic oscillation to be the LTT effect due to a low-luminosity K-type tertiary companion."444 As shown in Figure 1. our observations clearly indicate that the light curve morphology of CL Aur is not 3 Lyr type but rather very similar to that of Aleol. its class prototype.," As shown in Figure 1, our observations clearly indicate that the light curve morphology of CL Aur is not $\beta$ Lyr type but rather very similar to that of Algol, its class prototype."445 To understaud the geometrical structure and the physical parameters of the system. our BYRL light curves were solved simultaneously in a manner similar to those lor XX Cep (Lee et al.," To understand the geometrical structure and the physical parameters of the system, our $BVRI$ light curves were solved simultaneously in a manner similar to those for XX Cep (Lee et al."446 2007) aud GW Gem (Lee, 2007) and GW Gem (Lee447antenna-based phase errors; the corrections thus derived were applied to all the frequency bands thereby ensuring that the positional registration of the continuum emission. when imaged. was anchored to the position of the peak of the strong line emission.,"antenna-based phase errors; the corrections thus derived were applied to all the frequency bands thereby ensuring that the positional registration of the continuum emission, when imaged, was anchored to the position of the peak of the strong line emission."448 The continuum emission in Arp220 is problematic from a VLBI imaging perspective in (hat the structure has a very. dillerent character depending upon Che spatial resolution with which it is viewed., The continuum emission in Arp220 is problematic from a VLBI imaging perspective in that the structure has a very different character depending upon the spatial resolution with which it is viewed.449 As we have previously shown with MERLIN (Rovilos et al 2003). on arc-second scales we see all of the expected ~280mJv flix from Arp 220 (Baan. Wood llaschick 1932).," As we have previously shown with MERLIN (Rovilos et al 2003), on arc-second scales we see all of the expected $\sim$ 280mJy flux from Arp 220 (Baan, Wood Haschick 1982)."450 About 200 mJv of this is concentrated in two main components separated by 1 aresecond almost east-west. which are surrounded by more extended emission.," About 200 mJy of this is concentrated in two main components separated by 1 arcsecond almost east-west, which are surrounded by more extended emission."451 Lowever. our early VLBI observations (Sith et al..," However, our early VLBI observations (Smith et al.,"452 1998a) showed that on baselines bevond a [ew hundred kilometres most of the MEBRLIN-scale structure was resolved and we only see a few percent of the flux concentrated in a number of point sources., 1998a) showed that on baselines beyond a few hundred kilometres most of the MERLIN-scale structure was resolved and we only see a few percent of the flux concentrated in a number of point sources.453 Therelore. in producing the continuum image from (his dataset. which suffers from poor (u.v) coverage on short baselines. we must take care when imaging.," Therefore, in producing the continuum image from this dataset, which suffers from poor (u,v) coverage on short baselines, we must take care when imaging."454 We found that the best approach was to simply eliminate the short baselines and the dilfuse emission detected on them. by ignoring all data within a (uv) range of 5000kA. i.e. projected. baseline lengths of less than 900km.," We found that the best approach was to simply eliminate the short baselines and the diffuse emission detected on them, by ignoring all data within a (u,v) range of $\lambda$, i.e. projected baseline lengths of less than 900km."455 This results in only minimal loss of sensitivitv {ο compact features. because (he largest apertures in the telescope array tend to be al the periphery of the array. ancl tend not to participate in many short baselines.," This results in only minimal loss of sensitivity to compact features, because the largest apertures in the telescope array tend to be at the periphery of the array, and tend not to participate in many short baselines."456 Regions of the final image centered on the (wo nuclei of Arp 220 are shown in Figure 1., Regions of the final image centered on the two nuclei of Arp 220 are shown in Figure 1.457 The rms noise across the image varies somewhat with location due to dvnamic range limitations in areas with a high density of compact sources., The rms noise across the image varies somewhat with location due to dynamic range limitations in areas with a high density of compact sources.458 Llowever. we find (hat in the emptiest areas of the image the rms noise is 5.5 j(Jv/beam. and close to the NW region of compact structure il approaches 9 j(Jv/beam: this image is. therefore. the most sensitive VLBI image obtained to date bv some margin.," However, we find that in the emptiest areas of the image the rms noise is 5.5 $\mu$ Jy/beam, and close to the NW region of compact structure it approaches 9 $\mu$ Jy/beam: this image is, therefore, the most sensitive VLBI image obtained to date by some margin."459 Many. point sources are clearly visible., Many point sources are clearly visible.460 Due to the variable noise levels and the large number of [nint source candidates. objective criteria for identifving and measuring sources were recquired. as described in the next section.," Due to the variable noise levels and the large number of faint source candidates, objective criteria for identifying and measuring sources were required, as described in the next section."461 The locus of this investigation is (he population of point sources in the (wo nuclei ol Arp 220., The focus of this investigation is the population of point sources in the two nuclei of Arp 220.462 For a reliable ancl complete source list. we need objective and robust. criteria [or source detection.," For a reliable and complete source list, we need objective and robust criteria for source detection."463 A careful treatment is required. because the deconvolution step of image generation. via the CLEAN algorithm. is highlv non-linear in nature. contributing to strongly non-Gaussian noise statistics in the final image: another factor that may influence the statistics is residual poorly constrained diffuse continuum emission.," A careful treatment is required, because the deconvolution step of image generation, via the CLEAN algorithm, is highly non-linear in nature, contributing to strongly non-Gaussian noise statistics in the final image; another factor that may influence the statistics is residual poorly constrained diffuse continuum emission."464 The statistic we, The statistic we465applied monotonized central (MC) limiter to the conductive fluxes.,applied monotonized central (MC) limiter to the conductive fluxes.466 This method ensures that anisotropic conduction does not lead to negative temperatures in the presence of steep temperature gradients., This method ensures that anisotropic conduction does not lead to negative temperatures in the presence of steep temperature gradients.467" The three-dimensional computational domain was approximately 1 Mpc on each side, enclosing a large fraction of the cluster."," The three-dimensional computational domain was approximately 1 Mpc on each side, enclosing a large fraction of the cluster."468 The central regions of the cluster had an enhanced refinement level., The central regions of the cluster had an enhanced refinement level.469 The maximum spatial resolution for 6 levels of refinement was ~2.7h-! kpc., The maximum spatial resolution for 6 levels of refinement was $\sim 2.7h^{-1}$ kpc.470" The simulations were performed on a 384-processor cluster located at the Michigan Academic Computing Center at the University of Michigan in Ann Arbor and on theColumbia supercomputer at NASA We performed total of 16 runs including radiative cooling, anisotropic thermala conduction and self-gravitating particles to emulate the gas ""stirring"" by galaxies."," The simulations were performed on a 384-processor cluster located at the Michigan Academic Computing Center at the University of Michigan in Ann Arbor and on the supercomputer at NASA We performed a total of 16 runs including radiative cooling, anisotropic thermal conduction and self-gravitating particles to emulate the gas “stirring” by galaxies."471" We considered a uniform grid of parameters: 50, 100, 150, and 200 galaxies characterized by masses of (0.3, 0.6, 0.9, 1.2) x10? Mo."," We considered a uniform grid of parameters: 50, 100, 150, and 200 galaxies characterized by masses of (0.3, 0.6, 0.9, 1.2) $\times 10^{12}$ $_{\odot}$."472" With our cluster mass of 6.6x1014Mo, these parameters corresponds to a mass fraction in galaxies ranging from feal=2.2% to a maximum of [και=27%."," With our cluster mass of $6.6 \times 10^{14} \, {\rm M_{\odot}}$, these parameters corresponds to a mass fraction in galaxies ranging from $f_{\rm gal}=2.2\%$ to a maximum of $f_{\rm gal}=27\%$."473" For instance, for galaxies with mass 6x10!!Mo, our grid corresponds to feat=(4.3,8.3,12,15)96."," For instance, for galaxies with mass $6 \times 10^{11} \, {\rm M_{\odot}}$, our grid corresponds to $f_{\rm gal}=(4.3, 8.3, 12, 15) \%$."474" We also performed two control runs: one without the galaxies (and hence without stirring) to isolate the effect of heat buoyancy instability, and one without conduction to isolate the effect of dynamical friction heating by galaxies."," We also performed two control runs: one without the galaxies (and hence without stirring) to isolate the effect of heat buoyancy instability, and one without conduction to isolate the effect of dynamical friction heating by galaxies."475 Figure 3 (left panel) shows the evolution of the velocity dispersion measured within 100 kpc from the cluster center., Figure \ref{fig:velocities} (left panel) shows the evolution of the velocity dispersion measured within 100 kpc from the cluster center.476 Thin blue (red) lines are for 100 (200) galaxies respectively and for equally spaced masses ranging from 3x1011 to 1.2x1013., Thin blue (red) lines are for 100 (200) galaxies respectively and for equally spaced masses ranging from $3\times 10^{11}$ to $1.2\times 10^{12}$.477 The mass increases gradually from the lightest to the darkest color., The mass increases gradually from the lightest to the darkest color.478 The black dashed line is for the pure HBI case., The black dashed line is for the pure HBI case.479 The HBI case and lighter-colored curves are evolved for shorter times., The HBI case and lighter-colored curves are evolved for shorter times.480 These runs suffer from overcooling and the central temperatures reaches the low temperature threshold at which point the simulation is stopped., These runs suffer from overcooling and the central temperatures reaches the low temperature threshold at which point the simulation is stopped.481 The right panel in Figure 3 shows the median velocity within 100 kpc; the color coding corresponds to that in the left panel., The right panel in Figure \ref{fig:velocities} shows the median velocity within 100 kpc; the color coding corresponds to that in the left panel.482 It is clear from these figures that there is a clear trend for the velocity dispersion or the median velocity to increase with the typical galaxy mass., It is clear from these figures that there is a clear trend for the velocity dispersion or the median velocity to increase with the typical galaxy mass.483" A similar, albeit weaker, trend is seen for the galaxy number."," A similar, albeit weaker, trend is seen for the galaxy number."484" This is consistent with the findings of Kim (2007) who found in pure hydrodynamic simulations, that the gas velocity dispersion c scales as σοςNui?Mgai, where Ngai and Mgai are the number and mass of galaxies, respectively."," This is consistent with the findings of Kim (2007) who found in pure hydrodynamic simulations, that the gas velocity dispersion $\sigma$ scales as $\sigma\propto N_{\rm gal}^{1/2}M_{\rm gal}$, where $N_{\rm gal}$ and $M_{\rm gal}$ are the number and mass of galaxies, respectively."485" Note that a scaling Exοςc?NauMz is consistent with dynamical friction in the linear regime, since ExexMa for dynamical friction."," Note that a scaling ${\rm E}_{\rm k} \propto \sigma^{2} \propto N_{\rm gal} M_{\rm gal}^{2}$ is consistent with dynamical friction in the linear regime, since $\dot{\rm E}_{\rm k} \propto M_{\rm gal}^{2}$ for dynamical friction."486" As Ngai and Mga. increase, the cooling catastrophe is delayed, and is completely staved off at the upper envelope of these parameters."," As $N_{\rm gal}$ and $M_{\rm gal}$ increase, the cooling catastrophe is delayed, and is completely staved off at the upper envelope of these parameters."487" In this respect our MHD simulations differ markedly from those of ?,, who found that a cooling catastrophe was inevitable in purely hydrodynamic simulations, for all portions of parameter space."," In this respect our MHD simulations differ markedly from those of \citet{kim07}, who found that a cooling catastrophe was inevitable in purely hydrodynamic simulations, for all portions of parameter space."488 We explore these differences further in 8 refsection:heating.., We explore these differences further in \\ref{section:heating}.489" Note that in our case the velocity dispersion seems to increase more slowly than σοςN!/?,", Note that in our case the velocity dispersion seems to increase more slowly than $\sigma \propto N^{1/2}$.490" Besides the inclusion of MHD in our simulations, differing results could be due to a variety of factors, including the different assumed distribution of galaxies."," Besides the inclusion of MHD in our simulations, differing results could be due to a variety of factors, including the different assumed distribution of galaxies."491 Note that the stated number of galaxies are distributed over the entire cluster; the number of galaxies in the inner regions which, Note that the stated number of galaxies are distributed over the entire cluster; the number of galaxies in the inner regions which492with blue star formation regious aud a red bulge. with a color of g—r=1.2 lor the bulge. g—r—0.31 for the exponential disk. g—r=—0.63 for the star formatioi regions. aud the patelies oL dust. produce regious that are typically Ag—r)0.12 redder thau the surrounding dust-free regions.,"with blue star formation regions and a red bulge, with a color of $g-r=1.2$ for the bulge, $g-r=0.37$ for the exponential disk, $g-r=-0.63$ for the star formation regions, and the patches of dust produce regions that are typically $\Delta (g-r) \sim 0.12$ redder than the surrounding dust-free regions."493 Addiug sky aud source noise results in total signal-to-noise valies that vary trom 50.0 in the center of the bulge to 5-7 at the disk ball-lieht radiis. to nearly 0 where the outer parts of the disk [ade into the sky backgrouud.," Adding sky and source noise results in total signal-to-noise values that vary from 50.0 in the center of the bulge to 5-7 at the disk half-light radius, to nearly 0 where the outer parts of the disk fade into the sky background."494 For example. the bilee of themock galaxy has an r-baud flux = 1500 DN on average. where the typical HIE αι 1ux = 1250-180. wuch is comparable to the galaxies rom SDSS i1 Section [.," For example, the bulge of themock galaxy has an $r$ -band flux $=$ 1500 DN on average, where the typical HII $r$ -band flux $=$ 150-180, which is comparable to the galaxies from SDSS in Section 4."495 In tlle analysis 5eps below. the mock data is treaed as if it were real SDSS data.," In the analysis steps below, the mock data is treated as if it were real SDSS data."496 PSFs are Measure [rom point sources in the unage., PSFs are measured from point sources in the image.497 The PSF of each image is matched to the πα PSF. becausei is the broacest PSF.," The PSF of each image is matched to the $u$ -band PSF, because it is the broadest PSF."498 The background sky was ueasured [rom the corner ol the images ancl subtracted from the «lata., The background sky was measured from the corner of the images and subtracted from the data.499 Before this method cau be applied. the data must have PSFs matchect. all iustrunent artifacs removed. aud the galaxy identified.," Before this method can be applied, the data must have PSFs matched, all instrument artifacts removed, and the galaxy identified."500 The method proposed here uses the results of PCA to determine a quantitative relationship between the colors of different pixels. which are associated. during the smoothing process.," The method proposed here uses the results of PCA to determine a quantitative relationship between the colors of different pixels, which are associated during the smoothing process."501 [tis assumed that each pixel is a linear combination of basis spectra., It is assumed that each pixel is a linear combination of basis spectra.502 The flux at any pixel can be described as a linear combination of a set of weighted basis spectra., The flux at any pixel can be described as a linear combination of a set of weighted basis spectra.503 The normalized flux at pixe (x.v) can then be written as: Wwiere 7 is the number of bands (5 for SDSS). cr and y are spatial position iu this galaxy. aux A denotes the baud (Le. «griz). a;jGCr.5) is a eigenweight which varies as a fuuctiou of position ii the gaaxy. alid 6»; is the zthi basis at baud A.," The normalized flux at pixel (x,y) can then be written as: where $i$ is the number of bands (5 for SDSS), $x$ and $y$ are spatial position in this galaxy, and $\lambda$ denotes the band (i.e. $ugriz$ ), $a_i(x,y)$ is a eigenweight which varies as a function of position in the galaxy, and $e_{\lambda,i}$ is the $i$ th basis at band $\lambda$."504 Tle covariance matrix method is used to measure the eigenvectors aud eigeuweights., The covariance matrix method is used to measure the eigenvectors and eigenweights.505 The data are first normalized to the r-band., The data are first normalized to the $r$ -band.506 All pixels within 2 disk scale leneths aud. having a SNR ratio greate “than a miniununm value (discussed later) are included in a data matrix., All pixels within 2 disk scale lengths and having a SNR ratio greater than a minimum value (discussed later) are included in a data matrix.507 Using lower values lower than this includes pixels heavily influenced by background sky colors., Using lower signal-to-noise values lower than this includes pixels heavily influenced by background sky colors.508 The covariance matrix of this data matrix is calculated. aud then the eigenvectors and. eigenvalues of this covariance matrix are determined.," The covariance matrix of this data matrix is calculated, and then the eigenvectors and eigenvalues of this covariance matrix are determined."509 This is carried out using Python procedures in the NUMPY.LINALG library. where the eigenvectors are solved. using the APACIy routines dgeev aud zgeev 7.," This is carried out using Python procedures in the NUMPY.LINALG library, where the eigenvectors are solved using the APACK routines dgeev and zgeev ."510. Figure 2 showstie location of pixels within a subsection of the mock [n]oOalaxy image., Figure \ref{fig:PCAanalysis} shows the location of pixels within a subsection of the mock galaxy image.511 The figure, The figure512lt has previously been pointed out (Warner 1995a.b.c) that the theoretical reduction of AZ with increasing {δω together with the destablizing elfect of irradiation on the secondary (Wu. Wickramasinghe Warner 1995) produces stable high AL at the shortest periods. VY Scl behaviour at intermediate periods. and low AZ. probable long-interval large amplitude dwarf novae at the longest. periods.,"It has previously been pointed out (Warner 1995a,b,c) that the theoretical reduction of $\dot{M}$ with increasing $P_{orb}$, together with the destablizing effect of irradiation on the secondary (Wu, Wickramasinghe Warner 1995) produces stable high $\dot{M}$ at the shortest periods, VY Scl behaviour at intermediate periods, and low $\dot{M}$, probable long-interval large amplitude dwarf novae at the longest periods."513" Each of the recently discovered AM. CVn stars. including ""2003aw'. fits into this"," Each of the recently discovered AM CVn stars, including `2003aw', fits into this"514models given by equations (54)) ancl (58)) suggests a variability of acceleration efficiency depending on various choices of (he integrals of motion.,models given by equations \ref{model}) ) and \ref{constb}) ) suggests a variability of acceleration efficiency depending on various choices of the integrals of motion.515 Nevertheless. the former solution (54)) studied in the previous section will be interesting as a (vpical model revealing the high potentiality of MIID acceleration. which was also discussed by Okamoto(2002) in relation to pulsar winds (see also Michel(1969):Beeelman&Li (1994))).," Nevertheless, the former solution \ref{model}) ) studied in the previous section will be interesting as a typical model revealing the high potentiality of MHD acceleration, which was also discussed by \citet{ok02} in relation to pulsar winds (see also \cite{mc69,bl94}) )."516 By virtue of the form expression (54)) of the model. we can clearly understaiud (he following evolution of jet flows in the intermediate and asymptotic regions Lar bevond (he lieht exlinder. if the jet radius A extends (o an infinite distance: (1) The magnetic-energv dominated and magnetosonic outflows pass through the lieht evlinder surface (2=Rp) with the Alfvénn Mach number such that M?~1/E. (," By virtue of the closed-form expression \ref{model}) ) of the model, we can clearly understand the following evolution of jet flows in the intermediate and asymptotic regions far beyond the light cylinder, if the jet radius $R$ extends to an infinite distance: (1) The magnetic-energy dominated and sub-fast-magnetosonic outflows pass through the light cylinder surface $R=R_{\rm L}$ ) with the Alfvénn Mach number such that $M^{2}\sim1/E$. ("517"2) Then. the energy ratio E,/E, increases to 1/E?? al (he fast-magnetosonic point corresponding to the radius 227RE.","2) Then, the energy ratio $E_{k}/E_{m}$ increases to $1/E^{2/3}$ at the fast-magnetosonic point corresponding to the radius $R\sim R_{\rm L}E^{1/3}$."518 In the intermediate region (Rp«Rx RE) (the outllows can smoothly become super-Lust-magnetosonic. (, In the intermediate region $R_{\rm L} \ll R \leq R_{\rm L}E$ ) the outflows can smoothly become super-fast-magnetosonic. (5193) Ónuportantlv. the model claims (he realization of rough e«quipartütion between kinetic and magnetic energies al the radius of order of RLF. (,"3) Importantly, the model claims the realization of rough equipartition between kinetic and magnetic energies at the radius of order of $R_{\rm L}E$. ("5204) The further energy conversion toward a kinelic-enerev dominated state in logarithmic scales of £2 in the asvinptotic region RF) is also confirmed.,4) The further energy conversion toward a kinetic-energy dominated state in logarithmic scales of $R$ in the asymptotic region $R \gg R_{\rm L}E$ ) is also confirmed.521 Figure 7 is the summary of this jet solution., Figure \ref{fig:jet} is the summary of this jet solution.522 The [ull conversion of magnetic energv into kinetic one means (that. the asymptotic Lorentz [actor of bulk motion becomes equal to the total specific οποιον E of outflows injected near the central source., The full conversion of magnetic energy into kinetic one means that the asymptotic Lorentz factor of bulk motion becomes equal to the total specific energy $E$ of outflows injected near the central source.523" Hence. for kinetic-energy dominated jets observed with a huge bulk Lorentz factor 2. we can expect the rough equipartition £j~£I, to occur at the jel radius Ri=Rey | Note that the value of 5 (~ Ey) is roughly same order with £ |."," Hence, for kinetic-energy dominated jets observed with a huge bulk Lorentz factor $\gamma$, we can expect the rough equipartition $E_{k}\sim E_{m}$ to occur at the jet radius $R_{\rm jet}\equiv R_{\rm L}\gamma$ [ Note that the value of $\gamma$ $\sim E_k$ ) is roughly same order with $E$ ]."524 Though we have considered ideal MIID flows in (hiis paper. the observed jet activity. such as a prompt emission of radiation aud a ultra relativistic acceleration of electrons. should be due to dissipation of the power of bulk motion. for example. through formation of shocks.," Though we have considered ideal MHD flows in this paper, the observed jet activity, such as a prompt emission of radiation and a ultra relativistic acceleration of electrons, should be due to dissipation of the power of bulk motion, for example, through formation of shocks."525 Then. (he interesting hieh-enerey phenomena of jets will be observed. only after the kinetic energy of bulk motion begins to dominate. namely. the jet radius extends to this critical radius Zij4.," Then, the interesting high-energy phenomena of jets will be observed, only after the kinetic energy of bulk motion begins to dominate, namely, the jet radius extends to this critical radius $R_{\rm jet}$."526 The shock formed in the energy equipartition region would be distinct from shocks formed in the kinetically dominated asvinplolic region by observations., The shock formed in the energy equipartition region would be distinct from shocks formed in the kinetically dominated asymptotic region by observations.527 For the kinetically dominated flows. the compression ratio behind the shock is much higher. ancl flatter svnchrotron spectra. higher emissivities. etc.," For the kinetically dominated flows, the compression ratio behind the shock is much higher, and flatter synchrotron spectra, higher emissivities, etc."528 would be observed (Begelman&Li1994)., would be observed \citep{bl94}.529. In particular. [or AGN jets with ~10 (see. e.g.. Ghiselliniet al.," In particular, for AGN jets with $\gamma\sim 10$ (see, e.g., Ghiselliniet al."530 1993). we can estimate the crilical radius to be Rig~LOL.," 1993), we can estimate the critical radius to be $R_{\rm jet} \sim 10R_{\rm L}$."531 When magnetic fluxes for the jet connect to a rotating geometrically thin disk around a black hole. we can regard ο as Oir). where Oj(r) is the angular velocity for circular equatorial orbit in the Nery metric. which corresponds to the Keplerian angular velocity in the Newtonian case.," When magnetic fluxes for the jet connect to a rotating geometrically thin disk around a black hole, we can regard $\Omega_F$ as $\Omega_{\rm K}(r)$, where $\Omega_{\rm K}(r)$ is the angular velocity for circular equatorial orbit in the Kerr metric, which corresponds to the Keplerian angular velocity in the Newtonian case."532 If the foot points of most magnetic fluxes distribute near the inner part of the disk r~ y. the angular velocities of the magnetic field lines are roughly O(V)~On(Grays) (see. e.g.. Camenzind Ixrockenberger 1992).," If the foot points of most magnetic fluxes distribute near the inner part of the disk $ r \sim r_{\rm ms}$ , the angular velocities of the magnetic field lines are roughly $\Omega_F(\Psi) \sim \Omega_{\rm K}(r_{\rm ms})$ (see, e.g., Camenzind Krockenberger 1992)."533 Then. the," Then, the"534We now apply the above formalisin to derive Fisher matrices for specific cases of WL survevs aud their associated spectroscopic calibration surveys.,We now apply the above formalism to derive Fisher matrices for specific cases of WL surveys and their associated spectroscopic calibration surveys.535" Iu further sectious we vary the parameters of the photo-z errors and the spectroscopic survey and investigate the impact ou the accuracy of dark energy parameters derived frou cach survey,", In further sections we vary the parameters of the photo-z errors and the spectroscopic survey and investigate the impact on the accuracy of dark energy parameters derived from each survey.536 Following Maetal.(2006).. the fiducial galaxy redshift distribution »(:) is chosen to have the form Unuless otherwise stated we adopt a=2 and 3=1 aud fix zy such that the median redshift is tye=1.," Following \cite{Ma05}, the fiducial galaxy redshift distribution $n(z)$ is chosen to have the form Unless otherwise stated we adopt $\alpha=2$ and $\beta=1$ and fix $z_0$ such that the median redshift is $z_{\rm med} = 1$."537" The parametric model for 5(:) is determined by linear interpolation between My,=31 values n=ne"") at equally spaced redshitts between 0 aud 3.", The parametric model for $n(z)$ is determined by linear interpolation between $N_{\rm pz}=31$ values $n^i=n(z^i)$ at equally spaced redshifts between 0 and 3.538 Tn the Gaussian case as axstuned in Mactal. (2006).. we have The bias p; aud dispersion o. are functions of :.," In the Gaussian case as assumed in \cite{Ma05}, , we have The bias $z_{\rm bias}$ and dispersion $\sigma_z$ are functions of $z$."539 Iu reality. could bo far more complex than a single Gaussian.," In reality, $P(z_{\rm ph}|z)$ could be far more complex than a single Gaussian."540" We P(24,4]:)explore this complexity by assuuiug as the suni of Caussiaus.", We explore this complexity by assuming $P(z_{\rm ph}|z)$ as the sum of Gaussians.541" Using ΜΜ to P(ualz)describe Pnon|i). we have where is the normalization of the j""th Gaussian."," Using $N_{g}$ Gaussians to describe $P(z_{\rm ph}|z)$, we have where $C_j$ is the normalization of the $j^{th}$ Gaussian."542 Since we C; ΕΕ is normalized to unity. we have Ni=|].," Since we assume $P(z_{\rm ph}|z)$ is normalized to unity, we have $\sum_j C_j = 1$."543 We allow the biases μη) aud scatters fo be arbitrary fuuctious of redshift., We allow the biases $z_{\rm {bias;j}}(z)$ and scatters $\sigma_{z;j}(z)$ to be arbitrary functions of redshift.544 The redshitt G..(:)distribution of the tomographic bius defined by refequü can then be written as with where erfte) is the error functiou., The redshift distribution of the tomographic bins defined by \\ref{eq:ni} can then be written as with where ${\rm erf}(x)$ is the error function.545" Iu practice. we represeut the free functfious ipi) aud 0.,(:) by Tear interpolation between values at a discrete set of redshifts equally spaced from +=0 to 3."," In practice, we represent the free functions $z_{\rm{bias;j}}(z)$ and $\sigma_{z;j}(z)$ by linear interpolation between values at a discrete set of $N_{\rm pz}$ redshifts equally spaced from $z=0$ to 3."546" The photo-z NV,parameter set is hence the 2N,N,, values of the biases and dispersions{pyf of the Gaussiaus at these nodes.", The photo-z parameter set $\{p_\mu\}$ is hence the $2N_gN_{\rm pz}$ values of the biases and dispersions of the Gaussians at these nodes.547 With multiple Caussiaus. we can describe a wide variety of photo-z probability distributions μμ]r," With multiple Gaussians, we can describe a wide variety of photo-z probability distributions $P(z_{\rm ph}|z)$."548efügpzPDFes shows afew examples of ," \\ref{fig:pzPDFeg}549 shows afew examples of $P(z_{\rm ph}|z)$."550A wide variety of behaviors can be represented. Puul:).including ~catastrophic” outliers.," A wide variety of behaviors can be represented, including “catastrophic” outliers."551 Although catastrophic ploto-z errors could potentially have a big impact on what we can eet out of cosmic shear surveys (Amara.&Ποfregier2007).. we restrict ourselves to studying the core of Γον). iu this study.," Although catastrophic photo-z errors could potentially have a big impact on what we can get out of cosmic shear surveys \citep{Amara06}, we restrict ourselves to studying the core of $P(z_{\rm ph}|z)$ in this study."552" Maetal.(2006). show that M,,=31 between :=0 and 3 gives euough freedom to the photo-z parameters to destroy all tomographic information.", \cite{Ma05} show that $N_{\rm pz}=31$ between $z=0$ and 3 gives enough freedom to the photo-z parameters to destroy all tomographic information.553" Since we are giviug he photo-z even more freedom by allowing μμ) to ΡΟ ΡΕ Vy,=31 should be large enough."," Since we are giving the photo-z even more freedom by allowing $P(z_{\rm ph}|z)$ to be multiple Gaussians, $N_{\rm pz}=31$ should be large enough."554" Uuless stated otherwise, we use Vy,=31."," Unless stated otherwise, we use $N_{\rm pz}=31$."555 Thus. the tota iuauber of photo-z parameters is G2.Nye," Thus, the total number of photo-z parameters is $62 N_g$."556" The observables αςIE,ον determined ta bius of width op need not have the same bin width as the spacing o he nC) or the photo-z parameters."," The observables $n(z_{\rm ph}^i)$, determined in bins of width $\delta z_{\rm ph}$, need not have the same bin width as the spacing of the $n(z^i)$ or the photo-z parameters."557 In fact. they shoule vo mnore finely spaced.," In fact, they should be more finely spaced."558" We choose the size of óz,4 such hat further dividing it bv two docs not lead to auvmere information eains.", We choose the size of $\delta z_{\rm ph}$ such that further dividing it by two does not lead to anymore information gains.559 We find that 6:=0.0125 is smal enough for all the photo-z models explored iu this study., We find that $\delta z_{\rm ph} = 0.0125$ is small enough for all the photo-z models explored in this study.560 Iun this section wo investigate the size of the spectroscopic calibration sample required to limit photo-z systematics to some desired level., In this section we investigate the size of the spectroscopic calibration sample required to limit photo-z systematics to some desired level.561 Tn particular. we are interested im the increased demands that might result from giving the photo-z cistribution freedom to depart froma single-Gaussian form.," In particular, we are interested in the increased demands that might result from giving the photo-z distribution freedom to depart from a single-Gaussian form."562" We first demonstrate that.for a fixed fiducial photo-z model. the required calibration size increases with the nuniber of degrees of freedom (2.V,,) that we allow for deviations from the fiducial model."," We first demonstrate that,for a fixed fiducial photo-z model, the required calibration size increases with the number of degrees of freedom $2N_g$ ) that we allow for deviations from the fiducial model."563 This increase reaches an asviuptotic limit with Vy., This increase reaches an asymptotic limit with $N_g$.564 Secoud. we investigate how the required ρου varies as we allow the fiducial model to assunue nou-Ciaussiau shapes.," Second, we investigate how the required $N_{\rm spect}$ varies as we allow the fiducial model to assume non-Gaussian shapes."565 refequiNpriorCuuun aud ÁÀ-10 show that in the case of a Gaussian distribution. the ρου required to constrain the photo-z parameters is proportional to the square of the width of the distribution.," \\ref{eqn:NpriorGmu} and \ref{eqn:NpriorGsi} show that in the case of a Gaussian distribution, the $N_{\rm spect}$ required to constrain the photo-z parameters is proportional to the square of the width of the distribution."566 In the following. we hold the width (defined as the rms) of the fiducial photo-z distributions to be 0.05(1|2).," In the following, we hold the width (defined as the rms) of the fiducial photo-z distributions to be $0.05(1+z)$."567 Molding this fiducial width fixed means that auv variations we see are due oulv to variations in the of the photo-z probability distribution., Holding this fiducial width fixed means that any variations we see are due only to variations in the of the photo-z probability distribution.568" We use the error deeradations in ow, (that is. errors in wy relative to the error with perfect knowledee of the photo-z parameters) as the measure of dark energydegradations."," We use the error degradations in $w_{\rm a}$ (that is, errors in $w_{\rm a}$ relative to the error with perfect knowledge of the photo-z parameters) as the measure of dark energydegradations."569 The error degradatious iu wy? are about 30-50% lower and follow the same trend as that of a., The error degradations in $w_{\rm p}$ are about $30$ $50\%$ lower and follow the same trend as that of $w_{\rm a}$ .570 Roughly speaking. the figure of merit adopted bv the Dark Enerey Task Force(Albrechtetal.2006) will degrade as the square of the dark euergy degradation used here.," Roughly speaking, the figure of merit adopted by the Dark Energy Task Force\citep{DETF} will degrade as the square of the dark energy degradation used here."571 Iu this section we ase that the total spectroscopic galaxies are selected wuitormly iu Mauτοςκατ between 0 and 3., In this section we assume that the $N_{\rm spect}$ total spectroscopic galaxies are selected uniformly in redshift between 0 and 3.572" The left panel of plots the dark energy degradation versusthe size of the spectroscopic calibrationsample. when the photo-z error distribution has V,4= 1. 2. 3. and"," The left panel of \\ref{fig:fid1234G} plots the dark energy degradation versusthe size of the spectroscopic calibrationsample, when the photo-z error distribution has $N_g = 1$ $2$ , $3$ , and"573study on the «election of moous. we choose to cousider the simplest model.,"study on the detection of moons, we choose to consider the simplest model."574 In a future work we can consider a more realistic model. including orbital inclinations aud eccentricity for the moon.," In a future work we can consider a more realistic model, including orbital inclinations and eccentricity for the moon."575 Reearcing to ie planet eccentricity. followiug the strategy proposed by Ixippiug(2008).. tle model simulates elipical orbits ouly for the cases were ecceον was measured by other methods. suc1 as Facial veloci(les.," Regarding to the planet eccentricity, following the strategy proposed by \cite{Kipping2008}, the model simulates eliptical orbits only for the cases were eccentricity was measured by other methods, such as radial velocities."576 stellar iuout parameters to the model are: rotation period. radius. at(| —ass (the latter two i[un solar ilils).," Stellar input parameters to the model are: rotation period, radius, and mass (the latter two in solar units)."577 Tle rotaion period is use in the calcuation of the longitu«erp )osition of starspots 1[un SUCCESS]re |asits., The rotation period is used in the calculation of the longitude position of starspots in successive transits.578 Wie also have two liub carkenine parameters., We also have two limb darkening parameters.579" The plineary luput parameterW. are: orαἱ Ίος. ""acius. mass. orbital inclination alele. eccentricity aud JOSition of the periastrou."," The planetary input parameters are: orbital period, radius, mass, orbital inclination angle, eccentricity and position of the periastron."580 From le Orital period. the orbit's seii-major axisis caleulated by IvejleUs Third Law.," From the orbital period, the orbit's semi-major axis is calculated by Kepler's Third Law."581 Lastly. the moo parameers are orbital »eriod. radiIs. Nass. aud the augila“position of the moo1 in its or he starti1ig time of the first transit.," Lastly, the moon input parameters are orbital period, radius, mass, and the angular position of the moon in its orbit at the starting time of the first transit."582 For the rines. the inpil paraiueters consists of inner aud qai. twὉ Inclination aieles. one wihi respect to the plane of the sky aud the oler with the planets orbita ylane.," For the rings, the input parameters consists of inner and outer radii, two inclination angles, one with respect to the plane of the sky and the other with the planet's orbital plane."583 The transj»arency o‘the rings may also be fit., The transparency of the rings may also be fit.584" Τιe iporal resol=ion of the resitine light (""urve is chosen within the model. by adjusing the time ierval for each sitated pliotonetric caa poiut."," The temporal resolution of the resulting light curve is chosen within the model, by adjusting the time interval for each simulated photometric data point."585 Here we used a temporal resolutio of 32 s for fre compariso wil1 ColtoT d:via. and J min for hepey data.," Here we used a temporal resolution of 32 s for future comparison with CoRoT data, and 1 min for Kepler data."586 Usiue stellar aud plalelary paraljete*. the pkuietary orbit is caleulated.," Using stellar and planetary parameters, the planetary orbit is calculated."587 The line of transit is he pro.jection of this orbi O1 he stellar sur[face., The line of transit is the projection of this orbit on the stellar surface.588 A each time inerval. the program calculates the »osition of tlie planet all its moon.," At each time interval, the program calculates the position of the planet and its moon."589 The otal luminosity of the star-plauet-moon system is then determined by stumiumine he iHeLsities o- all the yixels iu the linage. where the pixels occupied wy the planet aix its moco luive zero value.," The total luminosity of the star–planet–moon system is then determined by summing the intensities of all the pixels in the image, where the pixels occupied by the planet and its moon have zero value."590 In the case of rings. the occupied pixels intensity is uultiplied by the rine traspa'eney factor.," In the case of rings, the occupied pixels intensity is multiplied by the ring transparency factor."591 At the end of the trait. the light curve is normalized o one from the ouside transit data points.," At the end of the transit, the light curve is normalized to one from the outside transit data points."592 Figure 1 shows two examples of trausits [rom the model., Figure \ref{fig:img_transitos} shows two examples of transits from the model.593 Figure l(a) shows he transit of a σασί and a moo rin front of the star. aL Figure 1(b) 5lows the transit of a platjet witli rings.," Figure \ref{fig:img_transitos_a} shows the transit of a planet and a moon in front of the star, and Figure \ref{fig:img_transitos_b} shows the transit of a planet with rings."594 Itis also possible to add spots ou the sellar surface. where each spot is moclelec by tliree extra j»arameters: radius. intensity. Owith respect to stellar cent‘al intensity). aud longit«le position : he stat ofthe first trausit.," It is also possible to add spots on the stellar surface, where each spot is modeled by three extra parameters: radius, intensity (with respect to stellar central intensity), and longitude position at the start of the first transit."595 The spot μαιde is considere Lto be that of the projected transit li, The spot latitude is considered to be that of the projected transit line.596 From t1e stellar rotation period. the spot longitude positio 10 Consectlive trausits is calculated.," From the stellar rotation period, the spot longitude position on consecutive transits is calculated."597 urther asstuuplion is that the spot position does not clanses)e within a sinele transit (usually lasts a few hours). which is reasonable when the stellar rotatioi period is much larger tlan the trai duration.," A further assumption is that the spot position does not change within a single transit (usually lasting a few hours), which is reasonable when the stellar rotation period is much larger than the transit duration."598 In this work. no attempt lias been mace for a dyuamical analysis of a planet with moous or," In this work, no attempt has been made for a dynamical analysis of a planet with moons or"599before. at JD~2454604.05. there is a small flare clearly visible in all three light curves.,"before, at $\rm JD \sim 2454604.05$, there is a small flare clearly visible in all three light curves."600 Mean source rates for the whole period are 1.63. 0.61. and 0.64 counts s! for pn. MOSI. and MOS2. respectively. with standard deviations of 0.15. 0.04. and 0.03 counts s7!.," Mean source rates for the whole period are 1.63, 0.61, and 0.64 counts $\rm s^{-1}$ for pn, MOS1, and MOS2, respectively, with standard deviations of 0.15, 0.04, and 0.03 counts $\rm s^{-1}$."601" This means fractional variations Εμ ofΤου,6%.. and 4%.. which. do not change significantly20ln ifκ. we exclude the last 6 hours."," This means fractional variations $F_{\rm var}$ of, and , which do not change significantly if we exclude the last 6 hours."602 Hence. we can conclude that the X-ray flux of BL Lacertae is mildly variable on an hour time scale.," Hence, we can conclude that the X-ray flux of BL Lacertae is mildly variable on an hour time scale."603 Figure 9. shows the broad-band SED of BL Lacertae in different brightness states., Figure \ref{sed_tot} shows the broad-band SED of BL Lacertae in different brightness states.604 The SED corresponding to 2008 May 16-17 includes the XMM-Newton UV and X-ray data analysed in refsecuum. f, The SED corresponding to 2008 May 16–17 includes the XMM-Newton UV and X-ray data analysed in \\ref{sec_xmm}.605" Inordertoavoidof setscausedbysourcevariability. Comptondominance.TRG ENSE WHE det Wo ο] ήipit, tbe EM epic)."," In order to avoid offsets caused by source variability, the OM spectrum was constructed with the last $W1$ data point, the first $M2$ (which is close to the last $M2$ ), and the first $W2$ datum of the corresponding light curves (see \\ref{om+epic}) )."606ThesedataindicateahardU V spectrum., These data indicate a hard UV spectrum.607T woSED. HAW Salta an ο eye 30. refseciwiftbecauseofthedif FerentspectralslopeintheX raxband.," Two SEDs in the figure refer to the Swift observations of 2008 August 25 and 30, which were chosen among those analysed in \\ref{sec_swift} because of the different spectral slope in the X-ray band."608T helow—frequency partofthesethreeS EDsisbuiltw LENIN LOM fFsecsaspYy. forAugust30near IRdatawerealsoavailable., The low-frequency part of these three SEDs is built with GASP optical and radio data (see \\ref{sec_gasp}) ); for August 30 near-IR data were also available.609Thestrongsourcevariabilitvintheopticalbangrequig Intheradiolggdag estie qe formationisqut Mga IRYlataaresimultaneouswiththesatelliteobservations., The strong source variability in the optical band requires that the optical (and near-IR) data are simultaneous with the satellite observations.610 uri vaste hista reed whether titπα —3days fromthesatelliteones. IR.optical. andU Vdatawerecorrected w," In the radio bands flux variations are slower, so that we used data taken within 2--3 days from the satellite ones, when simultaneous data were not available. Near-IR,"611hensimultaneousdatawerenotaveigblefortheGalacticextinction: IRandopticaldatawerealsocorrected f forthecontributionotheho: vise refsec.asp))," optical, and UV data were corrected for the Galactic extinction; the near-IR and optical data were also corrected for the contribution of the host galaxy (see \\ref{sec_gasp}) )."612" InAugust2008theFerniy-ray satellite detected BL Lacertae: the Fermi data we plottedin refsed,otwerederived fromAbdoet(2010a).", In August 2008 the Fermi $\gamma$ -ray satellite detected BL Lacertae; the Fermi data we plotted in \\ref{sed_tot} were derived from.613. The August 2008 SEDs indicate a faint. synchrotron-dominated state of the source that we fitted with the rotating helical jet model by1999.," The August 2008 SEDs indicate a faint, synchrotron-dominated state of the source that we fitted with the rotating helical jet model by,."614".2003..]vil99, This model has been used by to fit the broad-band SED of BL Lacertae in December 2007 — January 2008.", This model has been used by to fit the broad-band SED of BL Lacertae in December 2007 – January 2008.615 Their main finding was that the BL Lacertae broad-band SED cannot be explained by à single synchrotron component plus its self inverse-Compton emission., Their main finding was that the BL Lacertae broad-band SED cannot be explained by a single synchrotron component plus its self inverse-Compton emission.616 Indeed. the very strong historical X-ray variability requires an additional synchrotron (plus self inverse-Compton) component.," Indeed, the very strong historical X-ray variability requires an additional synchrotron (plus self inverse-Compton) component."617"ο). MoreoverLOreover.the UV excess suggests thermal contribution from the accretion disc (see 5.Wenoticerefsec,vor) VT"," Moreover,the UV excess suggests thermal contribution from the accretion disc (see \\ref{sec_uvot}) )."618 heS EDanalysedbyRaiteriet lacked simultaeqmonyot data. which made it impossible to constrain the emission of the high-energy component.," The SED analysed by lacked simultaneous $\gamma$ -ray data, which made it impossible to constrain the emission of the high-energy component."619" The 2008 August SED in refsed,otnowo ffersusthe possibilitvto perFormamoredetailedanalvsis.", The 2008 August SED in \\ref{sed_tot} now offers us the possibility to perform a more detailed analysis.620" In addition. we also display in refsed,otthebroad-bandS EDcorrespondingtothebigoutbursto f July 1997 he"," In addition, we also display in \\ref{sed_tot} the broad-band SED corresponding to the big outburst of July 1997, which showed a considerable inverse-Compton dominance."621X EAN MD , The X-ray spectra plotted in the figure are the result of the combined analysis of the ASCA and RXTE data by.622jovasd) aos Y which V 4 felad CORO: eifang AN up," Because of the very strong variability of the source in that period, the authors distinguished between a low state, which was well fitted by a single power law model, and a flare state, for which the best fit was obtained with a double power law model."623s instr, This last fit resulted in a very strong spectral.624ume, In July 1997 observations in the $\gamma$ -ray band were performed by CGRO.625nt onboard CGRO in ια . whilethose fromtheOS SEd," The data from the EGRET instrument onboard CGRO in \\ref{sed_tot} were taken from, while those from the OSSE detector were derived from the High Energy Astrophysics Science Archive Research (HEASARC)."626 frequencyin fromtheWE BTarchive: therangeo foptical flux , The low-frequency information is from the WEBT archive; the range of optical flux variation in the period is indicated.627fim was to se, This outburst state of the source was fitted with the same rotating helical model that we used to fit the faint state of 2008.628e it ii «eti Neige] configuration only., In performing the model fits our aim was to see whether it was possible to reproduce the high and low states by changing the geometrical configuration only.629 Moreover. we took into tdi: Hithe results by(2010)... who explained hdACC EIL. Lacertae optical andnear-IR variability in terms of variations of the Doppler boosting factor due to changes of the viewing angle of the emitting region.," Moreover, we took into account the results by, who explained the long-term BL Lacertae optical andnear-IR variability in terms of variations of the Doppler boosting factor due to changes of the viewing angle of the emitting region."630" The resulting model parameters are reported in Table 4.. while the corresponding fits are shown in refsed,ot."," The resulting model parameters are reported in Table \ref{modelfit}, while the corresponding fits are shown in \\ref{sed_tot}."631.Wealsoincludedblackbodyradiation f ronanacceretiondiscwithe and a temperature of ~16000K.The low-energy (radio-to-optical and related inverse-Compton) emission component comes from a helix portion with a pitch angle €= 2°. covering an angle a= 1807. The maximum Lorentz factor of the relativistic electrons is logyna.(0)= 3.7. while the bulk Lorentz factor of the plasma in the jet is T= 7.," We also included blackbody radiation from an accretion disc with a luminosity of $5 \times 10^{44} \rm \, erg \, s^{-1}$ and a temperature of $\sim 16000 \, \rm K$.The low-energy (radio-to-optical and related inverse-Compton) emission component comes from a helix portion with a pitch angle $\zeta=2\degr$ , covering an angle $a=180 \degr$ The maximum Lorentz factor of the relativistic electrons is $\log \gamma_{\rm max}(0)=3.7$ , while the bulk Lorentz factor of the plasma in the jet is $\Gamma=7$ ."632" The high-energy (UV—X-ray and related inverse-Compton) emission component comes from a helix portion with a pitch angle €= 8"". covering: an angle v= 3607.", The high-energy (UV–X-ray and related inverse-Compton) emission component comes from a helix portion with a pitch angle $\zeta=8\degr$ covering an angle $a=360 \degr$ .633 The maximum, The maximum634packages (FTOOLS version 4.2. NSPEC version 11.2).,"packages (FTOOLS version 4.2, XSPEC version 11.2)."635 Event files have been (hus filtered [or hieh-backeround time intervals and only events corresponding to patterns 0—12 (MOS 1&22) and 0—4 (pn) have been used (see Ehleetal.2001)): the net exposure times at the source position alter data cleaning are ~21.4 ks (MOSI. MOS2) and ~17.5 ks (pn).," Event files have been thus filtered for high-background time intervals and only events corresponding to patterns $0 - 12$ (MOS 2) and $0 - 4$ (pn) have been used (see \citealt{xmmhb}) ); the net exposure times at the source position after data cleaning are $\sim 21.4\:\,$ ks (MOS1, MOS2) and $\sim 17.5\:\,$ ks (pn)."636" The Newton.MOSI. MOS2 and pn images in the 0.5—10 keV energy range reveal a high signal-to-noise ratio (S/N~30 aud ~50 in the MOS and the pn. respectively) point- source within the positional error circle (~2"" radius) ofANJO447-0627."," The MOS1, MOS2 and pn images in the $0.5-10\:\,$ keV energy range reveal a high signal-to-noise ratio $S/N \sim 30$ and $\sim 50$ in the MOS and the pn, respectively) point-like source within the positional error circle $\sim 2^{\prime}$ radius) of."637. This is the only detected ancl visible X-ray source in the error circle: the X-ray position derived using the data is RA = 04:47:48.62. Dee = —06:28:12 (e21 away rom (he nominal position).," This is the only detected and visible X-ray source in the error circle; the X-ray position derived using the data is $\mbox{RA}$ = 04:47:48.62, $\mbox{Dec}$ = $-$ 06:28:12 $\sim 21\arcsec$ away from the nominal position)."638" Source counts were extracted from a circular region of radius 22.5"" for the MOS and 17.5"" or (he pn (this smaller radius for the pu is due to the proximity of à CCDs gap).", Source counts were extracted from a circular region of radius $22.5\arcsec$ for the MOS and $17.5\arcsec$ for the pn (this smaller radius for the pn is due to the proximity of a CCDs gap).639" Dackeround counts were extracted from a nearby source-[ree cireular region of ~42""—50"" radius.", Background counts were extracted from a nearby source-free circular region of $\sim 42\arcsec - 50\arcsec$ radius.640 The net count rates (0.5—10 keV οποιον range) are 0.051£0.002 counts ο. 0.054£0.002 counts | and 0.171+0.004 counts ! for MOSI. MOS?2 and pn. respectively: the source counts represents about. of the total counts in the source extraction region.," The net count rates $0.5 - 10\:\,$ keV energy range) are $0.051\pm6410.002\:\,$ counts $^{-1}$, $0.054\pm 0.002\:\,$ counts $^{-1}$ and $0.171\pm6420.004\:\,$ counts $^{-1}$ for MOS1, MOS2 and pn, respectively; the source counts represents about of the total counts in the source extraction region."643 No statistically significant source variability has been detected during the observation., No statistically significant source variability has been detected during the observation.644 To improve statistics. the MOSI and MOS2 data have been combined together. aud the MOS and pn spectra have been fitted simultaneously. keeping the relative normalization free.," To improve statistics, the MOS1 and MOS2 data have been combined together, and the MOS and pn spectra have been fitted simultaneously, keeping the relative normalization free."645 Source counts were binned so as to have at least 20 counts in each energy bin., Source counts were binned so as to have at least $20$ counts in each energy bin.646 We have also generated our own spectral response matrices al (he source position using ihe SAS tasks andrmfgen., We have also generated our own spectral response matrices at the source position using the SAS tasks and.647. All the models cliseussecl here have been filtered through the Galactic absorption column density along the line of sieht cuim5.6x107)τι Dickey&Lockman 1990)).," All the models discussed here have been filtered through the Galactic absorption column density along the line of sight $=5.6\times10^{20}\:\,$; \citealt{nh}) )."648 Unless otherwise stated. fit parameters are quoted in the of (2=0.214. see below). while the figures and the EWs are in the observer frame.," Unless otherwise stated, fit parameters are quoted in the rest-frame of $z=0.214$, see below), while the figures and the EWs are in the observer frame."649" A bright (RA = 04:47:48.5. Dee = —06:28:13: APM red magnitude = 17.7) optical source lies about 2"" from the X-ray position derived using the data."," A bright $\mbox{RA}$ = 04:47:48.5, $\mbox{Dec}$ = $-$ 06:28:13; APM red magnitude = 17.7) optical source lies about $2\arcsec$ from the X-ray position derived using the data."650 This object was observed spectroscopically at the TING on October 5th. 2002.," This object was observed spectroscopically at the TNG on October 5th, 2002."651 The optical spectrum. (ot reported here) covers the wavelength range ~3500—8000A (dispersion of 2.8 A//pixel) and clearly shows broad (FEWIIM 76000 kms 1) Mell. LL? and Ho lines as well as narrow (FWIIM <1000 kms I!) [OL4959.5007 A lines.," The optical spectrum (not reported here) covers the wavelength range $\sim 3500-8000\:\,$ (dispersion of $2.8\:\,$ /pixel) and clearly shows broad (FWHM $> 6000\:\,$ km $^{-1}$ ) MgII, $\beta$ and $\alpha$ lines as well as narrow (FWHM $< 1000\:\,$ km $^{-1}$ ) [OIII]4959,5007 $\ $ lines."652 The optical line properties and position, The optical line properties and position653"31.8°, —0.27?«b 0.27?) can be used to investigate this possibility.",", $-0.27\degr< b < 0.27\degr$ ) can be used to investigate this possibility."654 This area covers 196 Spitzer IRDCs in total and more than 80% of the clouds can be associated with CS emission., This area covers 196 Spitzer IRDCs in total and more than $\%$ of the clouds can be associated with CS emission.655 'The distances of these clouds can be calculated using the Reid et al. (, The distances of these clouds can be calculated using the Reid et al. (6562009) galactic rotation model.,2009) galactic rotation model.657 Figure 5 shows the distance distribution of the Spitzer IRDCs detected in CS., Figure \ref{cs_dist} shows the distance distribution of the Spitzer IRDCs detected in CS.658" In this figure, the IRDCs have been divided int to three size ranges."," In this figure, the IRDCs have been divided int to three size ranges."659" The 80 smallest IRDCs, those with Req (PF09) less than have a mean distance and standard deviation of 5.2 kpc and 0.8 kpc."," The 80 smallest IRDCs, those with $R_{eq}$ (PF09) less than have a mean distance and standard deviation of $5.2$ kpc and $0.8$ kpc."660" This is indistinguishable from the values of 5.3 kpc and 0.7 kpc and 5.3 kpc and 0.8 kpc for the IRDCs in the next two size ranges, 15""«Req30"", Req>30”, which contain 52 and 39 objects respectively."," This is indistinguishable from the values of $5.3$ kpc and $0.7$ kpc and $5.3$ kpc and $0.8$ kpc for the IRDCs in the next two size ranges, $15''<R_{eq}<30''$, $R_{eq}>30''$, which contain $52$ and $39$ objects respectively."661 These distributions therefore show no indication that large and small IRDCs have different distributions of distance., These distributions therefore show no indication that large and small IRDCs have different distributions of distance.662" The mass completeness limit for the IRDCs and fragments can be written as where R, is the smallest radius above which the sample is complete, d. is the distance within which the majority of the sources occur, and <Ny,Ἄς is the typical average column density of the structures with a radius R,."," The mass completeness limit for the IRDCs and fragments can be written as where $R_{\rm c}$ is the smallest radius above which the sample is complete, $d_{\rm c}$ is the distance within which the majority of the sources occur, and $<N_{\rm H_2}>_{\rm c}$ is the typical average column density of the structures with a radius $R_{\rm c}$."663 Figure 4((left) shows that about 9596 of the IRDCs in that plot have distances below 6 kpc and so we conservatively adopt d.=6 kpc., Figure \ref{dist}( (left) shows that about $\%$ of the IRDCs in that plot have distances below 6 kpc and so we conservatively adopt $d_{\rm c}=6$ kpc.664 Estimating Πο is less straighforward., Estimating $R_{\rm c}$ is less straighforward.665" The completeness limits of our survey are related to two parameters of the source extraction: Ng, the minimum column density amplitude of a source (which is related to sensitivity) from the boundary of a cloud to its peak; and the angular resolution, bboth for IRDCs and fragments."," The completeness limits of our survey are related to two parameters of the source extraction: $N_{\rm666 H_2}^{\rm amp}$, the minimum column density amplitude of a source (which is related to sensitivity) from the boundary of a cloud to its peak; and the angular resolution, both for IRDCs and fragments."667 In order to investigate how these contribute to the completeness limits we look at the distribution of average column density of IRDCs for objects of a given range of sizes as plotted in Fig. 6.., In order to investigate how these contribute to the completeness limits we look at the distribution of average column density of IRDCs for objects of a given range of sizes as plotted in Fig. \ref{pk_r}.668 We then plot the column density at the peak of these distributions as a function of cloud size., We then plot the column density at the peak of these distributions as a function of cloud size.669 This is also done for the fragments., This is also done for the fragments.670 Figure 7 shows these plots., Figure \ref{rad_comp} shows these plots.671 The plots show a similar structure for both the IRDCs and fragments., The plots show a similar structure for both the IRDCs and fragments.672" Up to some size, 55"" for the IRDCs and 9"" for the fragments, the peak of the column density distributions is constant."," Up to some size, $55\arcsec$ for the IRDCs and $9\arcsec$ for the fragments, the peak of the column density distributions is constant."673 Above these values it increases with increasing size., Above these values it increases with increasing size.674 This constant column density for small sizescales suggests that the sample is not fully probing the populations of objects at these, This constant column density for small sizescales suggests that the sample is not fully probing the populations of objects at these6751t has been well-known for many vears that core-dominated raclio-loud quasars appear optically brighter and bluer than Iobe-dominated: ones (e.g... Jackson ct 11989: Baker Llunstead 1995).,"It has been well-known for many years that core-dominated radio-loud quasars appear optically brighter and bluer than lobe-dominated ones (e.g., Jackson et 1989; Baker Hunstead 1995)."676 Since powerful radio cores are believed to result from. Doppler boosting when the radio jet is oriented. close to the line of sight. this implies a viewing angle dependence of the optical emission.," Since powerful radio cores are believed to result from Doppler boosting when the radio jet is oriented close to the line of sight, this implies a viewing angle dependence of the optical emission."677 Such a dependence could be caused by Doppler boosting of optical svnchrotron radiation (Jackson et 11989). anisotropic emission [ron an optically-thick acceretion. disc (Netzer 1985. 1987). or dust extinction ancl reddening of the continuum at large viewing angles (Baker 1997).," Such a dependence could be caused by Doppler boosting of optical synchrotron radiation (Jackson et 1989), anisotropic emission from an optically-thick accretion disc (Netzer 1985, 1987), or dust extinction and reddening of the continuum at large viewing angles (Baker 1997)."678 While the first mechanism is likeA o be important only in racdio-Ioud quasars. the others should also influence the appearance of radio-quiet quasars.," While the first mechanism is likely to be important only in radio-loud quasars, the others should also influence the appearance of radio-quiet quasars."679 Unfortunately. the techniques for. identifvine racdio-cquiet quasars (via optical and/or X-ray emission) are themselves allectecd by these mechanisms. and cannot therefore be used to produce samples suitable for studying them.," Unfortunately, the techniques for identifying radio-quiet quasars (via optical and/or X-ray emission) are themselves affected by these mechanisms, and cannot therefore be used to produce samples suitable for studying them."680 On the other hand. samples selected: at low radio frequency are especially useful since they consist. almost exclusively of objects which are viewed at random. orientations because their racio emission is both unshadowed and üunbeamed.," On the other hand, samples selected at low radio frequency are especially useful since they consist almost exclusively of objects which are viewed at random orientations because their radio emission is both unshadowed and unbeamed."681 In principle. therefore. one can learn much about viewing angle dependencies from Iow-frequeney radio saniples. with a view to understanding the physical mechanism(s) responsible.," In principle, therefore, one can learn much about viewing angle dependencies from low-frequency radio samples, with a view to understanding the physical mechanism(s) responsible."682 Alaking measurements in the region of Aj7Lum is also likely to prove fruitful. since the mechanisms responsible for emission on cither side of this wavelength are believec to be distinct.," Making measurements in the region of $\lambda_{\rm rest} \approx6831\,\mu$ m is also likely to prove fruitful, since the mechanisms responsible for emission on either side of this wavelength are believed to be distinct."684 At optical wavelengths (heneelorth define as Ayer< dum). the emission is thought to arise from the low-enereyv tail of an accretion disc spectrum. which coul be intrinsically anisotropic (Netzer 1987).," At optical wavelengths (henceforth defined as $\lambda_{\rm rest} < 1\,\mu$ m), the emission is thought to arise from the low-energy tail of an accretion disc spectrum, which could be intrinsically anisotropic (Netzer 1987)."685" The emission atinfrared. wavelengths: (Άγιοι lim) is believed. to be dominated by reprocessed thermal radiation [rom hot. dus on the inner walls of the obscuring ""torus. where any anisotropy. is almost certainty due to extrinsic cllects such as dust obscuration."," The emission atinfrared wavelengths $\lambda_{\rm rest} > 1\,\mu$ m) is believed to be dominated by reprocessed thermal radiation from hot dust on the inner walls of the obscuring `torus', where any anisotropy is almost certainly due to extrinsic effects such as dust obscuration."686 We can therefore hope to discriminate between intrinsic and extrinsic sources of anisotropy., We can therefore hope to discriminate between intrinsic and extrinsic sources of anisotropy.687 The 3€ sample of radio sources is an obvious place to start such investigations., The 3C sample of radio sources is an obvious place to start such investigations.688 Our earlier study of 2~1 3€ radio galaxies is described in Simpson. Rawlines Lacy (1999. hereafter SRL). and we present. here photometry and a preliminary statistical analysis of απ identically selected. sample of quasars.," Our earlier study of $z \sim 1$ 3C radio galaxies is described in Simpson, Rawlings Lacy (1999, hereafter SRL), and we present here photometry and a preliminary statistical analysis of an identically selected sample of quasars."689 Ln Section 2. we describe our observations and reduction method.," In Section 2, we describe our observations and reduction method."690 In Section 3. we present the results of our. photometry and. derive the optical and near-infrared. continuum spectral indices.," In Section 3, we present the results of our photometry and derive the optical and near-infrared continuum spectral indices."691 In Section 4. we look for correlations among the properties of the quasars in our sample. and investigate unusual objects.," In Section 4, we look for correlations among the properties of the quasars in our sample, and investigate unusual objects."692 We summarize our results in Section 5., We summarize our results in Section 5.693 Throughoutthispaper. we adopt Hy=5ükkmss + qo= 0.5. and A= 0.," Throughoutthispaper, we adopt $H_0 = 50$ $^{-1}$ $^{-1}$ , $q_0 =6940.5$ , and $\Lambda = 0$ ."695 Our, Our696confirmation of the power of the LM method to extract consistently distinct eroups in biased samples of a οἼνοιι stellar population.,confirmation of the power of the LM method to extract consistently distinct groups in biased samples of a given stellar population.697 Table 6 eives the umber of stars in our sample which are assigned to every crossed eroup CG) Le. to group C in IX aud CC in IRAS., Table \ref{tab_kincross} gives the number of stars in our sample which are assigned to every crossed group G(G') i.e. to group G in K and G' in IRAS.698 Iu Sect.5.2.1 our LPVs sample is shown to be representative of the LPVs population as far as the kinematics is concerned., In \ref{sec_kinbias} our LPVs sample is shown to be representative of the LPVs population as far as the kinematics is concerned.699 Thusgroup., Thus.700 Obviously such a consideration does not apply to the DIunuimosities (see Sect.5.2.2))., Obviously such a consideration does not apply to the luminosities (see \ref{sec_lumbias}) ).701 The assigned eroups are eiven in annex A (electronic table)., The assigned groups are given in annex A (electronic table).702 Table 6 eives the values of the axes of the velocity ellipsoids aud the scale height of each of the 7 crossed Ix and IRAS eroups., Table \ref{tab_kincross} gives the values of the axes of the velocity ellipsoids and the scale height of each of the 7 crossed K and IRAS groups.703 Civen that our sample is represcutative of the population in terms of kinciatics. as already. sccu in Sect. 5.2.1...," Given that our sample is representative of the population in terms of kinematics, as already seen in Sect. \ref{sec_kinbias},"704 we can use the kinematical values of table G as representative in terms of ealactic The relation between the mean kinematics of a galactic »pulation and its age allows us to estimate the rauge of ages of the eroups., we can use the kinematical values of table \ref{tab_kincross} as representative in terms of galactic The relation between the mean kinematics of a galactic population and its age allows us to estimate the range of ages of the groups.705 Furthermore. classical statistical studies of stars known to belong to different ealactic »opulatious aud of different moetallicity abundances allow us to add an estimate of the ranee of metallicity.," Furthermore, classical statistical studies of stars known to belong to different galactic populations and of different metallicity abundances allow us to add an estimate of the range of metallicity."706 By comparing the values in table 6 with the results on snematies and metallicity of the galactic populations by A\Ghalas and Binney (1981) aud by Stromercu (1987). we can deduce:," By comparing the values in table \ref{tab_kincross} with the results on kinematics and metallicity of the galactic populations by Mihalas and Binney (1981) and by Stromgren (1987), we can deduce:"707for? ete.,for etc.708 SinceQ.. we find the simple relation Both sides in this equation are functions of 2. although we liave uot iudicated tle dependence explicitly.," Since, we find the simple relation Both sides in this equation are functions of $\beta$, although we have not indicated the dependence explicitly."709 Iu any case. the relation coulirus our expectation that the mass accretion rate is iudependent of the spheres radius rin steady-state We also now see that the higher-order variables f.4. fο. ete.," In any case, the relation confirms our expectation that the mass accretion rate is independent of the sphere's radius $r$ in steady-state We also now see that the higher-order variables $f_{-1}$, $f_{-2}$, etc."710 play no part in cleterminine this rate., play no part in determining this rate.711 Now that we have tied the mass accretion rate to fo(0). we can inunediately rule out a subset of outer flow solutious as being unplivsical.," Now that we have tied the mass accretion rate to $f_0 (0)$, we can immediately rule out a subset of outer flow solutions as being unphysical."712 Figure 3. shows that. for1.12... fü(0) is negative. corresponding to atet mass efflux.," Figure \ref{fig:2ndfandg} shows that, for, $f_0 (0)$ is negative, corresponding to a net mass efflux."713 That such a situation is even. possible empliasizes once 1lore the need to extend tje. flow solution inward across the sotic surface., That such a situation is even possible emphasizes once more the need to extend the flow solution inward across the sonic surface.714 For this same cvoice of gοίπ]. the dotted curve in he lower panel of Figure 3. shows hatat).," For this same choice of $g_{-2} (\pi)$, the dotted curve in the lower panel of Figure \ref{fig:2ndfandg} shows that."715 Iileed. we have just found one example of a general result: the cdillerence agrees ln sien with fü(0).," Indeed, we have just found one example of a general result: the difference agrees in sign with $f_0 (0)$."716 We now show hat the two quantities are in [act ecual. apart from a multipicative factor.," We now show that the two quantities are in fact equal, apart from a multiplicative factor."717 Our proof starts with the [act that the leftμαμα side of the secouc-orcder equation (36)) is a perfect derivative., Our proof starts with the fact that the lefthand side of the second-order equation \ref{eqn:secondt2}) ) is a perfect derivative.718 Specifically. Turning to the righthand side of the same equation. we note first that sin@is an even function ol 0—2/2. while cos0 is an odd function.," Specifically, Turning to the righthand side of the same equation, we note first that ${\rm sin}\,\theta$is an even function of $\theta-\pi/2$, while ${\rm cos}\,\theta$ is an odd function."719 Since P depends only on sind. it has even syinmetry.," Since $\cal D$ depends only on ${\rm sin}\,\theta$, it has even symmetry."720 Iuspectiou shows that the righthaud side of equation (36)) has odd syiunetry., Inspection shows that the righthand side of equation \ref{eqn:secondt2}) ) has odd symmetry.721 If we now integrate equation (36)) [rom to 0. the righthaud side vanishes because of the odd syiumetry of the integrand.," If we now integrate equation \ref{eqn:secondt2}) ) from to 0, the righthand side vanishes because of the odd symmetry of the integrand."722 We find that Since aud 1.. we have which we recast as," We find that Since and , we have which we recast as"723directly computed from the IME.,directly computed from the IMF.724 The nature of the Type Ia SNe progenitors is still a matter of debate., The nature of the Type Ia SNe progenitors is still a matter of debate.725 Indeed. it is not vet clear which of the two main compctine models (DD. Then Tutukov. 1951) or (SD. Whelan Then 1973) applics for the SNIa precursors.," Indeed, it is not yet clear which of the two main competing models – (DD, Iben Tutukov 1984) or (SD, Whelan Iben 1973) – applies for the SNIa precursors."726 ere we assume the SD iunodel where Type Ia SNe are produced by C-ignition aud total disruption of a cold degenerate WD when this latter exceeds the Chandrasekhar wiass after mass transfer i a close binary svstem., Here we assume the SD model where Type Ia SNe are produced by C-ignition and total disruption of a cold degenerate WD when this latter exceeds the Chandrasekhar mass after mass transfer in a close binary system.727 We estimate the rate of Type Ia SNe according to the formalism by Ferrini et al. (, We estimate the rate of Type Ia SNe according to the formalism by Ferrini et al. (7281992: sec also Cregeio Reuzini 1983).,1992; see also Greggio Renzini 1983).729" In ou SD uodelliug. the free parameters ip;,¢ and og. are respectively the lower limit of the total mass of the binary svstems which can produce Type Ia SN. aud the fraction of the total mass of stars which belong to these systems."," In our SD modelling, the free parameters $m_{Binf}$ and $\alpha_{0}$, are respectively the lower limit of the total mass of the binary systems which can produce Type Ia SN, and the fraction of the total mass of stars which belong to these systems."730 They have been adjusted in order to reproduce the main properties of the solar neighbourhood., They have been adjusted in order to reproduce the main properties of the solar neighbourhood.731 We find that the couple ΕΡΕ.eg]=[3.0.05] correctly reproduces the data.," We find that the couple $[m_{Binf},{\alpha}_{0}]=[3,0.05]$ correctly reproduces the data."732 Tn order to assess the uncertainties iu the nature of the progenitors. we have also used the empirical paraiueterization of the rates of SN Iun as derived by Ciotti et al. (," In order to assess the uncertainties in the nature of the progenitors, we have also used the empirical parameterization of the rates of SN Ia as derived by Ciotti et al. ("7331991: hereafter C91) RiyaXUsxt* Osx,1991; hereafter C91) $R_{Ia}\propto \theta_{SN} t^{-s}$.734 isthe normalizing parameter., $\theta_{SN}$ is the normalizing parameter.735" The Ry, evolution is coutrolled by the free parameter s.", The $R_{Ia}$ evolution is controlled by the free parameter $s$.736 Iu. order to accouut for the Fe content in clusters of galaxies. COL concluded that s> 1.1.," In order to account for the Fe content in clusters of galaxies, C91 concluded that $s > 1.4$ ."737 Note that our standard modelling roughly corresponds to s~1.6., Note that our standard modelling roughly corresponds to $s \sim 1.6$.738 The other free parameter is the rise time which is fixed to ty5.y=0.05 i uuits of 15 Cyr (Reuzini et al.," The other free parameter is the rise time which is fixed to $t_{15,0}$ =0.05 in units of 15 Gyr (Renzini et al."739 1993)., 1993).740 Dv combining different photometric surveys such as the CFERS. up to + = 1 (Lilly et al.," By combining different photometric surveys such as the CFRS, up to $z$ = 1 (Lilly et al."741 1995) aud the IDF up to 2 = | (AD96). it has been possible to derive a picture of he star formation rate historv of the whole universe (M96. Madan et al.," 1995) and the HDF up to $z$ = 4 (M96), it has been possible to derive a picture of the star formation rate history of the whole universe (M96, Madau et al."742 1997: M97)., 1997; M97).743 To compute the evolution of the COSMIC SUpCrnova rates per comoving volume (CSNER). we nake use of the CSER as derived iu M97. (110dol MI).," To compute the evolution of the cosmic supernova rates per comoving volume (CSNR), we make use of the CSFR as derived in M97 (model M1)."744 We also introduce another CSER law ο account for a yosstble dust extinction correction αποςς M2)., We also introduce another CSFR law to account for a possible dust extinction correction (model M2).745 MI aud AI2 have he same shape at low : in agreemen with the ΣΕΤ derived from observations. but at higher τν we have allowed the CSFR to be higher iu model M2. than iu AIL in agreement with Pettini et al. (," M1 and M2 have the same shape at low $z$ in agreement with the CSFR derived from observations, but at higher $z$, we have allowed the CSFR to be higher in model M2 than in M1 in agreement with Pettini et al. ("7461997).,1997).747 These time-dependen CSFRs are then introduced as an input iu our code., These time-dependent CSFRs are then introduced as an input in our code.748 The amplitude is normalized in order to match the niüuositv density evolution in the UVcoutiuuuu with ie following coustants of proportionality between the ENjserved Dr and our derived SFR: (7.5<104.6.5« 1079] i units of W | AL. 3) 1 at (1500AL. 2800 A3).," The amplitude is normalized in order to match the luminosity density evolution in the UV–continuum with the following constants of proportionality between the observed $_{UV}$ and our derived SFR: $7.5\times 10^{19}, 6.5\times 10^{19}$ ) in units of W $^{-1}$ $_{\odot}$ $^{-1}$ $^{-1}$ at (1500, 2800 )."749" Before computing the SN rates we first have checked whether our simple model with a ""standard IME aud le paranueterization of the CSER is able to reproduce i6 observed volume averaged denusitv lunmuinosities p at different A."," Before computing the SN rates, we first have checked whether our simple model with a “standard” IMF and the parameterization of the CSFR is able to reproduce the observed volume averaged density luminosities $\rho_{\nu}$ at different ${\lambda}$."750 Figure 1 shows that the aerecmeut at longer A is satisfvius., Figure 1 shows that the agreement at longer ${\lambda}$ is satisfying.751 The L100 and 100600 luminosity deusities seen to hint at model M2 in agreement with the IR/subniuu backerouncd which also sugeests the preseuce of extinction at high + (Guiderdoni et al., The 4400 and 10000 luminosity densities seem to hint at model M2 in agreement with the IR/submm background which also suggests the presence of extinction at high $z$ (Guiderdoni et al.752 1997)., 1997).753 Iu this section. we use the CSER to derive the CSNR or both IL/Ib.c aud Ia Types.," In this section, we use the CSFR to derive the CSNR for both II/Ib,c and Ia Types."754 To our knowledec. this is he first time that the evolution of SN rates with redshift is predicted from a sclf&cousistent spectro-plotometric nodelling of galaxy evolution at Ligh : aud iudepeudoeutlv of the details of individual galaxy evolution.," To our knowledge, this is the first time that the evolution of SN rates with redshift is predicted from a self-consistent spectro-photometric modelling of galaxy evolution at high $z$ and independently of the details of individual galaxy evolution."755 As already uentionned. the direct measurement of SNe cau be used as an independent test for the cosmic star and metal onuation iu the universe.," As already mentionned, the direct measurement of SNe can be used as an independent test for the cosmic star and metal formation in the universe."756 In figure 2. we have plotted he predicted evolution of the CSNR per uuit of comoving voluue with redshift.," In figure 2, we have plotted the predicted evolution of the CSNR per unit of comoving volume with redshift."757 The Type II/To.c rate shows the sale shape as the instautancous CSER. that is the rise. oeak and drop from hieh redshift to the preseut time.," The Type II/Ib,c rate shows the same shape as the instantaneous CSFR, that is the rise, peak and drop from high redshift to the present time."758 This mcans that the SNIL/Tb.c rate cau be used as an independent tracer of the star formation rate.," This means that the SNII/Ib,c rate can be used as an independent tracer of the star formation rate."759" The Type Ta rate Ry, has a different shape from the CSFR while coincidently it has uecarly the same behaviour as the B-ο huuinositv.", The Type Ia rate $R_{Ia}$ has a different shape from the CSFR while coincidently it has nearly the same behaviour as the B-band luminosity.760 The most important poiut is the time delay we observe between the stellar (binary in this case) mth and the explosion time., The most important point is the time delay we observe between the stellar (binary in this case) birth and the explosion time.761 The occurrence of the SNla rate peak is shifted by a few Cars relatively to he CSFR peal., The occurrence of the SNIa rate peak is shifted by a few Gyrs relatively to the CSFR peak.762 However. at 2> 0.9. Type Ia SN can ο used as a probe of the past historv of the CSFR.," However, at $z \ge$ 0.9, Type Ia SN can be used as a probe of the past history of the CSFR."763" Firthermore. as can be seen from figure 3. Ry, lias very distinct shape (different normalization aud different iue of peak occurrence) depending on the adopted CSER shape. aud this difference is higher at higher :."," Furthermore, as can be seen from figure 3, $R_{Ia}$ has very distinct shape (different normalization and different time of peak occurrence) depending on the adopted CSFR shape, and this difference is higher at higher $z$."764 Therefore. mcasurements of SNIa at το 1 would be able ο discriniuate between the models.," Therefore, measurements of SNIa at $z {\sim}$ 1 would be able to discriminate between the models."765 Supernova rate has been measured in ucarby galaxies w several authors (Cappellarvo Turatto 1988. Ένας et al," Supernova rate has been measured in nearby galaxies by several authors (Cappellaro Turatto 1988, Evans et al."766 1989. C97).," 1989, C97)."767 At higher redshift. Pain aud his collaborators have receutlv reported the first ligh-: (5~ L1) rest-frame Type Ia SN rate with a value of ~ 152 h SNu.," At higher redshift, Pain and his collaborators have recently reported the first $z$ $z \sim$ 0.4) rest-frame Type Ia SN rate with a value of ${\sim}$ 0.82 $^{2}$ SNu."768 In order to compare the model to the observations. we have couverted the observed rate iu SNu (SNo/100 vr/ 10! L5.) into a ate iu | using due huuinosities as computed by our code.," In order to compare the model to the observations, we have converted the observed rate in SNu (SNe/100 yr/ $^{10}$ $_{B\odot}$ ) into a rate in $^{-1}$ using blue luminosities as computed by our code."769 From figure we can see that although the available imeasureimeut of the SNIa rate does uot allow one to cüscrininate )tween models MI and M2. the redshift is still low iid the statitics is poor.," From figure 3, we can see that although the available measurement of the SNIa rate does not allow one to discriminate between models M1 and M2, the redshift is still low and the statitics is poor."770 Measurementat 2~ 1 with he same statistics. would beein to discriminate between uodels. independently of the adopted model for Type Ia xogenitors.," Measurementat $z \sim$ 1 with the same statistics, would begin to discriminate between models, independently of the adopted model for Type Ia progenitors."771 By investigating different models for SNIa rates using COL empirical paraineterization (see figure, By investigating different models for SNIa rates using C91 empirical parameterization (see figure772of a sample of eight low redshift star foriiug galaxies.,of a sample of eight low redshift star forming galaxies.773" We take BC,=1.66 (A= 1600A)) computed by MIIC99 from a rauge of Starburst99 svuthetic starburst spectra (Leiherer et al.", We take $_{\rm stars} = 1.66$ $\lambda=1600$ ) computed by MHC99 from a range of Starburst99 synthetic starburst spectra (Leitherer et al.774 1999)., 1999).775 We lave obtained £u-UV fluxes for 29 of the 71 ealaxies in the sample from either the UV catalog of Marcum et al. (, We have obtained far-UV fluxes for 29 of the 71 galaxies in the sample from either the UV catalog of Marcum et al. (7762001) at ουν he homogenized UV. catalog of Rifatto et al. (,2001) at or the homogenized UV catalog of Rifatto et al. (7771995) at1650À.,1995) at.778. O ‘the 71 galaxies in this sample. 28 have 30 upper Innuits to the N-rav flux.," Of the 71 galaxies in this sample, 28 have $3\sigma$ upper limits to the X-ray flux."779 We accouut for this fact aud. compute the fit to the correlation by using the lear regression. techuique which computes coefficieuts based on Waplau-Meier residuals., We account for this fact and compute the fit to the correlation by using the linear regression technique which computes coefficients based on Kaplan-Meier residuals.780 If we consider only the far-IR cunission when approximating the bolometric lnnunosity (Lp; Lrig DCqua) we can use all 71 ealaxies to find the fit. shown by the solid line in the bottom panel of Figure 1.," If we consider only the far-IR emission when approximating the bolometric luminosity $_{\rm Bol}$ = $_{\rm FIR} \times$ $_{\rm dust}$ ) we can use all 71 galaxies to find the fit, shown by the solid line in the bottom panel of Figure 1."781 The error is lo., The error is $1\sigma$.782 Η we fit oulv the 29 with both far-IR aud far-UV data (of which 11 have N-rav upper lnuuits) using Equation 1 we compute the ft. which is represented by the dashed ie in the same figure.," If we fit only the 29 with both far-IR and far-UV data (of which 11 have X-ray upper limits) using Equation 1 we compute the fit, which is represented by the dashed line in the same figure."783 The fit becomes slightly steeper when the UV data are properly cousidered., The fit becomes slightly steeper when the UV data are properly considered.784 This is due to the act that the far-UV contributes measurably to Lp in svsteiis with lower ar-IBRO Dpuuiuosities. but has a uesligible effect ou Lp at higher far-IR luminosities.," This is due to the fact that the far-UV contributes measurably to $_{\rm Bol}$ in systems with lower far-IR luminosities, but has a negligible effect on $_{\rm Bol}$ at higher far-IR luminosities."785 The effect is lighliehted iu he top xuel of figure 1 where we plot the difference 08 Epp.gU] log Lyoyriny.," The effect is highlighted in the top panel of figure 1 where we plot the difference log $_{\rm Bol[FIR,FUV]}$ – log $_{\rm786Bol[FIR]}$."787 The 29 data points have COLL STOliped in bius of l dex ancl the mean diffCrence is plotte against log Lyviprimy., The 29 data points have been grouped in bins of 1 dex and the mean difference is plotted against log $_{\rm Bol[FIR]}$.788 The line is a simple chi-square Direar fit., The line is a simple chi-square linear fit.789 The trend is understood iuost sinplv as odd extinction. effect. where the £u-UV extinelon 1s positively correlated with Lp (IHeockman et al.," The trend is understood most simply as a extinction effect, where the far-UV extinction is positively correlated with $_{\rm Bol}$ (Heckman et al."790 1995)., 1998).791 We use the tιο correlation derived from the £u-UV axd far-IR (Equation 3) for predicting the LBG N-arav. output., We use the the correlation derived from the far-UV and far-IR (Equation 3) for predicting the LBG X-ray output.792 It is interesting to note that the mean observed N-rav huninosity of the LBC sample suggests a large mean Lp (~ Lem) of ~LOY L..," It is interesting to note that the mean observed X-ray luminosity of the LBG sample suggests a large mean $_{\rm793Bol}$ $\approx$ $_{\rm FIR}$ ) of $\sim 10^{12}$ $_{\odot}$."794 This is at the high Iuniuosity cud of the local starburst population., This is at the high luminosity end of the local starburst population.795 The assumption that we can model the far-IR to far-UV flux ratio as a function of UV reddening is the crucial compouent iu our techuique of predicting the X-lIAV OLudssloi of the LBC sample., The assumption that we can model the far-IR to far-UV flux ratio as a function of UV reddening is the crucial component in our technique of predicting the X-ray emission of the LBG sample.796 We first consider the clupirically derived UW reddening relation of ATIC99 (RX-.3) for local starbursts., We first consider the empirically derived UV reddening relation of MHC99 $\beta$ ) for local starbursts.797 The UV sxectral slope (2) is found from the photometric 1GUEIa color.," The UV spectral slope $\beta$ ) is found from the photometric $V_{\rm798606}-I_{\rm 814}$ color."799 We use equation 1L of MIIC99 which calibrates t1ο spectroscopically defined § to the, We use equation 14 of MHC99 which calibrates the spectroscopically defined $\beta$ to the800Although particles .from solar flares have been observed for 50 vears (Forbush1916)). there are still many unresolved questions about them acceleration sites and iechauisiis,"Although particles from solar flares have been observed for 50 years \cite{f46}) ), there are still many unresolved questions about their acceleration sites and mechanisms."801 Much receut research has focused1 onu a siniple classification scheme proposed by Pallavicini. Serio. Vaiana (1977).," Much recent research has focused on a simple classification scheme proposed by Pallavicini, Serio, Vaiana (1977)."802 In modern termuneloey. nupulsive solar flares are typically defined as those with a short ἐς hj) duration of X-ray cussion. while Sradual have a louger Acuy duratiou.," In modern terminology, “impulsive” solar flares are typically defined as those with a short $\lesssim 1$ h) duration of X-ray emission, while “gradual” flares have a longer X-ray duration."803 These two classes of Haresflares have been found to have several distinguishing ∢∢∖characteristics., These two classes of flares have been found to have several distinguishing characteristics.804" For DOexample,ON Norayvast eiission—eges impulsive flares is obrved to com from οςnupact frou. at low coronal heights ος10 au). while X- ου...from flares tend to arise frou broader raysregions or from graduallarge coronal loops up to ~107 kan above the photosphere."," For example, X-ray emission from impulsive flares is observed to come from compact regions at low coronal heights $\lesssim 10^4$ km), while X-rays from gradual flares tend to arise from broader regions or from large coronal loops up to $\sim10^5$ km above the photosphere."805" In addition. it is widely believed.⋅ that the main⋅ acceleration⋅ mechauisi⋅ for. particles. escaping. from. inpulsive. flares is. stochastic: (secoud-order. Fermi acceleration: bycistributious evroresonaut plasma waves: e.g... Tomer&Roth1992.Millerb""nas 1993)). while eradual fiflares are associatedted withwit! shock acceleration hieher iu the corona (e.g... Cane. 1986)). though not all observations support this paracdieni (e.g... Mazuretal. 1992))."," In addition, it is widely believed that the main acceleration mechanism for particles escaping from impulsive flares is stochastic acceleration (second-order Fermi acceleration by gyroresonant plasma waves; e.g., \cite{tr92,mv93}) ), while gradual flares are associated with shock acceleration higher in the corona (e.g., \cite{cea86,lr86}) ), though not all observations support this paradigm (e.g., \cite{mea92}) )."806 A more controversial issue is why solar cosmic ravs are observed a locations t are magueticalVv connected to solar longitudes far iatfrom the longitude of: theH flareqvo site., A more controversial issue is why solar cosmic rays are observed at locations that are magnetically connected to solar longitudes far from the longitude of the flare site.807uite 3n (1interplanetarynbars diffusionMua isle highly anisotropic. Decimhlibitius motion perpendicular to theEH magneticwastie fieldFa (o5...∖⊾ Palmer2.iv|αι 19531). foruv decadesMs x]it was asstuned that particles are transported to other solar longitudes within the solar corona. whence they escape fo travel along the interplanetary magnetic field to the observer.," Because interplanetary diffusion is highly anisotropic, inhibiting motion perpendicular to the magnetic field (e.g., \cite{p82}) ), for decades it was assumed that particles are transported to other solar longitudes within the solar corona, whence they escape to travel along the interplanetary magnetic field to the observer."808 The sctminal quantitative model of Reid (1961) assunes isotropic. two-dimensional. diffusion. gives specific for the coronalinjection rate of audparticles into the predictionsinterplanetary medium as a function of time aud coronal distauce from the flare site.," The seminal quantitative model of Reid (1964) assumes isotropic, two-dimensional, coronal diffusion, and gives specific predictions for the injection rate of particles into the interplanetary medium as a function of time and coronal distance from the flare site."809" It is not clear exactly what mechamisis eive rise to coronal diffusion. though Newkirk Weutzel (1978) preseuted the ""bird cage” model of rieidity-iudepeudenut coronal transport. in which flare-accelerated particles bonuce back and forth imde coronal loops aud occasionally trausfer to other loops at the footpoiuts."," It is not clear exactly what mechanisms give rise to coronal diffusion, though Newkirk Wentzel (1978) presented the “bird cage” model of rigidity-independent coronal transport, in which flare-accelerated particles bounce back and forth inside coronal loops and occasionally transfer to other loops at the footpoints."810 On the other hix. it has recently been proposed that for eradual flares. with are often associated with large.? iuterplauetary| shocks. |particles are freshly accelerated on open field Hues at differcut heliolongitudoes (Mason.Cloeckler.&Iovestadt198.Reames1990)).," On the other hand, it has recently been proposed that for gradual flares, with are often associated with large, interplanetary shocks, particles are freshly accelerated on open field lines at different heliolongitudes \cite{mgh84,r90}) )."811 It is further assumed that coronal transport does not occur. and that for iupulsive flares. which are usually not associated with interplanetary shocks. the narrower longitucinal dispersion (Reames.Cane.&vouRosenvinge1990)) is due to the spreading of maenetic field lines i the solar corona or the interplanetary inediuu.," It is further assumed that coronal transport does not occur, and that for impulsive flares, which are usually not associated with interplanetary shocks, the narrower longitudinal dispersion \cite{rea90}) ) is due to the spreading of magnetic field lines in the solar corona or the interplanetary medium."812" Therefore. it is no louger clear whether several decades worth of solar cosnüc rav observations were τοντο information about the solar Bare site, as previously assumed. or about accddleyation. at coronal or interplanetary. shocks."," Therefore, it is no longer clear whether several decades worth of solar cosmic ray observations were providing information about the solar flare site, as previously assumed, or about acceleration at coronal or interplanetary shocks."813 Further progress in the interpretation of the information emibocdied iu solar cosumic ταν observations thus depends on the resolution of this iso., Further progress in the interpretation of the information embodied in solar cosmic ray observations thus depends on the resolution of this issue.814 Receut observations which can shed much light ou this⋅⋅ issue concern the charge state distributions⋅⋅⋅ ofD. ious accelerated as a result of. solar flares., Recent observations which can shed much light on this issue concern the charge state distributions of ions accelerated as a result of solar flares.815 Such. for⋅⋅⋅ various⋅ Clements. provide a. rici source. acceleration of information ou the conditious of particle acceleration and escape from the corona., Such distributions for various elements provide a rich source of information on the conditions of particle acceleration and escape from the corona.816 Results of the ULEZEQ iustirmucnt on board theZSEE-3 spacecraft indicatcc that for several gradual flares. the charge states of 9 different clemeuts were not consistent with the sale temperature (Luluetal.198 D). but rather hac apparcut ionization temperatures of d to 84109 Wy. Iu coutrast. πο results for several less powertu iupulsive fares suggested that charge states of Si an Fe were characteristic: of sienificautlysig hotter plasma↴∙ as night. be expected ifqt particle. acceleration.. takes place at the site of a compact. impulsive flare (Luluetal. 1987)).," Results of the ULEZEQ instrument on board the spacecraft indicated that for several gradual flares, the charge states of 9 different elements were not consistent with the same temperature \cite{lea84}) ), but rather had apparent ionization temperatures of 1 to $\times10^6$ K. In contrast, summed results for several less powerful impulsive flares suggested that charge states of Si and Fe were characteristic of significantly hotter plasma, as might be expected if particle acceleration takes place at the site of a compact, impulsive flare \cite{lea87}) )."817 More receuth. measurements by three iustruments on board theSAALPEN mission (Octlikeretal.1995.LeskectMason 1995)) have eeneralle coufirmed the earlier results for gradual Hares. aud also provide measurements for more elements. with better statistical acemacy. aud for a broader chorey ranse.," More recently, measurements by three instruments on board the mission \cite{oea95,lea95,mea95}) ) have generally confirmed the earlier results for gradual flares, and also provide measurements for more elements, with better statistical accuracy, and for a broader energy range."818 A possible explanation for why the nean charges of energetic ious from gradual eveuts are not cousisteut with a single temperature was eiven y Alan and Waldron (19SG). who proposed that photoionization bv flare N-ravs changes the ionization equilibrimmn in the plasma from which ions are accelerated. and obtained good quantitative agreement with the apparent ionization temperatures for various eleiieuts.," A possible explanation for why the mean charges of energetic ions from gradual events are not consistent with a single temperature was given by Mullan and Waldron (1986), who proposed that photoionization by flare X-rays changes the ionization equilibrium in the plasma from which ions are accelerated, and obtained good quantitative agreement with the apparent ionization temperatures for various elements."819 Another possible explanation in terms of frecze-out temperatures was discussed by Mason et ((1995)., Another possible explanation in terms of freeze-out temperatures was discussed by Mason et (1995).820 Iu this report. we exanune effects ou the charge states of. energetic ⋅⋅ious associated. with. eradual solar," In this report, we examine effects on the charge states of energetic ions associated with gradual solar"821star clusters.,star clusters.822 Their ages range from 100 to 600 million years., Their ages range from 100 to 600 million years.823 However. only two of them (NGC 5715 and NGC 6268) are very voung. and their distances are compatible with the location of the Carina-Sagittarius arm (Russeil 2003).," However, only two of them (NGC 5715 and NGC 6268) are very young, and their distances are compatible with the location of the Carina-Sagittarius arm (Russeil 2003)."824 We can conclude hey formed inside the arm and are tracers of the arm. since with such ages they could not travel much away from their Unfortunately. we fail to find any young cluster located beyond the Carina-Sagittarius arm in the present sample.," We can conclude they formed inside the arm and are tracers of the arm, since with such ages they could not travel much away from their Unfortunately, we fail to find any young cluster located beyond the Carina-Sagittarius arm in the present sample."825 This clearly reflects the difficulty to penetrate the arm and to see further away because of the high density of dust and gas. unless absorption 10les allow to detect more distant clusters (see Vázzquez et al.," This clearly reflects the difficulty to penetrate the arm and to see further away because of the high density of dust and gas, unless absorption holes allow to detect more distant clusters (see Vázzquez et al."826 1995: Baume et al., 1995; Baume et al.827 2009)., 2009).828 The oldest clusters (Czernik. 38 and NGC 5715). are sarticularly interesting in. the context of cluster. dynamical evolution and dissolution models (Lamers et al 2005). since they could survive longer than the the typical open cluster life-time in a dense and hostile environment like the inner disk. where tidal orces and close encounters do not permit star clusters to survive ypically more than 100-200 Myrs (Wielen Not many clusters of this age or older are known o be located at these Galacto-centric distances (see tpwww.univie.ac.atwebda/navigation.html).," The oldest clusters (Czernik 38 and NGC 5715) are particularly interesting in the context of cluster dynamical evolution and dissolution models (Lamers et al 2005), since they could survive longer than the the typical open cluster life-time in a dense and hostile environment like the inner disk, where tidal forces and close encounters do not permit star clusters to survive typically more than 100-200 Myrs (Wielen Not many clusters of this age or older are known to be located at these Galacto-centric distances (see http://www.univie.ac.at/webda/navigation.html)."829 This combination of age and distance is extremely useful to investigate the Galactic disk radial abundance gradient in the inner disk (Magrini et al., This combination of age and distance is extremely useful to investigate the Galactic disk radial abundance gradient in the inner disk (Magrini et al.830 2010) and its evolution through Therefore. these two clusters are ideal targets for future spectroscopic follow-up to determine their metal abundances.," 2010) and its evolution through Therefore, these two clusters are ideal targets for future spectroscopic follow-up to determine their metal abundances."831 We acknowledge the staff ofΤΟ and LCO. in particular Edgardo Cosgrove and Patricio Pinto. for their valuable support during the runs.," We acknowledge the staff of CTIO and LCO, in particular Edgardo Cosgrove and Patricio Pinto, for their valuable support during the runs."832 The work of Α.Ε. Seleznev has been partly supported by the ESO Visiting Scientist Program., The work of A.F. Seleznev has been partly supported by the ESO Visiting Scientist Program.833 We are very grateful to Sandy Strunk. who carefully revised the paper and helped us to improve the language.," We are very grateful to Sandy Strunk, who carefully revised the paper and helped us to improve the language."834 This study made use of the SIMBAD and WEBDA databases., This study made use of the SIMBAD and WEBDA databases.8352009).,.836". The galaxy undergoes starburst upon final coalescence, and peaks in its bolometric luminosity."," The galaxy undergoes a starburst upon final coalescence, and peaks in its bolometric luminosity."837a Radiative transfer post-processing on these models have found that the model galaxy is then visible as a ULIRG 2007)., Radiative transfer post-processing on these models have found that the model galaxy is then visible as a ULIRG .838. The same gaseous inflows can drive sufficient black hole growth to result in optical quasar activity 2006)., The same gaseous inflows can drive sufficient black hole growth to result in optical quasar activity .839". Effects of the AGN feedback can be seen in both the warm infrared colours of the galaxy 2009),, as well as molecular outflows 2008b)."," Effects of the AGN feedback can be seen in both the warm infrared colours of the galaxy , as well as molecular outflows ."840". The truncation of the starburst by a combination of gas consumption and AGN feedback can render the galaxy observable as an E+A post-starburst before it evolves into into a dead early-type with colours comparable to those observed on the red-sequence 2008b,c).", The truncation of the starburst by a combination of gas consumption and AGN feedback can render the galaxy observable as an E+A post-starburst before it evolves into into a dead early-type with colours comparable to those observed on the red-sequence .841". The kinematic 2006c),, X-ray 2006a),, nuclear emission 2009),, and molecular disc properties of this merger remnant have all been studied and found to be comparable to those observed."," The kinematic , X-ray , nuclear emission , and molecular disc properties of this merger remnant have all been studied and found to be comparable to those observed."842" Similarly, the remnant lies on the fundamental plane 2008a)."," Similarly, the remnant lies on the fundamental plane ."843". In the far left panel of Figure 1,, we show the velocity-integrated brightness temperature map of the model disc galaxy."," In the far left panel of Figure \ref{figure:iso_d3_map}, we show the velocity-integrated brightness temperature map of the model disc galaxy."844" As expected, the central regions are the brightest, and the outer disc has little CO emission."," As expected, the central regions are the brightest, and the outer disc has little CO emission."845" In the top left panel of Figure 2,, we plot the emission- distribution of vvalues for the GMCs in our model disc galaxy and the fiducial model merger."," In the top left panel of Figure \ref{figure:xco_distribution}, we plot the emission-weighted distribution of values for the GMCs in our model disc galaxy and the fiducial model merger."846 We additionally plot the distribution of GMC physical properties in both the disc and merger., We additionally plot the distribution of GMC physical properties in both the disc and merger.847 We will return to this plot frequently throughout this section and the next., We will return to this plot frequently throughout this section and the next.848 The luminosity-weighted iin our model disc is ~4x wwith a relatively narrow dispersion.," The luminosity-weighted in our model disc is $\sim 4 \times849 10^{20}$ with a relatively narrow dispersion."850" The dispersion is narrow because the surface densities, kinetic temperatures and velocity dispersions of the model disc GMCs show fairly little variation."," The dispersion is narrow because the surface densities, kinetic temperatures and velocity dispersions of the model disc GMCs show fairly little variation."851" To remind the reader, the column densities in the GMCs in our disc galaxy are set to be the surface density of cold gas in the cell."," To remind the reader, the column densities in the GMCs in our disc galaxy are set to be the surface density of cold gas in the cell."852" When the GMC is unresolved in the simulation, we set the subgrid value of the surface density to Nctoua=100Mo/pe’."," When the GMC is unresolved in the simulation, we set the subgrid value of the surface density to $\Sigma_{\rm853 cloud} = 100 \ \msun/{\rm pc}^{2}$."854 This value was chosen to match the roughly constant surface density of Galactic molecular clouds., This value was chosen to match the roughly constant surface density of Galactic molecular clouds.855 Nearly all of the GMCs in the model disc take on this value for a surface density., Nearly all of the GMCs in the model disc take on this value for a surface density.856" The kinetic temperatures of GMCs in the disc have a relatively tight distribution near 10 K, as shown in Figure 2.."," The kinetic temperatures of GMCs in the disc have a relatively tight distribution near 10 K, as shown in Figure \ref{figure:xco_distribution}."857" Because the GMCs have a relatively low density compared with starbursts (the mass-weighted value is ~500 ?)), there is little coupling with the dust grains (which are a factor of a few hotter; Figure 2))."," Because the GMCs have a relatively low density compared with starbursts (the mass-weighted value is $\sim500$ ), there is little coupling with the dust grains (which are a factor of a few hotter; Figure \ref{figure:xco_distribution}) )."858" Thus the temperature is primarily determined by molecular line cooling, and heating by cosmic rays and the grain photoelectric effect."," Thus the temperature is primarily determined by molecular line cooling, and heating by cosmic rays and the grain photoelectric effect."859" The kinetic temperature helps to set the brightness temperature, though the two are not identical."," The kinetic temperature helps to set the brightness temperature, though the two are not identical."860" The emission-weighted brightness temperature for the merger(disc) are ~ 50(7)K. Finally, the distribution of velocity dispersions in the GMCs is fairly narrow."," The emission-weighted brightness temperature for the merger(disc) are $\sim 50$ (7)K. Finally, the distribution of velocity dispersions in the GMCs is fairly narrow."861" Recalling 2, the velocity dispersion of the clouds is taken by calculating the dispersion amongst the cell's nearest neighbours, with a subgrid model for unresolved clouds (Equation 7))."," Recalling \ref{section:methods}, the velocity dispersion of the clouds is taken by calculating the dispersion amongst the cell's nearest neighbours, with a subgrid model for unresolved clouds (Equation \ref{equation:sigma}) )."862" Because the disc is dynamically cold, the velocity dispersions are primarily set by the latter case."," Because the disc is dynamically cold, the velocity dispersions are primarily set by the latter case."863" This results in an emission-weighted velocity dispersion within GMCs in the model disc of ~31,, with a maximum of ~15 !."," This results in an emission-weighted velocity dispersion within GMCs in the model disc of $\sim 3$, with a maximum of $\sim15$ ."864". These values compare favourably with the velocity dispersions reported in the comprehensive survey of(1987),, and the more recent review by(2007)."," These values compare favourably with the velocity dispersions reported in the comprehensive survey of, and the more recent review by."865". We can ask why the simulated ffrom the model galaxy is comparable to the Galactic average, Xco&2—-—4xX10?9cm""?/K-kms! In principle this occurs because the physical conditions in the model. GMCs by and large match those of observed GMCs in the Milky Way."," We can ask why the simulated from the model galaxy is comparable to the Galactic average, $\xco \approx 2-4 \times 10^{20}866\xcounits.$ In principle this occurs because the physical conditions in the model GMCs by and large match those of observed GMCs in the Milky Way."867" In this sense, the fact that our model value for iin quiescent discs matches that of the Galaxy is by construction."," In this sense, the fact that our model value for in quiescent discs matches that of the Galaxy is by construction."868" However, there are two salient points here."," However, there are two salient points here."869" First, it is important to remember that we allow for the possibility that the galactic environment can set the physical conditions in the GMCs if the pressure is sufficiently high."," First, it is important to remember that we allow for the possibility that the galactic environment can set the physical conditions in the GMCs if the pressure is sufficiently high."870" The fact that the default value for the surface density and velocity dispersions in the clouds is typically used is a statement that the galactic environment in the model disc galaxy is not sufficiently extreme to cause significant changes in the surface densities,"," The fact that the default value for the surface density and velocity dispersions in the clouds is typically used is a statement that the galactic environment in the model disc galaxy is not sufficiently extreme to cause significant changes in the surface densities,"871"üre (av) Alfvénn waves at frequencies ||<<O,. where V4=Bebrmnym)1/2 is the Alfvénn speed and wp=Jt).","are (a) Alfvénn waves at frequencies $|\omega_{\rm R}|<<\Omega_{\rm p}$, where $V_{\rm A}=B(4\pi m_{\rm p}n_{\rm b})^{-1/2}$ is the Alfvénn speed and $\omega_{\rm R}=\Re (\omega )$."872 The dispersion relation (??)) accounts for four types of Alfvénn waves: forward and backward moving. right-handed and left-handed polarised: (b) Whistler waves ccatfeequeneiesbetweent MSwp<Q).," The dispersion relation \ref{alfven}) ) accounts for four types of Alfvénn waves: forward and backward moving, right-handed and left-handed polarised; (b) Whistler waves at frequencies between $\Omega_{\rm e}<\omega_{\rm R}<-\Omega_{\rm p}$."873 The dispersion relation (??)) describes right-handed polarised waves (because Ως 0) that propagate forward for negative &<ϐ and backward for positive &=>0., The dispersion relation \ref{whist}) ) describes right-handed polarised waves (because $\Omega_{\rm e}<0$ ) that propagate forward for negative $k<0$ and backward for positive $k>0$.874 All of these are stable («(i) 0) if the beam particles are absent., All of these are stable $\psi =\Im (\omega )=0$ ) if the beam particles are absent.875 As is explained in Achatz et al. (1990)), As is explained in Achatz et al. \cite{asl90}) )876 the beam protons and electrons. which under the given weak-beam condition do not affect each other. are however able to trigger each its instability which occurs if the resonance conditions can be satisfied.," the beam protons and electrons, which under the given weak-beam condition do not affect each other, are however able to trigger each its instability which occurs if the resonance conditions }=-kV can be satisfied."877 The time-dependent behaviour of the intensities ή:f) of the excited waves is given by (Lerche 1967.. Lee Ip 19875) vy1111 ccwhere thet growth rate «is where No=eckfwg is the index of refraction and further E;=Vl|ae whith pSQihe.," The time-dependent behaviour of the intensities $I(k,t)$ of the excited waves is given by (Lerche \cite{ler67}, Lee Ip \cite{li87}) ) , where the growth rate $\psi $ is where $N=ck/\omega_{\rm R}$ is the index of refraction and further $E_i=\sqrt{1+x_i^2}$ whith $x_i=\Omega _{\rm i,0}/kc$."878 To describe the influence of these excited waves on the beam particles we use the quasilinear Fokker-Planck equation (e.g. Schlicketser 1989)) for the resonant wave-particle interaction., To describe the influence of these excited waves on the beam particles we use the quasilinear Fokker-Planck equation (e.g. Schlickeiser \cite{sch89}) ) for the resonant wave-particle interaction.879 For Alfvénn waves and for Whistler waves the index of refraction Va“is large compared to unity.. so that the Lorentz force: associated with the magnetic field of the waves is much larger than the force associated with the electric field. so that on the shortest time scale these waves scatter the particles in pitch angle j/ but conserve their energy. re. they isotropise the beam particles.," For Alfvénn waves and for Whistler waves the index of refraction is large compared to unity, so that the Lorentz force associated with the magnetic field of the waves is much larger than the force associated with the electric field, so that on the shortest time scale these waves scatter the particles in pitch angle $\mu $ but conserve their energy, i.e. they isotropise the beam particles."880" The Fokker-Planck equation for the phase space density then reads ---cewhere EU the pitch angleVI Fokker-Planck coefficient is determined by the wave intensities 7, ccUsing the dispersion relation of Alfvénn waves (cpcVA). Eq. (185) "," The Fokker-Planck equation for the phase space density then reads =, where the pitch angle Fokker-Planck coefficient is determined by the wave intensities $I_{\rm n}$ Using the dispersion relation of Alfvénn waves $\omega_{\rm R}\simeq V_{\rm A}k$ ), Eq. \ref{Dmumu}) )"881indicates that bean protons and electrons resonate with waves at wavenumbers given by the inverse of their Larmor radii times yn kcQO;Deu(μι).," indicates that beam protons and electrons resonate with waves at wavenumbers given by the inverse of their Larmor radii times $\mu $, $k\simeq \Omega _i/\Gamma v\mu =(\mu R_i)^{-1}$."882 For protons these wavenumbers correspond to Alfvénn waves., For protons these wavenumbers correspond to Alfvénn waves.883 For electrons these wavenumbers correspond to Alfvénn waves if the bulk Lorentz factor is above μηΠο)=1836]¢). and Whistler waves for smaller Lorentz factors.," For electrons these wavenumbers correspond to Alfvénn waves if the bulk Lorentz factor is above $|\mu |(m_{\rm p}/m_{\rm e})=1836|\mu |$, and Whistler waves for smaller Lorentz factors."884 We concentrate here on the isotropisation by Alfvénn waves mainly for two reasons: (1) the bulk of the momentum of the inflowing interstellar particles is carried by the protons so that they are energetically more important than the electrons: (2) for Lorentz factors P»μμin; the isotropisation of electrons is also caused by scattering with Alfvénn waves: for smaller Lorentzfactors the mistake one makes 1n representing the Whistller dispersion relation still by the Alfvénn dispersion relation is relatively small., We concentrate here on the isotropisation by Alfvénn waves mainly for two reasons: (1) the bulk of the momentum of the inflowing interstellar particles is carried by the protons so that they are energetically more important than the electrons; (2) for Lorentz factors $\Gamma >|\mu | m_{\rm p}/m_{\rm e}$ the isotropisation of electrons is also caused by scattering with Alfvénn waves; for smaller Lorentzfactors the mistake one makes in representing the ler dispersion relation still by the Alfvénn dispersion relation is relatively small.885 For Alfvénn waves (??)) we obtain /VA..=20 which is positive for forward (hk> 0) moving waves and negative for backward (4< 0)movingwaves., For Alfvénn waves \ref{alfven}) ) we obtain ^2= which is positive for forward $k>0$ ) moving waves and negative for backward $k<0$ )movingwaves.886 According to Eq. (14)), According to Eq. \ref{psi}) )887 we obtain for the growth rate of forward (+) and backward (-) moving Alfvénn waves, we obtain for the growth rate of forward (+) and backward (-) moving Alfvénn waves with888Asteroseismology of solar type stars is a powerful tool which can lead to the precise determination of the stellar parameters when associated with spectroscopic observations.,Asteroseismology of solar type stars is a powerful tool which can lead to the precise determination of the stellar parameters when associated with spectroscopic observations.889 Special tools have been developped for this purpose. which help constraining the internal structure of the stars from the observable acoustic frequencies.," Special tools have been developped for this purpose, which help constraining the internal structure of the stars from the observable acoustic frequencies."890 In previous papers (e.g. Vauelair et al. 2008)).," In previous papers (e.g. Vauclair et al. \cite{vauclair08}) ),"891 we developped a systematic way of comparing models with observations., we developped a systematic way of comparing models with observations.892 First. for each set of chemical composition [Fe/H] and Y. we compute evolutionary tracks and for each mass the model which best fits the observed large separation is derived.," First, for each set of chemical composition [Fe/H] and Y, we compute evolutionary tracks and for each mass the model which best fits the observed large separation is derived."893 Then we keep only the model that also best fits the observed echelle diagram., Then we keep only the model that also best fits the observed echelle diagram.894 We finally have a set of models correctly fitting the seismic data. with different abundances.," We finally have a set of models correctly fitting the seismic data, with different abundances."895 Interestingly enough. all these models have the same mass and radius. and of course the same log g (see Vauclair et al.," Interestingly enough, all these models have the same mass and radius, and of course the same log g (see Vauclair et al."896 for the star « Hor)., \cite{vauclair08} for the star $\iota$ Hor).897 All these models are placed in a log g - log diagram. together with the observed spectroscopic boxes.," All these models are placed in a log $g$ - log diagram, together with the observed spectroscopic boxes."898 Finally we only keep the models which satisfy all the constraints. including seismology and spectroscopy.," Finally we only keep the models which satisfy all the constraints, including seismology and spectroscopy."899 In this way. precise values of the stellar parameters are found.," In this way, precise values of the stellar parameters are found."900 Here we test this method on the exoplanet-host star Jj Arae (HD 160691)., Here we test this method on the exoplanet-host star $\mu$ Arae (HD 160691).901 This G3 IV-V type star is at the centre of a four- system: The star's overmetallicity has been established by many groups of observers: Bensby et al. (2003)).," This G3 IV-V type star is at the centre of a four-planets system: The star's overmetallicity has been established by many groups of observers: Bensby et al. \cite{bensby03}) ),"902 Laws et al. (2003)).," Laws et al. \cite{laws03}) ),"903 Santos et al. (2004a)).," Santos et al. \cite{santos04a}) ),"904 Santos et al. (2004b)), Santos et al. \cite{santos04b}) )905 and Fischer Valenti (2005)) (see Table 1))., and Fischer Valenti \cite{fischer05}) ) (see Table \ref{tab1}) ).906 µ Arae was observed with the HARPS spectrometer at La Silla Observatory during eight nights in June 2004 to obtair radial velocity time series., $\mu$ Arae was observed with the HARPS spectrometer at La Silla Observatory during eight nights in June 2004 to obtain radial velocity time series.907 The analysis of these data led to the discovery of up to 43 frequencies that could be identified with p-modes of degrees £=0 to 3 (Bouchy et al. 2005))., The analysis of these data led to the discovery of up to 43 frequencies that could be identified with p-modes of degrees $\ell=0$ to 3 (Bouchy et al. \cite{bouchy05}) ).908 A detailec modelling was given by Bazot et al. (2005))., A detailed modelling was given by Bazot et al. \cite{bazot05}) ).909 They computed models with two different assumptions that could. expla the observed overmetallicity: overabundance of metals in the original interstellar cloud or accretion of planetary material onto the star., They computed models with two different assumptions that could explain the observed overmetallicity: overabundance of metals in the original interstellar cloud or accretion of planetary material onto the star.910 They tried to obtain evidence of the origi of i. Arae's overmetallicity., They tried to obtain evidence of the origin of $\mu$ Arae's overmetallicity.911 The results were not conclusive in that respect. as the differences were not large enough to decide between the two scenari.," The results were not conclusive in that respect, as the differences were not large enough to decide between the two scenarii."912 Later on. other evidences were obtained that the observed overmetallicity in. exoplanet host stars must be original. not due to accretion (see Castro et al. 2009)).," Later on, other evidences were obtained that the observed overmetallicity in exoplanet host stars must be original, not due to accretion (see Castro et al. \cite{castro09}) )."913 We computed new models. testing various values of the original metallicity and helium abundance.," We computed new models, testing various values of the original metallicity and helium abundance."914 We compared the parameters of these models with those obtained from spectroscopy. anc introduced the new value of the Hipparcos parallax. as given by van Leeuwen (2007)).," We compared the parameters of these models with those obtained from spectroscopy, and introduced the new value of the Hipparcos parallax, as given by van Leeuwen \cite{leeuwen07}) )."915 We also analysed models with overshooting at the edge of the stellar core., We also analysed models with overshooting at the edge of the stellar core.916 In some of these models. negative small separations appear so that there is a crossing point in the echelle diagram for the lines {= 0-¢=2.," In some of these models, negative small separations appear so that there is a crossing point in the echelle diagram for the lines $\ell=0$ - $\ell=2$."917 This effect was specially discussed in Soriano Vauclair (2008)). who showed that all solar type stars go through a stage where the small separations become negative in the observable range of frequencies.," This effect was specially discussed in Soriano Vauclair \cite{soriano08}) ), who showed that all solar type stars go through a stage where the small separations become negative in the observable range of frequencies."918 Here. we present a direct application of this theoretical effect. which is used to constrain core overshooting.," Here, we present a direct application of this theoretical effect, which is used to constrain core overshooting."919is solved with an explicit scheme. using a seconc-order finite dillerence approximation for the clissipative Duxes.,"is solved with an explicit scheme, using a second-order finite difference approximation for the dissipative fluxes."920" A star of mass Al,=0.8M. and radius 2,=2h. is located at the origin of the 3D spherical coordinate system (1.0.0). with the rotation axis coincident with the normal to the disk midplane."," A star of mass $M_* = 0.8 M_{\odot}$ and radius $R_* = 2 R_{\odot}$ is located at the origin of the 3D spherical coordinate system $(R,\theta,\phi)$, with the rotation axis coincident with the normal to the disk midplane."921 The rotation period of the star is assumed to be 9.2 days., The rotation period of the star is assumed to be 9.2 days.922 The initial unperturbed stellar atmosphere is approximately in equilibrium and consists of three components: the stellar magnetosphere. the extended stellar corona. and the quasi-Ixeplerian disk.," The initial unperturbed stellar atmosphere is approximately in equilibrium and consists of three components: the stellar magnetosphere, the extended stellar corona, and the quasi-Keplerian disk."923 The pre-Hare magnetosphere is assumed to be force-frec. with dipole topology and magnetic moment fry aligned with the rotation axis of the star.," The pre-flare magnetosphere is assumed to be force-free, with dipole topology and magnetic moment $\mu\rs{B}$ aligned with the rotation axis of the star."924 Thus the two field components in spherical coordinates are The magnetic moment is chosen in order to have a magnetic field strength of the order of 1 KG. at the stellar surface according to observations (7))., Thus the two field components in spherical coordinates are The magnetic moment is chosen in order to have a magnetic field strength of the order of 1 kG at the stellar surface according to observations \citealt{1999ApJ...510L..41J}) ).925 The initial corona and disk are set in order to satisfy mechanical equilibrium involving centrifugal. eravitational. and pressure eracicnt forces.," The initial corona and disk are set in order to satisfy mechanical equilibrium involving centrifugal, gravitational, and pressure gradient forces."926 In. particular. we adopted the initial conditions introduced. by 7/— (where the reader. is referred. to for. more details) and. describing the star-disk system in quiescent configuration.," In particular, we adopted the initial conditions introduced by \cite{2002ApJ...578..420R} (where the reader is referred to for more details) and describing the star-disk system in quiescent configuration."927" These conditions assume that initially the plasma is barotropic and that the disk and the corona are both isothermal with temperatures Zi, and 19. respectively."," These conditions assume that initially the plasma is barotropic and that the disk and the corona are both isothermal with temperatures $T\rs{d}$ and $T\rs{c}$, respectively."928" With these assumptions. the thermal pressure at any point of the spatial domain is given by where 2)22pkp/(pmg) is the initial pressureat the boundary between the disk aid the corona. pe is the mass density of the corona close to the disk truncation radius £24. ky is the Boltzmann constant. JF is the function kids a constant of the order o£1 that takes into account that the clisk is slightly non-Ixeplerian (we set &=1.01 as in 2)). 4, is the centrifugal potential written as and r=sind is the cevlindrical radius."," With these assumptions, the thermal pressure at any point of the spatial domain is given by where $P_0 = 2\rho\rs{c}k\rs{B}T\rs{c}/(\mu m\rs{H})$ is the initial pressureat the boundary between the disk and the corona, $\rho\rs{c}$ is the mass density of the corona close to the disk truncation radius $R\rs{d}$, $k\rs{B}$ is the Boltzmann constant, ${\cal F}$ is the function $k$ is a constant of the order of 1 that takes into account that the disk is slightly non-Keplerian (we set $k = 1.01$ as in \citealt{2002ApJ...578..420R}) ), $\Phi\rs{c}$ is the centrifugal potential written as and $r = R\sin\theta$ is the cylindrical radius."929 The mass density depends on the pressure given. in Eq., The mass density depends on the pressure given in Eq.930 19 at any point of the spatial domain The angular velocity of the plasma is given hy 1n our simulation. we assume the isothermal cisk to be cold (Z4=810 IX). dense (n4 ranges between 5.lot and 410712 em3 7) and to rotate with: angular velocity close to the Ixeplerian value Og.," \ref{ini_pres} at any point of the spatial domain The angular velocity of the plasma is given by In our simulation, we assume the isothermal disk to be cold $T\rs{d}931= 8\times 10^3$ K), dense $n\rs{d}$ ranges between $5\times 10^{11}$ and $4\times 10^{12}$ $^{-3}$ ), and to rotate with angular velocity close to the Keplerian value $\Omega\rs{K}$."932 The rotation axis of the disk (coincident with the rotation axis of the star) is aligned with the magnetic moment fry., The rotation axis of the disk (coincident with the rotation axis of the star) is aligned with the magnetic moment $\mu\rs{B}$.933 The clisk is initially truncated. by the stellar magnetosphere at the radius. / where the total gas pressure of the disk equals the magnetic pressure. p|pi?=Dfss: for the parameters. adopted here. Ay=2.86HR..," The disk is initially truncated by the stellar magnetosphere at the radius $R\rs{d}$ where the total gas pressure of the disk equals the magnetic pressure, $p+\rho u^2 = B^2/8\pi$; for the parameters adopted here, $R\rs{d} = 2.86\, R_*$."934" lo our simulation. the corotation radius Hao,=(CAL,/OyA924. where Ὃν ds the stellar angular velocity."," In our simulation, the corotation radius $R\rs{co} = (GM_*/\Omega_*^2)^{1/3}=9.2\, R_*$, where $\Omega_*$ is the stellar angular velocity."935" The corona is initially isothermal with ⊳⇁↙⋮∶≟↳∖∐∖⋜⋯∠⇂⋜∐↓∪∖∖⊽∠⇂∢⊾⊔⊳∖↓↿∙∖⇁∖∖⊽∐↓↕⊔⋅↓⋅⋜⋯⋏∙≟↓⊔⋏∙≟∣⋡⋖⋅∣∖∖⊽⋖⋅∢⊾⊔ ⊳ ⋠⊥⋠ ⋠ zc10°h and 10""(", The corona is initially isothermal with $T\rs{c} = 4$ MK and at low with $n\rs{c}$ ranging between $\approx 10^8$ and $10^9$ $^{-3}$.9360 E As shown in Eq. ld.," As shown in Eq. \ref{ini_omega},"937 we allow 1ο corona to be initially rotating with angular. velocity qual to the Ixeplerian rotation rate of the disk in order to jwe approximately equilibrium conditions and. reduce the ποσις of transients caused by the initial cdillerential rotation »etween the disk aid the corona (C?77?))," we allow the corona to be initially rotating with angular velocity equal to the Keplerian rotation rate of the disk in order to have approximately equilibrium conditions and reduce the effects of transients caused by the initial differential rotation between the disk and the corona \citealt{2002ApJ...578..420R,9382009A&A...508.1117Z}) )."939 Figure 1. shows the initial condition. together with rw numerical οἱ adopted. in| our simulation., Figure \ref{fig1_sup} shows the initial condition together with the numerical grid adopted in our simulation.940" Phe 'omputational domain extends between Ruin=ds (ie. 1e inner boundary coincides with the stellar surface) and Duas=14h, in the radial direction. and encompasses an angular sector goingH [rom. £u, Ὁ.ü to Cs=175""πρι inον the angular coordinate 6. and from óui,=0"" to Ons0=3603n in5 the angular coordinate ©."," The computational domain extends between $R\rs{min} = R_*$ (i.e. the inner boundary coincides with the stellar surface) and $R\rs{max} = 14\, R_*$ in the radial direction, and encompasses an angular sector going from $\theta\rs{min} = 5^0$ to $\theta\rs{max} = 175^0$ in the angular coordinate $\theta$, and from $\phi\rs{min} = 0^0$ to $\phi\rs{max} = 360^0$ in the angular coordinate $\phi$."941 Vhe inner and outer boundaries in € do not coincide with the rotation axis of the star-clisk system to avoid extremely small dO values. vastly increasing the computational cost.," The inner and outer boundaries in $\theta$ do not coincide with the rotation axis of the star-disk system to avoid extremely small $\delta \phi$ values, vastly increasing the computational cost."942 On the other hand. all the evolution relevant for this study never involve portions of the domain close to the star-disk rotation axis.," On the other hand, all the evolution relevant for this study never involve portions of the domain close to the star-disk rotation axis."943 The racial coordinate 2 has been discretized on a logarithmic grid. with the mesh size increasing with 7? (sec Fig. 1)).," The radial coordinate $R$ has been discretized on a logarithmic grid with the mesh size increasing with $R$ (see Fig. \ref{fig1_sup}) ),"944 giving a higher spatial resolution closer to the star as it is appropriate for simulations of accretion Lows to a star with a dipole field (2))., giving a higher spatial resolution closer to the star as it is appropriate for simulations of accretion flows to a star with a dipole field \citealt{2002ApJ...578..420R}) ).945" The radial grid is macle of Ny=50 points with a maximum resolution of AR=48«10"" em close to the star and a minimum resolution of AR=64102°t em close to the outer boundary."," The radial grid is made of $N\rs{R}946= 80$ points with a maximum resolution of $\Delta R = 4.8\times 10^9$ cm close to the star and a minimum resolution of $\Delta R = 6.4\times94710^{10}$ cm close to the outer boundary."948" ENThe angular coordinate € has been discretized. uniformly. with No=90 points. giving a resolution of A@=2""."," The angular coordinate $\theta$ has been discretized uniformly with $N\rs{\theta} = 90$ points, giving a resolution of $\Delta \theta = 2^o$."949" The angular coordinate © is nonuniform with the highest. resolution in an angular sector of 180"" placed where the faring loop and the stream evolve (see bottom panel in Fie. 13).", The angular coordinate $\phi$ is nonuniform with the highest resolution in an angular sector of $180^0$ placed where the flaring loop and the stream evolve (see bottom panel in Fig. \ref{fig1_sup}) ).950" The ó-grid is made of Αιξ110 points with a maximum resolution of Ao=2"" and a minimum resolution of 9"".", The $\phi$ -grid is made of $N\rs{\phi} = 110$ points with a maximum resolution of $\Delta \phi = 2^o$ and a minimum resolution of $9^0$ .951" ""Phe numerical erid is not static but tracks the hot loop and the stream as the calculation progresses. in such a way that the loop and the stream evolve in the portion of the domain with the highest spatial resolution (namely that with Ad= 2"")."," The numerical grid is not static but tracks the hot loop and the stream as the calculation progresses, in such a way that the loop and the stream evolve in the portion of the domain with the highest spatial resolution (namely that with $\Delta \phi = 2^o$ )."952 The boundary conditions at the stellar surface ain amount to assuming that the infalling material passes, The boundary conditions at the stellar surface $R\rs{min}$ amount to assuming that the infalling material passes953" 106AL, 2~2050,"," \citep{bl01,bl04a,bromm09}."954 1997).. (Yoshidaetal.2008).. (MacLow&Ilessen2001:MelIl&eeOstriker2007:Yorke2007).," $10^6\,M_\odot$ $z\sim 20-50$ \citep{htl96,tegmark97}. \citep{yoh08}. \citep{mk04,mo07,zy07}."955" With a few exceptions (CDukettal.al.200€ 2009).. simulations⋅lat.: oof thisus Initialtial collapse plisplase have shown no fragmentation (Abeletal.2002:BrounO'Shea&Novinan2007:Yoshidaetal. 2008).. leacing to the conclusion that the first stars formed im isolation and were extremely massive,"," With a few exceptions \citep{tao09}, , simulations of this initial collapse phase have shown no fragmentation \citep{abn02,bcl02,bl04b,yoshida06b,on07,yoh08}, leading to the conclusion that the first stars formed in isolation and were extremely massive."956 Iu contrast. studies of prescut-day star formation have eoncrally found fragmentation to occur shortly after the formation of the first protostar (IxXlessen&Burlk-al. 2010).," In contrast, studies of present-day star formation have generally found fragmentation to occur shortly after the formation of the first protostar \citep{kb00,bbb03,krumholz09,peters10}."957. Following up ou this result. simak particles were recently used in studies of the fragmentation of prunuordial gas in nünihalos (Clarketal.2008:Stacyetal.2010:Clark 2011a).," Following up on this result, sink particles were recently used in studies of the fragmentation of primordial gas in minihalos \citep{cgk08,sgb10,clark11a}."958. Thev found that the ietal-free eas clouds fragment strouglv. with the details of the process depending ou the degree of turbulence in the halo.," They found that the metal-free gas clouds fragment strongly, with the details of the process depending on the degree of turbulence in the halo."959 Focusing on the dynamical evolution of the hnieh-deusitv gas in the central regions of a nmünuihalo. Clarketal.(2011b) demoustrated that the protostellar disks around primordial stars accrete frou the iufalliug euvelope faster than they can trauster thoi mass outo the central object.," Focusing on the dynamical evolution of the high-density gas in the central regions of a minihalo, \citet{clark11b} demonstrated that the protostellar disks around primordial stars accrete from the infalling envelope faster than they can transfer their mass onto the central object."960: As a result. they rapidly: become unstable∙ and fragiient to⋅ buildaip binary. ironomisches or hieher-order multiple stella systems.," As a result, they rapidly become gravitationally unstable and fragment to build-up binary or higher-order multiple stellar systems."961 These results challeuge the idea that the first stars formed in isolation and eive rise to a number of new questions: What is the mass spectrum of DpPop III stars iu. groups?, These results challenge the idea that the first stars formed in isolation and give rise to a number of new questions: What is the mass spectrum of Pop III stars in groups?962""" [sit. differentfrom thepreviously preferred siugle mode of", Isit differentfrom thepreviously preferred single mode of963Some X-ray pulsars are kuown to have remarkable siilari vin their properties whicli are differeut frou other binary or isolated X-ray. pulsars (Mereshetti&audStella 1995)).,Some X-ray pulsars are known to have remarkable similarity in their properties which are different from other binary or isolated X-ray pulsars \cite{mere95b}) ).964 T1e properties conunon to most of these objects are a) pulse period iia small rauge of 5 12s.) monotonous spin down with P/P in the rauge of 5«1BOp.107 s ο) identical X-ray προςτι consstiie of steep power-law (PEE—25 1) aud black body coumponeut (KT —0.5 keV). d) stable N-rav 1iuuostyv (101y CreseS s 1j for vears. ο) fait or 1audentifed optical counterpart. and f) no evidence of orbital motion.," The properties common to most of these objects are a) pulse period in a small range of 5--12 s, b) monotonous spin down with ${\rm P} / \dot {\rm P}$ in the range of $5\times10^{11} - 1.3\times10^{13}$ s c) identical X-ray spectrum consisting of steep power-law $\Gamma = 3-4$ ) and black body component (kT $\sim 0.5$ keV), d) stable X-ray luminosity $^{34}-10^{36}$ ergs $^{-1}$ ) for years, e) faint or unidentified optical counterpart, and f) no evidence of orbital motion."965" The sources also have a galactic distribution. most of these are within |),<0.5* ivd all are probably young (~104 vr) because of their association with SNR or molecular clouds."," The sources also have a galactic distribution, most of these are within $|{b}|\leq 0.5^\circ$ and all are probably young $\sim 10^4$ yr) because of their association with SNR or molecular clouds."966 The objects in which all the properties meifioned. above have been observed areGl.. LE 2259|586. 5937.. IRNS 10 0axl 1E 015 (Ies 73).," The objects in which all the properties mentioned above have been observed are, 1E 2259+586, , 1RXS $-$ 400910 and 1E $-$ 045 (Kes 73)."967 Two more objects. AX 0300 etal. 1998)) aud RN 3125 (Ilaberlet 1996)). also probably Delong to the sade class but to establish their ANP cuxliacy. mnore X-ray observations are required to nieastY their pulse period variations. search for possibe pulse arrival tiue delay and investigate the flux staülity.," Two more objects, AX $-$ 0300 \cite{tori98}) ) and RX $-$ 3125 \cite{habe96}) ), also probably belong to the same class but to establish their AXP candidacy, more X-ray observations are required to measure their pulse period variations, search for possible pulse arrival time delay and investigate the flux stability."968 Classificatiou of au objec ax ANP ouly from some pro)erties sinilar to the ayove is not very firm., Classification of an object as AXP only from some properties similar to the above is not very firm.969 IU GT. probably a Dinary system. showed bot| spiu-u» and spin-cown (Chakrabartyetal. 1997)) andaIso has an optically brigh accretion disk (Mkdlecitehetal. 1981)).," 4U $-$ 67, probably a binary system, showed both spin-up and spin-down \cite{chak97}) ) and also has an optically bright accretion disk \cite{midd81}) )."970 Anoticr one object RA 0301 (Sclaventier 19911). does not have stable intensity aud pulsatious are aso not always detectable (Songetal. 1999)).," Another one object RX $-$ 0301 \cite{schw94}) ), does not have stable intensity and pulsations are also not always detectable \cite{song99}) )."971 Therefore. these two objects are 1ιο ANPs.," Therefore, these two objects are not AXPs."972 Considering the stroug similarity. betweei these hiucful of sources. it has been proposed that they have same plysical nature aud different. scenarios have been proposed to explain the observed properties.," Considering the strong similarity between these handful of sources, it has been proposed that they have same physical nature and different scenarios have been proposed to explain the observed properties."973 The xonminent models are a) accretion from low mass binary companion (Mereghoetti&audStella 1995)). b) suele neutron star accreting frou molecular cloud. or a product of common envelope evolution Zxvtkov object) of close high mass X-ray binaries iu which a solitary neutron star accretes matter from a fossil disk (wan Paradijs. Tau. van den Ieuvel 1995:: Ghosh. Aneclini. White 1997)). aud ο) extremely hieh magnetic Held neutron star radiating X-rays due to maeuetic field decay (Thompson&Duucanu 1996)).," The prominent models are a) accretion from low mass binary companion \cite{mere95b}) ), b) single neutron star accreting from molecular cloud, or a product of common envelope evolution $\rm \dot{Z}$ ytkov object) of close high mass X-ray binaries in which a solitary neutron star accretes matter from a fossil disk (van Paradijs, Taam, van den Heuvel \cite{vanp95a}; Ghosh, Angelini, White \cite{ghos97}) ), and c) extremely high magnetic field neutron star radiating X-rays due to magnetic field decay \cite{thom96}) )."974 Unlike the radio pulsars aud rotationally powered X-rav pulsars. iu the ANPs. the spin-down rate is not arge enough to power the observed. X-ray enission.," Unlike the radio pulsars and rotationally powered X-ray pulsars, in the AXPs, the spin-down rate is not large enough to power the observed X-ray emission."975 Auxmg the 90 or so known X-rav piIsars (Nagase1999)). direct evidence of binary natuvas known for nore than 35 sources (vauParadijs1995)).," Among the 90 or so known X-ray pulsars \cite{naga99}) ), direct evidence of binary nature is known for more than 35 sources \cite{vanp95b}) )."976 Iucludiug he inclizect evidences this uuuber ca o upto about 65 ae 7 pulsars are isolated stars m SNR aud are »owered by rotational energy losses., Including the indirect evidences this number can be upto about 65 and 7 pulsars are isolated stars in SNR and are powered by rotational energy losses.977 Iu the rest of the oulsars. m which uo binary signature is known. it is often due to lack of sufficient observation.," In the rest of the pulsars, in which no binary signature is known, it is often due to lack of sufficient observation."978 However. he 7 objects which are either ANPs or candidate ANPs (or 10 if we include he 3 Soft. Ciunnia-rav Repeaters in which pt]satious have beendetected). no ΗΝ signature has been fouud iu spite of extensive searches.," However, the 7 objects which are either AXPs or candidate AXPs (or 10 if we include the 3 Soft Gamma-ray Repeaters in which pulsations have been detected), no binary signature has been found in spite of extensive searches."979 The strong upper limit ou pulse arrival time delay that has been otained iu some of these sources strongly sugeests nobinary nature for the AND., The strong upper limit on pulse arrival time delay that has been obtained in some of these sources strongly suggests non-binary nature for the AXPs.980 Iu addition. the spin change behaviour of the AXNPs is also remarkably different from accreting pulsars (Bildstenetal. L997)).," In addition, the spin change behaviour of the AXPs is also remarkably different from accreting pulsars \cite{bild97}) )."981 Tu almost al the accreting N-rav pulsars. both spiu-up and spin-down episodes ie been secu which may be randomly distributed (ii persistent sources) or spiu-dowus in quiescence ollowed bv rapid spiuecups during bright trausieut phases Gu trausicut pulsars) or long monotonic spiu-up and spin-down episodes accompanied by spectral and Iuuinositv changes (e.g. LU 1626. 67. Y1&Vish-iac 19993).," In almost all the accreting X-ray pulsars, both spin-up and spin-down episodes have been seen which may be randomly distributed (in persistent sources) or spin-downs in quiescence followed by rapid spin-ups during bright transient phases (in transient pulsars) or long monotonic spin-up and spin-down episodes accompanied by spectral and luminosity changes (e.g. 4U $-$ 67, \cite{yiiv99}) )."982 The ANPs are i Linauy respect also similar to the N-ray counterparts of the Soft Camuna-ray Repeaters (SCR)., The AXPs are in many respect also similar to the X-ray counterparts of the Soft Gamma-ray Repeaters (SGR).983" The Nav spectral aud timing properties of these two type of ojects have strong similarities. half of the ANPs iud SGRs are associated with supernova τολμάές,"," The X-ray spectral and timing properties of these two type of objects have strong similarities, half of the AXPs and SGRs are associated with supernova remnants."984 This has lead to the suggestion that the ANPs are also maguetars in which the X-ray cussion is due to maguetic field decay. (Thompson&Duncau 1996)., This has lead to the suggestion that the AXPs are also magnetars in which the X-ray emission is due to magnetic field decay \cite{thom96}) ).985 The main difference between these two type of objects is the non detection of SCR bursts from the ANPs., The main difference between these two type of objects is the non detection of SGR bursts from the AXPs.986" However. considering the rarity of the SCR activity among the established SCRs (Ixouveliototetal. 1996)). the absence of bursts frou, ANPS is not a serious issuc."," However, considering the rarity of the SGR activity among the established SGRs \cite{kouv96}) ), the absence of bursts from AXPs is not a serious issue."987 From a relatively voung age of the ANP. IE 015 in the supernova renuuui Ives 73. Gotthelf. Vasisht. Dotani (1999)) propose that inthe evolutionary track. the ANPs are au carly quiescent state of the SCRs.," From a relatively young age of the AXP, 1E $-$ 045 in the supernova remnant Kes 73, Gotthelf, Vasisht, Dotani \cite{gott99}) ) proposed that inthe evolutionary track, the AXPs are an early quiescent state of the SGRs."988 Stability of the X- endssion properties (spin-down rate. ΠΠ PAectral shape aud pulse shape aud fraction) is usually mentioncd as one important aspect of the ANXPs," Stability of the X-ray emission properties (spin-down rate, luminosity, spectral shape and pulse shape and fraction) is usually mentioned as one important aspect of the AXPs"989covered by our photometric survey is 2.0 arcmin?.,covered by our photometric survey is 2.9 $^2$.990 Compared to the combined areas of the survevs by Wilkingetal.(2001) and Aspinetal.(1994).. which surveved both the north and south components of the cluster and covered 79 arcmin? and 81 arcmin? respectively. our photometric survey. thus covers z of the area of NGC 1333.," Compared to the combined areas of the surveys by \citet{wil03} and \citet{as94}, which surveyed both the north and south components of the cluster and covered 79 $^2$ and 81 $^2$ respectively, our photometric survey thus covers $\approx$ of the area of NGC 1333."991 Total integration times for the fields with 3 roll angles. including dithers and multiple roll angles were 383 s [or F160W. 766 s lor F110W and 4608 s lor G141 per field.," Total integration times for the fields with 3 roll angles, including dithers and multiple roll angles were 383 s for F160W, 766 s for F110W and 4608 s for G141 per field."992 Data reduction was carried out. using a combination of IRAF and eustom IDL routines., Data reduction was carried out using a combination of IRAF and custom IDL routines.993 The methods described below closely follow those used in W.Liuetal.(2003)., The methods described below closely follow those used in \citet{liu03}.994. The images were dark and sky subtracted using combined dark plus skv Games created with the routine NICSKYDARI in the NICRED package for ΗΑΕ (AleLeod1997)., The images were dark and sky subtracted using combined dark plus sky frames created with the routine NICSKYDARK in the NICRED package for IRAF \citep{mc97}.995.. Cosmic rays and bad pixels were located and removed using the routine FULLEITDAM. by searching for discontinuities in the fIux in each pixel over time.," Cosmic rays and bad pixels were located and removed using the routine FULLFITBAM, by searching for discontinuities in the flux in each pixel over time."996" Alter clark and sky subtraction and cosmic ray reduction. there did not appear to be anv bias offsets or “pedestal effects"" between (he different quadrants of the NICMOS chip. which is a common problem experienced with NICMOS (e.g. W.Liuetal. (2003)))."," After dark and sky subtraction and cosmic ray reduction, there did not appear to be any bias offsets or “pedestal effects” between the different quadrants of the NICMOS chip, which is a common problem experienced with NICMOS (e.g. \citet{liu03}) )."997 This is because our fields were not verv crowded. ancl did. not show significant nebulositv. as (he pedestal is also dependent on the total charge of the quadrant.," This is because our fields were not very crowded and did not show significant nebulosity, as the pedestal is also dependent on the total charge of the quadrant."998 The only field that shows any nebulosity is $2., The only field that shows any nebulosity is S2.999 Finally. (he images were fIatfielded using the routine NICFLATTEN and the appropriate epoch on-orbit fLlatfields from the Space Telescope Science Center Institute (STSclI) website.," Finally, the images were flatfielded using the routine NICFLATTEN and the appropriate epoch on-orbit flatfields from the Space Telescope Science Center Institute (STScI) website."1000 The dithers were (hen combined using the IDL software IDP3 (Lytleetal.1999).. aligning the different frames by. their World Coordinate System as well as resampling the images by a [actor of 2 using a bi-cubic sinc {ο increase the resolution of the images.," The dithers were then combined using the IDL software IDP3 \citep{ly99}, aligning the different frames by their World Coordinate System as well as resampling the images by a factor of 2 using a bi-cubic sinc to increase the resolution of the images."1001 The three roll angles were combined using IDP3. first rotating the images to align in angle and then shifting the images to correct for offsets in x and v. For both of these routines. παν was conserved.," The three roll angles were combined using IDP3, first rotating the images to align in angle and then shifting the images to correct for offsets in x and y. For both of these routines, flux was conserved."1002 All sources presented in the Color-Magnitude Diagram (CMD) (See., All sources presented in the Color-Magnitude Diagram (CMD) (Sec.1003 3.3) were detected in both the FIGOW and FLIOW filters using the IRAF routine DAOFIND with a 10 6 detection threshold., 3.3) were detected in both the F160W and F110W filters using the IRAF routine DAOFIND with a 10 $\sigma$ detection threshold.1004 Even with the low amount of nebulosity every image still had. some spurious detections which were removed by visual inspection., Even with the low amount of nebulosity every image still had some spurious detections which were removed by visual inspection.1005 We took advantage of the three roll angles for each field to help identily false detections., We took advantage of the three roll angles for each field to help identify false detections.1006 If à source appeared in the sanie position lor each roll angle. it was assunied (o be genuine.," If a source appeared in the same position for each roll angle, it was assumed to be genuine."1007 False detections were located al (heir position on the chip for only one roll angle., False detections were located at their position on the chip for only one roll angle.1008 Since the fields were not crowded. the photometry was performed using the IRAE routine APPIIOT.," Since the fields were not crowded, the photometry was performed using the IRAF routine APPHOT."1009 The optimal aperture size was caleulated to be 8 pixels in radius with the, The optimal aperture size was calculated to be 8 pixels in radius with the1010After recombining at :~1000. the interealactic nedium (IGM) remains neutral until the formation of he first luminous objects.,"After recombining at $z \sim 1000$, the intergalactic medium (IGM) remains neutral until the formation of the first luminous objects."1011 Standard ACDM structure orluation models sugeest that the first objects with sufficient mass to lonize substantially the ICAL form at 215 (Cuediu Ostriker 1997)., Standard $\Lambda$ CDM structure formation models suggest that the first objects with sufficient mass to ionize substantially the IGM form at $z \le 15$ (Gnedin Ostriker 1997).1012 It is currently not shown whether these first luminous objects are active star formuue galaxies or accreting massive black holes. mut eiveu the current observations of both ealaxics and huninous AGN at i:>6 (In et al.," It is currently not known whether these first luminous objects are active star forming galaxies or accreting massive black holes, but given the current observations of both galaxies and luminous AGN at $z > 6$ (Hu et al."1013 2002: Fan et al., 2002; Fan et al.1014 2001). it seems Likely that a mixture of sources exists.," 2001), it seems likely that a mixture of sources exists."1015 Rejonization of the IGAL proceeds slowly at fist. with cach object essentially contained witlin its own Stromeren sphere.," Reionization of the IGM proceeds slowly at first, with each object essentially contained within its own Stromgren sphere."1016 Eveutually these spheres overlap such that ionizing photons from a given object contribute to unmiltiple spheres. resulting in a run-away process.," Eventually these spheres overlap such that ionizing photons from a given object contribute to multiple spheres, resulting in a run-away process."1017 This overlap phase is known as fast relonization. and in the discussion below we call this phase the “epoch of reionization (Loeb Barkana 2001: Cinediu 2000: 2002).," This overlap phase is known as 'fast' reionization, and in the discussion below we call this phase the 'epoch of reionization' (Loeb Barkana 2001; Gnedin 2000; 2002)."1018 The recent detection of Lya absorption bv the neutral IGM. in two :—6 QSOs. as predicted. by Cunu Petersou (1965). has revolutionized our uuderstaudiue of cosmic relonization bv placing coustraiuts ou the x»cli of reionization (Becker ct al.," The recent detection of $\alpha$ absorption by the neutral IGM in two $z \sim 6$ QSOs, as predicted by Gunn Peterson (1965), has revolutionized our understanding of cosmic reionization by placing constraints on the epoch of reionization (Becker et al."1019 2002: Djoreovski al., 2002; Djorgovski et al.1020 2002: Peutericci et al., 2002; Pentericci et al.1021 2002)., 2002).1022 Using these constraints in the context of ACDM structure formation models. (ποσα (2002) fixes the redshift of fast reionization at zo6.2. although adnüttedlv the possibility. of cosnüc variance makes this estimate uncertain (Ih et al.," Using these constraints in the context of $\Lambda$ CDM structure formation models, Gnedin (2002) fixes the redshift of fast reionization at $z \sim 6.2$, although admittedly the possibility of cosmic variance makes this estimate uncertain (Hu et al."1023 2002)., 2002).1024 The picture then is one of a egrav-age! between the epoch of formation of the first luminous objects auk the epoch of fast reionization., The picture then is one of a 'gray-age' between the epoch of formation of the first luminous objects and the epoch of fast reionization.1025 During the erav-age the TGAL is predominantly neutral. with pockets of 10nizec eas around the first hnuuinous objects," During the gray-age the IGM is predominantly neutral, with pockets of ionized gas around the first luminous objects."1026 The neutra IGAL is opaque at rest waveleugths shorter than Ένα such that study of objects. and the IGAL during this age will be limited to observatious at wavelengths longer than about ljan. One method for studying structures duriug this erav-age is through the IIT 21c1à line of neutral bydrogen.," The neutral IGM is opaque at rest wavelengths shorter than $\alpha$ such that study of objects, and the IGM, during this age will be limited to observations at wavelengths longer than about $\mu$ m. One method for studying structures during this gray-age is through the HI 21cm line of neutral hydrogen."1027 Mauy eroups have considered observations of TT 21cm cussion frou the neutral IGAL during the erayv-age. or absorption against the microwave backeround (cepeuding ou the III excitation temperature: Scott Rees 1990: Baela. Nath. Padmanabhan 1997: Tozzi ct al.," Many groups have considered observations of HI 21cm emission from the neutral IGM during the gray-age, or absorption against the microwave background (depending on the HI excitation temperature; Scott Rees 1990; Bagla, Nath, Padmanabhan 1997; Tozzi et al."1028 2000: Shaver et al., 2000; Shaver et al.1029 1999: They et al 2002: Bharachwaj Sethi 2002)., 1999; Iliev et al 2002; Bharadwaj Sethi 2002).1030 ILowever. even iu the most optimistic case of thermal noise lanited observations. the seusitivitv of large area future radio telescopes. such as the Square Kilometer Array (SIVA). πράος that such studies will be lmuited to large scale structure (MIT) >Lol ALL).," However, even in the most optimistic case of thermal noise limited observations, the sensitivity of large area future radio telescopes, such as the Square Kilometer Array (SKA), implies that such studies will be limited to large scale structure (M(HI) $> 10^{12}$ $_\odot$ )."1031 Tucliding systematic errors. such as source confusion (di Matteo et al.," Including systematic errors, such as source confusion (di Matteo et al."1032 2001). could deerade substantially this limit.," 2001), could degrade substantially this limit."1033 Iu this paper we consider stuclving the neutral ICAL during the erav-aee via III 21e1à absorption studies toward diserete radio sources., In this paper we consider studying the neutral IGM during the gray-age via HI 21cm absorption studies toward discrete radio sources.1034 The nuaportaut point is that while the ΙΝΕ is opaque to the Lya line. the weakness of the maeuetic lyperfiue transition mukes the IGAL transhicent to HIT 210 absorption.," The important point is that while the IGM is opaque to the $\alpha$ line, the weakness of the magnetic hyperfine transition makes the IGM translucent to HI 21cm absorption."1035 Coluun density scusitivity for absorption studies is set only by the surface brightucss of the background source. thereby allowing absorption studies to probe to orders-ofanaenitude lower masses than can be detected in emission.," Column density sensitivity for absorption studies is set only by the surface brightness of the background source, thereby allowing absorption studies to probe to orders-of-magnitude lower masses than can be detected in emission."1036 We use models of structure formation and cosmic relonization to predict the ITE 21cm optical depth (7) of the ICAL curing the erav-age., We use models of structure formation and cosmic reionization to predict the HI 21cm optical depth $\tau$ ) of the IGM during the gray-age.1037 We then address the questions: Ave future radio telescopes adequate to detect these signals?, We then address the questions: Are future radio telescopes adequate to detect these signals?1038 And are there high redshift radio sources of sufficient luminosity to permit these CXudies?, And are there high redshift radio sources of sufficient luminosity to permit these studies?1039 Iu a future paper we will consider iu more detail the physical iuformation about the rich structures in the pre-veionization ICAL that can be obtained through HI 21e absorption observations., In a future paper we will consider in more detail the physical information about the rich structures in the pre-reionization IGM that can be obtained through HI 21cm absorption observations.1040" We assume Πυ 10]1iuns | Mpe+t. Qa,=0.35. and Q4=0.65,"," We assume $H_0$ = 70 km $^{-1}$ $^{-1}$, $\Omega_M = 0.35$, and $\Omega_\Lambda = 0.65$."1041 Our analvsis relies ou the simulations of Cuediu (2000. 2002).," Our analysis relies on the simulations of Gnedin (2000, 2002)."1042 These simulations iuclude the three nain physical iueredieuts required to model neutra ivdrogeu absorption in the redshifted 21 cm line: Inhomogeneous small-scale structure of the universe. radiative transter. aud accurate treatinent of the leve »opulatious in atomic lbydroecu.," These simulations include the three main physical ingredients required to model neutral hydrogen absorption in the redshifted 21 cm line: inhomogeneous small-scale structure of the universe, radiative transfer, and accurate treatment of the level populations in atomic hydrogen."1043 The αναος of gas and dark matter is followed in a quasi-Lagraneian ashion wwith high resolution) with the Softee: Lagrangian Uvdrodvuamics (“SLI) code (Caedin 1995. (ιο Bertschinecr 1996).," The dynamics of gas and dark matter is followed in a quasi-Lagrangian fashion with high resolution) with the Softened Lagrangian Hydrodynamics (“SLH”) code (Gnedin 1995, Gnedin Bertschinger 1996)."1044 The radiative rauster is modeled with the newly developed Optically Thin Variable Eddington Tensor (OTVET) approximation (ποσα Abel 2001)., The radiative transfer is modeled with the newly developed Optically Thin Variable Eddington Tensor (OTVET) approximation (Gnedin Abel 2001).1045" Finally, we include all the effects that couple gas kinetic temperature to the spin temperature of the atomic hydrogen: Lya pumping aud collisions with electrous aud neutral atoms."," Finally, we include all the effects that couple gas kinetic temperature to the spin temperature of the atomic hydrogen: $\alpha$ pumping and collisions with electrons and neutral atoms."1046 The, The1047very unlikely.,very unlikely.1048 The infrared component is ~15 times brighter than our estimate of the stellar luminosity in Table 3.., The infrared component is $\sim$ 15 times brighter than our estimate of the stellar luminosity in Table \ref{evol_t}.1049 This would imply that the central star is significantly more luminous and that what we observe in the optical is the stellar spectrum severely attenuated by dust., This would imply that the central star is significantly more luminous and that what we observe in the optical is the stellar spectrum severely attenuated by dust.1050 Then the central star. reddened by dust. would be much hotter than estimated from the photometry.," Then the central star, reddened by dust, would be much hotter than estimated from the photometry."1051 Yet. as noted in Sect. 3..," Yet, as noted in Sect. \ref{spectr},"1052 our spectrum suggests an early M type star. thus consistent with the photometric results.," our spectrum suggests an early M type star, thus consistent with the photometric results."1053 Besides. the large present luminosity and the presence of hot nearby dust would inevitably indicate that V4332 Ser has recently experienced another outburst.," Besides, the large present luminosity and the presence of hot nearby dust would inevitably indicate that V4332 Sgr has recently experienced another outburst."1054 With the effective radius of the blackbody component estimated above (~400 Re) and adopting a rather lower limit of 10 kin/s for the expansion velocity. one finds that dust would have been lost a year before May-September 2003.," With the effective radius of the blackbody component estimated above $\sim$ $R_{\sun}$ ) and adopting a rather lower limit of 10 km/s for the expansion velocity, one finds that dust would have been lost a year before May-September 2003."1055 However. as discussed in Sect. 4.3..," However, as discussed in Sect. \ref{analys_fr},"1056 the infrared excess started developing in 1999., the infrared excess started developing in 1999.1057 Besides. there is no other observational indication that the object recently suffered from a secondary outburst.," Besides, there is no other observational indication that the object recently suffered from a secondary outburst."1058 On the contrary. as can be seen from Fig. 4..," On the contrary, as can be seen from Fig. \ref{evol_f},"1059 the object has been following the long term decline after the 1994 eruption., the object has been following the long term decline after the 1994 eruption.1060 The lack of any significant evolution in the optical magnitudes between May and September 2003 is also an evidence for the long time scale evolution of the object in 2003., The lack of any significant evolution in the optical magnitudes between May and September 2003 is also an evidence for the long time scale evolution of the object in 2003.1061 BVAO4 have recently considered that V4332 Ser might be a young object surrounded by a protostellar circumstellar disc., BVA04 have recently considered that V4332 Sgr might be a young object surrounded by a protostellar circumstellar disc.1062 This hypothesis has been invoked to explain the origin of the water-ice band observed in the infrared spectrum., This hypothesis has been invoked to explain the origin of the water-ice band observed in the infrared spectrum.1063 It also allows BVAOA to suggest. following the ideas of Soker Tylenda (2003)) and Retter Marom (2003). that the 1994 outburst of V4332 Sgr might have been due to infall of an inner planet due to tidal interaction with the disc.," It also allows \cite{bva} to suggest, following the ideas of Soker Tylenda \cite{soktyl}) ) and Retter Marom \cite{retmar}) ), that the 1994 outburst of V4332 Sgr might have been due to infall of an inner planet due to tidal interaction with the disc."1064" According to BVAO4 the presently observed emission spectrum and the infrared excess. which appeared in 1999, might have been due to interaction of the 1994 ejecta with the circumstellar disc."," According to \cite{bva} the presently observed emission spectrum and the infrared excess, which appeared in 1999, might have been due to interaction of the 1994 ejecta with the circumstellar disc."1065 However. an interaction of this kind should have happen during the 1994 eruption or soon after it.," However, an interaction of this kind should have happen during the 1994 eruption or soon after it."1066 Assuming an expansion velocity of 100 km/s (MWT99)) the distance of 400 Reo is reached in a month., Assuming an expansion velocity of 100 km/s \cite{martini}) ) the distance of 400 $R_{\sun}$ is reached in a month.1067 Thus it is not clear why the IR excess might have appeared ~5 years later and why it might have been brightening with time. as observed.," Thus it is not clear why the IR excess might have appeared $\sim$ 5 years later and why it might have been brightening with time, as observed."1068 The idea of a protostellar disc can however be used to explain the increasing IR excess. although in a different way to that in BVAOA..," The idea of a protostellar disc can however be used to explain the increasing IR excess, although in a different way to that in \cite{bva}."1069 In the case of an inner planet and an external aceretion disc. tidal forces transfer angular momentum from the orbital motion of the planet to the inner edge of the disc.," In the case of an inner planet and an external accretion disc, tidal forces transfer angular momentum from the orbital motion of the planet to the inner edge of the disc."1070 As a result the inner edge of the dise is kept at a certain distance from the planet orbit. determined by à balance between the rate of the angular momentum gain from the planet and the rate in which the angular momentum is transported outward in the disc.," As a result the inner edge of the disc is kept at a certain distance from the planet orbit, determined by a balance between the rate of the angular momentum gain from the planet and the rate in which the angular momentum is transported outward in the disc."1071 When the planet gets accreted by the central star the eruption can disturb the disc. possibly destroying its inner regions.," When the planet gets accreted by the central star the eruption can disturb the disc, possibly destroying its inner regions."1072 After the event the inner disc does not gain angular momentum any more., After the event the inner disc does not gain angular momentum any more.1073 As a result the inner edge will be approaching the star and the inner dise will become hotter and hotter., As a result the inner edge will be approaching the star and the inner disc will become hotter and hotter.1074 Thus we can speculate that just after the 1994 eruption the inner disc in the V4332 Sgr system was far away. cold and thus radiating only in the far IR. while at present the inner disc regions are closer to the star. thus they are hotter and dominating the observed brightness in the KLM bands.," Thus we can speculate that just after the 1994 eruption the inner disc in the V4332 Sgr system was far away, cold and thus radiating only in the far IR, while at present the inner disc regions are closer to the star, thus they are hotter and dominating the observed brightness in the $KLM$ bands."1075 There are. however. observational facts which do not favour the hypothesis that V4332 Ser is à young object.," There are, however, observational facts which do not favour the hypothesis that V4332 Sgr is a young object."1076 First. the position of V4332 Ser in galactic coordinates is /=13°63. b=-9340.," First, the position of V4332 Sgr in galactic coordinates is $l = 13\fdg63$, $b = -9\fdg40$."1077 According to catalogues of Sharpless (1959)). Lynds (1965)). Avedisova (2002)) and Russeil (2003)) there is no HII or star forming region closer than 5° from the position of V4332 Ser.," According to catalogues of Sharpless \cite{sharp}) ), Lynds \cite{lynds}) ), Avedisova \cite{aved}) ) and Russeil \cite{russ}) ) there is no HII or star forming region closer than $5\degr$ from the position of V4332 Sgr."1078 Also from the CO map of Dame et al. (2001.. ," Also from the CO map of Dame et al. \cite{dame}, ,"1079Fig 2) one can see that the object Hes well outside any significant CO emission., Fig 2) one can see that the object lies well outside any significant CO emission.1080" Second. MWT99.. from their high resolution spectra near H,. have derived a radial velocity of V4332 Ser to be —180 km/s. A similar value. te. —160 km/s. can be obtained from the observed positions of all the emission lines listed in their Table 3."," Second, \cite{martini}, from their high resolution spectra near $_\alpha$, have derived a radial velocity of V4332 Sgr to be $-$ 180 km/s. A similar value, i.e, $-$ 160 km/s, can be obtained from the observed positions of all the emission lines listed in their Table 3."1081" This result has been derived from outburst spectra so interpretations other than the radial velocity of the object (although very unlikely) can be considered (see MWT99)),", This result has been derived from outburst spectra so interpretations other than the radial velocity of the object (although very unlikely) can be considered (see \cite{martini}) ).1082 However. the observed wavelengths of the Cal line in our Table 2. still gives a similar radial velocity. te -- 140 km/s. The Galactic rotation curve (see e.g. Brand Blitz 1993)) predicts. for the position of V4332 Ser and a distance of 1.8 kpe. Visr= km/s which ts equivalent to a heliocentric radial velocity of ~+3 km/s. A CO radial velocity map of Dame at al. (2001.," However, the observed wavelengths of the CaI line in our Table \ref{spectr_t}1083 still gives a similar radial velocity, i.e $-$ 140 km/s. The Galactic rotation curve (see e.g. Brand Blitz \cite{bb93}) ) predicts, for the position of V4332 Sgr and a distance of 1.8 kpc, $V_{LSR} = +13$ km/s which is equivalent to a heliocentric radial velocity of $\sim +3$ km/s. A CO radial velocity map of Dame at al. \cite{dame},"1084 Fig., Fig.1085 3) shows. at the galactic longitude of V4332 Ser. Vise between O and +140 km/s. Thus V4332 Ser does not follow the Galactic rotation which is notexpected to be the case for a young object.," 3) shows, at the galactic longitude of V4332 Sgr, $V_{LSR}$ between 0 and +140 km/s. Thus V4332 Sgr does not follow the Galactic rotation which is notexpected to be the case for a young object."1086 Formation of a circumstellar dise is. however. also possible in the merger scenario proposed by Soker Tylenda (2003)," Formation of a circumstellar disc is, however, also possible in the merger scenario proposed by Soker Tylenda \cite{soktyl}) )"1087rest-frame UV luminosity.,rest–frame UV luminosity.1088" No measure of the dependence of the EW distribution as a function of Myy has been obtained, however."," No measure of the dependence of the EW distribution as a function of $M_{UV}$ has been obtained, however."1089" We model the EW distribution assuming that at EW>0 it is represented by a Gaussian centered on EW=0 with an additional constant tail up to150A,, and at EW<0 by a constant level down to some EW,;,;, value, and null below."," We model the EW distribution assuming that at $>0$ it is represented by a Gaussian centered on EW=0 with an additional constant tail up to, and at $<0$ by a constant level down to some $_{min}$ value, and null below."1090 We take the width of the Gaussian and the two tails to reproduce the results of V09 and S10 at different rest-frame magnitudes., We take the width of the Gaussian and the two tails to reproduce the results of V09 and S10 at different rest–frame magnitudes.1091" Specifically, we derive from the bright galaxies in V09 a standard deviation for the Gaussian of10A,, and assume that it is constant at all magnitudes."," Specifically, we derive from the bright galaxies in V09 a standard deviation for the Gaussian of, and assume that it is constant at all magnitudes."1092" We then divide our sample in two luminosity bins (—20.5«Μυν and —20.5«Muy —19.5) and adjust the two tails in order to reproduce the fraction of galaxies with EW>50À ggiven by S10 and the fraction of galaxies with EW 5 and EW>20 ((for the two bins, respectively), as given by the V09 data."," We then divide our sample in two luminosity bins $-20.5< M_{UV}$ and $-20.5<M_{UV}<-19.5$ ) and adjust the two tails in order to reproduce the fraction of galaxies with $EW>50$ given by S10 and the fraction of galaxies with EW $>5$ and $>20$ (for the two bins, respectively), as given by the V09 data."1093" The resulting distributions are shown in Figure 3 for the two magnitude bins, and are reasonably similar in shape to the EW distribution at z5—6 (S07), and show a moderate evolution from the z~3—5 (Shapleyetal.2003, one."," The resulting distributions are shown in Figure \ref{EW_sim} for the two magnitude bins, and are reasonably similar in shape to the EW distribution at $z\simeq1094 5-6$ (S07), and show a moderate evolution from the $z\simeq 3-5$ \citep[][D10]{Shapley2003} one."1095 We then D10)compute the probability of detecting N lines at a given S/N in our sample of 7 objects., We then compute the probability of detecting $N$ lines at a given S/N in our sample of 7 objects.1096" For each object we randomly extract a redshift from the C10 distribution, we compute the corresponding Myy from the observed Y band magnitude (taking into account the IGM absorption at that redshift), and we then randomly extract an EW from the corresponding distribution."," For each object we randomly extract a redshift from the C10 distribution, we compute the corresponding $M_{UV}$ from the observed $Y$ band magnitude (taking into account the IGM absorption at that redshift), and we then randomly extract an EW from the corresponding distribution."1097 If the EW is larger than the minimum detectable EW at the corresponding wavelength (Fig.1)) for a given S/N we conclude that the object would be detected., If the EW is larger than the minimum detectable EW at the corresponding wavelength \ref{EW_lambda}) ) for a given S/N we conclude that the object would be detected.1098" We assume ffor the line, asfound at z=6.9 by Iyeetal.(2006) (see also Fig.1))."," We assume for the line, asfound at $z=6.9$ by \citet{Iye2006} (see also \ref{EW_lambda}) )."1099" Clearly, intrinsically broader lines would be harder to detect."," Clearly, intrinsically broader lines would be harder to detect."1100" We perform this exercise 10° times over the whole sample, requiring S/N»10 for the detection (larger than the S/N of the possible detection in G2.11408), and we finally obtain the probability distribution shown in the lower panel of Fig. 3.."," We perform this exercise $10^5$ times over the whole sample, requiring $>10$ for the detection (larger than the S/N of the possible detection in 1408), and we finally obtain the probability distribution shown in the lower panel of Fig. \ref{EW_sim}."1101" Under these assumptions, the probability of detecting no line in our sample is very small, about2%,, while the typical number of that we should have detected is between 2 and 4."," Under these assumptions, the probability of detecting no line in our sample is very small, about, while the typical number of that we should have detected is between 2 and 4."1102" We also find a low probability of having 1 detection at 5, as found in our (4%))sample."," We also find a low probability ) of having 1 detection at $>5$, as found in our sample."1103" The same probability S/N>adopting the S07 distribution would be much smaller (c 107%), because of the substantial tail of objects with large EW."," The same probability adopting the S07 distribution would be much smaller $\simeq 10^{-3}$ ), because of the substantial tail of objects with large EW."1104" Even using the Shapleyetal.(2003) distribution, which has a lower fraction of high EW objects, the probability is still rather low (996))."," Even using the \citet{Shapley2003}1105 distribution, which has a lower fraction of high EW objects, the probability is still rather low )."1106" We conclude that, with all the obvious caveats due to the small size of our sample and to possible observational mishaps, the lack of prominent lines in our sample is statistically significant."," We conclude that, with all the obvious caveats due to the small size of our sample and to possible observational mishaps, the lack of prominent lines in our sample is statistically significant."1107" On practical level, our results show how challenging it is to obtaina large samples of spectroscopically confirmed galaxies at z>6.5 with current instrumentation, especially if one aims at reaching the level of completeness (>> needed to robustly measure the luminosity function."," On a practical level, our results show how challenging it is to obtain large samples of spectroscopically confirmed galaxies at $z>6.5$ with current instrumentation, especially if one aims at reaching the level of completeness $\gg 50\%$ ) needed to robustly measure the luminosity function."1108" 50%)Our observations imply that this goal will have to wait until a future generation of instruments is available, either 8m telescopes equipped with multi-object spectrographs more efficient in the z and Y bands or, more likely, the new generation of telescopes, such as theTelescope or 20-40m ground-based facilities."," Our observations imply that this goal will have to wait until a future generation of instruments is available, either 8m telescopes equipped with multi-object spectrographs more efficient in the $z$ and $Y$ bands or, more likely, the new generation of telescopes, such as the or 20-40m ground-based facilities."1109" Nonetheless, our analysis appears to show that the failure to detect prominent in our sample is not only due to the insufficiency of current instrumentation."," Nonetheless, our analysis appears to show that the failure to detect prominent in our sample is not only due to the insufficiency of current instrumentation."1110 One possibility is that a significant fraction of the candidates are lower redshift interlopers., One possibility is that a significant fraction of the candidates are lower redshift interlopers.1111" We test this possibility by extrapolating to z~7 the observed contamination in spectroscopic samples at z~4, 5 and 6 (V09, Table 4 of B, V and i dropouts), which increases with redshift."," We test this possibility by extrapolating to $z\sim 7$ the observed contamination in spectroscopic samples at $z\sim 4$, 5 and 6 (V09, Table 4 of B, V and i dropouts), which increases with redshift."1112" We assume that amongst our z—band dropout sample the fraction of contaminants could be ~25%,, i.e., 2 out of 7 candidates."," We assume that amongst our z--band dropout sample the fraction of contaminants could be $\sim$, i.e., 2 out of 7 candidates."1113" This estimate may be pessimistic, given the excellent photometric quality of the Hawk-I and WFC3 data, and the more careful cleaning of lower z interlopers compared to the V09 samples."," This estimate may be pessimistic, given the excellent photometric quality of the Hawk–I and WFC3 data, and the more careful cleaning of lower $z$ interlopers compared to the V09 samples."1114" However, the contaminant population may be changing at higher redshifts, and different and previously unstudied galaxy types may be entering the selection window."," However, the contaminant population may be changing at higher redshifts, and different and previously unstudied galaxy types may be entering the selection window."1115" Ignoring these uncertainties, we repeated the Monte Carlo simulation for all possible choices of 5 candidates from our 7, finding that the probability of detecting no line at S/N>10 is still rather low, being typically 896,, and only in one case reaching (this range depends on which candidates are excluded from the sample)."," Ignoring these uncertainties, we repeated the Monte Carlo simulation for all possible choices of 5 candidates from our 7, finding that the probability of detecting no line at $>10$ is still rather low, being typically , and only in one case reaching (this range depends on which candidates are excluded from the sample)."1116 Another explanation for the paucity of detections, Another explanation for the paucity of detections111716.,16.1118 This model may have plausible values of the plivsical paraueters., This model may have plausible values of the physical parameters.1119 It shows a modest temperature and a rich chemistry., It shows a modest temperature and a rich chemistry.1120 Evideuth. the familar galactic tracers CS. SO. IICO!. CSIL. CN. ΠΝΟ aud ΠΟΝ should be useful for determining the nature of the intracluster material iu cases where the parameters are similar to those of Model 16.," Evidently, the familar galactic tracers CS, SO, $^{+}$ , $_{2}$ H, CN, HNC and HCN should be useful for determining the nature of the intracluster material in cases where the parameters are similar to those of Model 16."1121 Figure 2. compares the chemistries predicted by two Models. 14 aud 18. which both have low values of ¢ but differ stronelv inJ£ (1«1077 ere ὃν 1 aud," Figure \ref{fig:2} compares the chemistries predicted by two Models, 14 and 18, which both have low values of $\zeta$ but differ strongly in$H$ $1\times 10^{-22}$ erg $^{-3}$ $^{-1}$ and"1122location of in a void region of the galaxy distribution. is the main question addressed in the analvsis.,"location of in a void region of the galaxy distribution, is the main question addressed in the analysis."1123 The obtained results and the final conclusions are as follows:, The obtained results and the final conclusions are as follows:1124"The LMCC exists as eight files. one for each hour in RA from 20"" to 4"".","The LMCC exists as eight files, one for each hour in RA from $^{\mbox{\small h}}$ to $^{\mbox{\small h}}$."1125 Each file contains 60. subdirectories corresponding to the minutes of RA. and the light-motion curves are stored in these directories based on their mean RA coordinates.," Each file contains 60 subdirectories corresponding to the minutes of RA, and the light-motion curves are stored in these directories based on their mean RA coordinates."1126 The LMCC contains 3700548 light-motion curves. 2807047 of which have at least 20 epochs.," The LMCC contains 3700548 light-motion curves, 2807047 of which have at least 20 epochs."1127 The files (~29.5 Gb compressed) may be obtained by web download fromhttpz/fdus., The files $\sim$ 29.5 Gb compressed) may be obtained by web download from.1128sdss.org/value-added/stripe-82public. Light-motion curve plotting tools written in IDL variabilitv/SDS8may also82 be downloaded from the same website., Light-motion curve plotting tools written in IDL may also be downloaded from the same website.1129 A single light-motion curve is stored as an tile with a name constructed from the unweighted mean position of the corresponding object., A single light-motion curve is stored as an file with a name constructed from the unweighted mean position of the corresponding object.1130 The light-motion curve tile contains a header line describing the column meanings. followed by exactly five rows for each epoch tone row for each wave band) in strict time order.," The light-motion curve file contains a header line describing the column meanings, followed by exactly five rows for each epoch (one row for each wave band) in strict time order."1131 All five wave band measurements are included for completeness. even though it is possible that at any one epoch. up to four wave band measurements may not satisfy the quality criteria described in Section 2.3.," All five wave band measurements are included for completeness, even though it is possible that at any one epoch, up to four wave band measurements may not satisfy the quality criteria described in Section 2.3."1132 In Table 3 we deseribe the columns that make up a light-motion curve from the LMCC., In Table \ref{tab:lmc} we describe the columns that make up a light-motion curve from the LMCC.1133 Figure 7 shows some clear examples of photometric variability and motion from the LMCC., Figure \ref{fig:example} shows some clear examples of photometric variability and motion from the LMCC.1134 Figure 7(a). presents the lighteurve in + (upper points) and g Gower points) of the large- long-period variable star SDSS J220514.58+000845.7. most likely a Mira variable (Watkinsetal. 2008)).," Figure \ref{fig:lc} presents the lightcurve in $r$ (upper points) and $g$ (lower points) of the large-amplitude long-period variable star SDSS J220514.58+000845.7, most likely a Mira variable \citealt{wat2008}) )."1135 Figure 7(b) presents the motion curve of the Known ultracool white dwarf SDSS J224206.19+004822.7 (Kilié.etal. 2006)., Figure \ref{fig:pm} presents the motion curve of the known ultracool white dwarf SDSS J224206.19+004822.7 \citealt{kil2006}) ).1136 Both panels illustrate the dramatic increase in temporal sampling produced by the start of the SDSS-IT Supernova Survey in 2005., Both panels illustrate the dramatic increase in temporal sampling produced by the start of the SDSS-II Supernova Survey in 2005.1137 The HLC supplies a set of 229 derived quantities for euch light-motion curve in the LMCC., The HLC supplies a set of 229 derived quantities for each light-motion curve in the LMCC.1138 These quantities are aimed describing the mean magnitudes. photometric variability and astrometric motion of the objects in the LMCC. and they are calculated using only light-motion curve entries that satisfy the quality constraints from Section 2.3.," These quantities are aimed at describing the mean magnitudes, photometric variability and astrometric motion of the objects in the LMCC, and they are calculated using only light-motion curve entries that satisfy the quality constraints from Section 2.3."1139 Those quantities in the HLC related to photometry are described in Table 4.. while those related o astrometry are described in Table 5..," Those quantities in the HLC related to photometry are described in Table \ref{tab:hlc1}, while those related to astrometry are described in Table \ref{tab:hlc2}."1140 In Table 4.. if a tag name is associated with a 5-element array. hen the 5 values represent the described quantity for each of the tive SDSS wave bands in the order η. g. r. £ and z.," In Table \ref{tab:hlc1}, if a tag name is associated with a 5-element array, then the 5 values represent the described quantity for each of the five SDSS wave bands in the order $u$, $g$, $r$, $i$ and $z$."1141 When a certain wave band has insufficent “good” light-motion curve entries o calculate a particular quantity. a value of zero is stored (this also applies to Table 59).," When a certain wave band has insufficent “good” light-motion curve entries to calculate a particular quantity, a value of zero is stored (this also applies to Table \ref{tab:hlc2}) )."1142" For instance. the first value in the array is set to zero for any light-motion curves with no ""good"" entries for the # band."," For instance, the first value in the array is set to zero for any light-motion curves with no “good” entries for the $u$ band."1143 All quantities in Table + with at the end of the tag name are calculated using a do-clip algorithm that rejects only the worst outlier at any one iteration. and terminates when no more outliers are identified.," All quantities in Table \ref{tab:hlc1} with at the end of the tag name are calculated using a $\sigma$ -clip algorithm that rejects only the worst outlier at any one iteration, and terminates when no more outliers are identified."1144 Similarly. all quantities in Table + withITER at the end of the tag name are calculatedusing the iterative procedure described in Stetson(1996) to dynamically reweight data points based on the size of the residuals from the mean.," Similarly, all quantities in Table \ref{tab:hlc1} with at the end of the tag name are calculatedusing the iterative procedure described in \citet{ste1996} to dynamically reweight data points based on the size of the residuals from the mean."1145 Both these sets, Both these sets1146in the last section.,in the last section.1147 After performing this procedure ou the Afr205 and Afr21 samples (both real aud. mock). we smooth the former with /=10 aud the latter with 7=15AIpe.," After performing this procedure on the $Mr205$ and $Mr21$ samples (both real and mock), we smooth the former with $l=10$ and the latter with $l=15$."1148. These choices maximize the vohune covered while keeping the sampling rate lnieh enough that flament contamination is uuder 25 per cout (see 5.2))., These choices maximize the volume covered while keeping the sampling rate high enough that filament contamination is under $\sim 25$ per cent (see \ref{subsubsec:FilClustRemove}) ).1149" We run the flament finderon the 3/7205 and A£21 ealaxy samples using €=1071 aud C=5071|, respectively, and /v=1."," We run the filament finderon the $Mr205$ and $Mr21$ galaxy samples using $C=40$ and $C=50$ respectively, and $K=1$."1150 The resulting filaments are shown in Fie. 9.., The resulting filaments are shown in Fig. \ref{fig:DataFils}.1151 After removing filaments shorter thaw a sinoothing leugth. the algorithm finds [89 filameuts in Afr205. having a total leneth per unit volume of LOSLO? 5? 2 (£210h Mpc)). while in AL721. 226 filaments are found with a total leusth per unit volune of 7.6«101 4? ? (04—15 Mpc)).," After removing filaments shorter than a smoothing length, the algorithm finds $489$ filaments in $Mr205$, having a total length per unit volume of $1.9 \times 10^{-3}$ $h^{2}$ $^{-2}$ $l=10$ ), while in $Mr21$, $226$ filaments are found with a total length per unit volume of $7.6 \times 10^{-4}$ $h^{2}$ $^{-2}$ $l=15$ )."1152 For comparison. the mock A7205 catalogue contains 151 fiiuneuts with a total leugth per unit volume of 1.7«1027 433 per (121075 Afpc)) aud the mock AZ£21 catalogue contains 235 filamcuts with a total leneth per unit vohune of 8.2.10 152 Mpc? 215p Mpc)).," For comparison, the mock $Mr205$ catalogue contains $451$ filaments with a total length per unit volume of $1.7 \times 10^{-3}$ $h^{2}$ $^2$ $l=10$ ) and the mock $Mr21$ catalogue contains $235$ filaments with a total length per unit volume of $8.2 \times 10^{-4}$ $h^{2}$ $^2$ $l=15$ )."1153 Thus. the προς density of filaments in the sinulatious closely matches that in the real universe.," Thus, the number density of filaments in the simulations closely matches that in the real universe."1154 We found iu 1l. that. above two smootlhiug leugths. dark matter filaments had an exponential leneth distribution that very closely matched that fouud im a Gaussian random field with the same power spectrum.," We found in \ref{subsec:FilGauss} that, above two smoothing lengths, dark matter filaments had an exponential length distribution that very closely matched that found in a Gaussian random field with the same power spectrum."1155 This suggests that. even if the filaments in thedata are iu a different stage of their evolution (i.c.. having differeut ox) than those in the simulations. the leueth distributions should be the same between the two.," This suggests that, even if the filaments in thedata are in a different stage of their evolution (i.e., having different $\sigma_8$ ) than those in the simulations, the length distributions should be the same between the two."1156 This does appear to be the case. as shown iu Fig. 10..," This does appear to be the case, as shown in Fig. \ref{fig:DataLength}."1157 More interesting is the similarity of the width distributions of filament elements. shown in Fie. 11..," More interesting is the similarity of the width distributions of filament elements, shown in Fig. \ref{fig:DataWidth}."1158" Iu the SDSS, we find mean filament widths of 5.5 aand 8.1 oou 10 aand 15 sunoothing scales; with standard deviations of Linh aaud L1Mpe. respectively."," In the SDSS, we find mean filament widths of $5.5$ and $8.4$ on $10$ and $15$ smoothing scales, with standard deviations of $1.1$ and $1.4$, respectively."1159 As was demonstrated in Fie. 7..," As was demonstrated in Fig. \ref{fig:EvolveAll},"1160" filament clement width distributions broaden and shift to naller widths as non-linear evolution proceeds,", filament element width distributions broaden and shift to smaller widths as non-linear evolution proceeds.1161 À laree discrepancy in. for example. σς between the simulations and real data should produce filament populations that are at different stages of non-linear evolution aud have different width distributions.," A large discrepancy in, for example, $\sigma_8$ between the simulations and real data should produce filament populations that are at different stages of non-linear evolution and have different width distributions."1162 As such. Fig.," As such, Fig."1163 11. sugeests that the SDSS filaments are both consistent with the standard model aud. cousistcut with the set of cosmological parameters used iu the simulation., \ref{fig:DataWidth} suggests that the SDSS filaments are both consistent with the standard model and consistent with the set of cosmological parameters used in the simulation.1164 This paper develops and uses an algorithm called the Siioothed. Major Axis Filament Finder to identify individual filaments iu laree-scale structure., This paper develops and uses an algorithm called the Smoothed Major Axis Filament Finder to identify individual filaments in large-scale structure.1165 Iu short. it uses the local eigeuvectors of the deusitv secoud-derivative field to define the filament axis and trace individual filaments.," In short, it uses the local eigenvectors of the density second-derivative field to define the filament axis and trace individual filaments."1166 Filament euds are defined as points at which the rate of chauge of the axis of structure exceeds a specified threshold (see 2))., Filament ends are defined as points at which the rate of change of the axis of structure exceeds a specified threshold (see \ref{subsec:Method}) ).1167 In a ACDM cosmmological simulation. this definition produces filament BHuples that are cousisteut with our visual pression of structure on a particular scale. are complete with few duplicate detections 0.21). andare robust to sparse suupling 3. 1)).," In a $\Lambda$ CDM cosmological simulation, this definition produces filament samples that are consistent with our visual impression of structure on a particular scale, are complete with few duplicate detections\ref{subsubsec:Ct}) ), andare robust to sparse sampling \ref{subsec:SparseFils}) )."1168 Iu addition to the smoothing scale. the flament finder takes the input parameters C. the maxinuun aneular rate of change of the filament axis. aud A. the width of filament removal in units of the smoothing leneth.," In addition to the smoothing scale, the filament finder takes the input parameters $C$ , the maximum angular rate of change of the filament axis, and $K$, the width of filament removal in units of the smoothing length."1169 Using Gaussian smoothing. the “best” values of these input paramctors ave Co= 30. 10. and 507 on 5. 10. and 15 smoothing scales. respectively. and AV=1 for all smoothing scales.," Using Gaussian smoothing, the `best' values of these input parameters are $C=30$ , $40$ , and $50$ on $5$ ,$10$ , and $15$ smoothing scales, respectively, and $K=1$ for all smoothing scales."1170 After we collapse, After we collapse11712NXMMi J225036.9me|573154 (hereafter NATAL J2250|5731). was found which d à characteristic repeating shape on a period of ~174 min (Figure 1)).,"2XMMi J225036.9+573154 (hereafter XMM J2250+5731), was found which showed a characteristic repeating shape on a period of $\sim$ 174 min (Figure \ref{light}) )."1172 was found in the field. of €107.5-1.5 which was observed on 23rd Jan 2007., was found in the field of G107.5-1.5 which was observed on 23rd Jan 2007.1173 The EPIC detectors were cach configured. in full window mode ancl used. the mecium filter., The EPIC detectors were each configured in full window mode and used the medium filter.1174 The field was observed for a total of 32.9 ksec in the EPIC pn detector and 34.5 ksec in both IZPIC MOS detectors., The field was observed for a total of 32.9 ksec in the EPIC pn detector and 34.5 ksec in both EPIC MOS detectors.1175 The source was just outside the field of view of the Optical Monitor., The source was just outside the field of view of the Optical Monitor.1176 Since the source was towards edge of the EPIC detectors. and was a nearby (28) ) theX-ray source. NATAL J225087.9extracted|57312there. which appears to be an active late-type star. we the data from the archive and re-extracted the N-rav light curves and spectra of5731.," Since the source was towards the edge of the EPIC detectors, and there was a nearby $^{''}$ ) X-ray source, XMM J225037.9+573127, which appears to be an active late-type star, we extracted the data from the archive and re-extracted the X-ray light curves and spectra of."1177 The data were processed. usingNALAI-Nowtlou SAS vS.Q.1 (released. Oct 2008)., The data were processed using SAS v8.0.1 (released Oct 2008).1178 Only X-ray events which were eracecl as PATTERN=0-4 and FLAG=O0 were used., Only X-ray events which were graded as =0-4 and =0 were used.1179 Events were extracted from a circular aperture with 10 radius centred on the source. with background events being extracted [rom source [ree areas on the same chip as the source.," Events were extracted from a circular aperture with $^{''}$ radius centred on the source, with background events being extracted from source free areas on the same chip as the source."1180 The background: cata were scaled. to give the same area as the source extraction. area ancl subtracted from. the source area. (, The background data were scaled to give the same area as the source extraction area and subtracted from the source area. (1181We estimate that the nearby source. NMM J225037.9[573127 contributes around 1.5 percent of the lux below 2keV. and a negligible amount at energies above dkeV).,"We estimate that the nearby source XMM J225037.9+573127 contributes around 1.5 percent of the flux below 2keV, and a negligible amount at energies above 4keV)."1182 To ensure that the spectra were correctly Hux calibrated we produced detector spectral response files and ancillary files using the SAS tasksrmfgen anclarfgen respectively., To ensure that the spectra were correctly flux calibrated we produced detector spectral response files and ancillary files using the SAS tasks and respectively.1183 We extracted light curves of in the 0.2LokeV. 0.2LOkeV. 2IO0keV and 4LOkeV energy bands from the EPIC pn. EPIC MOSI and EPIC MOS2 detectors using the method. described. above.," We extracted light curves of in the 0.2--10keV, 0.2–1.0keV, 2–10keV and 4–10keV energy bands from the EPIC pn, EPIC MOS1 and EPIC MOS2 detectors using the method described above."1184 We then obtained. a combined light curve for cach cnerey band. by adding the separate Light curves., We then obtained a combined light curve for each energy band by adding the separate light curves.1185 Each light curve shows a distinctive sharp drop in intensity every 174 min., Each light curve shows a distinctive sharp drop in intensity every 174 min.1186 This is due to the secondary star eclipsing the accretion region(s) on the white dwarf and represents the binary orbital period., This is due to the secondary star eclipsing the accretion region(s) on the white dwarf and represents the binary orbital period.1187 The observation covers 3 eclipses., The observation covers 3 eclipses.1188 We used the standard. Lomb-Scargle power spectrum analysis to search for periods in the data (Figure 2))., We used the standard Lomb-Scargle power spectrum analysis to search for periods in the data (Figure \ref{power}) ).1189 The error on the period was then determined using a bootstrap approach incorporating the generation of synthetic light curves., The error on the period was then determined using a bootstrap approach incorporating the generation of synthetic light curves.1190 We find that the period. is 0.12103 LO0OIS clays (=174.2+2.6 mins)., We find that the period is $\pm$ 0.0018 days $\pm$ 2.6 mins).1191 We folded the light curve in each of the 4 energy bands on this period and show these light curves in Figure 1.., We folded the light curve in each of the 4 energy bands on this period and show these light curves in Figure \ref{light}.1192 We have phased the data so that the eclipse. MM is total in cach energy band. defines ó200.," We have phased the data so that the eclipse, which is total in each energy band, defines $\phi$ =0.0."1193" is relatively faint in reaching a peak of NN| CUs in the combined IZPIC ""iDUUMlOkeV ma""light curve. although this count rate has not for the source being Lar oll-axis."," is relatively faint in X-rays, reaching a peak of $\sim$ 0.08 ct/s in the combined EPIC 0.2–10keV light curve, although this count rate has not been corrected for the source being far off-axis."1194nucertaity is mtroduced by the οἳtects of the umber of lavers ina stack: iu preliminary studies. the dc conductivity perpendicular to the laver plane was shown to increase distinctly with tje nuniber of lavers (see Fer et al. (2008))).,"uncertainty is introduced by the effects of the number of layers in a stack: in preliminary studies, the dc conductivity perpendicular to the layer plane was shown to increase distinctly with the number of layers (see Fei et al. \cite{fei}) )."1195 Assundug this is also true in the optical range. and by analogy with our reatinent of eL. we shall therefore take €| to be similarly proportional to s. which would entail a chauge of the wuning ratio. f. in the DrugeenmanOO formula roni 1/3 to s/3 (to a factor). eusuriug that the contribution of the parallel o)larization is null for erapheue.," Assuming this is also true in the optical range, and by analogy with our treatment of $\epsilon\perp$, we shall therefore take $\epsilon\parallel$ to be similarly proportional to $s$, which would entail a change of the mixing ratio, $f$, in the Bruggeman formula from 1/3 to $s/3$ (to a factor), ensuring that the contribution of the parallel polarization is null for graphene."1196" Iu the abseuce of consensual laboratory ueasurements of €| for uxulti-laver erapheuc. we let ourselves be euided again x the IS feature measurements: we adopt the dielectric function tabulated wv Draiue (1985).. aud take f=5ο, which eives a better fit to observations."," In the absence of consensual laboratory measurements of $\epsilon\parallel$ for multi-layer graphene, we let ourselves be guided again by the IS feature measurements: we adopt the dielectric function tabulated by Draine \cite{dra}, and take $f=s/6$, which gives a better fit to observations."1197 Figure 2 shows the extinction featre for s=0 (1.0. graphene). 0.1. 0.2. 0.1. 0.6 and 1 (e. close to exaphite).," Figure 2 shows the extinction feature for $s=0$ (i.e. graphene), 0.1, 0.2, 0.4, 0.6 and 1 (i.e. close to graphite)."1198 Clearly. the extinction feature cau be coutroled through s. ic. the average umber of lavers per stack iu a erain while the peak frequency shift remains within tli| observational linüts.," Clearly, the extinction feature can be controled through $s$, i.e. the average number of layers per stack in a grain while the peak frequency shift remains within the observational limits."1199 Note. however. that this frequency is 6 % lueher than the observed 5.7 eV (1.6 jan Ly," Note, however, that this frequency is 6 $\%$ higher than the observed 5.7 eV (4.6 $\mu$ $^{-1}$ )."1200 But this is, But this is

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