source,target
 In relation to the present study. tidal ellects like orbital circularisation inject internal energy into the planet. aud herefore delay its contraction.," In relation to the present study, tidal effects like orbital circularisation inject internal energy into the planet, and therefore delay its contraction."
 7 have estimated. the amount of tidal energy absorbed by the planet curing orbital circularisation. and show that it is large enough to alfect the rylanctary raclius.," \citet{jac08} have estimated the amount of tidal energy absorbed by the planet during orbital circularisation, and show that it is large enough to affect the planetary radius."
 We repeat our analysis of tical scales. identifying aancts with anomalously large. radii. by comparison of heir observed. position in the mass-racius ciagrams with heoretical mocels from. ?..," We repeat our analysis of tidal scales, identifying planets with anomalously large radii, by comparison of their observed position in the mass-radius diagrams with theoretical models from \citet{bar08}."
 Phe result is shown in bie. 4..," The result is shown in Fig. \ref{bloat},"
 on the same plane as Fig. 2.., on the same plane as Fig. \ref{tides}.
 Η tidal elfects. were related. το radius excess. some correlation would be expected between the strength of tidal forces and the excess radius.," If tidal effects were related to radius excess, some correlation would be expected between the strength of tidal forces and the excess radius."
 Correlation does not. prove, Correlation does not prove
"where g(p.rg)=ara|rp)21/2""|. and ro=FfC (0.1) where R. is the distance of blob to the center.","where $g(\mu,r_0)=\left[\left(1-r_0^2+r_0^2\mu^2\right)^{1/2}-r_0\mu\right]^{-2}$, and $r_0=R_{\gam}/\rblr \in (0,1)$ , where $R_{\gam}$ is the distance of blob to the center."
 Figure 2 shows the augular distribution in blob comoving frame., Figure 2 shows the angular distribution in blob comoving frame.
 It can be secu that the geometry effect of reflecting mirror isotropizes the radiation at some degrees. but the beamine effect still dominates.," It can be seen that the geometry effect of reflecting mirror isotropizes the radiation at some degrees, but the beaming effect still dominates."
 It is still a good approxination that the radiation is beamed with a cone of solid angle z/I?., It is still a good approximation that the radiation is beamed with a cone of solid angle $\pi/\Gam^2$.
 Thus the pair opacity cau be written as (Could Schrédder 1967) where e.=2/(1gle. and the photon-photou cross section σ.. reads where 2! is the speed of the electron and positron iu the center momentum frame 3={lL2/e;e.(1li., Thus the pair opacity can be written as (Gould Schrédder 1967) where $\eps_c=2/(1-\mu)\epsg$ and the photon-photon cross section $\siggg$ reads where $\beta$ is the speed of the electron and positron in the center of momentum frame $\beta=\left[1-2/\epsg \eps_s(1-\mu)\right]^{1/2}$.
" Performing ofthe integral we have where A(q) reads with ly (ye)nFGdo,οηνis the. integral of. μι) andover the eutire broad line region. which can be evaluated as with oq=7/2I). 0»=arcoosflyAy. and Oy=aresinv1pe."," Performing the integral we have where $A(q)$ reads with and $A_1(\mu)=\int_0^1f(\mu,r_0)dr_0$ is the integral of $f(\mu,r_0)$ over the entire broad line region, which can be evaluated as with $\phi_1=\pi/2-\theta_0$, $\phi_2=\arccos \sqrt{1-\mu^2}-\theta_0$, and $\theta_0=\arcsin \sqrt{1-\mu^2}$."
 The function fy) is plotted in Fig 3., The function $A(q)$ is plotted in Fig 3.
 Since the beamed radiation field reduces the effective cross section of plioton-plioton interaction by a factor 1/(2D)7. it would be couvenieut to check our approximation by the quantity (2D?Atq).," Since the beamed radiation field reduces the effective cross section of photon-photon interaction by a factor $1/(2\Gam)^2$, it would be convenient to check our approximation by the quantity $(2\Gam)^2A(q)$."
 We can easilv fud that it is close to 0.30.1 when q~1.7. sugeesting our approximation is accurate chough.," We can easily find that it is close to $\sim$ 0.4 when $q \sim 1.7$, suggesting our approximation is accurate enough."
 It should be pointed out that the prescut treatincuts of reflected. svuchrotrou radiation cau be couvenicutly extended to the inclusion ofthe radiation from the secondary electrons if we further study the pair cascade im the future., It should be pointed out that the present treatments of reflected synchrotron radiation can be conveniently extended to the inclusion of the radiation from the secondary electrons if we further study the pair cascade in the future.
 The last two subsections are devoted to theinfernal absorption of TeV photons. the developments of pair cascade due to the present mechanisua will be treated In a preparing paper (Wane. Zhou Cheng 2000).," The last two subsections are devoted to the absorption of TeV photons, the developments of pair cascade due to the present mechanism will be treated in a preparing paper (Wang, Zhou Cheng 2000)."
 However it would be useful to compare the dimensions aud radiation of the pair cloud due to theinterne? absorption and the pair halo sugeested by Abaronian. Coppi. Voolk (1991). who argue the formation of pair halo due to the interaction of TeV photons from ACNs with infrared photous of cosmological background. radiation.," However it would be useful to compare the dimensions and radiation of the pair cloud due to the absorption and the pair halo suggested by Aharonian, Coppi, Veolk (1994), who argue the formation of pair halo due to the interaction of TeV photons from AGNs with infrared photons of cosmological background radiation."
 Thisο absorption produces pairs. which are quickly isotropized by an ambicut random magnetic field. forming a extended halo of pairs with typical dimension of (2>1Mpc).," This absorption produces pairs, which are quickly isotropized by an ambient random magnetic field, forming a extended halo of pairs with typical dimension of $R>1$ Mpc)."
 Without specific mechanisin we know that the time scale of halo formation is of about 109 vr., Without specific mechanism we know that the time scale of halo formation is of about $10^6$ yr.
 Usually this absorption is regarded as the main mechanism of deficiency of TeV emissiou from ECRET-loud blazars (Stecker de Jager 1998)., Usually this absorption is regarded as the main mechanism of deficiency of TeV emission from -loud blazars (Stecker de Jager 1998).
 Let us simply estimate the scale of pair halo before it is isotropized by the aimbicut maguetic field., Let us simply estimate the scale of pair halo before it is isotropized by the ambient magnetic field.
 Asstuning the intergalactic magnetic field B=1Ὁ Gauss. then the mean free path of pair electrous im halo reacls The initial halo is of such a dimension. which is much larger than that of absorption case.," Assuming the intergalactic magnetic field $B=10^{-9}$ Gauss, then the mean free path of pair electrons in halo reads The initial halo is of such a dimension, which is much larger than that of absorption case."
 Abarouian. Coppi. Volk (1991) have suggested. some signatures of such an exteuded halo. especially for the light curves iu ligh euergv bands (Coppi Aharonian 1990).," Aharonian, Coppi, Volk (1994) have suggested some signatures of such an extended halo, especially for the light curves in high energy bands (Coppi Aharonian 1999)."
 Auvway this is much larger than that of the present pair cloud., Anyway this is much larger than that of the present pair cloud.
" Thus it is easier to distinguish the two cases,", Thus it is easier to distinguish the two cases.
 We have set a new constraint on the very high enerev Cluission iu terii of observable quantities., We have set a new constraint on the very high energy emission in term of observable quantities.
 As the applications of the present model. we would like to address sole properties of very high cnerey from blazars.," As the applications of the present model, we would like to address some properties of very high energy from blazars."
 The broadbaud contimmiun of blazars show attractive features which indicate the different processes powering the objects., The broadband continuum of blazars show attractive features which indicate the different processes powering the objects.
" The ratio L./L,, of >-ray huuiuositv to optical in flat spectrum radio quasars (FSRQs) is quite different frou that in BL Lacs (Doudi Chiscllini 1995).", The ratio $L_{\gam}/L_{\rm op}$ of $\gam$ -ray luminosity to optical in flat spectrum radio quasars (FSRQs) is quite different from that in BL Lacs (Dondi Ghisellini 1995).
" Comastzi et al (1997) confirmed this result in a lore larger samples and found this mean ratio is roushlv of unity iu DL Lacs aud £./£L.,,%30. uamely hasc0.03 in FSROs."," Comastri et al (1997) confirmed this result in a more larger samples and found this mean ratio is roughly of unity in BL Lacs and $L_{\gam}/L_{\rm op}\approx 30$, namely $l_{\rm rsc}\approx 0.03$ in FSRQs."
" Chisellini et al (1993).using the classical limut of SSC model. show that there is a svstemiatical difference i Doppler factors D between BL Lacs and corc-dominated quasars, (dlogD;=0.12 for BL Lacs aud dogD)=ΟΤΙ for core-domiunatec quasars."," Ghisellini et al (1993),using the classical limit of SSC model, show that there is a systematical difference in Doppler factors $\cd$ between BL Lacs and core-dominated quasars, $\langle \log \cd \rangle=0.12$ for BL Lacs and $\langle \log \cd \rangle=0.74$ for core-dominated quasars."
 These differences have Όσοι. confirmed by Cuuijosa Daly (1996) who assiune that the particles aud magnetic field are in equipartition., These differences have been confirmed by Güiijosa Daly (1996) who assume that the particles and magnetic field are in equipartition.
 This differeuce would lead to more prominent difference of reflected svuchrotron plotou energv deusitv. sugeesting a differcut mechanisui in these objects.," This difference would lead to more prominent difference of reflected synchrotron photon energy density, suggesting a different mechanism in these objects."
" The two systematically differcut features m 7A, and Doppler factor 2 strongly sueecst that the different mechaisiu of 5-ray radiation may operate in these objects.", The two systematically different features in $\lrsc$ and Doppler factor $\cd$ strongly suggest that the different mechanism of$\gam$ -ray radiation may operate in these objects.
 From eq.(15) it is believed that the deficiency of TeV Cluission im radio-loud quasars may be due to the present mechanism., From eq.(15) it is believed that the deficiency of TeV emission in radio-loud quasars may be due to the present mechanism.
"lu this section we preseut some analytical calculations which clearly demoustrate what Kind of combinations vetween amplitudes aud phases of the CMD signal in the V.OW bands aud phases of foregrounds are represented iu the dA, estimator.","In this section we present some analytical calculations which clearly demonstrate what kind of combinations between amplitudes and phases of the CMB signal in the V, W bands and phases of foregrounds are represented in the $d^{\Delta}_{\l,m}$ estimator."
 As was mentioned in Section 1. lis cstimator is desigued as a luear cstimator of the oase difference εςνε. if the phase difference is s«nuall.," As was mentioned in Section 1, this estimator is designed as a linear estimator of the phase difference $\Phi_{\l+\Delta,m}-\Phi_{\l,m}$, if the phase difference is small."
" Let us introduce the model of the signal at each baud Dm[2]=cs|F,Γη]v where ο ds. frequencyd independent. CMD sigual. aud F, Di.is the sui over all kinds. of foregrouuds for cach baud j (svuchrotron. frec-frec. dust endssion etc.)."," Let us introduce the model of the signal at each band $a^{(j)}_{\l,m}=c_{\l,m}+F^{(j)}_{\l,m}$, where $c_{\l,m}$ is frequency independent CMB signal and $F^{(j)}_{\l,m}$ is the sum over all kinds of foregrounds for each band $j$ (synchrotron, free-free, dust emission etc.)."
" According to the investigation above ou the foreground models. it ix realized that without the ILC signal the d, estimation of the foregrounds. especially for V aud W bauds. correspoucds to the signal the power of which is significantly simaller then that of the CAIB Iu terms of moduli aud phases of the foreerouuds at cach frequency baud where δι aud εν are the phases of foreground. aud the CAIB. respectively,"," According to the investigation above on the foreground models, it is realized that without the ILC signal the $d^{\Delta}_{\l,m}$ estimation of the foregrounds, especially for V and W bands, corresponds to the signal the power of which is significantly smaller then that of the CMB In terms of moduli and phases of the foregrounds at each frequency band where $\Phi_{\l,m}$ and $\xi_{\l,m}$ are the phases of foreground and the CMB, respectively."
" Aud from Eq.(31)) we eet aud practically speaking. we have $,,,—&Aim"," And from \ref{dd0}) ) we get and practically speaking, we have $\Phi_{\l,m}=\Phi_{\l+\Delta,m}$."
 Thus taking the le correlation iuto account. we can conclude that it reflects directly the high correlation of the phases of the foregrounds. determined by the GF.," Thus, taking the $4n$ correlation into account, we can conclude that it reflects directly the high correlation of the phases of the foregrounds, determined by the GF."
 Moreover. if any foreground cleaned CAIB maps derived from) different imethods display the ly correlation of phases. if would be evident that foreground residuals still determine the statistical properties of the derived signal.," Moreover, if any foreground cleaned CMB maps derived from different methods display the $4n$ correlation of phases, it would be evident that foreground residuals still determine the statistical properties of the derived signal."
" One of the basic ideas for comparison of phases of two sjenals is to define the following trigonometric moments for the phases $, aud Wy, as: where («€(0, "," One of the basic ideas for comparison of phases of two signals is to define the following trigonometric moments for the phases $\xi_{\l^{'},m}$ and $\Psi_{\l,m}$ as: where $\l \le \l^{'}$."
We appv these trigonometric moments to investigate the pliase correlations for ΤΟΠ FCM and WEM., We apply these trigonometric moments to investigate the phase correlations for TOH FCM and WFM.
" For that we simpy substitute (= in Eq.(33)). and define €),, as the pvase of FOCAL and V, as that of WEN."," For that we simply substitute $\l= \l^{'}$ in \ref{def2}) ), and define $\xi_{\l,m}$ as the phase of FCM and $\Psi_{\l,m}$ as that of WFM."
 The result of he caleulatiouns is preseuted iu Fie.13.., The result of the calculations is presented in \ref{comp}.
 From Fig.13 it can be clearly ποσα that the PCM has strong Af=| correlations starting from (.~10 which rapidly increase for 6>LO. while for WEM these correlations are significautlv damped. especially at low iultipole rauge 6<10.," From \ref{comp} it can be clearly seen that the FCM has strong $\Delta \l=4$ correlations starting from $\l\simeq 40$ which rapidly increase for $\l > 40$, while for WFM these correlations are significantly damped, especially at low multipole range $\l \le 40$."
 ITowever. the dA estimator allow us to clarify the properties of pliase. correlations for low multipole rauge.," However, the $d^{\Delta}_{\l,m}$ estimator allow us to clarify the properties of phase correlations for low multipole range."
 The idea is to apply dA estimator to FCM and WEAL and to compare the power spectra of the signals obtained before and after that.," The idea is to apply $d^{\Delta}_{\l,m}$ estimator to FCM and WFM, and to compare the power spectra of the signals obtained before and after that."
 According to the definition of dA estimator. the power spoectru of the sienalis given by Eq.(28)). which now has the form," According to the definition of $d^{\Delta}_{\l,m}$ estimator, the power spectrum of the signal is given by \ref{pp1}) ), which now has the form"
(2001).. Fanetal.(2003)..and Whiteetal.(2003) (Table 3)).,", \citet{Fan03},and \citet{White03} (Table \ref{tab:transmissiondata}) )."
 Note that we did not include the J1044-0125 (2=5.74) data in Table 3. because it is already incorporated into the compilation of SongailaandCowie(2002)., Note that we did not include the J1044-0125 $z=5.74$) data in Table \ref{tab:transmissiondata} because it is already incorporated into the compilation of \citet{SC2002}.
. In order to combine (hese data. we must take into account the fact that at redshifts 2S5.6. the uncertainties in the transmission data are dominated by intrinsic scatter. while ab higher5 redshifts. (he uncertainties ave dominated by measurement errors.," In order to combine these data, we must take into account the fact that at redshifts $z\lesssim 5.6$, the uncertainties in the transmission data are dominated by intrinsic scatter, while at higher redshifts, the uncertainties are dominated by measurement errors."
 Since our model predicts the transmission. our likelihood. function must properly account Lor both nmeasurenient error as well as the scatter.," Since our model predicts the transmission, our likelihood function must properly account for both measurement error as well as the scatter."
 For a compilation of transmissions. such as that by SongailaandCowie(2002).. the contribution of the data D.=[(z;.T;)] to the likelihood is simply The contribution to the likelihood from the compiled data is estimated (his wav.," For a compilation of transmissions, such as that by \citet{SC2002}, the contribution of the data $D_c = \{(z_j, {\cal T}_j)\}$ to the likelihood is simply The contribution to the likelihood from the compiled data is estimated this way."
" For the individual transmission data. especially in the case where both measurement error and scatter are important. the simplest wav (o do Chis is to estimate scatter 9,444 aul acd in e(quadrature wilh the measurement error ey, to obtain the total uncertainty in the mean for each data point."," For the individual transmission data, especially in the case where both measurement error and scatter are important, the simplest way to do this is to estimate scatter $\sigma_{\rm scatter}$ and add in quadrature with the measurement error $\sigma_{\rm meas}$ to obtain the total uncertainty in the mean for each data point."
" In particular. for individual transmission data D;=σι7;)] with known (Gaussian) errors Tyea, ALC o,uos probabilitv of a mean transmission as a function ol redshift7. is Note that if the measurement error is negligible. (hen Equation 22 reduces (ο Equation (21))."," In particular, for individual transmission data $D_i = \{(z_j,T_j)\}$ with known (Gaussian) errors $\sigma_{\rm meas}$ and $\sigma_{\rm scatter}$, probability of a mean transmission as a function of redshift${\cal T}_z$ is Note that if the measurement error is negligible, then Equation \ref{eq:likelihoodindividual} reduces to Equation \ref{eq:likelihoodcompiled}) )."
" To estimate e,45. we use (he SongailaandCowie(2002) compilation [ουςS5.6 (interpolated). and use the data themselves (binned) to determine o,,,,, al higher redshilis."," To estimate $\sigma_{\rm scatter}$, we use the \citet{SC2002}
 compilation for $z\lesssim 5.6$ (interpolated), and use the data themselves (binned) to determine $\sigma_{\rm scatter}$ at higher redshifts."
 These results are also labluatecl in Table 3.., These results are also tabluated in Table \ref{tab:transmissiondata}.
 The total likelihood function. Z. then is given by the product of the two separate likelihoods.," The total likelihood function, $\cal L$, then is given by the product of the two separate likelihoods."
 For this initial comparison. we fix the other cosmological parameters to their best fit values: n= 0.99. O45?= 0.024. and /= 0.72.," For this initial comparison, we fix the other cosmological parameters to their WMAP-only best fit values: $n=0.99$ , $\Omega_b h^2=0.024$ , and $h=0.72$ ."
 We use the WALAP-only results, We use the WMAP-only results
The low surface brightness of the nebulositv. and the presence of faint [field stars superimposed upon it and nearby. combine to make a morphological interpretation dillicult.,"The low surface brightness of the nebulosity, and the presence of faint field stars superimposed upon it and nearby, combine to make a morphological interpretation difficult."
 Phe image of taken from the ESO data shows the nebulosity clearly but it is unclear whether some of the brighter patehes are stars or enhanced nebular emission., The image of taken from the ESO data shows the nebulosity clearly but it is unclear whether some of the brighter patches are stars or enhanced nebular emission.
 However. because the filter bandpass is ellectively à subset of the broader I. filter. the stellar continuum emission present in the images can be largely removed using an Ro band image.," However, because the filter bandpass is effectively a subset of the broader R filter, the stellar continuum emission present in the images can be largely removed using an R band image."
 The final result shows that the emission is smoother than would appear from the image alone. though two slightly brighter features are weakly visible to the SSW.," The final result shows that the emission is smoother than would appear from the image alone, though two slightly brighter features are weakly visible to the SSW."
 Initial analyses of the nebula used CCD imaging from the SAAO 1.0m provided the ‘first look’., Initial analyses of the nebula used CCD imaging from the SAAO 1.0m provided the 'first look'.
 Though images were obtained in all filters. only the image shows nebulosity.," Though images were obtained in all filters, only the image shows nebulosity."
 Though emission is present (Lable 2)). it is apparently too weak to register in the image. due both to the Large amount of continuum also transmitted. and to the poor response of the CCD in the blue.," Though emission is present (Table \ref{tab:nebulalinelist}) ), it is apparently too weak to register in the image, due both to the large amount of continuum also transmitted and to the poor response of the CCD in the blue."
 llowever these cata were to some extent superseded by the archival images from the ESO 2.2m telescope due to their superior depth and. resolution., However these data were to some extent superseded by the archival images from the ESO 2.2m telescope due to their superior depth and resolution.
 Images in//o.. continuum. OLI] and SH]. were dearchived.. but again only clearly clisplaved nebulosity (Figure. 8))," Images in, continuum, [OIII] and [SII] were dearchived, but again only clearly displayed nebulosity (Figure \ref{fig:xtej0111halpha}) )."
 Continua were removed for both the SL] anc OLLI] images., Continua were removed for both the [SII] and [OIII] images.
 Unfortunately only continuum images were available to model the continuum component: this was not a problem with SH] (at AGTIT.G?31 ο," Unfortunately only continuum images were available to model the continuum component; this was not a problem with [SII] (at $\lambda$ 6717,6731 c.f."
 at. AG563) because of the minimal spectral separation. but the OLLI] A5007 continuum was not adequately represented. by the continuiun image. thus leaving significant stellar images and preventing the detection of the small OLLI] nebula (Section 3.3.4)).," at $\lambda6563$ ) because of the minimal spectral separation, but the [OIII] $\lambda5007$ continuum was not adequately represented by the continuum image, thus leaving significant stellar images and preventing the detection of the small [OIII] nebula (Section \ref{section:linedistribution}) )."
 The ENIM. of was measured both in the un-subtracted ancl continuume-subtracted. OLLI] images and found to be consistent with the other field stars. , The FWHM of was measured both in the un-subtracted and continuum-subtracted [OIII] images and found to be consistent with the other field stars. [
SLL] appears to show extremely weak emission coincicent withdla.,SII] appears to show extremely weak emission coincident with.
 Thus analyses of the structure of the nebula in the light of these species were undertaken spectroscopically (see below).Using the maximum diameter D of ata distance of 63.1pc one arrives at à diameter of 6.1pc for the nebulosity., Thus analyses of the structure of the nebula in the light of these species were undertaken spectroscopically (see below).Using the maximum diameter $D$ of at a distance of 63.1pc one arrives at a diameter of 6.1pc for the nebulosity.
"We begin our study of velocity asymmetries by evaluating the effects of using different values forN, the number of particles used to define the mean velocity of the core in equations (3) to (5).","We begin our study of velocity asymmetries by evaluating the effects of using different values for$N$, the number of particles used to define the mean velocity of the core in equations (3) to (5)."
 Large values of N result in less measurement noise but an overly large effective core for low-mass halos., Large values of $N$ result in less measurement noise but an overly large effective core for low-mass halos.
 Some compromise is thus required for these systems., Some compromise is thus required for these systems.
 In Fig., In Fig.
 2 we show cumulative distributions for our estimates of the square of the velocity offset between the core and the bulk of the main subhalo in each of our objects (equation (3) with Voun taken to be Vinain)-, \ref{fig:fig2} we show cumulative distributions for our estimates of the square of the velocity offset between the core and the bulk of the main subhalo in each of our objects (equation (3) with $\vec{V}_{\rm bulk}$ taken to be $\vec{V}_{\rm main}$ ).
 The four panels refer to our four different halo mass ranges and the four curves in each panel refer to different values of Ν., The four panels refer to our four different halo mass ranges and the four curves in each panel refer to different values of $N$.
 Here and below we divide the estimate for each halo by Ven.=GMa200/r2o0 in order to make it easier to compare results for the different mass ranges., Here and below we divide the estimate for each halo by $V_{200}^2 = G M_{200}/r_{200}$ in order to make it easier to compare results for the different mass ranges.
" In the two bottom panels and for the three lowest N values in the upper right panel, the curves coincide within the noise for all but thesmallest velocity offsets."," In the two bottom panels and for the three lowest $N$ values in the upper right panel, the curves coincide within the noise for all but thesmallest velocity offsets."
 This shows that for these N the central region for which the core velocity is estimated is small enough to be considered to move as a unit., This shows that for these $N$ the central region for which the core velocity is estimated is small enough to be considered to move as a unit.
 In the upper left panel and for the N=1000 curve in the upper right panel a trend towards less extreme offsets for larger N is visible., In the upper left panel and for the $N=1000$ curve in the upper right panel a trend towards less extreme offsets for larger $N$ is visible.
 This is because these halos have small enough masses (~3000 particles on average in the upper left panel) that increasing N washes out a significant part of the core motion., This is because these halos have small enough masses $\sim 3000$ particles on average in the upper left panel) that increasing $N$ washes out a significant part of the core motion.
" On the other hand, differences between the curves at small velocity offset clearly show the effects of small-N noise in our estimates of core velocity."," On the other hand, differences between the curves at small velocity offset clearly show the effects of $N$ noise in our estimates of core velocity."
" These are significant for N=100 but appear acceptably small for N>200, at least as judged from the curvesfor higher mass halos which appear converged at large velocity offset."," These are significant for $N=100$ but appear acceptably small for $N\ge 200$, at least as judged from the curvesfor higher mass halos which appear converged at large velocity offset."
 In the following we adopt N=200 as the best compromise between these competing effects., In the following we adopt $N=200$ as the best compromise between these competing effects.
" In the left panel of Fig. 3,,"," In the left panel of Fig. \ref{fig:fig3},"
 we replot the N=200 curves of Fig., we replot the $N=200$ curves of Fig.
 2 on top of each other for easier comparison., \ref{fig:fig2} on top of each other for easier comparison.
" There is a clear systematic trend for more massive halos to have larger velocity asymmetries, in direct analogy to the trend found above for position asymmetries."," There is a clear systematic trend for more massive halos to have larger velocity asymmetries, in direct analogy to the trend found above for position asymmetries."
" More than a quarter of all cluster halos have core velocities which differ from the mean halo value by at least of V200 (i.e. by velocities greater than about 300 km/s), whereas only a few percent of Milky Way halos have such a large offset."," More than a quarter of all cluster halos have core velocities which differ from the mean halo value by at least of $V_{200}$ (i.e. by velocities greater than about 300 km/s), whereas only a few percent of Milky Way halos have such a large offset."
 The typical offset for low mass halos is small and only about of them have offsets exceeding 0.2V200 (i.e. greater than about 40 km/s)., The typical offset for low mass halos is small and only about of them have offsets exceeding $0.2 V_{200}$ (i.e. greater than about 40 km/s).
 The right panel of Fig., The right panel of Fig.
" 3 shows identical curves, except that the offset is now calculated with respect to the barycentric motion of the halo."," \ref{fig:fig3} shows identical curves, except that the offset is now calculated with respect to the barycentric motion of the halo."
" The resulting distributions are almost indistinguishable from those in the left panel, showing that effects due to substructures and to the definition of the halo boundary are too small to be significant for these statistics."," The resulting distributions are almost indistinguishable from those in the left panel, showing that effects due to substructures and to the definition of the halo boundary are too small to be significant for these statistics."
 An interesting question is whether the relative motions we measure are due to non-equilibrium effects in the outer part of the halos or whether they also reflect significant motions of the core with respect to intermediate halo regions., An interesting question is whether the relative motions we measure are due to non-equilibrium effects in the outer part of the halos or whether they also reflect significant motions of the core with respect to intermediate halo regions.
 Presumably motions of the latter type are more likely to relate to observable galaxy distortions such as warps or lopsidedness., Presumably motions of the latter type are more likely to relate to observable galaxy distortions such as warps or lopsidedness.
" In Fig. 4,,"," In Fig. \ref{fig:fig4},"
 we address this question by measuring velocity offsets of the core relative to different regions of the halo., we address this question by measuring velocity offsets of the core relative to different regions of the halo.
 The four panels here refer to halos in each of our four mass ranges., The four panels here refer to halos in each of our four mass ranges.
" The three curves in each panel give the cumulative offset distributions for core velocities calculated relative to all particles within rooo, relative to all particles within 0.5rooo, and relative to all particles within 0.25r200."," The three curves in each panel give the cumulative offset distributions for core velocities calculated relative to all particles within $r_{200}$ , relative to all particles within $0.5
r_{200}$ , and relative to all particles within $0.25 r_{200}$ ."
" As expected, typical offsets go down in all cases as the"," As expected, typical offsets go down in all cases as the"
Examination of the MOS spectra for X1 also show the same absorption.,Examination of the MOS spectra for X1 also show the same absorption.
" The absorption feature cannot be explained by a neutron star atmosphere model alone, but perhaps may be indicative of metals present on the surface of the neutron star due to active accretion (Rutledgeetal."," The absorption feature cannot be explained by a neutron star atmosphere model alone, but perhaps may be indicative of metals present on the surface of the neutron star due to active accretion \citep{Rutledge02}."
" However, further observations are necessary to 2002)..determine whether the source of the absorption is inherent to the source and any possible physical significance it may have."," However, further observations are necessary to determine whether the source of the absorption is inherent to the source and any possible physical significance it may have."
" X1 is well separated from the other X-ray objects in Figure 6,, lying within the area for population I, X-ray transient and qLMXB, objects."," X1 is well separated from the other X-ray objects in Figure \ref{xcmd}, lying within the area for population I, X-ray transient and qLMXB, objects."
 There is a possible UV counterpart for X1 located 1/445 (1.30) from the X-ray position., There is a possible UV counterpart for X1 located 45 $\sigma$ ) from the X-ray position.
" The UV source is located in the region blue-ward of the main sequence in Figure 5bb, also known as the UV-excess region."," The UV source is located in the region blue-ward of the main sequence in Figure \ref{cmds}b b, also known as the UV-excess region."
" This region is primarily dominated by objects with accretion disk emission (Dieball and is a plausible location in a UV CMD for a qLMXB2007),, counterpart."," This region is primarily dominated by objects with accretion disk emission \citep{Dieball07}, and is a plausible location in a UV CMD for a qLMXB counterpart."
" Although this is also the location in a UV CMD where CV populations are found, we note that the lack of coherent periodicity as well as the absence of a hard component to the X-ray spectrum makes it highly unlikely that X1 is an intermediate polar CV, so we tentatively classify X1 as acandidate qLMXB."," Although this is also the location in a UV CMD where CV populations are found, we note that the lack of coherent periodicity as well as the absence of a hard component to the X-ray spectrum makes it highly unlikely that X1 is an intermediate polar CV, so we tentatively classify X1 as acandidate qLMXB."
" The nearest optical source has V= 16.4, B—V=1.5 and is located 1556 away from ΧΙ (Figure 3)).", The nearest optical source has $V=16.4$ $B - V=1.5$ and is located 56 away from X1 (Figure \ref{findingcharts}) ).
" The UV counterpart and nearby optical source are separated by 0/334, which is a —2c separation given the UV and optical source position errors, making the association of the two rather unlikely."," The UV counterpart and nearby optical source are separated by 34, which is a $\sim$ $\sigma$ separation given the UV and optical source position errors, making the association of the two rather unlikely."
" In addition, the optical source is far from the main sequence of NGC 6819, indicating that it is not a cluster member, and its red color is difficult to reconcile with UV emission."," In addition, the optical source is far from the main sequence of NGC 6819, indicating that it is not a cluster member, and its red color is difficult to reconcile with UV emission."
 The most plausible scenario is that the UV source is associated with X1 (separated by 1.30) and the optical source is an unrelated field star., The most plausible scenario is that the UV source is associated with X1 (separated by $\sigma$ ) and the optical source is an unrelated field star.
" This nearby bright star makes detection of the true optical counterpart difficult, and it is likely that the true optical counterpart to X1 is hidden in the glare from this bright source."," This nearby bright star makes detection of the true optical counterpart difficult, and it is likely that the true optical counterpart to X1 is hidden in the glare from this bright source."
" Although we do not detect an optical counterpart, the X-ray to optical luminosity limit for X1 also shows it to be in the range of X-ray transient objects (Figure 7))."," Although we do not detect an optical counterpart, the X-ray to optical luminosity limit for X1 also shows it to be in the range of X-ray transient objects (Figure \ref{xopt}) )."
" We note that the lack of an optical counterpart is not surprising as qLMXBs can be extremely optically faint (Heinkeetal.2003),, with the added complication of a nearby bright source making detection of an optical counterpart very difficult."," We note that the lack of an optical counterpart is not surprising as qLMXBs can be extremely optically faint \citep{Heinke03}, with the added complication of a nearby bright source making detection of an optical counterpart very difficult."
" The X-ray spectrum, X-ray color, and limit of the X-ray to optical luminosity ratio all point to X1 as a candidate quiescent low mass X-ray binary."," The X-ray spectrum, X-ray color, and limit of the X-ray to optical luminosity ratio all point to X1 as a candidate quiescent low mass X-ray binary."
" In addition, the absence of a hard component to the X-ray spectrum makes it highly unlikely that X1 is a dwarf nova or intermediate polar CV."," In addition, the absence of a hard component to the X-ray spectrum makes it highly unlikely that X1 is a dwarf nova or intermediate polar CV."
" With this candidate qLMXB classification, X1 may be the first system of its kind found in an open cluster environment."," With this candidate qLMXB classification, X1 may be the first system of its kind found in an open cluster environment."
 Higher spatial resolution X-ray and optical data are needed to confirm the counterpart identification and source classification., Higher spatial resolution X-ray and optical data are needed to confirm the counterpart identification and source classification.
 We note that vandenBergetal.(2004) discovered a highly variable and soft X-ray source in M67 (CX 2)., We note that \citet{vandenBerg04} discovered a highly variable and soft X-ray source in M67 (CX 2).
" 'They lack a confident classification due to the variable nature of the source, but detect system parameters that fall between expected values for a black hole qLMXB and a neutron star qLMXB."," They lack a confident classification due to the variable nature of the source, but detect system parameters that fall between expected values for a black hole qLMXB and a neutron star qLMXB."
" However, vandenBergetal. emphasize the need for more information before final classification."," However, \citeauthor{vandenBerg04} emphasize the need for more information before final classification."
" Source X2 has a confident UV source counterpart at Q""779 (0.74c) and a possible optical counterpart at a distance of 17224 (1.150).", Source X2 has a confident UV source counterpart at 79 $\sigma$ ) and a possible optical counterpart at a distance of 24 $1.15\sigma$ ).
" The UV counterpart to X2 is located in the blue UV-excess region of Figure 5bb, an area dominated by objects with accretion disk emission such as CVs (Dieballetal.2007)."," The UV counterpart to X2 is located in the blue UV-excess region of Figure \ref{cmds}b b, an area dominated by objects with accretion disk emission such as CVs \citep{Dieball07}."
". The possible optical counterpart is located in the “gap” region between the main-sequence and white dwarf cooling sequence (Figure 5aa), an area also dominated by CV systems."," The possible optical counterpart is located in the “gap” region between the main-sequence and white dwarf cooling sequence (Figure \ref{cmds}a a), an area also dominated by CV systems."
" Based on the X-ray color and luminosity of X2 (see Figure 6)) it lies on the boundary between population I and II objects, indicating the possibility that this object is a CV candidate or qLMXB."," Based on the X-ray color and luminosity of X2 (see Figure \ref{xcmd}) ) it lies on the boundary between population I and II objects, indicating the possibility that this object is a CV candidate or qLMXB."
" Its X-ray to optical luminosity ratio (Figure 7)) is more indicative of a CV than an LMXB, therefore we classify X2 as a CV candidate."," Its X-ray to optical luminosity ratio (Figure \ref{xopt}) ) is more indicative of a CV than an LMXB, therefore we classify X2 as a CV candidate."
" We fit the X-ray spectrum of X2 with an APEC thermal plasma model with an absorption component from the neutral hydrogen column density to the cluster, resulting in a best fit temperature of kT —6.224 keV, in good agreement with expected dwarf novae temperatures (Bycklingetal.2010;Fertigetal. 2011)."," We fit the X-ray spectrum of X2 with an APEC thermal plasma model with an absorption component from the neutral hydrogen column density to the cluster, resulting in a best fit temperature of kT $=6.2^{+2.1}_{-1.0}$ keV, in good agreement with expected dwarf novae temperatures \citep{Byckling10, Fertig11}."
". The pn spectrum and best-fit model with residuals for X2 is shown in Figure 10,, binned to 15 counts perbin for plotting purposes only."," The pn spectrum and best-fit model with residuals for X2 is shown in Figure \ref{X2spec}, binned to 15 counts perbin for plotting purposes only."
 The error and residual calculations are the same as in Figure 9.., The error and residual calculations are the same as in Figure \ref{X1spec}.
" The background-subtracted light curve for X2, with 500s bins, is shown in Figure 11.."," The background-subtracted light curve for X2, with 500s bins, is shown in Figure \ref{X2lc}."
 The appropriately scaled background light curve is shown in grey., The appropriately scaled background light curve is shown in grey.
 Errors are calculated using the same method as in Figure 11.., Errors are calculated using the same method as in Figure \ref{X2lc}.
" 'The average number of counts per bin is 5.8, or a rate of 0.012 counts s~!, shown with the dashed line."," The average number of counts per bin is 5.8, or a rate of 0.012 counts $^{-1}$ , shown with the dashed line."
" We find a bin with 13 counts,or a rate of 0.026 counts s~, as can be seen at ~10500ss. The probability of finding a bin"," We find a bin with 13 counts,or a rate of 0.026 counts $^{-1}$ , as can be seen at $\sim$ s. The probability of finding a bin"
ins ?7?7.. m which we also compare our abundances will literature results.,"in \ref{sec:abundances}, in which we also compare our abundances with literature results."
 In section we analvze the scatter in the N/O plateau., In section \ref{sec:analysis} we analyze the scatter in the N/O plateau.
 Finally. section ?? gives à stummary of our work.," Finally, section \ref{sec:conclusion} gives a summary of our work."
 In a companion paper to (his one (Henry et al., In a companion paper to this one (Henry et al.
 2006). chemical evolution models are combined with Monte Carlo techniques to further interpret the plateau morphology.," 2006), chemical evolution models are combined with Monte Carlo techniques to further interpret the plateau morphology."
 The N and O ionic abundances for our sample of low metallicity svstems were determined directly. [rom published optical emission-lines using the general expression: where the fraction on the left is the number density ratio of the N or O ion relative to IL. the first fraction on the right is the line flux ratio of a nebular [orbidden feature of ion X! relative to the strength. of 112. is the electron densitv 7). is the electron temperature of the region where the relevant ion is emitting (Ix). a4i is the effective recombination coefficient of IL3 which includes radiative and three-body processes (cm? 1). V ds the fraction of ions N! with an electron in the upper level of the transition of interest. Af is the corresponding spontaneous de-excitation rate coefficient 1). and the last term is the wavelength ratio of the line of interest and Ilo (A).," The N and O ionic abundances for our sample of low metallicity systems were determined directly from published optical emission-lines using the general expression: where the fraction on the left is the number density ratio of the N or O ion relative to $^+$, the first fraction on the right is the line flux ratio of a nebular forbidden feature of ion $^i$ relative to the strength of $\beta$, is the electron density $^{-3}$ ), is the electron temperature of the region where the relevant ion is emitting (K), $\alpha^{eff}_{H\beta}$ is the effective recombination coefficient of $\beta$ which includes radiative and three-body processes $^3$ $^{-1}$ ), $\chi^u$ is the fraction of ions $^i$ with an electron in the upper level of the transition of interest, $^u_l$ is the corresponding spontaneous de-excitation rate coefficient $^{-1}$ ), and the last term is the wavelength ratio of the line of interest and $\beta$ )."
 Note that equation (1)) is based on the assumption that IL? arises [rom recombination., Note that equation \ref{eq:ionabun}) ) is based on the assumption that $\beta$ arises from recombination.
 The contribution to II lines from collisional excitation was neglected because the excitation potentials of II levels are much veher (han (he average thermal equilibrium temperature that characterizes II 1 regions (Osterbrock.D.E.&Ferland.G.J.2006)., The contribution to H lines from collisional excitation was neglected because the excitation potentials of H levels are much higher than the average thermal equilibrium temperature that characterizes H II regions \citep{osterbrock06}.
. In addition. equation (1)) assumes that nebular orbidden lines originate from collisionally excited levels.," In addition, equation \ref{eq:ionabun}) ) assumes that nebular forbidden lines originate from collisionally excited levels."
 Collisional excitation is significant in this case because the low-lving energy levels of the relevant ions are of the order of ATT. However. according to Rubin(1986).. recombinations of 7 can excite the nebular doublet AA3126.3129. used to compute.," Collisional excitation is significant in this case because the low-lying energy levels of the relevant ions are of the order of T. However, according to \cite{rubin86}, recombinations of $^{+2}$ can excite the nebular doublet $\lambda\lambda3726, 3729$, used to compute."
. The effect of this process will be analvzed in the future., The effect of this process will be analyzed in the future.
 Although the presence of the nebular He II 4686 emission-line in the spectra of several metal-poor svstems implies the presence of unobserved ΤΗ . the contribution of this ion to the total oxveen abundance amounts to only a few percent according to our photoionization models.," Although the presence of the nebular He II 4686 emission-line in the spectra of several metal-poor systems implies the presence of unobserved $^{+3}$ $^+$, the contribution of this ion to the total oxygen abundance amounts to only a few percent according to our photoionization models."
 Therefore. we obtained the oxvgen abundance by assuming that," Therefore, we obtained the oxygen abundance by assuming that"
The recent release of Kepler data reporting more than 1200 planet eaudidates transiting (heir host star confirms (hat the core accretion scenario for forming planets is ubiquitous.,The recent release of Kepler data reporting more than 1200 planet candidates transiting their host star confirms that the core accretion scenario for forming planets is ubiquitous.
 Indeed.," Indeed,"
were taken [rom the experimental compilation in the NIST database.,were taken from the experimental compilation in the NIST database.
 For the remainder. the theoretical values of Agearwal Wkeenan were adopted.," For the remainder, the theoretical values of Aggarwal Keenan were adopted."
 Test) calculations inelucling higher-Iving levels. such as those arising from the 3s ?3p4s configuration. were found to have a negligible elfeet on the theoretical line ratios considered in this paper.," Test calculations including higher-lying levels, such as those arising from the $^{2}$ $^{4}$ 4s configuration, were found to have a negligible effect on the theoretical line ratios considered in this paper."
 The electron impact excitation cross sections employed in the present paper are those calculated using the fully relativistic Dirac code by Aggarwal lIxeenan (2005)., The electron impact excitation cross sections employed in the present paper are those calculated using the fully relativistic Dirac code by Aggarwal Keenan (2005).
" For Einstein A-coellicicnts. Agegarwal Keenan (2004) cmplovecl the fully relativistic code to generate results for all transitions among the 54 fine-structure levels of⋅ the ⋅⊐⋅ "" 3873p. i383p"". VEN Le3d. and. MM3d configurations⋅. ofFex."," For Einstein A-coefficients, Aggarwal Keenan (2004) employed the fully relativistic code to generate results for all transitions among the 54 fine-structure levels of the $^{2}$ $^{5}$ , $^{6}$, $^{2}$ $^{4}$ 3d and $^{5}$ 3d configurations of."
. Subsequently. Aegearwal Ixeenan (2005) extended this work to include the additional 36 levels arising from," Subsequently, Aggarwal Keenan (2005) extended this work to include the additional 36 levels arising from"
using new calibrations to improve the accuracy.,using new calibrations to improve the accuracy.
 This position. listed in Table 7.. excludes the ROSAT source ax a possible counterpart. and thus confirms that ASCA iudecd detected a source in the cluster.," This position, listed in Table \ref{tabpos}, excludes the ROSAT source as a possible counterpart, and thus confirms that ASCA indeed detected a source in the cluster."
" Tu the three low-reddened clustersCeu. 66397 aud 66752 we have detected a total of 17 dim N-rav sources, of which 5 are well outside the core."," In the three low-reddened clusters, 6397 and 6752 we have detected a total of 17 dim X-ray sources, of which 5 are well outside the core."
 The N-ray huninosities of these sources are listed in reftxhuu.. and plotted in roffxlun..," The X-ray luminosities of these sources are listed in \\ref{txlum}, and plotted in \\ref{fxlum}."
 The interpretation of roffixhun iuust be made with some care., The interpretation of \\ref{fxlum} must be made with some care.
 First. sources outside the core may not belong to the cluster: the faiutest core source in niuav be a fore- or background source.," First, sources outside the core may not belong to the cluster; the faintest core source in may be a fore- or background source."
 Second. the conversion of observed countrate to huninositv depends ou the assumed spectimm. aud. frour PSPC observations we know that different sources have different spectral paranueters (Johnston et 11991).," Second, the conversion of observed countrate to luminosity depends on the assumed spectrum, and from PSPC observations we know that different sources have different spectral parameters (Johnston et 1994)."
 For exaniple. the kkeV. black body spectra used for he sources iu eoives a ihigher fux for the same countrate than an assumed κο bronasstrahluueOo spectra would Ooeive.," For example, the keV black body spectrum used for the sources in gives a higher flux for the same countrate than an assumed keV bremsstrahlung spectrum would give."
 The wenistrahluung spectrum is used for the three other clusters., The bremsstrahlung spectrum is used for the three other clusters.
 Third. the detection limuts iu 66397. 66752 aud are higher iu the cores. where the »oiut spread fictions of sources overlap. than outside the core.," Third, the detection limits in 6397, 6752 and are higher in the cores, where the point spread functions of sources overlap, than outside the core."
 Such a difference is not presen oereCon., Such a difference is not present in.
 Fourth. we show the average bhuuinositv. and several sources are shown to be variable.," Fourth, we show the average luminosity, and several sources are known to be variable."
 With these poiuts in mind. we note from roffxhun that im all clusters except possibly the most huuimous sources appear be in the cluster core.," With these points in mind, we note from \\ref{fxlum} that in all clusters except possibly the most luminous sources appear to be in the cluster core."
 The main difference between aand the other clusters is that the collision frequency in lis so low that oue expects no low-1nass X-ray. binaries iu it. and that most cataclvsiiic variables in it will be evolved roni primordial binaries (Verbuut ADMevlan 1988. Davies 1997).," The main difference between and the other clusters is that the collision frequency in is so low that one expects no low-mass X-ray binaries in it, and that most cataclysmic variables in it will be evolved from primordial binaries (Verbunt Meylan 1988, Davies 1997)."
 Iu addition. the lius segregation iu this cluster is very low.," In addition, the mass segregation in this cluster is very low."
 Thus in there is no marked difference between the core aud the regions outside the core., Thus in there is no marked difference between the core and the regions outside the core.
 Iu cach cluster we detect sources down to the detection nuit: this sugesests---- hat Wore scusitive observations will detect more sources., In each cluster we detect sources down to the detection limit; this suggests that more sensitive observations will detect more sources.
 Iu the πο... of 66397 and 66752. the ctection of nore source will also require better imagine. so that the [αι sources cau be detected against the brighter ones.," In the cores of 6397 and 6752 the detection of more source will also require better imaging, so that the faint sources can be detected against the brighter ones."
 We do not detect a difference between he hunuinosities of sources in the collapsed elobular cluster 66397. aud the much less concentrated elobular cluster 66752., We do not detect a difference between the luminosities of sources in the collapsed globular cluster 6397 and the much less concentrated globular cluster 6752.
 On the other haud. the hieghlv concentrated cluster [7 Tuc contaius three sources which are an order of magnitude brighter than the brightest sources iun 66397 and 66752.," On the other hand, the highly concentrated cluster 47 Tuc contains three sources which are an order of magnitude brighter than the brightest sources in 6397 and 6752."
Thus external constraints on expansion velocity evolution. wip (oz =1.7 trom the SNAP spectrograph. will be needed (o assess the efficacy of using photometric redshifts from the 5Ne Ia themselves for cosmology.,"Thus external constraints on expansion velocity evolution, up to $z=1.7$ from the SNAP spectrograph, will be needed to assess the efficacy of using photometric redshifts from the SNe Ia themselves for cosmology."
 As described above. calibrated photometric redshifts from the host galaxies are sulficient Lo satislv the bias limit lor the recshilt range z>1.7.," As described above, calibrated photometric redshifts from the host galaxies are sufficient to satisfy the bias limit for the redshift range $z>1.7$."
 Independently determined photometric redshifts from the SN lighteurves could then play a role in resolving remaining ambieuities., Independently determined photometric redshifts from the SN lightcurves could then play a role in resolving remaining ambiguities.
 The accuracy of (he photometric redshifts is certainly sullicient for triggering targeted programs. uusing JWST.," The accuracy of the photometric redshifts is certainly sufficient for triggering targeted follow-up programs, using JWST."
 Categorizing the supernovae into (vpes is an essential element of supernova. astrophysical. and cosmological studies.," Categorizing the supernovae into types is an essential element of supernova, astrophysical, and cosmological studies."
 Photometric seeregation of supernovae can be performed wilh some success based on lighteurve and color evolution., Photometric segregation of supernovae can be performed with some success based on lightcurve and color evolution.
 Early variations of this approach have been presented in Poznanskietal.(2002);Gal-Yam(2004):RiessοἱTonry (2004).," Early variations of this approach have been presented in \citet{poznanski02,galyam04,riess04,barris04}."
. In general. when restirame UV data are available Type II SNe are seen {ο be UV-bright whereas opacity [vom iron-group elements in SNe Ia suppresses the UV brightness.," In general, when restframe UV data are available Type II SNe are seen to be UV-bright whereas opacity from iron-group elements in SNe Ia suppresses the UV brightness."
 In addition. the lighteurves of SNe IP are quite distinct lom those of other SN twpes.," In addition, the lightcurves of SNe IIP are quite distinct from those of other SN types."
 Using a Monte Carlo lighteurve simulation appropriate to the color coverage. cadence. and depth ol SNAP we confirm that for the S/N achievable for SNe la at 2=3 the shape of the D-band lighteurve distinguishes between Type Ia and Type LIP 5Ne.," Using a Monte Carlo lightcurve simulation appropriate to the color coverage, cadence, and depth of SNAP we confirm that for the $S/N$ achievable for SNe Ia at $z=3$ the shape of the B-band lightcurve distinguishes between Type Ia and Type IIP SNe."
 Disünguishing Types Ib and Ie from Type Ia is more difficult. especially in the [ace of an uncertain redshift and dust exGuelion.," Distinguishing Types Ib and Ic from Type Ia is more difficult, especially in the face of an uncertain redshift and dust extinction."
 SNe Ib/c are generally redder (han SNe Ia (wilh resUrame B-V color about 0.5 magnitudes redder (Poznauskietal. 2002)))., SNe Ib/c are generally redder than SNe Ia (with restframe B-V color about 0.5 magnitudes redder \citep{poznanski02}) ).
 The color evolution is different as well: Type Ib/c have similar pre- ancl post-iaxinmumn colors while Type Ia become τοον after (heir maximum brightness is reached., The color evolution is different as well: Type Ib/c have similar pre- and post-maximum colors while Type Ia become redder after their maximum brightness is reached.
 The colors ancl magnitude can be used to largely break (he degeneracy between dust reddening and SN (wpe once a full lighteurve is obtained., The colors and magnitude can be used to largely break the degeneracy between dust reddening and SN type once a full lightcurve is obtained.
 The precise multi-band. photometry afforded by SNAP can greatly improve (he power of such techuiiques., The precise multi-band photometry afforded by SNAP can greatly improve the power of such techniques.
 The largest and most complete sample on which this technique has been tested to date is (he Supernova Legacy Survey., The largest and most complete sample on which this technique has been tested to date is the Supernova Legacy Survey.
 Sullivanetal.(200Ga) use 4-color lighteurve photometry to reject all ten SNe II while rejecting only one SN Ia in a sample of 85 spectroscopically- high-redshift SNe., \citet{sullivan06} use 4-color lightcurve photometry to reject all ten SNe II while rejecting only one SN Ia in a sample of 85 spectroscopically-classified high-redshift SNe.
 ILowever. Sullivanetal.(200G6a) do not demonstrate their ability to reject SNe Ib/c and they do not comment on whether or not all the SNe II were HP. or might have included SNe IL.," However, \citet{sullivan06} do not demonstrate their ability to reject SNe Ib/c and they do not comment on whether or not all the SNe II were IIP, or might have included SNe IIL."
whereas the lower density material is hotter and produces ligh ionization enuission.,whereas the lower density material is hotter and produces high ionization emission.
the value that gives the best results on the following discussion of the cosmological parameters.,the value that gives the best results on the following discussion of the cosmological parameters.
 The initial guess for the cosmological weight function wl!) is not very important for the determination of the shear. because a wrong choice would only lead to an increase of the errors (at least in the weak lensing limit: see comment after Eq. (C4))).," The initial guess for the cosmological weight function $w^{[0]}$ is not very important for the determination of the shear, because a wrong choice would only lead to an increase of the errors (at least in the weak lensing limit; see comment after Eq. \ref{K}) ))."
 This behavior ts well verified in the simulations., This behavior is well verified in the simulations.
" At the end of the first inner iteration. the mass density obtained. 5/4], does not differ significantly from the determinations obtained at the end of the following cycles. wl! (except for a factor arising from the global sealing invariance)."," At the end of the first inner iteration, the mass density obtained, $\kappa^{[0,I]}$, does not differ significantly from the determinations obtained at the end of the following cycles, $\kappa^{[j, I]}$ (except for a factor arising from the global scaling invariance)."
 After the first inner cycle. the mass and the shear maps are close to the best ones that we can hope to have.," After the first inner cycle, the mass and the shear maps are close to the best ones that we can hope to have."
 This suggests that the number of inner iterations 7 could decrease after the first cycle., This suggests that the number of inner iterations $I$ could decrease after the first cycle.
 Thus we could start with J=5 for j= 0. and then let /=3 or even [=2 for j>0.," Thus we could start with $I=5$ for $j=0$ , and then let $I=3$ or even $I=2$ for $j > 0$."
 The reduction of 7 can lead to a significant reduction of machine-time., The reduction of $I$ can lead to a significant reduction of machine-time.
 The following step ts the determination of the cosmological weight., The following step is the determination of the cosmological weight.
 Simulations show that the number of outer iterations needed to obtain a good estimation of the cosmological weight is very low. say .7=3. because the method ts able to give the correct shear map after the first inner loop.," Simulations show that the number of outer iterations needed to obtain a good estimation of the cosmological weight is very low, say $J = 3$, because the method is able to give the correct shear map after the first inner loop."
 Simulations also show that this property. strictly expected only in the weak lensing limit. is in fact valid (at least approximately) in the general case. provided the lens is not too “strong.”," Simulations also show that this property, strictly expected only in the weak lensing limit, is in fact valid (at least approximately) in the general case, provided the lens is not too “strong.”"
 The estimated cosmological weight w(:) (solid lines in Fig., The estimated cosmological weight $\hat w(z)$ (solid lines in Fig.
 4) is smooth on the characteristic scale of IW.(+.<7) (see Eq. (C7))):," 4) is smooth on the characteristic scale of $W_z(z, z')$ (see Eq. \ref{<w>}) ));"
 because of this. 1t remains at a finite value atτα.," because of this, it remains at a finite value at $z
= z_\mathrm d$."
 Moreover. as expected. the error on the cosmological weight increases slightly for + near στ and more for high values of :. where the number of galaxies decreases (see Fig.," Moreover, as expected, the error on the cosmological weight increases slightly for $z$ near $z_\mathrm d$ and more for high values of $z$, where the number of galaxies decreases (see Fig."
 3)., 3).
 In particular. the smoothing of the cosmological weight is very important near the cluster. at +— tq. Where the true cosmological weight vanishes abruptly.," In particular, the smoothing of the cosmological weight is very important near the cluster, at $z \simeq z_\mathrm d$ , where the true cosmological weight vanishes abruptly."
 This clearly indicates that neither the limit (OL)> lfor:>x nor the limit for 5.>» can be used to break the global scaling invariance., This clearly indicates that neither the limit $w(z) \rightarrow 1$ for $z \rightarrow \infty$ nor the limit for $z\rightarrow z_\mathrm d^+$ can be used to break the global scaling invariance.
 Figure 4+ shows the reconstruction of the cosmological weight in the case of an Einstein-de Sitter universe: different choices for € and O4 lead to similar results., Figure 4 shows the reconstruction of the cosmological weight in the case of an Einstein-de Sitter universe; different choices for $\Omega$ and $\Omega_{\Lambda}$ lead to similar results.
 Figure 4 should be compared with Fig., Figure 4 should be compared with Fig.
 3 which shows the expected mean value (w(:) and the related error for an Einstein-de Sitter universe with NN—.—10000 galaxies.," 3 which shows the expected mean value $\langle w \rangle (z)$ and the related error for an Einstein-de Sitter universe with $N = 10\, 000$ galaxies."
 Figure 3 clarifies the general properties discussed above and. in more detail. in Appendix D. for the expected error on «w(:) and suggests that our method should be able to constrain significantly the cosmological parameters.," Figure 3 clarifies the general properties discussed above and, in more detail, in Appendix D, for the expected error on $w(z)$ and suggests that our method should be able to constrain significantly the cosmological parameters."
 In fact. the error or w(2) for:c2 is sufficiently small to distinguish different consmological models even if the measured weight are affected by the global scaling invariance.," In fact, the error on $w(z)$ for $z \simeq 2$ is sufficiently small to distinguish different consmological models even if the measured weight are affected by the global scaling invariance."
 From the estimation w(:) of the cosmological weight we can obtain information on the cosmological parameters as explained in Appendix D. For a description of the results. we plot the contours of the (» function of Eq. (D6))," From the estimation $\hat w(z)$ of the cosmological weight we can obtain information on the cosmological parameters as explained in Appendix D. For a description of the results, we plot the contours of the $\ell_2$ function of Eq. \ref{ell2}) )"
 in a square domain |0.1]«10.1 of the O-O4 plane.," in a square domain $[0,1] \times
[0, 1]$ of the $\Omega$ $\Omega_\Lambda$ plane."
 Figures 5 and 6 show the confidence regions obtained for various confidence levels CL in different cosmological models and with different numbers of source galaxies., Figures 5 and 6 show the confidence regions obtained for various confidence levels $\CL$ in different cosmological models and with different numbers of source galaxies.
 From diagrams of this type we argue that 510000 can be considered as a lower bound for the applicability of our method.," From diagrams of this type we argue that $N = 10\,000$ can be considered as a lower bound for the applicability of our method."
 These figures also show that. unless an exceedingly high number of source galaxies is available. point estimation 15 notvery meaningful.," These figures also show that, unless an exceedingly high number of source galaxies is available, point estimation is notvery meaningful."
 In fact. these examples show that the minimum of the4? function can occur quite far from the true values of the cosmological parameters.," In fact, these examples show that the minimum of the$\chi^2$ function can occur quite far from the true values of the cosmological parameters."
 On the other hand.," On the other hand,"
maeuitucde listed tu Table 2:: the effective surface brielituess is then clefinecl as the average surface brightuess within the Reg elliptical aperture.,magnitude listed in Table \ref{tab:results}; the effective surface brightness is then defined as the average surface brightness within the $_{\rm eff}$ elliptical aperture.
 We note that while these values are model iudepencdent. they can be sensitive to systematic errors introduced by shallow imaging observations since tle radii are defined based ou the observed apparent iuagnitucdes.," We note that while these values are model independent, they can be sensitive to systematic errors introduced by shallow imaging observations since the half-light radii are defined based on the observed apparent magnitudes."
 Several of the galaxies tn this sample have surface photometry reported in the literature., Several of the galaxies in this sample have surface photometry reported in the literature.
 Recent observations with partially overlapping samples include. VCC 513. 917. 1036. 1261. anne 1308 reported in Gehaetal.(2003):: VCC 965. 1036. 1122. and 1308 reported in Pieriui(2002):: VCC 1036 aud 1261 reported in Barazzaetal.(2003): aud VCC 513 aud 1713 reported in (1997).," Recent observations with partially overlapping samples include VCC 543, 917, 1036, 1261, and 1308 reported in \citet{GGvM03}; VCC 965, 1036, 1122, and 1308 reported in \citet{P02}; VCC 1036 and 1261 reported in \citet{BBJ03}; and VCC 543 and 1743 reported in \citet{D97}."
. In egeneral. there is reasouable agreement[we between the literature values and those reporte iere. particularly given the variety of filters aud telescopes used by these studies.," In general, there is reasonable agreement between the literature values and those reported here, particularly given the variety of filters and telescopes used by these studies."
 However. there are iotable differeuces between the values reported lere aud the magnitudes aud effective racii liste in Celiaetal.(2003).," However, there are notable differences between the values reported here and the magnitudes and effective radii listed in \citet{GGvM03}."
. With the exception of VCC 1261. it is likely that many of these differeuces can be attributed to dillering image deptlis aud to incomplete surface photometry as a result of the restricted field-otCview of HST unagiug observatious.," With the exception of VCC 1261, it is likely that many of these differences can be attributed to differing image depths and to incomplete surface photometry as a result of the restricted field-of-view of HST imaging observations."
 For VCC 1261. the large discrepaucy between he effective radius reported here aud that reported in Gelaetal.(2003) (inore than a factor of 2) cannot be reconciled by mere observatioual differences: the present. value agrees with the surface shotometry of VCC 1261 reported in Barazzaetal.(2003)..," For VCC 1261, the large discrepancy between the effective radius reported here and that reported in \citet{GGvM03} (more than a factor of 2) cannot be reconciled by mere observational differences; the present value agrees with the surface photometry of VCC 1261 reported in \citet{BBJ03}."
" To enable comparison with galaxies of other morphological types. the outer isophotes (r > 10"")) of the B-band surface brightness profiles were fit to an exponential model: where μή) is the observed surface brightness at semi-major axis r. ji! is the extrapolated centra surlace brightuess. aud a is the exponential scale leneth."," To enable comparison with galaxies of other morphological types, the outer isophotes (r $>$ ) of the B-band surface brightness profiles were fit to an exponential model: where $\mu(r)$ is the observed surface brightness at semi-major axis $r$, $\mu^0$ is the extrapolated central surface brightness, and $\alpha$ is the exponential scale length."
 The face-on central surface brightuess. TM was calculated by applying a Galactic extinction correction aud a line-of-sight correction ol —2.5log(cos7). where 7 is the inclination derived [rom the observed axial ratios and a this disk approximation.," The face–on central surface brightness, $\mu_B^{0,c}$, was calculated by applying a Galactic extinction correction and a line–of–sight correction of $-2.5~{\rm log~(cos}~i)$, where $i$ is the inclination derived from the observed axial ratios and a thin disk approximation."
 While au exponential distribution does uot necessarily imply a clisk-like stellar cistributiou. the simplicity of this fuuctioual form permits a robust analysis of the light distributioi of the outer regious of a galaxy. regardless of the actual shape of the galaxy. (disk or spherokl).," While an exponential distribution does not necessarily imply a disk-like stellar distribution, the simplicity of this functional form permits a robust analysis of the light distribution of the outer regions of a galaxy, regardless of the actual shape of the galaxy (disk or spheroid)."
 The exponeutial fitsenable a direct. comparison between the structural parameters of dE aud dl galaxies (Figure 2))., The exponential fitsenable a direct comparison between the structural parameters of dE and dI galaxies (Figure \ref{fig:struct}) ).
 The top pauel of Figure 2. slows the scale leneth as a fuuctiou of absolute umaguitude while the bottom panel shows the face-ou central surface brightness as a fuuction of absolute magnitude.," The top panel of Figure \ref{fig:struct}
 shows the scale length as a function of absolute magnitude while the bottom panel shows the face–on central surface brightness as a function of absolute magnitude."
 For comparison. the structural parameters for dwarf irregular galaxies frou vanZee(2000) are also shown.," For comparison, the structural parameters for dwarf irregular galaxies from \citet{vZ00} are also shown."
 Despite the fact that this comparisou sample may uot be ideal. since the dls were selected to be isolated galaxies and the dEs are cluster members. these figures confirm that the structural parameters of cls ancl dEs are quite similar 1983)..," Despite the fact that this comparison sample may not be ideal, since the dIs were selected to be isolated galaxies and the dEs are cluster members, these figures confirm that the structural parameters of dIs and dEs are quite similar \citep[as first discussed in][]{LF83}. ."
 In particular. both the dE aud dE samples contain low sur(ace brightuess galaxies," In particular, both the dI and dE samples contain low surface brightness galaxies"
"the cluster core, because low-mass stars have small radii and therefore smaller cross sections for collisions.","the cluster core, because low-mass stars have small radii and therefore smaller cross sections for collisions."
 Fig., Fig.
 4 depicts the number of stellar collisions in the unsegregated cluster models., \ref{fig:coll} depicts the number of stellar collisions in the unsegregated cluster models.
" In clusters less concentrated than r;,z 0.1 pc hardly any collisions occur.", In clusters less concentrated than $r_h \approx$ 0.1 pc hardly any collisions occur.
 The reason is the low stellar density of these models in combination with the large inspiral times of the massive stars., The reason is the low stellar density of these models in combination with the large inspiral times of the massive stars.
" In clusters with half-mass radii larger than τη=0.1 pc, the inspiral times are larger than a few 0.1 Myr, meaning that massive stars do not have time to accumulate in the center while still in the accretion phase."," In clusters with half-mass radii larger than $r_h = 0.1$ pc, the inspiral times are larger than a few 0.1 Myr, meaning that massive stars do not have time to accumulate in the center while still in the accretion phase."
" Instead, they reach the center only after a few Myr, at which point they have arrived already on the main-sequence and have much smaller radii and collision cross sections."," Instead, they reach the center only after a few Myr, at which point they have arrived already on the main-sequence and have much smaller radii and collision cross sections."
" However, even in the most massive and concentrated cluster with N—104 starsand initial half-mass radius Τη=0.033 pc only 41 stars collide with each other."," However, even in the most massive and concentrated cluster with $N=10^4$ starsand initial half-mass radius $r_h=0.033$ pc only 41 stars collide with each other."
" For a Kroupa mass function with stars up to 100 Mo, 90 collisions between 10 Mo stars are necessary to create the missing stars if the initial mass function only extends up to 15 Μο."," For a Kroupa mass function with stars up to 100 $_\odot$, 90 collisions between 10 $_\odot$ stars are necessary to create the missing stars if the initial mass function only extends up to 15 $_\odot$."
" The number of collisions in the unsegregated models are therefore too small compared to the required number of collisions, especially for clusters which have final half-light radii compatible with observed clusters, i.e. clusters starting with rj>0.1 pc."," The number of collisions in the unsegregated models are therefore too small compared to the required number of collisions, especially for clusters which have final half-light radii compatible with observed clusters, i.e. clusters starting with $r_h > 0.1$ pc."
" We note that in the unsegregated runs most collisions happen after massive stars have reached their final mass, in the N=1031 stars, ry=0.033 pc run for example only 3 collisions happen in the pre-MS phase while roughly half of the collisions happen after 1 Myr, i.e. by the time the cluster has already become gas free and collisions would possibly be observable through their collision products or flashes of bright emission."," We note that in the unsegregated runs most collisions happen after massive stars have reached their final mass, in the $N=10^4$ stars, $r_h=0.033$ pc run for example only 3 collisions happen in the pre-MS phase while roughly half of the collisions happen after 1 Myr, i.e. by the time the cluster has already become gas free and collisions would possibly be observable through their collision products or flashes of bright emission."
 We next discuss the evolution of the mass-segregated clusters., We next discuss the evolution of the mass-segregated clusters.
 In the segregated clusters energy and final mass of stars were correlated such that proto-stellar cores which will become the highest mass stars have the lowest energies., In the segregated clusters energy and final mass of stars were correlated such that proto-stellar cores which will become the highest mass stars have the lowest energies.
 The high mass stars therefore form in the cluster center and do not have to spiral into the center by dynamical friction., The high mass stars therefore form in the cluster center and do not have to spiral into the center by dynamical friction.
" Furthermore, the cluster core contracts during the accretion process due to mass increase, further facilitating stellar collisions."," Furthermore, the cluster core contracts during the accretion process due to mass increase, further facilitating stellar collisions."
 Table 1 and Fig., Table 1 and Fig.
 5 show the number of collisions in the segregated models., \ref{fig:collmseg} show the number of collisions in the segregated models.
 It can be seen that the number of collisions which occur in the runs are now about a factor 10 higher than in the unsegregated models., It can be seen that the number of collisions which occur in the runs are now about a factor 10 higher than in the unsegregated models.
" In particular in the most concentrated models starting with τη=0.033 pc, the number of collisions is now sufficiently high to build up a main-sequence of high-mass stars."," In particular in the most concentrated models starting with $r_h=0.033$ pc, the number of collisions is now sufficiently high to build up a main-sequence of high-mass stars."
" A closer look at the data also shows that in the mass segregated models the collisions happen at earlier times on average, for the star cluster with N=104 and τι=0.033 pc cluster, for example, roughly half of all collisions happen in the pre-MS stage."," A closer look at the data also shows that in the mass segregated models the collisions happen at earlier times on average, for the star cluster with $N=10^4$ and $r_h=0.033$ pc cluster, for example, roughly half of all collisions happen in the pre-MS stage."
" We note that these results are not in contradiction to results obtained by Ardietal. (2008),, who found that primordial mass segregation does not lead to a significant"," We note that these results are not in contradiction to results obtained by \citet{ardietal2008}, , who found that primordial mass segregation does not lead to a significant"
mass (again. the flow has a narrow accretion column. since it is almost isothermal).,"mass (again, the flow has a narrow accretion column, since it is almost isothermal)."
 The eravilational attraction of the more massive star is stronger on the accretion line. and most of the mass flows toward (he massive companion.," The gravitational attraction of the more massive star is stronger on the accretion line, and most of the mass flows toward the massive companion."
 | treat the more massive companion. a treatment holds for equal mass components as well.," I treat the more massive companion, a treatment holds for equal mass components as well."
" The angular momentum of mass element residing near the center of mass relative to the center of the accreting star is Ja,Εμ where wyο7/e1y is the angular velocity of the binary svstem."," The angular momentum of mass element residing near the center of mass relative to the center of the accreting star is $j_{\rm cm}=\omega_{12} a_1^2$, where $\omega_{12}=[G (M_{12})]^{1/2}/a_{12}^{3/2}$ is the angular velocity of the binary system."
" For an aceretion disk to form. (his should be larger than the specilie angular momentum on the equator of the accreting star jj=(GM,)!."," For an accretion disk to form, this should be larger than the specific angular momentum on the equator of the accreting star $j_1=(G M_{b1} R_1)^{1/2}$."
 The condition Τους This condition can be met by a large fraction of binary svstems with mass ratio of q~1., The condition reads This condition can be met by a large fraction of binary systems with mass ratio of $q \simeq 1$.
 For q=Ll. a main sequence accretor. with Ry2R.. requires ay21048... and there are many binary svstems with LOR.Say1AU.," For $q=1$, a main sequence accretor, with $R_1 \simeq R_\odot$, requires $a_{12} \gtrsim 10 R_\odot$, and there are many binary systems with $10 R_\odot \lesssim a_{12} \lesssim 1 \AU$."
 However. for ¢=0.5. and q=0.3. the condition reads (yo> 54424. and ay>2071484. respectively.," However, for $q=0.5$, and $q=0.3$, the condition reads $a_{12} > 54 {R_1}$ , and $a_{12} > 271 {R_1}$, respectively."
 For q<1 the condition reads diy>qRy., For $q \ll 1$ the condition reads $a_{12} > q^{-4} {R_1}$.
 Therefore. for ¢<0.5 only a small number of svstems with accreting main sequence stars are expected to form accretion disks.," Therefore, for $ q \lesssim 0.5$ only a small number of systems with accreting main sequence stars are expected to form accretion disks."
 For accreting WDs. the mass ratio can be as small as q0.1.," For accreting WDs, the mass ratio can be as small as $ q \sim 0.1$."
 IHlowever. such svstems are rare and live for a short (ime. since a WD mass is ον0.5.U.. and a companion e10 times as massive. will evolve [ast off the main sequence.," However, such systems are rare and live for a short time, since a WD mass is $\sim 0.5 M_\odot$, and a companion $\sim 10$ times as massive, will evolve fast off the main sequence."
 Overall. (he formation of an accretion disk requires (hat the mass ratio in most systems be g20.3 (note that q<1 by definition).," Overall, the formation of an accretion disk requires that the mass ratio in most systems be $q \gtrsim 0.3$ (note that $q \leq 1$ by definition)."
 If the mass is accreted directly from the center of mass of the accreting binary svstem. the accretion disk will be in the binary-orbital plane. i.e.. perpendicular to the orbital plane wilh (he mass-losing star.," If the mass is accreted directly from the center of mass of the accreting binary system, the accretion disk will be in the binary-orbital plane, i.e., perpendicular to the orbital plane with the mass-losing star."
 Some inclination is expected. though. since mass will start flowing toward the accreting star before reaching the center of mass.," Some inclination is expected, though, since mass will start flowing toward the accreting star before reaching the center of mass."
 In any case. jets. if blown. will be in the orbital plane of the triple-star svsten. or close to it.," In any case, jets, if blown, will be in the orbital plane of the triple-star system, or close to it."
 In this case. each of the two stars in (he accreting binary svstem accretes mass in (urn. when it reaches (he accretion column up-stream side.," In this case, each of the two stars in the accreting binary system accretes mass in turn, when it reaches the accretion column up-stream side."
 The star in(urn. sav star Mg. “cleans”," The star inturn, say star $M_{b1}$ , “cleans”"
“species” of Π.Ι) .,“species” of $_2$ $^+$.
" This approximation is expected to be valid at temperatures of 10 to 20 Ix. since E, is LOL Is for the first excited state of o-H3D. ."," This approximation is expected to be valid at temperatures of 10 to 20 K, since $_{\rm u}$ is 104 K for the first excited state of $o$ $_2$ $^+$."
" The o-H3D partition functiou is caleulated from where the cegeneracies. g, aud gr. are equal to 3."," The $o$ $_2$ $^+$ partition function is calculated from where the degeneracies, $g_u$ and $g_l$, are equal to 3."
" Assumiug LTE (Equation 1). excitation temperatures iu the range 10 to 20 Ix aud the partition [uuction appropriate for these temperatures. N(o-H5D .) limits are calculated [rom the observed upper limit on the HeD J—14,05-14,4 liue flix."," Assuming LTE (Equation 1), excitation temperatures in the range 10 to 20 K and the partition function appropriate for these temperatures, $o$ $_2$ $^+$ ) limits are calculated from the observed upper limit on the $_2$ $^+$ $_{\rm1,0}$ $_{\rm1,1}$ line flux."
 As shown in Figure 3a. the limits on the disk average N(o-HoD ) range from 1.3—2.7x10P em? (with the highest value corresponding to the lowest temperature).," As shown in Figure 3a, the limits on the disk average $o$ $_2$ $^+$ ) range from $1.3-2.7\times10^{12}$ $^{-2}$ (with the highest value corresponding to the lowest temperature)."
 The limit ou the total N(HeD ). Le. o—p. depends on the o-HeD colt density aud the temperature dependent o/p ratio.," The limit on the total $_2$ $^+$ ), i.e. $o + p$, depends on the $o$ $_2$ $^+$ column density and the temperature dependent $o/p$ ratio."
 This ratio has been modeled by Sipilaetal.(2010) [rom Lto 20 Ix for a molecular hydrogene cdeusity of 10° ? aud au interstellar eerain distribution., This ratio has been modeled by \citet{Sipila10} from 4 to 20 K for a molecular hydrogen density of $^6$ $^{-3}$ and an interstellar grain distribution.
 The derived o/p1 ratio is 0.5—0.1 for temperatures of 10-20 Ix. with the lowest ratio at 13-17 Ix. While the assume erain size distribution aud deusity are more appropriate for dense clouds thau disk midplaues. the erain size cdistributiou mainly affects the temperature profile. and Sipilaetal.(2010). [iud that the Ha3D o/pis nearly constant with increasingD> density.," The derived $o/p$ ratio is 0.5–0.1 for temperatures of 10–20 K, with the lowest ratio at 13–17 K. While the assumed grain size distribution and density are more appropriate for dense clouds than disk midplanes, the grain size distribution mainly affects the temperature profile, and \citet{Sipila10} find that the $_2$ $^+$ $o/p$ is nearly constant with increasing density."
 Thus we expect the values calculated with these asstunptious sliould be valid for cisk iuidplaue couditious., Thus we expect the values calculated with these assumptions should be valid for disk midplane conditions.
 Figure 3b shows the result of applying the literature o/p ratios to the previously calculated o-HeD colt density limits to obtain limits on N(H2D ) as a functiou of temperature., Figure 3b shows the result of applying the literature $o/p$ ratios to the previously calculated $o$ $_2$ $^+$ column density limits to obtain limits on $_2$ $^+$ ) as a function of temperature.
 These ⋅⋅ e ⋅ ∐∐↕∐⊳∖↓⋅⋜↕∐∩≺↵↥⋅∩⋯⊓∩−≻↥≍∐≻∟∢∙⋯−⋅⊺∐≺↵↽∐⋅≺↵⊳∖≺, These limits range from $4$ to $21\times10^{12}$ $^{-2}$.
"↵∐∢∙≺↵∩↥⋜↕↽≻≺↲⋜↕↕⊆∖⇁⋜↕↥⋯↵⋜↕↕↥∙⋝⊾↓⊳∖⋜⊔⇂∐⋅≺↵∢∙↕∢∙∩∐⊳∖≺↲≺↽⋯↵∐∢∙≺↵o,2 2 . ye - ⋅ of the [act that the limit on N(o-H3D ) decreases with temperature. while the o/p ratio has a minununm at 13-17 Is. The limit on the total columu density of ious in the midplaue. ΝΟ Ha. D, ). depeuds ou N(H3D } aud the ratio of N(H2D ) to N(OS7 Hy ,D, )."," The presence of a peak value at 13 K is a direct consequence of the fact that the limit on $o$ $_2$ $^+$ ) decreases with temperature, while the $o/p$ ratio has a minimum at 13–17 K. The limit on the total column density of ions in the midplane, $\sum$ $_{3-x}$ $_x^+$ ), depends on $_2$ $^+$ ) and the ratio of $_2$ $^+$ ) to $\sum$ $_{3-x}$ $_x^+$ )."
 Like the o/p ratio. this ratio is expected to vary with deusity and temperature. as well as other euviroumental properties such as grain size (Floweretal.2001:Ceccarelli&Dominik2005:Sipilà2010).," Like the $o/p$ ratio, this ratio is expected to vary with density and temperature, as well as other environmental properties such as grain size \citep{Flower04,Ceccarelli05,Sipila10}."
. This ratio will also depend ou whether or uot CO aud Ne are frozen out. completely. or if low abundauces of these species can be maintained in the gas-phase through non-thermal clesorption (AseusioRamosetal.," This ratio will also depend on whether or not CO and $_2$ are frozen out completely, or if low abundances of these species can be maintained in the gas-phase through non-thermal desorption \citep{AsensioRamos07}."
"2007).. Casellietal.(2008) modeled the ratios of all deuterated Η., isotopologues for molecular hydrogen density 10? ? aud interstellar erains as a [tection of temperature.", \citet{Caselli08} modeled the ratios of all deuterated $_3^+$ isotopologues for molecular hydrogen density $^5$ $^{-3}$ and interstellar grains as a function of temperature.
 As mentioned previously. the," As mentioned previously, the"
or YSOs.,or YSOs.
 The median OCLE-UI color of these blue stars i (V.PDσ0. nae., The median OGLE-III color of these blue stars is $(V-I)\simeq0.1$ mag.
δέ. These are stars with red spectra sometimes showing molecular absorption bauds., These are stars with red spectra sometimes showing molecular absorption bands.
" We identified 68 red sources, where SIMDAD classifies Las infrared. sources (one YSO from and three 211255 sources). d as a PN. 1 as a YSO. 1 as an asvinptotic eiut branch star. and 1 asa galaxy."," We identified 68 red sources, where SIMBAD classifies 4 as infrared sources (one YSO from and three 2mass sources), 1 as a PN, 1 as a YSO, 1 as an asymptotic giant branch star, and 1 as a galaxy."
 The median III color of these sources is (VD)z1.1 nae., The median OGLE-III color of these sources is $(V-I)\simeq1.1$ mag.
 The vield of our survey is determined by a combination of contamination and depth. where it is dificult to fully characterize the effects of depth because of the large variation in integration time created by the weather.," The yield of our survey is determined by a combination of contamination and depth, where it is difficult to fully characterize the effects of depth because of the large variation in integration time created by the weather."
 For diseussion. we simply combine the four fields (Table 2)).," For discussion, we simply combine the four fields (Table \ref{tab:selectionresults}) )."
 We know from the surface density of candidates compared to quasars 9011... 2006)). shown in Figure 8.. that the level of contanunation is high. but much of this is bv desigu jecause of the large απο of available fibers.," We know from the surface density of candidates compared to quasars , ), shown in Figure \ref{fig:cumul}, that the level of contamination is high, but much of this is by design because of the large number of available fibers."
 Tn Figure d aud Table 2.. we preseut our vields frou his observing run. divided by the selection method.," In Figure \ref{fig:VennCand} and Table \ref{tab:selectionresults}, we present our yields from this observing run, divided by the selection method."
 The wid-IR-sclected sources were divided iuto several classes (QSO-Aa. QSO-Ab. ete:," The mid-IR-selected sources were divided into several classes (QSO-Aa, QSO-Ab, etc.;"
 see Section 2.1))., see Section \ref{sec:mIRsel}) ).
 As expected he purest class. QSO-Aa. has the highest coufiiinationrate. atδ," As expected the purest class, QSO-Aa, has the highest confirmationrate, at."
"ν, Next. the two classes QSO-Àb aud QSO- had coufinnation rates of and respectively."," Next, the two classes QSO-Ab and QSO-Ba had confirmation rates of and , respectively."
— 7.7masxX5.7mas. ad a position angle (PA) of —31 for 1.6 Εν. and 3.6masx3.0mos al a PA of —13° for 4.9 Gllz.,"= $7.7~mas\times5.7~mas$, at a position angle (PA) of $-31\degr$ for 1.6 GHz, and $3.6~mas\times3.0~mas$ at a PA of $-13\degr$ for 4.9 GHz."
 We have detected. (wo distinct radio components al 1.6 Gllz (marked 1 and 2 in Figure 1))., We have detected two distinct radio components at 1.6 GHz (marked 1 and 2 in Figure \ref{fig:vlbi}) ).
 There is a suggestion of a weak third component to the north of component 1., There is a suggestion of a weak third component to the north of component 1.
 New sensitive radio observations are required to confirm this feature., New sensitive radio observations are required to confirm this feature.
" The source position is ~0.58"" away [rom the listed optical host galaxy position of R.A. 19h 081m 16.3708. decl."," The source position is $\sim0.58\arcsec$ away from the listed optical host galaxy position of R.A. 19h 08m 16.370s, decl."
 50d 55m 59.588 (7)., 50d 55m 59.58s \citep{Clements81}.
 We identify components l. 2 and 3 to be the core. the jet. aud apossible counterjet. respectively.," We identify components 1, 2 and 3 to be the core, the jet, and a counterjet, respectively."
" The ""core-jet structure extends to 70.3 parsec. αἱ a DA, ol ~25""."," The “core-jet” structure extends to $\sim$ 0.8 parsec, at a P.A. of $\sim25\degr$."
" The peak surlace brightness of the core. and the total radio flux densitv of the ""core-jet structure are 0.46 mJv + and 0.80 mJy. respectively."," The peak surface brightness of the core, and the total radio flux density of the “core-jet” structure are 0.46 mJy $^{-1}$ and 0.80 mJy, respectively."
 The peak intensities and positions of the radio components. estimated using AIPS tasks JAIFIT and IMDIST. are listed in Table 1..," The peak intensities and positions of the radio components, estimated using AIPS tasks JMFIT and IMDIST, are listed in Table \ref{tabparam}."
 In order to test the credibility of component 3. we estimated a surface densitv of spurious noise peaks in the image by dividing the number of noise peaks having a signal-to-noise ratio (S/N) similar to component 3 (£e. S/N 22.6). bv (he entire image.," In order to test the credibility of component 3, we estimated a surface density of spurious noise peaks in the image by dividing the number of noise peaks having a signal-to-noise ratio (S/N) similar to component 3 $i.e.,$ S/N $>$ 2.6), by the entire image."
" We found around. 230 noise peaks in an image of size 0.4""x0.4"". resulting in a noise peak surface density of ~1370 peaks 7."," We found around 230 noise peaks in an image of size $0.4\arcsec\times0.4\arcsec$, resulting in a noise peak surface density of $\sim1370$ peaks $^{-2}$."
 Considering then a region of size (15masx15 mas). centered around the peak source enission. as (he region where a noise peak could be mistaken [or source emission (the core-]et distance being ~5 mes). we estimated thal 0.3 noise peaks could be expected in this region.," Considering then a region of size $15~mas\times15~mas$ ), centered around the peak source emission, as the region where a noise peak could be mistaken for source emission (the core-jet distance being $\sim5~mas$ ), we estimated that 0.3 noise peaks could be expected in this region."
 Therefore. there is a chance that component 3 is a noise peak.," Therefore, there is a chance that component 3 is a noise peak."
" The 1.4 GHz peak flux density of the VLA A-array core (size ~ 1.5"") is ~14 mJv (?)..", The 1.4 GHz peak flux density of the VLA A-array core (size $\sim1.5\arcsec$ ) is $\sim14$ mJy \citep{HotaSaikia06}.
 This implies that only about of the VLA flux density is detected by the VLBA., This implies that only about of the VLA flux density is detected by the VLBA.
 This is similar to what is observed in the Sevlert galaxy NGC 4151 (28%:??)..," This is similar to what is observed in the Seyfert galaxy NGC 4151 \citep[$\sim$8\%;][]{Pedlar93,Ulvestad98}."
 We believe. however. that had self-calibration worked in NGC 6764. its [raction of VLBA to VLA αν density would have been higher.," We believe, however, that had self-calibration worked in NGC 6764, its fraction of VLBA to VLA flux density would have been higher."
 Nevertheless. it appears that there is either a lot of diffuse emission on scales of tens or hundreds of parsecs which are not visible to the VLBA. or the VLBA is nol sensitive to the diffuse radio emission on parsec scales (this appears to be less likely in the case of NGC 6764 considering the large fraction of missing flux). or a combination ol both (see Orienti Prieto 2010 for a discussion on (he missing diffuse emission in VLBI observations).," Nevertheless, it appears that there is either a lot of diffuse emission on scales of tens or hundreds of parsecs which are not visible to the VLBA, or the VLBA is not sensitive to the diffuse radio emission on parsec scales (this appears to be less likely in the case of NGC 6764 considering the large fraction of missing flux), or a combination of both (see Orienti Prieto 2010 for a discussion on the missing diffuse emission in VLBI observations)."
 There appears to be a tentative detection of components 1 and 2 at 4.9 GlIz., There appears to be a tentative detection of components 1 and 2 at 4.9 GHz.
 However. (his detection is also in need of new confirmatory observations. as (he radio peak in components," However, this detection is also in need of new confirmatory observations, as the radio peak in components"
"the loss of opacity leading to smaller NLTE corrections. the opposite being true for $,»B,.","the loss of opacity leading to smaller NLTE corrections, the opposite being true for $S^l_\nu > B_\nu$."
" We see that for the lower level of the weaker line considered in the figure has J.>B,.. whilst STsB,. which leads to overionization of that level and greater departures than the stronger line and greater NLTE abundance corrections,"," We see that for the lower level of the weaker line considered in the figure has $\bar J_{\nu} > B_{\nu}$ , whilst $S^l_\nu \approx B_\nu$, which leads to overionization of that level and greater departures than the stronger line and greater NLTE abundance corrections."
 The effect of H collisions ts in general to reduce the spread of departure coefficients and drive populations towards LTE values., The effect of H collisions is in general to reduce the spread of departure coefficients and drive populations towards LTE values.
 This reduction in the spread of departure coefhcients comes from the coupling of bound states., This reduction in the spread of departure coefficients comes from the coupling of bound states.
 The increase of H collisions gradually reduces the departures from LTE through the atmosphere as shown in Fig. 1: , The increase of H collisions gradually reduces the departures from LTE through the atmosphere as shown in Fig. \ref{fig:departplots}; ;
with an increasing Sy the slope in the departure coefficient profile becomes shallower., with an increasing $\rm S_{H}$ the slope in the departure coefficient profile becomes shallower.
 In Fig., In Fig.
 | it is interesting to see that the rise i b; at around τει3—2.5 for the levels below 1.83 eV becomes smaller with increasing Sy., \ref{fig:departplots} it is interesting to see that the rise in $b_{i}$ at around $\tau_{\rm 5000} \approx -2.5$ for the levels below 1.83 eV becomes smaller with increasing $S_{\rm H}$.
 This could in fact mean an increase in NLTE departures for some levels for increasing Sy. rather than H collidions driving conditions towards LTE which is normally the case.," This could in fact mean an increase in NLTE departures for some levels for increasing $S_{\rm H}$, rather than H collidions driving conditions towards LTE which is normally the case."
 This rise is most likely caused by increased recombination in the upper (infrared) levels followed by a cascade of electrons down to lower levels., This rise is most likely caused by increased recombination in the upper (infrared) levels followed by a cascade of electrons down to lower levels.
 Exactly how this is affected by the increase in Sy is not yet knowr and requires further study., Exactly how this is affected by the increase in $\rm S_{H}$ is not yet known and requires further study.
 The decrease in level population at 73000.< ] causes a drop in opacity for all lines., The decrease in level population at $\tau_{\rm 5000} <$ 1 causes a drop in opacity for all lines.
 As a result of this. the lines form deeper in the atmosphere than in LTE.," As a result of this, the lines form deeper in the atmosphere than in LTE."
 In Fig. 3.. ," In Fig. \ref{fig:depthform}, ,"
"we clearly see this effect. where we show the continuum optical depth 73666 at which the line optical depth r, = 2/3."," we clearly see this effect, where we show the continuum optical depth $\tau_{\rm 5000}$ at which the line optical depth $\tau_{\nu}$ = 2/3."
" We also see that there is an increasingly large logarithmic optical depth difference. Alogrsooo(7,=2/3). between the formation of weak lines in NLTE and LTE. up to = 50mA.. after which the difference becomes constant."," We also see that there is an increasingly large logarithmic optical depth difference, $\Delta\log{\tau_{\mathrm{5000}}(\tau_\nu=2/3)}$, between the formation of weak lines in NLTE and LTE, up to $\approx$ 50, after which the difference becomes constant."
 With a decrease in opacity compared. to LTE. there needs to be an increase of abundance to match the equivalent width of a given line in NLTE.," With a decrease in opacity compared to LTE, there needs to be an increase of abundance to match the equivalent width of a given line in NLTE."
 Opacity is not the only variable affected by NLTE. the source function can also be affected.," Opacity is not the only variable affected by NLTE, the source function can also be affected."
 However. it is the dominant force in driving the NLTE departures within the Fe atom.," However, it is the dominant force in driving the NLTE departures within the Fe atom."
 In Fig. 4..," In Fig. \ref{fig:Chi-WEQHD140283},"
 we plot the abundance correction versus equivalent width for the star HD140283., we plot the abundance correction versus equivalent width for the star HD140283.
 We see that there is a positive correction for the different values of Su., We see that there is a positive correction for the different values of $\rm S_{H}$ .
 There ts a clear trend with equivalent width., There is a clear trend with equivalent width.
 It is how this translates to trends with excitation energy y that will affect Tay: if the abundance corrections only shifted the mean abundance without depending on y then the derived Το would not change., It is how this translates to trends with excitation energy $\chi$ that will affect $T_{\rm eff}$ : if the abundance corrections only shifted the mean abundance without depending on $\chi$ then the derived $T_{\rm eff}$ would not change.
 Through Fig., Through Fig.
 | to Fig., \ref{fig:departplots} to Fig.
 4. the general effects of NLTE on line formation can be seen., \ref{fig:Chi-WEQHD140283} the general effects of NLTE on line formation can be seen.
 The depletion of level populations (Fig. 1) , The depletion of level populations (Fig. \ref{fig:departplots}) )
leads to a lower opacity and shiftsthe depth of formation to deeper levels (Fig. 3))., leads to a lower opacity and shiftsthe depth of formation to deeper levels (Fig. \ref{fig:depthform}) ).
 This also means that a higher abundance is needed within NLTE. leading to positive abundance corrections (Fig. 4)).," This also means that a higher abundance is needed within NLTE, leading to positive abundance corrections (Fig. \ref{fig:Chi-WEQHD140283}) )."
 However. there is a competing effect in some cases where the source function deviates from the Planck function (Fig. 2).," However, there is a competing effect in some cases where the source function deviates from the Planck function (Fig. \ref{fig:sjbplots}) ),"
 which. in the case of the strong lines. compensates for the level depletion and decreases the abundance correction. as is seen in Fig. 4..," which, in the case of the strong lines, compensates for the level depletion and decreases the abundance correction, as is seen in Fig. \ref{fig:Chi-WEQHD140283}. ."
 In order to determine a new Zr for a star. we first need to calculate NLTE corrections for the LTE abundances derived in Paper 1. Abundance corrections of the form ΑνΜα - ApriMuri] are caleulated and applied to the LTE abundances from Paper I to generate NLTE abundances on the same scale as that paper. rather than using solely the new NLTE analysis.," In order to determine a new $T_{\rm eff}$ for a star, we first need to calculate NLTE corrections for the LTE abundances derived in Paper I. Abundance corrections of the form $A_{\rm NLTE,{\sc MULTI}}$ $-$ $A_{\rm LTE,{\sc MULTI}}$ are calculated and applied to the LTE abundances from Paper I to generate NLTE abundances on the same scale as that paper, rather than using solely the new NLTE analysis."
 This procedure is used so as to tie this work to the previous results. thus allowing the limitations of the LTE assumptions in that work to be seen.," This procedure is used so as to tie this work to the previous results, thus allowing the limitations of the LTE assumptions in that work to be seen."
 To do this. a grid of results for a range of abundances is created with increments of 0.02 dex.," To do this, a grid of results for a range of abundances is created with increments of 0.02 dex."
 The abundance values covered by this grid depend on the spread of abundances from individual lines in each star., The abundance values covered by this grid depend on the spread of abundances from individual lines in each star.
 gives an LTE and NLTE equivalent width for each abundance in this σης., gives an LTE and NLTE equivalent width for each abundance in this grid.
 A first step is to determine what WL; from the grid corresponds to the LTE abundance derived in Paper I (?).., A first step is to determine what $W_{\rm LTE}$ from the grid corresponds to the LTE abundance derived in Paper I \citep{Hosfordetal2009}.
 This is done for all Fe lines that are measured in the star., This is done for all Fe lines that are measured in the star.
" The NLTE abundance inferred for a line is the abundance that corresponds to this W, within the grid of NLTE results.", The NLTE abundance inferred for a line is the abundance that corresponds to this $W_{\lambda}$ within the grid of NLTE results.
 The correction is then calculated as AA(Fe) = A(Fednyy — A(Fe)ngi., The correction is then calculated as $\Delta A(\rm Fe)$ = $A(\rm Fe)_{NLTE}$ $-$ $A(\rm Fe)_{LTE}$.
 Fig., Fig.
 5 shows the corrections for the star HD140283 calculated for the three different Sy values: Sy = 0. 0.001 and I.," \ref{Fig:abnd-chi/ew-HD140283} shows the corrections for the star HD140283 calculated for the three different $\rm S_{H}$ values: $\rm S_{H}$ = 0, 0.001 and 1."
 We see a trend in the abundance correction with y. where we have values. from least square fits. of: The non-zero coefficient of y implies that a 7. correction is needed.," We see a trend in the abundance correction with $\chi$, where we have values, from least square fits, of: The non-zero coefficient of $\chi$ implies that a $T_{\rm eff}$ correction is needed."
 The values for Sy = 0 and Sy = 0.001 are very similar and imply that 7. corrections for these two values will be very similar., The values for $\rm S_{H}$ = 0 and $\rm S_{H}$ = 0.001 are very similar and imply that $T_{\rm eff}$ corrections for these two values will be very similar.
 We therefore decided that corrections for only Su = 0 and 1 would be calculated. Sy = 0 representing the maximal NLTE corrections and Sy = | representing the full Drawinian magnitude of neutral H collisions.," We therefore decided that corrections for only $\rm S_{H}$ = 0 and 1 would be calculated, $\rm S_{H}$ = 0 representing the maximal NLTE corrections and $\rm S_{H}$ = 1 representing the full Drawinian magnitude of neutral H collisions."
" To test the corrections. we compared synthetic profiles from the NLTE abundance with the observed profile. and compared measured Wy's with NLTE W,’s fromMULTI. obtained from an abundance given by Ap; + ΔΑ."," To test the corrections, we compared synthetic profiles from the NLTE abundance with the observed profile, and compared measured $W_{\lambda}$ 's with NLTE $W_{\lambda}$ 's from, obtained from an abundance given by $A_{\rm LTE}$ + $\Delta A$."
 The synthetic profiles are convolved with a Gaussian whose width is allowed to vary from line to line., The synthetic profiles are convolved with a Gaussian whose width is allowed to vary from line to line.
 This represents the macroturbulent and instrumental broadening. the latter calculated by fittingσι Gaussian profiles to ThAr lines in IRAF and found to be é 100mA.," This represents the macroturbulent and instrumental broadening, the latter calculated by fitting Gaussian profiles to ThAr lines in IRAF and found to be $\sim$ 100."
. We found that the profiles match the observed line reasonably well. and that measured and calculated W's are comparable. with a standard deviation of 2.3mA.," We found that the profiles match the observed line reasonably well, and that measured and calculated $W_{\lambda}$ 's are comparable, with a standard deviation of 2.3."
. This gives us confidence that the corrections are realistic within the framework of the atomic model used., This gives us confidence that the corrections are realistic within the framework of the atomic model used.
 These corrections were then applied to theWIDTH6LTE abundances used in Paper [ andnew plots of y versus A(Fe) were plotted., These corrections were then applied to theWIDTH6LTE abundances used in Paper I andnew plots of $\chi$ versus$A$ (Fe) were plotted.
 We then nulledtrends in this plot to constrain Z4;NLTE) by recalculating the LTE abundances using the radiative transfer program WIDTH6 (Kurucz Furenlid 1978) exactly as in ? and reapplying the NLTE corrections.derived here from for the original LTEparameters.," We then nulledtrends in this plot to constrain $T_{\rm eff}$ (NLTE) by recalculating the LTE abundances using the radiative transfer program WIDTH6 (Kurucz Furenlid 1978) exactly as in \citet{Hosfordetal2009} and reapplying the NLTE corrections,derived here from for the original LTEparameters."
infalling masses are self-similar.,infalling masses are self-similar.
" Therefore, when the abundance ratios and gas fraction are examined (Figs."," Therefore, when the abundance ratios and gas fraction are examined (Figs."
" and 6)), the three series of models (M8, M9, and M10) are overlapping."," \ref{Fig:nowdOFeH} and \ref{Fig:nowdmuY}) ), the three series of models (M8, M9, and M10) are overlapping."
 In Fig., In Fig.
 7 we plot the mass-metallicity relations predicted by models without wind., \ref{Fig:nowdMsZ} we plot the mass-metallicity relations predicted by models without wind.
" We run models for three different infall masses (105, 109, 10'!°Mo))."," We run models for three different infall masses $10^{8}$, $10^{9}$, $10^{10}$ )."
" The effects of different numbers of bursts (n=1,3, 7), different durations (d= 0.03,0.1,0.3) and different SFEs (e=0.2,0.5,1.0, 2.0) as functions of galactic mass are shown."," The effects of different numbers of bursts $n=1, 3, 7$ ), different durations $d=0.03, 0.1, 0.3$ ) and different SFEs $\epsilon=0.2, 0.5, 1.0,
2.0$ ) as functions of galactic mass are shown."
 It is evident from Fig., It is evident from Fig.
" 7 that our models can very well reproduce the M-Z relation even without galactic wind but just assuming an increase of the number, or duration of bursts, or the efficiency of SF."," \ref{Fig:nowdMsZ} that our models can very well reproduce the M-Z relation even without galactic wind but just assuming an increase of the number, or duration of bursts, or the efficiency of SF."
" If the galactic wind has the same chemical composition as the well-mixed ISM, i.e. w;—1 for all the elements, we call it “normal wind""."," If the galactic wind has the same chemical composition as the well-mixed ISM, i.e. $w_i=1$ for all the elements, we call it “normal wind”."
" In Fig. 8,"," In Fig. \ref{Fig:wdZmuY},"
" we show the the evolutionary tracks predicted by models with normal wind; abundance ratios of log(C/O) vs. 12+log(O/H) and [O/Fe] vs. [Fe/H] are on the left side, while u,—Z and Y—Z relations on the right side."," we show the the evolutionary tracks predicted by models with normal wind; abundance ratios of log(C/O) vs. 12+log(O/H) and [O/Fe] vs. [Fe/H] are on the left side, while $\mu-Z$ and $Y-Z$ relations on the right side."
" The models have the same total infall mass (Ming=10? M5)) and same bursts sequence (t=1/3/5/7/9/11/13 Gyr, with d—0.1 Gyr for each burst), but different wind efficiencies (Aw=0,0.2,0.5,1.0)."," The models have the same total infall mass $M_{inf}=10^9$ ) and same bursts sequence $t=1/3/5/7/9/11/13$ Gyr, with $d=0.1$ Gyr for each burst), but different wind efficiencies $\lambda_w=0, 0.2, 0.5, 1.0$ )."
" Oxygen is produced by massive stars, therefore no oxygen will be ejected into the ISM after star formation ceases."," Oxygen is produced by massive stars, therefore no oxygen will be ejected into the ISM after star formation ceases."
" On the other hand, elements, such as C and N produced by low- and intermediate-mass stars, and Fe mainly produced by SN Ia explosion, are continuously polluting the ISM after the star formation stops, owing to their long lifetime."," On the other hand, elements, such as C and N produced by low- and intermediate-mass stars, and Fe mainly produced by SN Ia explosion, are continuously polluting the ISM after the star formation stops, owing to their long lifetime."
" Therefore, the decrease of the mass of gas (i.e., H and He) and the o-elements lost with the wind will result in a dramatic increasing of the abundance of the “time-delayed? elements."," Therefore, the decrease of the mass of gas (i.e., H and He) and the $\alpha$ -elements lost with the wind will result in a dramatic increasing of the abundance of the ``time-delayed'' elements."
" The stronger the wind, the higher the abundance of C or Fe relative to O predicted by the models."," The stronger the wind, the higher the abundance of C or Fe relative to O predicted by the models."
" 'The main effect of normal winds is to decrease the gas fraction with a smaller effect on the O/H abundance, as we can see from the 12+log(O/H)-y relation (upper right panel of Fig. 8))."," The main effect of normal winds is to decrease the gas fraction with a smaller effect on the O/H abundance, as we can see from the $\mu$ relation (upper right panel of Fig. \ref{Fig:wdZmuY}) )."
 This means that models with normal wind cannot explain the whole spread in O/H observed at a given µ for these galaxies., This means that models with normal wind cannot explain the whole spread in O/H observed at a given $\mu$ for these galaxies.
 There are two possible reasons for that., There are two possible reasons for that.
" One is the same wind efficiency (i.e., wiAw) for both oxygen and hydrogen in the normal wind, so that both O and H decrease at the same time."," One is the same wind efficiency (i.e., $w_i\lambda_w$ ) for both oxygen and hydrogen in the normal wind, so that both O and H decrease at the same time."
 The other one is the short infall time scale assumed (τ1 Gyr)., The other one is the short infall time scale assumed $\tau=1$ Gyr).
" In this case, no primordial gas falls into the galaxies to dilute the ISM at late evolutionary times."," In this case, no primordial gas falls into the galaxies to dilute the ISM at late evolutionary times."
" Therefore, we also developed a model with long infall time scale (r—10 Gyr, magenta dotted lines in Fig. 8))."," Therefore, we also developed a model with long infall time scale $\tau=10$ Gyr, magenta dotted lines in Fig. \ref{Fig:wdZmuY}) )."
 It is clear that in this model the metallicity decreases in the interburst time., It is clear that in this model the metallicity decreases in the interburst time.
" Actually, the infall of primordial gas (i.e., H and He) results in a lower mass loss rate of H and He than metals, similar to the metal-enhanced wind case which will be further discussed in the next section."," Actually, the infall of primordial gas (i.e., H and He) results in a lower mass loss rate of H and He than metals, similar to the metal-enhanced wind case which will be further discussed in the next section."
" A very strong normal wind (e.g. A,> 0.5) seems unlikely in late-type dwarf galaxies since it would lose a large amount of gas, and hence it would predict a too low gas fraction, as it is evident in Fig. 8.."," A very strong normal wind (e.g. $\lambda_w>0.5$ ) seems unlikely in late-type dwarf galaxies since it would lose a large amount of gas, and hence it would predict a too low gas fraction, as it is evident in Fig. \ref{Fig:wdZmuY}."
" In the lower right panel of Fig. 8,,"," In the lower right panel of Fig. \ref{Fig:wdZmuY},"
" the predicted Y vs. (O/H) relation is shown and it is consistent with the observational data at the low metallicity, because the wind does not develop yet when the galaxy is still very metal poor."," the predicted Y vs. (O/H) relation is shown and it is consistent with the observational data at the low metallicity, because the wind does not develop yet when the galaxy is still very metal poor."
" However, after the wind, an increase of the helium abundance as well as of the abundances of elements produced on long timescale occurs, especially in the case of a strong wind which produces a very small final gas fraction."," However, after the wind, an increase of the helium abundance as well as of the abundances of elements produced on long timescale occurs, especially in the case of a strong wind which produces a very small final gas fraction."
is defined as those objects blue-wards (in both plots) of the lines in Figure 11..,is defined as those objects blue-wards (in both plots) of the lines in Figure \ref{samcut}.
 The lines eut the top of the white dwarf locus at (13) but allow slightly redder objects with higher RPMs into the sample., The lines cut the top of the white dwarf locus at $(B_{\rm J}-R)\sim1.2$ but allow slightly redder objects with higher RPMs into the sample.
 The white dwarf locus is unambiguous wards of (£3)H)-12., The white dwarf locus is unambiguous blue-wards of $(B_{\rm J}-R)\sim1.2$.
 Every object in the sample must appear in at least 15 stacks in each passhanc., Every object in the sample must appear in at least 15 stacks in each passband.
 Phe positional data as a function of time have been scrutinised lor every object selected as a white dwarf candidate. and those with dubious motions rejected.," The positional data as a function of time have been scrutinised for every object selected as a white dwarf candidate, and those with dubious motions rejected."
 While such a process may seem rather arbitrary. il was necessary (0 incorporate this screening stage in the sample extraction. because simple automated rejection algorithms such as the 30 rejection routine used here cannot be guaranteed. to eliminate spurious motions.," While such a process may seem rather arbitrary, it was necessary to incorporate this screening stage in the sample extraction because simple automated rejection algorithms such as the $\rm 3\sigma$ rejection routine used here cannot be guaranteed to eliminate spurious motions."
 Some examples are shown in Figures 12. and. 13.., Some examples are shown in Figures \ref{Rexamples} and \ref{Bexamples}.
 All three objects shown successfully satisfied: all the survey criteria., All three objects shown successfully satisfied all the survey criteria.
 Phe object plotted at the top of Figures 12 and 13. (IXN27) shows a clear. genuine motion in both x and v in both passbands anc was included. in the final sample without hesitation.," The object plotted at the top of Figures \ref{Rexamples} and \ref{Bexamples} (KX27) shows a clear, genuine motion in both x and y in both passbands and was included in the final sample without hesitation."
 The middle object (IRX18) has arger positional uncertainties and a smaller overall motion. out still shows consistent. smooth motions and was also included.," The middle object (KX18) has larger positional uncertainties and a smaller overall motion, but still shows consistent, smooth motions and was also included."
 “Phe final object shows evidence of large non-linear deviations in the last four epochs of the x measures. in roth passbands., The final object shows evidence of large non-linear deviations in the last four epochs of the x measures in both passbands.
 Although the bad-point rejection algorithm vas removed at least one datum [rom each x plot (as shown bv the multiple straight line fits). this object shows no evidence of proper motion based on the first. 16 data »oints and certainly cannot be considered a reliable proper motion object candidate.," Although the bad-point rejection algorithm has removed at least one datum from each x plot (as shown by the multiple straight line fits), this object shows no evidence of proper motion based on the first 16 data points and certainly cannot be considered a reliable proper motion object candidate."
 This object. along with 9 others. were rejected. from. the final WD. sample.," This object, along with 9 others, were rejected from the final WD sample."
 These rejected objects tended either to have large olfsets from a positional distribution otherwise consistent with zero motion at either the first or last few epochs. as is the case with the rejected object. described above: or the positional measures had an unusually laree scatter around the mean position. indicating the error in position was larger than the objects magnitude," These rejected objects tended either to have large offsets from a positional distribution otherwise consistent with zero motion at either the first or last few epochs, as is the case with the rejected object described above; or the positional measures had an unusually large scatter around the mean position, indicating the error in position was larger than the objects magnitude"
"as ""background"" and treated in LTE.",as “background” and treated in LTE.
 The statistical equilibrium equations in the RH-code are solved under the assumption of LTE populations within a single vibrational level since it would be time-consuming to solve them for each individual vibrational-rotational level (as we have to deal with more than 2000 transitions simultaneously)., The statistical equilibrium equations in the RH-code are solved under the assumption of LTE populations within a single vibrational level since it would be time-consuming to solve them for each individual vibrational-rotational level (as we have to deal with more than 2000 transitions simultaneously).
 Such an approximation is commonly used in NLTE molecular calculations (cf.mann1989:Uitenbroek 2000).. because oscillator strengths of pure rotational radiative transitions are negligibly small and radiative processes can contribute to the rotational population balance only via two-step processes (Raman seattering for the ground electronic state and emission followed by absorption for the exited state).," Such an approximation is commonly used in NLTE molecular calculations \citep[cf.][]{thompson1973, mountlinsky1974, ayreswiedemann1989, uitenbroek2000}, because oscillator strengths of pure rotational radiative transitions are negligibly small and radiative processes can contribute to the rotational population balance only via two-step processes (Raman scattering for the ground electronic state and emission followed by absorption for the exited state)."
" We have included calculations of electronic-vibrational molecular transitions with the fine structure into the RH-code in order to use it for the CN violet system computation,", We have included calculations of electronic-vibrational molecular transitions with the fine structure into the RH-code in order to use it for the CN violet system computation.
 Although the modified code can be used for a more general case. for the CN violet system we consider only diagonal vibrational bands as the Franck-Condon factors of the non-diagonal bands are much lower than those of the diagonal.," Although the modified code can be used for a more general case, for the CN violet system we consider only diagonal vibrational bands as the Franck-Condon factors of the non-diagonal bands are much lower than those of the diagonal."
 With the RH-code we compute the opacities and intensity. neglecting polarization.," With the RH-code we compute the opacities and intensity, neglecting polarization."
 Then. we use these as the input for the second code (hereafter POLY-code) written by Fluri&Stenflo(2003) and Flurietal.(2003)... which iteratively solves the polarized radiative transfer equation. taking into aecount the Hanle effect and assuming that opacities obtained in the RH-code remain unchanged (which ts à good approximation since the degree of polarization in our calculations is always lower than )).," Then, we use these as the input for the second code (hereafter POLY-code) written by \citet{fluristenflo2003} and \citet{flurietal2003}, which iteratively solves the polarized radiative transfer equation, taking into account the Hanle effect and assuming that opacities obtained in the RH-code remain unchanged (which is a good approximation since the degree of polarization in our calculations is always lower than )."
 The POLY-code was adjusted for dealing with many blended lines as initially it was designed for treating only a few non-overlaping atomic lines. while in molecular bands even a narrow spectral region can contain several hundred lines. which have to be computed simultaneously.," The POLY-code was adjusted for dealing with many blended lines as initially it was designed for treating only a few non-overlaping atomic lines, while in molecular bands even a narrow spectral region can contain several hundred lines, which have to be computed simultaneously."
 As in most NLTE-codes. the line source function in the RH-code is calculated as a function of NLTE populations. temperature and frequency (seeUitenbroek2000).," As in most NLTE-codes, the line source function in the RH-code is calculated as a function of NLTE populations, temperature and frequency \citep[see][]{uitenbroek2000}."
. The source function deviations from the LTE are defined via population departure coefficients (the ratio between the LTE and NLTE populations)., The source function deviations from the LTE are defined via population departure coefficients (the ratio between the LTE and NLTE populations).
 Such a formalism ts sufficient for calculations of the intensity field but does not provide any direct information about the contribution of scattering processes to the line., Such a formalism is sufficient for calculations of the intensity field but does not provide any direct information about the contribution of scattering processes to the line.
 However. for polarization calculations it is important to know the exact balance between scattering and thermal processes. since polarization is produced only by scattering.," However, for polarization calculations it is important to know the exact balance between scattering and thermal processes, since polarization is produced only by scattering."
 Therefore in the POLY-code the source function is calculated as à sum of scattering and thermal parts (see Eq. (10))). (10)., Therefore in the POLY-code the source function is calculated as a sum of scattering and thermal parts (see Eq. \ref{eq:lineS}) )).
 Defining these coefficients is a standard problem for a two-level system. which is described in many textbooks (e.g.Mihalas1970).," Defining these coefficients is a standard problem for a two-level system, which is described in many textbooks \citep[e.g.][]{mihalas1970}."
. However. in the case of the CN violet system the situation becomes much more complicated as we have to consider a lot of levels. many of which are collistonally and radiatively coupled with each other.," However, in the case of the CN violet system the situation becomes much more complicated as we have to consider a lot of levels, many of which are collisionally and radiatively coupled with each other."
 To find the branching coefficients let us consider the prehistory of the photon which was emitted in the /-th line. when molecule changed from the upper to the lower state.," To find the branching coefficients let us consider the prehistory of the photon which was emitted in the -th line, when molecule changed from the upper to the lower state."
" If the excited state was collisionaly populated (it could be both the collisional excitation from a lower level or de-excitation from a higher energy level) a thermal photon. which contributes As we solve the statistical equilibrium equations only for vibrational levels. the coefficients 6), and 0, are also calculated only for vibrational levels of the excited electronic state BX (adding up all significant radiative and. collisional processes. populating the considered. vibrational level)."," If the excited state was collisionaly populated (it could be both the collisional excitation from a lower level or de-excitation from a higher energy level) a thermal photon, which contributes As we solve the statistical equilibrium equations only for vibrational levels, the coefficients $\delta_{\rm th}$ and $\delta_{\rm sc}$ are also calculated only for vibrational levels of the excited electronic state $B {}^2 \Sigma$ (adding up all significant radiative and collisional processes, populating the considered vibrational level)."
 We assume that the same fractions of thermal and scattered photons óy(v) and ON) apply to every rotational level within a given rotational band (v.v).," We assume that the same fractions of thermal and scattered photons $\delta_{\rm th}({\rm v})$ and $\delta_{\rm sc}({\rm v})$ apply to every rotational level within a given rotational band (v,v)."
 This seems to be à good approximation às the rotational energy is low in comparison with the electronic energy of the excited BE state., This seems to be a good approximation as the rotational energy is low in comparison with the electronic energy of the excited $B {}^2 \Sigma$ state.
" In the case of a two-level system the branching coethcients €"" are equal to the coefficients dy and ὃς.", In the case of a two-level system the branching coefficients $\varepsilon_{\rm th}^i$ are equal to the coefficients $\delta_{\rm th} $ and $\delta_{\rm sc} $.
 However in our case of a multi-level system this equality is not valid anc the branching coethcients depend not only on the coefficients On and o... but rather on the whole balance of populating anc depopulating rates.," However in our case of a multi-level system this equality is not valid and the branching coefficients depend not only on the coefficients $\delta_{\rm th} $ and $\delta_{\rm sc} $, but rather on the whole balance of populating and depopulating rates."
 The following algorithm was employed to compute coefficients., The following algorithm was employed to compute .
 Using the NLTE populations of the /; and, Using the NLTE populations of the $l_i$ and
]t is now generally agreed. that accretion cliscs are driven mainly by magnetic torques (Shakura Sunvaev. 1973) and that the magnetic fields in such discs are maintained by local cvnamo processes.,"It is now generally agreed that accretion discs are driven mainly by magnetic torques (Shakura Sunyaev, 1973) and that the magnetic fields in such discs are maintained by local dynamo processes."
 What is not. understood however. despite considerable theoretical clforts. is how such dvnamo processes work. and exactly where and in what form the accretion energv is released (Ixing ct al..," What is not understood however, despite considerable theoretical efforts, is how such dynamo processes work, and exactly where and in what form the accretion energy is released (King et al.,"
. 2007: Blackman. 2010).," 2007; Blackman, 2010)."
 Given the Πο of this problem. we should look for observational evidence ενwhich bears on it.," Given the difficulty of this problem, we should look for observational evidence which bears on it."
 We have unique insight into the remains of one accretion clisc the protosolar nebula., We have unique insight into the remains of one accretion disc – the protosolar nebula.
 Lt is therefore worth asking if the presentclay solar system ollers evidence which can constrain the way energy is released. in accretion disces., It is therefore worth asking if the present–day solar system offers evidence which can constrain the way energy is released in accretion discs.
 Among the oldest. solid objects in the solar svstenm ave chondrules. the round. grains present in the majority of meteorites.," Among the oldest solid objects in the solar system are chondrules, the round grains present in the majority of meteorites."
 These must have formed. as molten droplets in space before being accreted., These must have formed as molten droplets in space before being accreted.
 The need to heat them sulliciently. implies a connection with energy release in the protosolar nebula., The need to heat them sufficiently implies a connection with energy release in the protosolar nebula.
 We consider here the overall energeties required of the heating process., We consider here the overall energetics required of the heating process.
 We argue that since most of the meteoritic material has been subject to this heating. the energy. source required. must. be quite. widespread: aud substantial. and must. therefore. involve the major local source of energy. Le. disc accretion.," We argue that since most of the meteoritic material has been subject to this heating, the energy source required must be quite widespread and substantial, and must therefore involve the major local source of energy, i.e. disc accretion."
. We find that. around 10 per cent of the accretion. energy ds required., We find that around 10 per cent of the accretion energy is required.
 LL so. it is evident that the existence of chondrules: provides funclamental information about energy release in accretion clises.," If so, it is evident that the existence of chondrules provides fundamental information about energy release in accretion discs."
 This paper is organised as follows., This paper is organised as follows.
 In Section 2 we give a brief introduction to chondrules and. ideas about. their formation., In Section 2 we give a brief introduction to chondrules and ideas about their formation.
 In Section 3 we consider the elobal energeties of the heating process., In Section 3 we consider the global energetics of the heating process.
 A likely mechanism for the provision of transient. heating within the solar nebula involves strong shocks (Deschl Connolly. 2002. Connolly et ab.," A likely mechanism for the provision of transient heating within the solar nebula involves strong shocks (Deschl Connolly, 2002, Connolly et al.,"
 2006)., 2006).
 In Section 4. we brielly present. the model of chondrule formation through shock heating suggested. by Deschl Connolly (2002) but argue that their proposed. mechanism [or generating these shocks by gravitational instabilitios within the dise requires dise properties which do not sit easily with our current understanding of cise evolution.," In Section 4, we briefly present the model of chondrule formation through shock heating suggested by Deschl Connolly (2002) but argue that their proposed mechanism for generating these shocks by gravitational instabilities within the disc requires disc properties which do not sit easily with our current understanding of disc evolution."
 In Section 5 we describe the dise properties which appear likely to hold at the time of chondrule formation., In Section 5 we describe the disc properties which appear likely to hold at the time of chondrule formation.
 The need. for significant clisc dissipation in the form. of shocks. together with evidence for magnetic fields within the nebula at that time (Section. 6) lead. us to a picture of a dise dvnamo (Section 7).," The need for significant disc dissipation in the form of shocks, together with evidence for magnetic fields within the nebula at that time (Section 6) lead us to a picture of a disc dynamo (Section 7)."
 This cülfers markedly from the results of current numerical simulations., This differs markedly from the results of current numerical simulations.
 Our picture draws on the ανπαπιο model of Bagealey et al. (, Our picture draws on the dynamo model of Baggaley et al. (
20092. b) and involves thin [lux ropes and reconnection. analogous to mocels of Haring on the solar surface.,"2009a, b) and involves thin flux ropes and reconnection, analogous to models of flaring on the solar surface."
 Section S is a discussion., Section 8 is a discussion.
 Alost meteorites are chondrites (more than 75 per cent. Sears Dodd. 1988: HIutchinson 2004). and the most abundant constituent of the majority of chondritio meteorite groups are chondrules (Crossman et al..," Most meteorites are chondrites (more than 75 per cent, Sears Dodd, 1988; Hutchinson 2004), and the most abundant constituent of the majority of chondritic meteorite groups are chondrules (Grossman et al.,"
 1988)., 1988).
 Chondrules are small particles of silicate material (typically around a millimetre in size. Weishere et al.," Chondrules are small particles of silicate material (typically around a millimetre in size, Weisberg et al.,"
 2006) that experienced. melting before incorporation into chondritie meteorite parent bodies, 2006) that experienced melting before incorporation into chondritic meteorite parent bodies
We observed six of the eight members of our sample in 1999 and 2000. using the VLBA. which is described by Napieretal.(1993): the two other RQQs had been observed previously with the VLBA (Blundell&Beasley1998;Dlundellοἱal.2003).,"We observed six of the eight members of our sample in 1999 and 2000, using the VLBA, which is described by \citet{nap93}; the two other RQQs had been observed previously with the VLBA \citep{blu98,blu03}."
. J12192-0638 and J1353+6345 also were imaged by Blunclell&Beasley(1998).. but we observed them al multiple lrequencies to derive spectral inlormation useful for constraining (heir emission processes.," J1219+0638 and J1353+6345 also were imaged by \citet{blu98}, but we observed them at multiple frequencies to derive spectral information useful for constraining their emission processes."
 Table 5 summarizes the VLBA observations., Table \ref{tab:obs} summarizes the VLBA observations.
 In most cases. all 10 VLBA antennas participated successhully.," In most cases, all 10 VLBA antennas participated successfully."
 However. data and/or antennas occasionally were removed cue to snowsloris. raclio-lrequency interference. or instrumentation Lailures.," However, data and/or antennas occasionally were removed due to snowstorms, radio-frequency interference, or instrumentation failures."
 All target quasars were below the flux-densitv threshold needed to solve for instrumental and atmospheric effects. so each was pliase-relerenced (Beasley&Conway1995) (0 a nearby strong compact radio source.," All target quasars were below the flux-density threshold needed to solve for instrumental and atmospheric effects, so each was phase-referenced \citep{bea95} to a nearby strong compact radio source."
 Phase-reference cvcle limes were 45 minutes. and total times ranged Irom 1 to 4 hours per frequency. band.," Phase-reference cycle times were 4–5 minutes, and total on-source times ranged from 1 to 4 hours per frequency band."
 Charged-particle effects were corrected with global ionospheric using AIPS. the Astronomical Image Processing System (Greisen2003).," Charged-particle effects were corrected with global ionospheric using AIPS, the Astronomical Image Processing System \citep{gre03}."
. Amplitudes were calibrated using standard gain files as well as svslem lemperatures measured al 12 minute intervals. (hen checked by observations of simple strong sources.," Amplitudes were calibrated using standard gain files as well as system temperatures measured at 1–2 minute intervals, then checked by observations of simple strong sources."
 Global clock offsets were found. Irom observations of strong sources. while atmospheric aud electronic drilts were calibrated using the local phase-reference sources.," Global clock offsets were found from observations of strong sources, while atmospheric and electronic drifts were calibrated using the local phase-reference sources."
 Imaging of J13164-0051 [iled. since the distance of 87 [rom the phase-velerencing source prevented successIul atmospheric calibration.," Imaging of J1316+0051 failed, since the distance of $8^\circ$ from the phase-referencing source prevented successful atmospheric calibration."
 The other quasars had. plase+velerence sources between 1.57 and 4.57 from the target., The other quasars had phase-reference sources between $1.5^\circ$ and $4.5^\circ$ from the target.
 Although their peak flux densities initially were reduced by imperfect atmospheric calibration. all had enough signal for sell-calibration. which largely eliminated the image degradation.," Although their peak flux densities initially were reduced by imperfect atmospheric calibration, all had enough signal for self-calibration, which largely eliminated the image degradation."
 Final images were made using (wo different data weishtüngs. pure “natural” weighting which maximizes the sensitivity. aud a compromise between natural and “uniform” weighting. which provides a good combination of sensitivity and resolution.," Final images were made using two different data weightings, pure “natural” weighting which maximizes the sensitivity, and a compromise between natural and “uniform” weighting, which provides a good combination of sensitivity and resolution."
 1219-0005. J13534-6345. and J1436-4-584T were observed at multiple frequencies and are unresolved at all bands. while JQ046-2-0104. also was unresolved ab its only observed. [requencey. of 4.99 GIIz.," J1219+0638, J1353+6345, and J1436+5847 were observed at multiple frequencies and are unresolved at all bands, while J0046+0104 also was unresolved at its only observed frequency of 4.99 GHz."
 JOS044+6459 contains a core wilh weak extended emission at 4.99 Gllz. but a prominent jet al 1.67 Ην.," J0804+6459 contains a core with weak extended emission at 4.99 GHz, but a prominent jet at 1.67 GHz."
 The 4.99 GIIz VLBA images of the quasars are shown in Figure 1. and a 1.67 GlIz image also is displaved [or 0504-60-00.," The 4.99 GHz VLBA images of the quasars are shown in Figure 1, and a 1.67 GHz image also is displayed for J0804+6459."
 Table 4. lists radio positions. flux densities. powers. and brightness temperatures [or each quasar.," Table \ref{tab:tb} lists radio positions, flux densities, powers, and brightness temperatures for each quasar."
 For all except J0804-2-6459. (these results come [rom Gaussian [its to the core," For all except J0804+6459, these results come from Gaussian fits to the core"
"periods. at a significance level below:SOC. were found: Ds— 2p»: D,—23p»= 3119 aud Ds=2p.","periods, at a significance level below, were found: $\rm P_3 ~ = ~ 2 P_2$ ; $\rm P_4 ~ = ~ 2/3 P_2
~ = ~ 344^d$ ; and $\rm P_5 ~ = ~ 2 P_4$."
 The data for V Uva have several properties which cause problems for this analysis., The data for V Hya have several properties which cause problems for this analysis.
 The time interval spanned by the data is only 2.1 Py. and the uaenitude estimates in the deep stellay iunuinuuni are often poorly determuned and/or lower limits.," The time interval spanned by the data is only 2.4 $\rm P_1$, and the magnitude estimates in the deep stellar minimum are often poorly determined and/or lower limits."
 The third woblem is iat the amplitudes of the variations are cdiffereut. which oeicreases the range of the data and hence the “noise” or individual period determinations.," The third problem is that the amplitudes of the variations are different, which increases the range of the data and hence the “noise” for individual period determinations."
 We therefore refined ur estimates of the periods of V Ia as follows., We therefore refined our estimates of the periods of V Hya as follows.
 First. ENutial estimates of the amplitude. period aud phase for 1e long-terii variation (G1607) were determined from the olded. data.," First, initial estimates of the amplitude, period and phase for the long-term variation $6160^d$ ) were determined from the folded data."
 This variation was approxiuate by a sine wave aud subtracted., This variation was approximated by a sine wave and subtracted.
 The residual data were heu folded o redetermine the periods., The residual data were then folded to redetermine the periods.
 The helt curve or P» was determined. aud subtracted.," The light curve for $\rm P_2$ was determined, and subtracted."
 No significant further periods were found in the residual data., No significant further periods were found in the residual data.
 The amplitude. period and phase of πιαπα lieht for Py aud P» are given iu Table 1.," The amplitude, period and phase of maximum light for $\rm P_1$ and $\rm
P_2$ are given in Table 1."
 The resulting folded light curves. for Py=6160 with the 529.19 variation subtracted aud for Py=529.| with the 61609 variation subtracted. are shown in Fies.," The resulting folded light curves, for $\rm P_1 ~ = ~ 6160^d$ with the $\rm 529.4^d$ variation subtracted and for $\rm P_2 ~ = ~ 529.4^d$ with the $\rm 6160^d$ variation subtracted, are shown in Figs."
 3 aud L., \ref{p6160} and \ref{p529}.
 The plotted πο curves are averaged iuto 50 bins per evcle aud the evele repeated several tines., The plotted light curves are averaged into 50 bins per cycle and the cycle repeated several times.
 As discussed above. the minima are not well determined.," As discussed above, the minima are not well determined."
 Fig.lL shows a fairly typical \Mra-type light curve. to first order a sine-wave variation but with a slow rise and more rapid decline.," \ref{p529} shows a fairly typical Mira-type light curve, to first order a sine-wave variation but with a slow rise and more rapid decline."
 The loug-period variation (Fig.3)) is nof approxinatelv sinusoidal however: it resembles the light curve of an eclipsing binary star. but with a far longer period aud duration.," The long-period variation \ref{p6160}) ) is not approximately sinusoidal however; it resembles the light curve of an eclipsing binary star, but with a far longer period and duration."
 The data of Mavall (1965)) show a sinular light curve shape. aud cover five previous loue-period eveles. with P4.=620082-1008. in good agreeinent with the more recent data.," The data of Mayall \cite{mayall}) ) show a similar light curve shape, and cover five previous long-period cycles, with $\rm P_1 ~ = ~ 
6200^d \pm 400^d$, in good agreement with the more recent data."
 The long-period variation in V να thus appears to be very regular., The long-period variation in V Hya thus appears to be very regular.
" Finally, we note that the amplitude of the 17-veur variation. 3.5!"" agrees with that eiven by Ἱνμοίορον et al. ("," Finally, we note that the amplitude of the 17-year variation, $\rm 3.5^m$, agrees with that given by Kholopov et al. ("
1985) aud is naller than the amplitude of 5ο sugeested by Mavall. which is actually he total range of variation due to both periodicities.,"1985) and is smaller than the amplitude of $\rm 5^m - 6^m$ suggested by Mayall, which is actually the total range of variation due to both periodicities."
 What is the origin of the GOOQT variation of V. Ilva?, What is the origin of the $\rm 6000^d$ variation of V Hya?
 The ereat depth of the unin. arc V Tva’s large mass loss rate (Paper I). suggest obscuration by dust. analogous to the ejection of obscuring material by R CrB stars: however. the dinmuninge of R CrB stars is irregular.," The great depth of the minimum, and V Hya's large mass loss rate (Paper I), suggest obscuration by dust, analogous to the ejection of obscuring material by R CrB stars; however, the dimming of R CrB stars is irregular."
 The reeularity of V να G000' variation support. rather. a dyaiunical origin for the variation.," The regularity of V Hya's $\rm 6000^d$ variation support, rather, a dynamical origin for the variation."
 We sugeest that V να is an eclipsing binary. t that he eclipse is caused not oa stellar companion but bv cireunistellar dust.," We suggest that V Hya is an eclipsing binary, but that the eclipse is caused not by a stellar companion but by circumstellar dust."
 Siuiluw phenomena are seen in a small ummber of stars of widely different spectral types and at very different evolutionary stages., Similar phenomena are seen in a small number of stars of widely different spectral types and at very different evolutionary stages.
 One such star is the FO supergiaut € Aur. which undergoes an eclipseevery27.1 vears.," One such star is the F0 supergiant $\epsilon$ Aur, which undergoes an eclipseevery27.1 years."
 There Is no secondary eclipse. and the primary eclipse is of long duration. (22 mouths) showing that the secondary. camnot," There is no secondary eclipse, and the primary eclipse is of long duration, (22 months) showing that the secondary cannot"
1e results o£ from measurements of the X-ray. luminosity function of galaxy clusters within z«0.7.,the results of from measurements of the X-ray luminosity function of galaxy clusters within $z<0.7$.
 Figure 1. shows the redshift distribution (solid line) for clusters detected above the Spectrum-RG/eCROSLLA X-ray Hux limit with mass-weighted temperatures στου 5keV. A sky coverage. fa=0.5 ds assumed.," Figure \ref{fig:distr} shows the redshift distribution (solid line) for clusters detected above the Spectrum-RG/eROSITA X-ray flux limit with mass-weighted temperatures $kT_{2500}>5$ keV. A sky coverage, $f_{\rm sky}=0.5$ is assumed."
 Approximately 5000 clusters meet these criteria from. which. following our observing strategy. 4000 will be observed by short snapshots.," Approximately $5000$ clusters meet these criteria from which, following our observing strategy, $4000$ will be observed by short snapshots."
 Assuming that ~1/8 of these clusters will also meet the relaxation criteria based on X-ray morphology(??).. a sample of ~500 hot. N-rav luminous. dynamically relaxed clusters can be defined.," Assuming that $\sim 1/8$ of these clusters will also meet the relaxation criteria based on X-ray morphology, a sample of $\sim 500$ hot, X-ray luminous, dynamically relaxed clusters can be defined."
 Taking snapshot observations of the available ~5000 clusters instead of 4000. and assuming that 1/8 of these clusters are relaxed. we will obtain a sample of ~625 ως targets.," Taking snapshot observations of the available $\sim 5000$ clusters instead of $4000$, and assuming that $\sim 1/8$ of these clusters are relaxed, we will obtain a sample of $\sim 625$ $f_{\rm gas}$ targets."
 This allows us to either use a larger sample of clusters. assume an even more conservative ratio ol relaxed. clusters. or select a different redshift’ distribution for the faa. sample of ~500 clusters.," This allows us to either use a larger sample of clusters, assume an even more conservative ratio of relaxed clusters, or select a different redshift distribution for the $f_{\rm gas}$ sample of $\sim 500$ clusters."
 In Section 5.3... we cliscuss the latter case.," In Section \ref{redshift_distr}, we discuss the latter case."
 For comparison purposes. Figure 1. also shows (dashed curve) the redshift distribution for the case of a luminosity limit of Li73.3510-2 tinthe 0.1.2.4 band (dashed line: no temperature cut is imposed).," For comparison purposes, Figure \ref{fig:distr} also shows (dashed curve) the redshift distribution for the case of a luminosity limit of $L_{\rm i}> 3.35\times 10^{44}h_{\rm 70}^{-2}$ in the $0.1-2.4$ band (dashed line; no temperature cut is imposed)."
 Ehe effect of the X-ray flux limit on the distribution is evident. towards the highest redshifts (2~ 1.5) in this case., The effect of the X-ray flux limit on the distribution is evident towards the highest redshifts $z\sim1.5$ ) in this case.
 lt is clear from that figure that the temperature and luminosity cutslead to cilferent redshift) distributions, It is clear from that figure that the temperature and luminosity cutslead to different redshift distributions.
 In the case of the temperature cut (solid line). the redshi distribution peaks around z~0.65 and relatively [ew clusters are found at z21.5.," In the case of the temperature cut (solid line), the redshift distribution peaks around $z\sim0.65$ and relatively few clusters are found at $z>1.5$."
 For the case of the luminosity cut. (dashed. line). the distribution peaks around 2~ and has many more clusters in the redshift range 1<2<2.," For the case of the luminosity cut (dashed line), the distribution peaks around $z\sim 1$, and has many more clusters in the redshift range $1<z<2$ ."
 Lis important to note. however. that a recshi distribution. weighted towards higher redshifts does not necessarily imply tighter constraints on dark energy.," It is important to note, however, that a redshift distribution weighted towards higher redshifts does not necessarily imply tighter constraints on dark energy."
" For the DIETE FoM criterion. constraints around the pivot redshi are important: for the fas. experiment 2,~(0.25 (see Figure 5))."," For the DETF FoM criterion, constraints around the pivot redshift are important; for the $f_{\rm gas}$ experiment $z_{\rm p}\sim 0.25$ (see Figure \ref{fig:evol}) )."
 In Section 5.3. we further discuss the elfect that using dillercnt redshift distributions has on the dark energy constraints., In Section \ref{redshift_distr} we further discuss the effect that using different redshift distributions has on the dark energy constraints.
 We generate mock f; measurements for 500 clusters with the redshift distribution appropriate for the case of the temperature cut solid. curve. Figure D: in accordance with the selection criteria. used for current. fii work (?)]].," We generate mock $f_{\rm gas}$ measurements for 500 clusters with the redshift distribution appropriate for the case of the temperature cut [solid curve, Figure \ref{fig:distr}; in accordance with the selection criteria used for current $f_{\rm gas}$ work ]."
 bor each cluster. we assign à statistical error in the fii measurements of 5 per cent.," For each cluster, we assign a statistical error in the $f_{\rm gas}$ measurements of $\sim 5$ per cent."
 We have also generated a set of mock measurements for the case of 250 clusters observed with fj; measurements accurate to 3.5 per cent.," We have also generated a set of mock measurements for the case of $250$ clusters observed with $f_{\rm
 gas}$ measurements accurate to $3.5$ per cent."
 This latter data set ds used. to study the impact on the dark energy constraints in the case that the fraction of suitably relaxed clusters is less than 1/8 at high redshilts., This latter data set is used to study the impact on the dark energy constraints in the case that the fraction of suitably relaxed clusters is less than $1/8$ at high redshifts.
" We stress that the predicted redshift distribution. which peaks around z~0.65 in the case of the temperature cut. has already. been probed. at. least partially, over the luminosity and temperature range of interest. by the ALAC'S survey(2): ALACS covers the redshift range 0.3<z«0.7 to a lux limit of fia,=LOI inthe0.1 2.4keV band."," We stress that the predicted redshift distribution, which peaks around $z\sim 0.65$ in the case of the temperature cut, has already been probed, at least partially, over the luminosity and temperature range of interest, by the MACS survey; MACS covers the redshift range $0.3<z<0.7$ to a flux limit of $F_{\rm
 lim}=10^{-12}$ in the $0.1-2.4\keV$ band."
 For ALACS. approximately 1/4 clusters are found to be sullicienthy relaxed. for ως work(?).," For MACS, approximately $1/4$ clusters are found to be sufficiently relaxed for $f_{\rm gas}$ work."
. Pherefore. our asstunption that 1/8 clusters detected in a future X-ray survey and meeting the X-ray [lux and leniperature criteria will be suitably relaxed. appears reasonable.," Therefore, our assumption that $\sim 1/8$ clusters detected in a future X-ray survey and meeting the X-ray flux and temperature criteria will be suitably relaxed, appears reasonable."
 Moreover. as discussed in Section 5.1. for the case of the 250-cluster sample (i.e. assuming that only ~1/16 clusters are relaxed) and using a similar total observing time to obtain individua fons measurements to 3.5 per cent accuracy. we obtain very similar dark energv constraints (see Table 4)).," Moreover, as discussed in Section 5.1, for the case of the 250-cluster sample (i.e. assuming that only $\sim 1/16$ clusters are relaxed) and using a similar total observing time to obtain individual $f_{\rm gas}$ measurements to $\sim 3.5$ per cent accuracy, we obtain very similar dark energy constraints (see Table \ref{tab:models}) )."
 A final important. point regards contaminating poin sources: for ALACS clusters. the fraction of the measure Ol24keV X-ray flux arising from contaminating poin sources is small. twpically of order a per cent (Mantz ο al.," A final important point regards contaminating point sources: for MACS clusters, the fraction of the measured $0.1-2.4\keV$ X-ray flux arising from contaminating point sources is small, typically of order a per cent (Mantz et al."
 2008: this is also the case for the hottest. Ad.2 Ske. relaxed clusters at. lower redshifts.)," 2008; this is also the case for the hottest, $kT_{\rm e}\gsim5$ keV, relaxed clusters at lower redshifts.)"
 Therefore. we do no expect our target clusters. which have comparable X-ray temperatures ancl luminosities. to be severely. allectecl by contaminating point sources. especially at zz 1.," Therefore, we do not expect our target clusters, which have comparable X-ray temperatures and luminosities, to be severely affected by contaminating point sources, especially at $z\lesssim 1$ ."
 This alleviates the instrumental requirements on the point spread function., This alleviates the instrumental requirements on the point spread function.
 An instrument with capabilities similar to the baseline characteristics listed in Table 1 should be capable of making significant strides in dark energy. work., An instrument with capabilities similar to the baseline characteristics listed in Table 1 should be capable of making significant strides in dark energy work.
The rest-frame far-infrared (farv-UR) thermal emission from dust. grains heated. by various sources the. cilluse interstellar radiation Ποιά (SRI) in galaxies. sites of active star formation. and a central active galactic nucleus (AGN) can dominate the spectral energy. distribution (SED) of galaxies (Soifer Neugebauer. 1991: Sanders AMirabel. 1996).,"The rest-frame far-infrared (far-IR) thermal emission from dust grains heated by various sources – the diffuse interstellar radiation field (ISRF) in galaxies, sites of active star formation, and a central active galactic nucleus (AGN) – can dominate the spectral energy distribution (SED) of galaxies (Soifer Neugebauer 1991; Sanders Mirabel 1996)."
 The most luminous galaxy apparent in the Universe (APALOOS279|5255: Lewin et 11998). emits approximately GOpper cent of its bolometric Luminosity in the far-LR.o waveband. while low-redshilt: galaxies with blue optical colours that were detected. by theLRAS satellite also release about GOpper cent of their total bolometric luminosity as thermal raciation from cust (Alazzarella Balzano 1986).," The most luminous galaxy apparent in the Universe 08279+5255; Irwin et 1998) emits approximately per cent of its bolometric luminosity in the far-IR waveband, while low-redshift galaxies with blue optical colours that were detected by the satellite also release about per cent of their total bolometric luminosity as thermal radiation from dust (Mazzarella Balzano 1986)."
 Even the most quiescent spiral galaxies such as the Alilky Way emit. of order pper cent of their total luminosity [rom dust. (Reach et 11995: Alton ct 11998: Dale et 22001: Dale Llelou 2002)., Even the most quiescent spiral galaxies such as the Milky Way emit of order per cent of their total luminosity from dust (Reach et 1995; Alton et 1998; Dale et 2001; Dale Helou 2002).
 Dust emission remains important at high redshifts., Dust emission remains important at high redshifts.
 Phe most distant quasi-stellar objects (QSOs) (Benford ct 11999: Carilli ct 22001: Isaak et 22002) ancl more typical. but still very luminous galaxies detected in submillimetre(submnm) wave surveys (Blain et 22002: Smail et 22002) emit strongly ad rest-Framoe far-LR wavelengths.," The most distant quasi-stellar objects (QSOs) (Benford et 1999; Carilli et 2001; Isaak et 2002) and more typical, but still very luminous galaxies detected in submillimetre(submm) wave surveys (Blain et 2002; Smail et 2002) emit strongly at rest-frame far-IR wavelengths."
 As compared with the rich variety. of features. in the SEDs of galaxies at near-It.. optical and ultraviolet wavelengths. the far-LR SED is simple. dominated. by a smooth pseudo-thermal continuum. emission spectrum.," As compared with the rich variety of features in the SEDs of galaxies at near-IR, optical and ultraviolet wavelengths, the far-IR SED is simple, dominated by a smooth pseudo-thermal continuum emission spectrum."
 At most about pper cent of the emitted energy is associated with spectral lines from atomic fine-structure and molecular rotational transitions (Malhotra et 11997: Lubman οἱ 11998: Combes. Maoli Omont 1099: Blain οἱ 22000).," At most about per cent of the emitted energy is associated with spectral lines from atomic fine-structure and molecular rotational transitions (Malhotra et 1997; Luhman et 1998; Combes, Maoli Omont 1999; Blain et 2000)."
 The mid-Lllt spectra of galaxies [rom 10 to fam are expected. to be significantly. more complex. especially," The mid-IR spectra of galaxies from 10 to $\mu$ m are expected to be significantly more complex, especially"
random sampling with replacement from the observed light curve.,random sampling with replacement from the observed light curve.
 This analysis shows that the probability of measuring a random signal with 20.1796 amplitude at the orbital period is296., This analysis shows that the probability of measuring a random signal with $\geq$ amplitude at the orbital period is.
". Based on Figures 3 and 4 and our statistical analysis, the doppler boosting signal is likely detected in J1630."," Based on Figures 3 and 4 and our statistical analysis, the doppler boosting signal is likely detected in J1630."
" However, the boosting signal does not provide any new physical constraints on the properties of this binary."," However, the boosting signal does not provide any new physical constraints on the properties of this binary."
" Along with J0106 and J0651, J1630 is only the third detached WD binary known to have a period shorter than an hour."," Along with J0106 and J0651, J1630 is only the third detached WD binary known to have a period shorter than an hour."
 All three were discovered in the last year as part of the ELM Survey., All three were discovered in the last year as part of the ELM Survey.
" Previously, all known systems with P«1 hr were interacting AM CVn systems."," Previously, all known systems with $P<1$ hr were interacting AM CVn systems."
 We do not see any evidence of mass transfer (no emission lines and no obvious excess continuum) in the three targets mentioned above., We do not see any evidence of mass transfer (no emission lines and no obvious excess continuum) in the three targets mentioned above.
 'The primary WDs in J0106 and J0651 both show ellipsoidal variations due to their larger size and shorter orbital periods., The primary WDs in J0106 and J0651 both show ellipsoidal variations due to their larger size and shorter orbital periods.
 These variations are extremely useful for constraining the inclination angle of the systems., These variations are extremely useful for constraining the inclination angle of the systems.
" However, J1630 does not show any significant ellipsoidal variations, and we only have an upper limit on the inclination angle due to the lack of eclipses."," However, J1630 does not show any significant ellipsoidal variations, and we only have an upper limit on the inclination angle due to the lack of eclipses."
 For the companion to avoid detection in the SDSS photometry and our spectroscopy implies that it is 210x fainter than the visible WD., For the companion to avoid detection in the SDSS photometry and our spectroscopy implies that it is $\ge10\times$ fainter than the visible WD.
" For an edge on orbit, such a companion would have M=0.30Mo,, Ter<7500 K, and R«€0.021 Rociteppanei07.."," For an edge on orbit, such a companion would have $M=0.30$, $T_{\rm eff}\leq7500$ K, and $R\leq0.021$ \\citep{panei07}."
 A total eclipse would be <70% deep and last for about 1.8 minutes., A total eclipse would be $\leq$ deep and last for about 1.8 minutes.
 T'he lack of eclipses in the photometry constrain the inclination angle to i<82°., The lack of eclipses in the photometry constrain the inclination angle to $i\leq82^{\circ}$.
" Hence, J1630 is best explained by a binary system containing a 0.30 WWD with a M 20.30 WWD companion at a separation of 20.32Ro."," Hence, J1630 is best explained by a binary system containing a 0.30 WD with a $M\geq$ 0.30 WD companion at a separation of $\geq$ 0.32."
. 'The two WDs in the J1630 binary will merge in «31 Myr due to gravitational wave radiation (Landau&Lifshitz 1958)., The two WDs in the J1630 binary will merge in $\leq$ 31 Myr due to gravitational wave radiation \citep{lan58}.
". After J0651, J1630 is currently the second quickest WD merger system known."," After J0651, J1630 is currently the second quickest WD merger system known."
" When the mass transfer starts, J1630 will most likely have unstable mass transfer due to the mass ratio of its components being close to unity (Marshetal. 2004)."," When the mass transfer starts, J1630 will most likely have unstable mass transfer due to the mass ratio of its components being close to unity \citep{marsh04}."
". However, the merger outcome is uncertain because of the unknown inclination angle and the companion mass."," However, the merger outcome is uncertain because of the unknown inclination angle and the companion mass."
" If the companion is another He-core WD, the merger will likely create a single He-burning subdwarf in 23-31 Myr."," If the companion is another He-core WD, the merger will likely create a single He-burning subdwarf in 23-31 Myr."
" If the companion is a more massive carbon/oxygen core WD, the system will merge in <23 Myr to form a rapidly rotating massive WD (seeKilicetal.2010,forotherpossibilites).."," If the companion is a more massive carbon/oxygen core WD, the system will merge in $\leq$ 23 Myr to form a rapidly rotating massive WD \citep[see][for other possibilites]{kilic10}."
 Short period binary WDs are important gravitational wave sources., Short period binary WDs are important gravitational wave sources.
" For example, the 12 minute orbital period binary J0651 should be detected by LISA within its first week of operation (Brownetal.2011c;Nelemans 2004)."," For example, the 12 minute orbital period binary J0651 should be detected by LISA within its first week of operation \citep{brown11c,nelemans04}."
". Similarly, the previously discovered 39 min orbital period system J0106 may be detected by LISA after 1 yr of observations."," Similarly, the previously discovered 39 min orbital period system J0106 may be detected by LISA after 1 yr of observations."
" With an orbital period similar to J0106, J1630 is also a promising candidate for detection."," With an orbital period similar to J0106, J1630 is also a promising candidate for detection."
" For an average inclination angle of 60? and model-dependent distance of 0.7 kpc, we expect the gravitational wave strain at Earth logh——22.0 at a frequency logv (Hz) — —3.08 "," For an average inclination angle of $^{\circ}$ and model-dependent distance of 0.7 kpc, we expect the gravitational wave strain at Earth $\log h = -22.0$ at a frequency $\log \nu$ (Hz) = $-3.08$ "
accreted.,accreted.
 We find that this masimaun density always mereases with ο and its functional dependence is clearly greater than linear., We find that this maximum density always increases with $a$ and its functional dependence is clearly greater than linear.
 The specific values we obtain can be found in the corresponding captions of Figs., The specific values we obtain can be found in the corresponding captions of Figs.
 1 to L., 1 to 4.
 Typical density cuhancements in the post-shock region (BL coordinates) with respect to the asvinptotic density rauge in between 1.65 (=0) aud 1.77 (0=0.99) (in logarithmic scale, Typical density enhancements in the post-shock region (BL coordinates) with respect to the asymptotic density range in between $1.65$ $a=0$ ) and $1.77$ $a=0.99$ ) (in logarithmic scale).
 Now we turn to the description of the flow morphology for different values of 5. the fluid adiabatic exponent.," Now we turn to the description of the flow morphology for different values of $\gamma$, the fluid adiabatic exponent."
 The accretion patterus for models 6 (5=1/3). Lis =5/3) and 7 (5=2) are depicted in Fig.," The accretion patterns for models 6 $\gamma=4/3$ ), 4 $\gamma=5/3$ ) and 7 $\gamma=2$ ) are depicted in Fig."
 5., 5.
 Once more. the variable we show in this figure is the logaritlin of the scaled restauass density.," Once more, the variable we show in this figure is the logarithm of the scaled rest-mass density."
 Clearly visible in this plot are the larger shock- opening angeles for the larger values of >., Clearly visible in this plot are the larger shock opening angles for the larger values of $\gamma$.
" This is explained by the eulianced values of the pressure inside the shock ""cone as 5 mereases;"," This is explained by the enhanced values of the pressure inside the shock “cone"" as $\gamma$ increases."
 We already noticed this behaviour in the non-rotating simmlatious performed in FI98a.b. Now. the larger values of 5. combined with the rapid rotation of the black hole («= 0.99). wrap the upper shock wave around the accretor.," We already noticed this behaviour in the non-rotating simulations performed in FI98a,b. Now, the larger values of $\gamma$, combined with the rapid rotation of the black hole $a=0.99$ ), wrap the upper shock wave around the accretor."
 This effect is more pronounced for the larger 5 values., This effect is more pronounced for the larger $\gamma$ values.
 We also note that the lower shock wave is less affected by the increase in 5., We also note that the lower shock wave is less affected by the increase in $\gamma$.
 While it still opens to larger angeles. the existing rotational flow counteracts the effects of the pressure force keeping its position almost unchanged.," While it still opens to larger angles, the existing rotational flow counteracts the effects of the pressure force keeping its position almost unchanged."
 The cuhancement of the pressure in the post-shock zone is responsible for the so-called “drag” force experienced by the accretor., The enhancement of the pressure in the post-shock zone is responsible for the so-called “drag” force experienced by the accretor.
 We notice here that the rotating black hole is redistributing the high pressure area. with non-trivial effects on the nature of the drag force.," We notice here that the rotating black hole is redistributing the high pressure area, with non-trivial effects on the nature of the drag force."
 Whereas in the Sclavarzschild case the drag force is alligned with the flow lines. poiutiug in the upstream direction. in the Ier case we notice a distinct asviuuetry between the co-rotating aud counterrotating side of the flow.," Whereas in the Schwarzschild case the drag force is alligned with the flow lines, pointing in the upstream direction, in the Kerr case we notice a distinct asymmetry between the co-rotating and counter-rotating side of the flow."
 The pressure cubancement is predominantly on the counterrotating side., The pressure enhancement is predominantly on the counter-rotating side.
 In Fie., In Fig.
 6 this observation is made more precise with the exaiination of the pressure profile. at the imuermost radius. for the ο=0.99 case.," 6 this observation is made more precise with the examination of the pressure profile, at the innermost radius, for the $a=0.99$ case."
 Three differcut ~ values are illustrated. showing the strong dependence of the pressure asvuuuetry on the adiabatic iudex.," Three different $\gamma$ values are illustrated, showing the strong dependence of the pressure asymmetry on the adiabatic index."
 This is particularly clear in the limiting case >=2 (dashed line)., This is particularly clear in the limiting case $\gamma=2$ (dashed line).
 We observe a pressure difference of alinost two orders of magnitude. aloug the axis normal to the asviuuptotic flow direction.," We observe a pressure difference of almost two orders of magnitude, along the axis normal to the asymptotic flow direction."
 The implication of this asviuuetry is that a rotating hole moving accross the interstellar media (or accreting from a wind). will experience. ou top of the drag force. a “litt” force. normal to its direction of motion (to the wind direction).," The implication of this asymmetry is that a rotating hole moving accross the interstellar medium (or accreting from a wind), will experience, on top of the drag force, a “lift” force, normal to its direction of motion (to the wind direction)."
" It is interesting to note that this effect bears a strong superficial reseimiblauce to the so-called ""Maeuus? effect. Ίνοι, the experience of lift forces by rotating bodies inunersed im a stream flow."," It is interesting to note that this effect bears a strong superficial resemblance to the so-called “Magnus” effect, i.e., the experience of lift forces by rotating bodies immersed in a stream flow."
 There. the lift force is due to the increased speed of the flow ou the co-rotating side (due to friction with the object). aud the Increase of pressure on the couuter-rotating side Gvhich follows inunuediatelv from the Bernoulli equation).," There, the lift force is due to the increased speed of the flow on the co-rotating side (due to friction with the object), and the increase of pressure on the counter-rotating side (which follows immediately from the Bernoulli equation)."
 We note that the direction of the lift. in relation to the scuse of rotation. agrees in both contexts.," We note that the direction of the lift, in relation to the sense of rotation, agrees in both contexts."
 We caution though that the underline causes nay be very different., We caution though that the underlying causes may be very different.
 Iu the black hole case the flow is 8upersouic aud there is no bouudary laver., In the black hole case the flow is supersonic and there is no boundary layer.
 Completing the study of the broad morphology of the dow aud its dependence on the black hole spin. we exteud the value of & above AZ. choosing. iu particular. &=LIAL GQuodel 5).," Completing the study of the broad morphology of the flow and its dependence on the black hole spin, we extend the value of $a$ above $M$, choosing, in particular, $a=1.1M$ (model 5)."
 This case correspouds to accretion onto asinyilarity. Although from the theoretical point of view such objects are believed not to exist in Nature (according to thehypothesis. allphysical singularities formed bv the gravitational collapse of nousiugular. asviuptotically flat initial data. must be hidden from the exterior world mside an event horizon) we nouctheless decided to perform such a computation. in order to assess the behaviour of the code in this regime. aud to explore the extrapolation of previous simulations.," This case corresponds to accretion onto a. Although from the theoretical point of view such objects are believed not to exist in Nature (according to the, all singularities formed by the gravitational collapse of nonsingular, asymptotically flat initial data, must be hidden from the exterior world inside an event horizon) we nonetheless decided to perform such a computation, in order to assess the behaviour of the code in this regime, and to explore the extrapolation of previous simulations."
 The resulting morphology for this simulation. using NS coordinates. is plotted in Fie.," The resulting morphology for this simulation, using KS coordinates, is plotted in Fig."
 7., 7.
 We are showing isocontours of the logarithm of the rest-imass density iu a region extending LAL in the ος and y directions from the sineularity., We are showing isocontours of the logarithm of the rest-mass density in a region extending $4M$ in the $x$ and $y$ directions from the singularity.
 Iu this situation. there is an ambiguity as to where to place the iuner boundary of the domain.," In this situation, there is an ambiguity as to where to place the inner boundary of the domain."
 The closer one gets tor=0 the «ποσο the eravity becomes (with imfuite tidal forces at the singularitv)., The closer one gets to $r=0$ the stronger the gravity becomes (with infinite tidal forces at the singularity).
 This introduces important resolution requirements on the mumierical code., This introduces important resolution requirements on the numerical code.
 For this reason we chose ων=M. iu accordance with the location of the inner boundary in the maximal case à=M.," For this reason we chose $r_{min}=M$, in accordance with the location of the inner boundary in the maximal case $a=M$."
 As can be seen from Fig., As can be seen from Fig.
" 7 the flow morphology for this uocel follows the previous treud found for lower values of à (nodels 1 to D: the shock appears slightly more wrapped (around the 7—AL circle) aud the aNd rest-mass density in the rear part of the acceretor mereases,", 7 the flow morphology for this model follows the previous trend found for lower values of $a$ (models 1 to 4): the shock appears slightly more wrapped (around the $r=M$ circle) and the maximum rest-mass density in the rear part of the accretor increases.
 We compute the accretion rates of mass. radial momentum aud angular moment.," We compute the accretion rates of mass, radial momentum and angular momentum."
 The procedure of, The procedure of
1900+500 km/s. respectively.,"$1900\pm 500$ km/s, respectively."
 The rest of the lines are not resolved and likely have smaller velocity widths., The rest of the lines are not resolved and likely have smaller velocity widths.
 We made lightcurves of the dispersed spectrum from each observation in. the 0.5-10.0 keV band., We made lightcurves of the dispersed spectrum from each observation in the 0.5–10.0 keV band.
 The lighteurves are presented in Figure 5., The lightcurves are presented in Figure 5.
 Strong variability is present in. all observations on the timescale of few«100 s. but particularly in the first and second observations.," Strong variability is present in all observations on the timescale of $few \times 100$ s, but particularly in the first and second observations."
 This variability is confirmed to also be present in the lightcurves of the simultaneousRXTE lightcurves ofH 1743-322 (see 22 and 33)., This variability is confirmed to also be present in the lightcurves of the simultaneous lightcurves of H $-$ 322 (see 2 and 3).
 A “dip” feature may be present in the third observation., A “dip” feature may be present in the third observation.
 A preliminary examination of otherRXTE observations in the publie archive revealed much stronger dipping activity typical of dipping black hole binaries viewed at high inclinations (see. e.g.. Kuulkers et 11998).," A preliminary examination of other observations in the public archive revealed much stronger dipping activity typical of dipping black hole binaries viewed at high inclinations (see, e.g., Kuulkers et 1998)."
 To explore whether the absorption lines vary within the dip. we made spectra of observation 3 from before the dip and within the dip (see Figure 5). and fit the spectrum in the manner noted above.," To explore whether the absorption lines vary within the dip, we made spectra of observation 3 from before the dip and within the dip (see Figure 5), and fit the spectrum in the manner noted above."
 Again. the only lines apparent are the Fe XXV and Fe XXV absorption lines found in the time-averaged spectrum.," Again, the only lines apparent are the Fe XXV and Fe XXVI absorption lines found in the time-averaged spectrum."
 The resultsI of fits to the spectrum prior to the dip and within the dip are given in Table 4. and the spectra are shown in Figure 6.," The results of fits to the spectrum prior to the dip and within the dip are given in Table 4, and the spectra are shown in Figure 6."
 Both absorption lines are stronger within the dip., Both absorption lines are stronger within the dip.
 While the statistical significance of the variability in the Fe XXVI line is marginal. the variability in the strength of the Fe XXV line is clearly significant.," While the statistical significance of the variability in the Fe XXVI line is marginal, the variability in the strength of the Fe XXV line is clearly significant."
 Indeed. in observation 3. the Fe XXV line is not clearly detected in the pre-dip spectrum.," Indeed, in observation 3, the Fe XXV line is not clearly detected in the pre-dip spectrum."
 Next. we investigated whether the absorption lines in observation |. and 4 vary on the few«100 s flaring variability timescale.," Next, we investigated whether the absorption lines in observation 1, and 4 vary on the $few \times 100$ s flaring variability timescale."
 We calculated the mean count rate for observations I. 2. and 4 (see 55). produced lists of time intervals for count rates above and below the mean rate. and extracted spectra from the time intervals above and below the mear rates.," We calculated the mean count rate for observations 1, 2, and 4 (see 5), produced lists of time intervals for count rates above and below the mean rate, and extracted spectra from the time intervals above and below the mean rates."
 Those spectra were also fitted in the manner described above: the spectra are shown in Figure 6., Those spectra were also fitted in the manner described above; the spectra are shown in Figure 6.
 As was the case with the time-averaged spectrum of observation 2. the count-rate selected spectra of observation. 2. show no absorptior lines.," As was the case with the time-averaged spectrum of observation 2, the count-rate selected spectra of observation 2 show no absorption lines."
 In observation 4. the strength of the lines does not vary significantly between the high rate and low rate spectra.," In observation 4, the strength of the lines does not vary significantly between the high rate and low rate spectra."
 Ii observation 1. however. the absorption lines are marginally stronger in the low rate spectra than in the high rate spectra.," In observation 1, however, the absorption lines are marginally stronger in the low rate spectra than in the high rate spectra."
 This raises the interesting possibility that the absorption may vary on a timescale of few«100 s. To establish this more clearly. we created a list of time intervals in which the source count rate was more than 5 counts/s above the mean rate. and more than 5 counts/s below the mean rate. and extracted spectra from these time intervals.," This raises the interesting possibility that the absorption may vary on a timescale of $few \times 100$ s. To establish this more clearly, we created a list of time intervals in which the source count rate was more than 5 counts/s above the mean rate, and more than 5 counts/s below the mean rate, and extracted spectra from these time intervals."
 The resultant spectra are shown in Figure 7., The resultant spectra are shown in Figure 7.
 The Fe XXV absorption line ts clearly stronger in the low count rate spectrum than in the high count rate spectrum. and the variation is statistically significant (see Table 4).," The Fe XXV absorption line is clearly stronger in the low count rate spectrum than in the high count rate spectrum, and the variation is statistically significant (see Table 4)."
 To interpret the iron absorption lines measured in the time-averaged spectra in more detail. we computed line profiles from a spherical wind with a photoionization code.," To interpret the iron absorption lines measured in the time-averaged spectra in more detail, we computed line profiles from a spherical wind with a photoionization code."
 We used the atomic physics packages from the X-ray illuminated accretion disk models of Raymond (1993)., We used the atomic physics packages from the X-ray illuminated accretion disk models of Raymond (1993).
 In this case. since the gas is optically thin. it is only necessary to consider a single slab.," In this case, since the gas is optically thin, it is only necessary to consider a single slab."
 We used the spectral parameters taken from Table 5 for each of the 4 observations and specify a density. slab thickness and distance from the central object. to model a 1-d. optically-thin wind.," We used the spectral parameters taken from Table 5 for each of the 4 observations and specify a density, slab thickness and distance from the central object, to model a 1-d, optically-thin wind."
 The code iterates to find a self-consistent temperature and ionization state. then computes equivalent widths assuming a Doppler profile with widths ranging from 100 to 2000 kms.," The code iterates to find a self-consistent temperature and ionization state, then computes equivalent widths assuming a Doppler profile with widths ranging from 100 to 2000 $\rm
km~s^{-1}$."
 It assumes 1onization equilibrium. which 15 a good approximation for the densities and dynamical times estimated below in all cases.," It assumes ionization equilibrium, which is a good approximation for the densities and dynamical times estimated below in all cases."
 From the output of the code we find the parameters that match the observed Fe XXV and Fe XXVI absorption line equivalent widths., From the output of the code we find the parameters that match the observed Fe XXV and Fe XXVI absorption line equivalent widths.
 The Fe XXV and Fe XXVI lines cannot originate in exactly the same gas. because the line widths and line shifts differ.," The Fe XXV and Fe XXVI lines cannot originate in exactly the same gas, because the line widths and line shifts differ."
 A plausible physical picture might consist of denser. slower clumps containing more Fe XXV embedded in a less dense gas where Fe XXVI dominates.," A plausible physical picture might consist of denser, slower clumps containing more Fe XXV embedded in a less dense gas where Fe XXVI dominates."
 However. a modest density contrast of order 2 (approximately) is sufficient to shift the balance between Fe XXV and Fe XXVL so an average density that produces both lines is meaningful. and we do not have enough measured parameters to justify the additional free parameters of a more complex model (the line widths are uncertain).," However, a modest density contrast of order 2 (approximately) is sufficient to shift the balance between Fe XXV and Fe XXVI, so an average density that produces both lines is meaningful, and we do not have enough measured parameters to justify the additional free parameters of a more complex model (the line widths are uncertain)."
 We therefore choose densities to match the observed equivalent widths and Doppler velocity widths intermediate between those of Fe XXV and Fe XXVI., We therefore choose densities to match the observed equivalent widths and Doppler velocity widths intermediate between those of Fe XXV and Fe XXVI.
 There is à maximum radius at which a slab can plausibly be causing the observed absorption., There is a maximum radius at which a slab can plausibly be causing the observed absorption.
 The Fe XXVI/XXV line flux ratio gives an ionization parameter (&=Ly/m7). and for any given distance from the central source r. a specific density 7118 required.," The Fe XXVI/XXV line flux ratio gives an ionization parameter $\xi = L_{X} / nr^{2}$ ), and for any given distance from the central source $r$, a specific density $n$ is required."
 A specific slab thickness. which cannot exceed r. is required to produce the measured line column densities. so that Nx:nr.," A specific slab thickness, which cannot exceed $r$, is required to produce the measured line column densities, so that $N
\leq nr$."
 Because the density that gives the tonization parameter scales as 7/7. the slab thickness scales as 777 (therefore Nj7! ). and there is a maximum radius at which a slab of thickness r can reproduce the lines observed.," Because the density that gives the ionization parameter scales as $r^{-2}$, the slab thickness scales as $r^{-2}$ (therefore $N~r^{-1}$ ), and there is a maximum radius at which a slab of thickness $r$ can reproduce the lines observed."
 We call this parameter Fray., We call this parameter $r_{max}$.
 The lines can be formed at smaller r. and if the lines are formed in a wind with ~~ density profile. the density at ως is correspondingly smaller.," The lines can be formed at smaller $r$, and if the lines are formed in a wind with $r^{-2}$ density profile, the density at $r_{max}$ is correspondingly smaller."
 Table 5 shows the derived parameters., Table 5 shows the derived parameters.
 We can also compute the mass loss rate in a wind of any velocity. because 7/77 is constant.," We can also compute the mass loss rate in a wind of any velocity, because $n r^{-2}$ is constant."
 The velocity shifts in Table 3 are generally uncertain. so We present mass loss rates for a reference wind speed of 300 kmsl.," The velocity shifts in Table 3 are generally uncertain, so we present mass loss rates for a reference wind speed of 300 $\rm km~s^{-1}$."
 The table shows that the lack of absorption lines in the second observation cannot entirely be attributed to a higher ionizing flux., The table shows that the lack of absorption lines in the second observation cannot entirely be attributed to a higher ionizing flux.
 Although the ionizing flux in this observation is about double that of the first observation. an order of magnitude lower column density of absorbing material is needed to match the upper limits to the equivalent widths.," Although the ionizing flux in this observation is about double that of the first observation, an order of magnitude lower column density of absorbing material is needed to match the upper limits to the equivalent widths."
 An interesting result in Table 5 ts the small value of ως for the 4th observation. which results primarily from the small ionizing flux.," An interesting result in Table 5 is the small value of $r_{max}$ for the 4th observation, which results primarily from the small ionizing flux."
 The lines do not show a significant Doppler shift. but they arise in a white dwarf-size region near the central source.," The lines do not show a significant Doppler shift, but they arise in a white dwarf-size region near the central source."
" In the first and third observations. the absorbing regions have a size comparable to a few solar radit, which is similar to the separation of the components m several binary black hole systems."," In the first and third observations, the absorbing regions have a size comparable to a few solar radii, which is similar to the separation of the components in several binary black hole systems."
 The second interesting aspect of the table is the indication that the observed absorption lines are formed in a wind with a mass loss rate of order 2«107?M..yr! — comparable to the inferred mass accretion rate., The second interesting aspect of the table is the indication that the observed absorption lines are formed in a wind with a mass loss rate of order $2 \times 10^{-8}~M_{\odot}~{\rm yr}^{-1}$ — comparable to the inferred mass accretion rate.
 If we assume a distance of 8.5 kpe. take the highest inferred 0.5—10.0 keV luminosity of Ly26.8\10°8eres! and assume an efficiency of in the equation Ly=aya. We obtain my.=7.6«1075es!.," If we assume a distance of 8.5 kpc, take the highest inferred 0.5–10.0 keV luminosity of $L_{X} = 6.8
\times 10^{38}~{\rm erg}~{\rm s}^{-1}$ and assume an efficiency of in the equation $L_{X} = \eta \dot{m}_{acc} c^{2}$, we obtain $\dot{m}_{acc} = 7.6 \times 10^{18}~{g}~{\rm s}^{-1}$."
 Taking the highest measured 0.5-10.0 keV luminosity to be a lower limit on the true Eddington luminosity. the highest inferred mass loss in the wind Giving=1.7«1015es ) corresponds to an outflow rate that is of the Eddington mass aceretion rate (for a unity filling factor).," Taking the highest measured 0.5–10.0 keV luminosity to be a lower limit on the true Eddington luminosity, the highest inferred mass loss in the wind $\dot{m}_{wind} = 1.7 \times 10^{18}~{\rm g}~{\rm s}^{-1}$ ) corresponds to an outflow rate that is of the Eddington mass accretion rate (for a unity filling factor)."
 Note. however. that the fraction depends crucially on the filling factor. the energy range used. and the," Note, however, that the fraction depends crucially on the filling factor, the energy range used, and the"
are comparalivelv small because of the low75.,are comparatively small because of the low.
. The dashed lime encloses the points for the same model with of the dust mass in a constant density: the scattered points are lower (there is more scattering) because there are no almost empty. paths through the nebula., The dashed line encloses the points for the same model with of the dust mass in a constant density; the scattered points are lower (there is more scattering) because there are no almost empty paths through the nebula.
 The dot-dashed and dotted lines are the boundaries of the points in Figures |. and 2.. which ave lor (e.gy.D.τι) = (0.6. 0.6. 2.6. 2).," The dot-dashed and dotted lines are the boundaries of the points in Figures \ref{fig1} and \ref{fig2}, which are for $a,\,g,\,D,\,\tau_0)$ = (0.6, 0.6, 2.6, 2)."
 The long dashed lines enclose (le very. wide boundaries for (a.g.D.τι) = (0.8. 0.35. 2.3. 4).," The long dashed lines enclose the very wide boundaries for $a,\,g,\,D,\,\tau_0)$ = (0.8, 0.85, 2.3, 4)."
observations. However. we will see that a minimun « of 70.5 is required to produce enough scattered light.," However, we will see that a minimum $a$ of $\sim$ 0.5 is required to produce enough scattered light."
 Figure 4 shows theaveraged(ru). (Teen))) values for all of the 21 different initial seeds that we tried.," Figure \ref{fig4} shows the, ) values for all of the 21 different initial seeds that we tried."
 Figure dashowspurelyhierarchicalmodels: 4b. models with of the dust in a constant distribution.," Figure \ref{fig4}$ $a$ shows purely hierarchical models; \ref{fig4}$ $b$, models with of the dust in a constant distribution."
 The parameters are (he same as in Figures | and 2: 0 = gy = 0.6. = 2.," The parameters are the same as in Figures \ref{fig1} and \ref{fig2}: $a$ = $g$ = 0.6, = 2."
 The points are as belore., The points are as before.
 The open squares are models with D = 2.3 instead of 2.6., The open squares are models with $D$ = 2.3 instead of 2.6.
 Within each panel. the D = 2.3 models are higher because the radiation. both scaltered ancl stellar. can escape more easily [rom the more stronely clamped structure.," Within each panel, the $D$ = 2.3 models are higher because the radiation, both scattered and stellar, can escape more easily from the more strongly clumped structure."
 The lines show uniform models. with albedos marked in 4a.," The lines show uniform models, with albedos marked in $a$."
 The uniform model appropriate to the actual albedo assumed in the models (α = 0.6) is just below the bottom of the figure., The uniform model appropriate to the actual albedo assumed in the models $a$ = 0.6) is just below the bottom of the figure.
 The dashed line is the uniform model with « = 0.5 and g = 0.4 instead of 0.6., The dashed line is the uniform model with $a$ = 0.5 and $g$ = 0.4 instead of 0.6.
 The differences are nol large in comparison to the effects of the other parameters., The differences are not large in comparison to the effects of the other parameters.
" The ellects on hierarchical models of changingSills, 5gg are similar.", The effects on hierarchical models of changing $g$ are similar.
 Perhaps the most striking difference between the two panels is the lower values of [for hierarchical models with of the dust in a uniform component., Perhaps the most striking difference between the two panels is the lower values of for hierarchical models with of the dust in a uniform component.
 The increase of scattering from the dust between the clumps causes the decrease in and greatly reduces the differences between models with dillerent spatial distributions of dust., The increase of scattering from the dust between the clumps causes the decrease in and greatly reduces the differences between models with different spatial distributions of dust.
 In either panel. ihe points from various viewing angles of (he (wo values of D are completely intertwined.," In either panel, the points from various viewing angles of the two values of $D$ are completely intertwined."
 Itellection nebulae are poor diagnostics of D as well as other properties of the ISM., Reflection nebulae are poor diagnostics of $D$ as well as other properties of the ISM.
 The three open circles in Figure 4 are with (μου values of the initial seed., The three open circles in Figure \ref{fig4} are with three values of the initial seed.
 Two of these values show the extrema in((74).. (744))) among the 21 initial seeds that we tested.," Two of these values show the extrema in, ) among the 21 initial seeds that we tested."
 The üghiness of the mean optical depths of the models shows the importance of hierarchical geometry. as opposed to simple Gvo-phase chuups.," The tightness of the mean optical depths of the models shows the importance of hierarchical geometry, as opposed to simple two-phase clumps."
 The of the mnocdels are similar to those of the hierarchical models with uniform dust., The of the models are similar to those of the hierarchical models with uniform dust.
 The contrast of both with purely hierarchical models illustrates the importance ofdust-Iree regions (in real space. possibly caused by extensions of hot. low-density material into the nebulae).," The contrast of both with purely hierarchical models illustrates the importance ofdust-free regions (in real space, possibly caused by extensions of hot, low-density material into the nebulae)."
" Figure 4 shows what we meant bv saving thal Figure 1 was produced by a typical hierarchical model | one with a tvpical (7,4).", Figure \ref{fig4} shows what we meant by saying that Figure \ref{fig1} was produced by a typical hierarchical model – one with a typical .
. Those with large have the star embedded within dusty material. and a relatively low central dustdensity leads to a low (Tox).," Those with large have the star embedded within dusty material, and a relatively low central dustdensity leads to a low ."
broad-band (grit) magnitudes of each individual. galaxy with one of a set of BCOS model template spectra. from which the g and i-band. k-corrections are then measured.,"broad-band $ugriz$ ) magnitudes of each individual galaxy with one of a set of BC03 model template spectra, from which the $g$ and $r$ -band k-corrections are then measured."
 We also compare with two simpler models in which the same k-corrections are assumed. for all the I2/80s: firstly. our evolving DC03 model 2 (age 12 Gyr τς1 Gyr) and secondly. the non-evolving for ellipticals from Coleman. Wu and Weedman (1980. hereafter. CYWWN).," We also compare with two simpler models in which the same k-corrections are assumed for all the E/S0s: firstly, our evolving BC03 model 2 (age 12 Gyr $\tau=1$ Gyr) and secondly, the non-evolving for ellipticals from Coleman, Wu and Weedman (1980, hereafter CWW)."
 Figure 13. shows the dilferent AyA against redshift (Table 1&22 list. the different &-corrections used in the Figure)., Figure \ref{kgkr} shows the different $k_g-k_r$ against redshift (Table 2 list the different $k$ -corrections used in the Figure).
 Firstly. note that k-corrections computed from: and magnitudes dilfer (primarily at z< 0.15) by [ess than 0.01 mag.," Firstly, note that k-corrections computed from and magnitudes differ (primarily at $z < 0.15$ ) by less than 0.01 mag."
 t z«0.15 the spectra and Blanton-Rowcis Ayare all below our DC Model 4. but the uncorrected Ay crosses all the models to reach Model 1: at z>0.3.," At $z<0.15$ the spectra and Blanton-Roweis $k_g-k_r$are all below our BC Model 4, but the uncorrected $k_g-k_r$ crosses all the models to reach Model 1 at $z>0.3$."
 1C COLTLOCCIOCC Vevp 18 ο... Ugh Lrectsht alc 16 331011OD-T00WOCLIS /i OEC LONOCD SN and 5Sive ο Leas evolution.," The corrected $k_g-k_r$ is lower at high redshift, and the Blanton-Roweis $k_g-k_r$ are lower still, and give the least evolution."
" The CWW £A,Ay is always more positive than that from the spectra. and so gives an overestimate of r) and too-blue rest-frame colours."," The CWW $k_g-k_r$ is always more positive than that from the spectra, and so gives an overestimate of $\Delta(g-r)$ and too-blue rest-frame colours."
 Neither the corrected or uncorrected Ay follow any of the BCOS models over the whole redshift range., Neither the corrected or uncorrected $k_g-k_r$ follow any of the BC03 models over the whole redshift range.
 This is to some extent because the mean luminosity increases with redshift., This is to some extent because the mean luminosity increases with redshift.
 Po examine the effect of this. Figure 14 shows mean Ay with the galaxies divided into four intervals of AL).," To examine the effect of this, Figure \ref{kgkr-fixedM} shows mean $k_g-k_g$ with the galaxies divided into four intervals of $M_r$ ."
 Ata given redshift the more luminous galaxies do tend to have a slightly more positive kyAy., At a given redshift the more luminous galaxies do tend to have a slightly more positive $k_g-k_r$.
 However. the galaxy Ay within each of the four luminosity intervals still do not closely follow any of the BCOS mocdels. tending to be lower at 0.05«z0.2. indicating systematic. dilferences al some wavelengths between the spectra and. this set of models.," However, the galaxy $k_g-k_r$ within each of the four luminosity intervals still do not closely follow any of the BC03 models, tending to be lower at $0.05<z<0.2$, indicating systematic differences at some wavelengths between the spectra and this set of models."
 We compared the spectroscopically-based. €MIS with the CAMs from photometric colours &-corrected. using only model-based or template spectra., We compared the spectroscopically-based CMRs with the CMRs from photometric colours $k$ -corrected using only model-based or template spectra.
 Phe top panel in Figure 15 shows the effect. of fitting Blanton Roweis (2007) À-corrections tofiber magnitudes. and using the observed g r.," The top panel in Figure \ref{cmr-br} shows the effect of fitting Blanton Roweis (2007) $k$ -corrections to magnitudes, and using the observed $g-r$."
 The CAIR is best-fit with The slope is very similar to that in the bottom panel ofFigure 10.. but with these rather than the spectra-based. &-the evolution is 0.092 weaker.," The CMR is best-fit with The slope is very similar to that in the bottom panel ofFigure \ref{cmr-fiber}, , but with these rather than the spectra-based $k$ -corrections,the evolution is $0.09z$ weaker."
 A similar analysis of the colour (Blanton Roweis A-corrections from, A similar analysis of the colour (Blanton Roweis $k$ -corrections from
In some applications. inclicding the present one. it is more convenient to work in terms of matter and Iuminositw densities.,"In some applications, including the present one, it is more convenient to work in terms of matter and luminosity densities."
 This is especially Cie case when thedensity. j. is known for a given band. rather than (he mass-to-light ratio.," This is especially the case when the, $j$, is known for a given band, rather than the mass-to-light ratio."
 We have which gives The luminositw densitw j for à particular band of the survey can be calculated for example from the luminosity funetion 9$(L) of galaxies in this band citealt De93)) Note that using the Relation (9)) in ((1)). we get the linear-theorv velocity measured [rom the {lux dipole as where g denotes the acceleration of the LG. in units of velocity.," We have which gives The luminosity density $j$ for a particular band of the survey can be calculated for example from the luminosity function $\Phi(L)$ of galaxies in this band \\citealt{Pe93}) ) Note that using the Relation \ref{eq:flux dipole}) ) in \ref{eq:v.and.g}) ), we get the linear-theory velocity measured from the flux dipole as where $\tilde{\bmg}$ denotes the acceleration of the LG, in units of velocity."
" The term 57,5r; is the [αν dipole moment of sources down to the zero flux over the whole skv.", The term $\sum_i S_i\hat{\bmr}_i$ is the flux dipole moment of sources down to the zero flux over the whole sky.
 The universal Iuminosity density j. measured Irom a fair sample of galaxies in the given band with known apparent Iuminosities and redshifts. is proportional to fy). which means (hat the overall result does not depend on the IIubble constant.," The universal luminosity density $j$, measured from a fair sample of galaxies in the given band with known apparent luminosities and redshifts, is proportional to $H_0$, which means that the overall result does not depend on the Hubble constant."
 Realistic galaxy catalogs will never reach down to zero flux. irrespectivelv of the used wavelength.," Realistic galaxy catalogs will never reach down to zero flux, irrespectively of the used wavelength."
 On the contrary. survevs are usually/ic-lanited.. which means that the number of observed sources. Nis finite.," On the contrary, surveys are usually, which means that the number of observed sources, $N$, is finite."
 For that reason. in the following we will denote the flux dipole of a finite sample as d: Note that the clustering dipole calculated for a finite. [Iux-limited sample may be a biased estimator of the peculiar acceleration of the Local Group. ((2)).," For that reason, in the following we will denote the flux dipole of a finite sample as $\bmd$: Note that the clustering dipole calculated for a finite, flux-limited sample may be a biased estimator of the peculiar acceleration of the Local Group, \ref{eq:g.theor}) )."
 This can be overcome by, This can be overcome by
"themselves as part of a synchrotron self-Compton (SSC) approach (see, e.g., Markoff et al.","themselves as part of a synchrotron self-Compton (SSC) approach (see, e.g., Markoff et al."
" 2001, Liu Melia 2001, Eckart et al."," 2001, Liu Melia 2001, Eckart et al."
 2004)., 2004).
" It has also been suggested that both the NIR and the X-ray components may be due to the same synchrotron process (see, e.g., Yuan et al."," It has also been suggested that both the NIR and the X-ray components may be due to the same synchrotron process (see, e.g., Yuan et al."
 2004)., 2004).
 But the analysis of several bright flares detected over the past few years has all but ruled out these earlier proposals (see Dodds-Eden et al., But the analysis of several bright flares detected over the past few years has all but ruled out these earlier proposals (see Dodds-Eden et al.
" 2009, Trap et al."," 2009, Trap et al."
 2011)., 2011).
 Single component synchrotron models (with a particle distribution dN(y)eΥ5dy) are problematic because the high-energy electrons required to generate X-ray synchrotron emission have very short cooling times (much shorter than a typical X-ray flare duration).," Single component synchrotron models (with a particle distribution $dN(\gamma)
\propto {\gamma}^{-p}\,d\gamma$ ) are problematic because the high-energy electrons required to generate X-ray synchrotron emission have very short cooling times (much shorter than a typical X-ray flare duration)."
" Thus, even a continuous injection to replenish the energetic population cannot sustain the same power-law index p at both low and high energies."," Thus, even a continuous injection to replenish the energetic population cannot sustain the same power-law index $p$ at both low and high energies."
" If instead the X-rays are submm photons inverse-Compton scattered by the electrons producing the NIR emission, the largest permissible size of the quiescent radio-producing region is about 0.27Rs (Dodds-Eden et al."," If instead the X-rays are submm photons inverse-Compton scattered by the electrons producing the NIR emission, the largest permissible size of the quiescent radio-producing region is about $0.27\,R_S$ (Dodds-Eden et al."
" 2009, Trap et al."," 2009, Trap et al."
" 2011), far smaller than the measured FWHM size (53.7R5; Doeleman et al."," 2011), far smaller than the measured FWHM size $\approx
3.7\,R_S$; Doeleman et al."
 2008) of Sgr A* at 1.3 mm., 2008) of Sgr A* at 1.3 mm.
" This uncomfortably tight restriction (see also Liu Melia 2001) is compounded by several other difficulties, but even the size issue on its own already rules out the “external” Compton scattering scenario."," This uncomfortably tight restriction (see also Liu Melia 2001) is compounded by several other difficulties, but even the size issue on its own already rules out the “external"" Compton scattering scenario."
" Attempts at fitting an SSC spectrum to the combined NIR/X-ray data are equally problematic because such a model requires low electron energies (with y~ 10-15) and unrealistically strong magnetic fields (B>1,000 G) and very large particle densities (much larger than those inferred for the quiescent emission)."," Attempts at fitting an SSC spectrum to the combined NIR/X-ray data are equally problematic because such a model requires low electron energies (with $\gamma
\sim$ 10–15) and unrealistically strong magnetic fields $B>1,000$ G) and very large particle densities (much larger than those inferred for the quiescent emission)."
 We are thus left with the following rather tightly constrained indicators., We are thus left with the following rather tightly constrained indicators.
" During a typical flare, there is little if any detectable emission at 11.88 jm, implying that the flare emission spectrum (characterized by the power density vF,) must rise from the MIR towards the NIR."," During a typical flare, there is little if any detectable emission at $11.88\;\mu$ m, implying that the flare emission spectrum (characterized by the power density $\nu F_\nu$ ) must rise from the MIR towards the NIR."
 This is consistent with the spectral index a~0.6 described above., This is consistent with the spectral index $\alpha\sim 0.6$ described above.
" It is clear, therefore, that the electron population producing the L’-band flare has a different distribution of energies than that associated with the submm bump (see, e.g., Melia 2007), so a NIR flare cannot simply be a small change in the overall properties of the steady radio-submm emitting region."," It is clear, therefore, that the electron population producing the $L^\prime$ -band flare has a different distribution of energies than that associated with the submm bump (see, e.g., Melia 2007), so a NIR flare cannot simply be a small change in the overall properties of the steady radio-submm emitting region."
" However, it is unrealistic to expect a power-law particle distribution such as this to maintain the same power-law index p at all energies."," However, it is unrealistic to expect a power-law particle distribution such as this to maintain the same power-law index $p$ at all energies."
" Synchrotron energy losses, not to mention the escape time from the acceleration region, both depend on the particle energy (see, e.g., Liu et al."," Synchrotron energy losses, not to mention the escape time from the acceleration region, both depend on the particle energy (see, e.g., Liu et al."
 2006)., 2006).
" It is well known (see, e.g., Pacholezyk 1970) that while energetic electrons are injected continuously into the system bythe acceleration process, the emitted steady-state photon spectrum has an index a=(3—p)/2 (with p the particle index) up to a “cooling break"" frequency v;, steepening to"," It is well known (see, e.g., Pacholczyk 1970) that while energetic electrons are injected continuously into the system bythe acceleration process, the emitted steady-state photon spectrum has an index $\alpha=(3-p)/2$ (with $p$ the particle index) up to a “cooling break"" frequency $\nu_b$ , steepening to"
phenomena ever discovered iu the corona.,phenomena ever discovered in the corona.
 The wave is nostly MIID. kink node propagating outwards along he thin plaza sheet., The wave is mostly MHD kink mode propagating outwards along the thin plasma sheet.
 The restoring force supporting he wavy notion is provided bv the magnetic feld of he streamer structure. which is ecucrated by the large streamer deflection upon the CALE impact.," The restoring force supporting the wavy motion is provided by the magnetic field of the streamer structure, which is generated by the large streamer deflection upon the CME impact."
 The energy received from the inipact is carried outwards by tlie wave rturbation., The energy received from the impact is carried outwards by the wave perturbation.
 Consequently. the amplitude of the wave rear the sun declines rapidly with time. aud only a few xeriods of the wave are observable.," Consequently, the amplitude of the wave near the sun declines rapidly with time, and only a few periods of the wave are observable."
 The wave period is estimated to be about 1 hour. the wavelength varies from 2to LR... the wave amplitude is a few tens of solar radii. and the plase speed is about 300 to 500 kins .," The wave period is estimated to be about 1 hour, the wavelength varies from 2 to 4 $_\odot$, the wave amplitude is a few tens of solar radii, and the phase speed is about 300 to 500 km $^{-1}$."
 There exists a general trend for the phase speed to decrease with increasing helioceutric distance., There exists a general trend for the phase speed to decrease with increasing heliocentric distance.
 Tuteractious between CAIE and παΟΠΟΥ5 are frequently observed. especially during the active phase of solu cycles.," Interactions between CME and streamers are frequently observed, especially during the active phase of solar cycles."
 Usually. such interactions result in apparent deflections of interacting streamers (Ce. Tndhausen ot al.," Usually, such interactions result in apparent deflections of interacting streamers (e.g., Hundhausen et al.,"
 1987: Sime IIuudhnauseu. 1987: Sheeley et al.," 1987; Sime Hundhausen, 1987; Sheeley et al.,"
 2000)., 2000).
 We emphasize that the streamer wavy motion. reported im the present study. is a direct consequence of a streamer deflection.," We emphasize that the streamer wavy motion, reported in the present study, is a direct consequence of a streamer deflection."
 Nevertheless. as revealed from a prelaminary overview of the long-erm LASCO observations. in onlv α ταν siall yaction of the deflection events the streamer exhibits wavelike phenomena.," Nevertheless, as revealed from a preliminary overview of the long-term LASCO observations, in only a very small fraction of the deflection events the streamer exhibits wavelike phenomena."
 In other words. most CME-diveu deflections. even very fast aud strong. are not followed κα streniner wavy motion.," In other words, most CME-driven deflections, even very fast and strong, are not followed by a streamer wavy motion."
 Therefore. there exist ain strict conditions for streamer waves to be excited| κα CAME-streamer deflection.," Therefore, there exist certain strict conditions for streamer waves to be excited by a CME-streamer deflection."
 Two observational eatures of the July 6 event can help us evaluate the relevant conditious., Two observational features of the July 6 event can help us evaluate the relevant conditions.
" Firstly. it is found that the CME ""OIree region lies ou the flank side of the closed loops colmprising the streamer. that meaus the CATE does not originate from beneath the streamer structure. aud the ejecta can collide with the streamer frou the fanuk sido."," Firstly, it is found that the CME source region lies on the flank side of the closed loops comprising the streamer, that means the CME does not originate from beneath the streamer structure, and the ejecta can collide with the streamer from the flank side."
 Secondly. the CME is a fast eruption with a speed of c1300 liu |. which has two consequences favoring he excitation of the streamer wave.," Secondly, the CME is a fast eruption with a speed of $\ge 1300$ km $^{-1}$, which has two consequences favoring the excitation of the streamer wave."
 One is that a aster eruption results iu a strouger inmupiugenieut on he nearby streamer and a consequent larger deflection of the σος structure from its equilibria position. he other is that the ejecta moves out of the corona in a relatively short time. and leaves cnough time or the streamer wave to develop.," One is that a faster eruption results in a stronger impingement on the nearby streamer and a consequent larger deflection of the streamer structure from its equilibrium position, the other is that the ejecta moves out of the corona in a relatively short time, and leaves enough time for the streamer wave to develop."
 Otherwise if the eruption is nof fast enoush. the deflected streamer nay simply moves backwards along with the cjecta. and uo wavy motions result.," Otherwise if the eruption is not fast enough, the deflected streamer may simply moves backwards along with the ejecta, and no wavy motions result."
 To observe one example of such a case. one inav check the online LASCO observations of the iuteraction eveut between a CALE and a streamer in the uortheaster quadrant dated ou July 9th. 2001.," To observe one example of such a case, one may check the online LASCO observations of the interaction event between a CME and a streamer in the northeastern quadrant dated on July 9th, 2004."
 Sheeley et al. (, Sheeley et al. (
2000) also presents LASCO exmuples of strong streamer deflection. events without accompanying apparcut streamer wavy motions.,2000) also presents LASCO examples of strong streamer deflection events without accompanying apparent streamer wavy motions.
 It should be noted that a more complete understanding of the excitation coucitious of the streamer wave cau only be obtained from observational investigations ou much more similar events and from elaborate theoretical modelling eudeavors., It should be noted that a more complete understanding of the excitation conditions of the streamer wave can only be obtained from observational investigations on much more similar events and from elaborate theoretical modelling endeavors.
 As inentioned in the introduction section. a well developed typical streamer consists of the main bocly. which is a bunch of closed field. arcades confining high density coronal plasimas. and a dense plasma sheet within which a lone thin current sheet i$ emibedded.," As mentioned in the introduction section, a well developed typical streamer consists of the main body, which is a bunch of closed field arcades confining high density coronal plasmas, and a dense plasma sheet within which a long thin current sheet is embedded."
 The intersection of the closed streamer main body and he open plasina sheet gives the streamer cusp. which is ecnerally thought to be below 2 to 2.5 R.. w(DAY close to the bottom of the LASCO ο) ΕΟΝ.," The intersection of the closed streamer main body and the open plasma sheet gives the streamer cusp, which is generally thought to be below 2 to 2.5 $R_\odot$, very close to the bottom of the LASCO C2 FOV."
 After he impact from a CAIE. the streamer deflects away roni its original equilibria position.," After the impact from a CME, the streamer deflects away from its original equilibrium position."
 The cousequeut restoring mmotion may excite the wavelike oscillations ILOweating along the plasma sheet., The consequent restoring motion may excite the wavelike oscillations propagating along the plasma sheet.
 Therefore. the σοςuetry supporting the discussed streamer wave motion can be simplified as a long slender plasiua slab extend o infinity with the lower end attaching to the streamer Cus» Which bounces back auc forth iLa quasi-periodici liamer.," Therefore, the geometry supporting the discussed streamer wave motion can be simplified as a long slender plasma slab extending to infinity with the lower end attaching to the streamer cusp which bounces back and forth in a quasi-periodic manner."
 The oscillations are observed to be genera Yaisverse to the nouinal direction of he maenetic ficd., The oscillations are observed to be generally transverse to the nominal direction of the magnetic field.
 The manifestation and the geometry of the pleLOMla are very simular to that of the well-suown sank uxxle deποσα from a sleπο. magnetic sla) except clue iu a spherical expanding 9eometrv (Roberts. 1981: Edwin Roberts. 1982).," The manifestation and the geometry of the phenomena are very similar to that of the well-known kink mode deduced from a slender magnetic slab except being in a spherical expanding geometry (Roberts, 1981; Edwin Roberts, 1982)."
" It is therefore μιeeesttCOONxL tiat the wave phenomenon discussed. ia this study rTOepreseuts the kink node. wuch Is. lu a 1uore eenera ποσο, a type of ast mnagnetosonide waves propagaimgo iu ali inhomogencous maguctized plasuia cuviromment."," It is therefore suggested that the wave phenomenon discussed in this study represents the kink mode, which is, in a more general sense, a type of fast magnetosonic waves propagating in an inhomogeneous magnetized plasma environment."
 It is interesting o notice that the morphology of the sYOOluecr wave discussed above is very similar to a raditional Chinese daice named as Colored Belt Dance? which is performed w dancers holding one cud of a long belt iu color., It is interesting to notice that the morphology of the streamer wave discussed above is very similar to a traditional Chinese dance named as 'Colored Belt Dance' which is performed by dancers holding one end of a long belt in color.
 Au important exteion to the coronal wave study is to develop diagnostic techniques of plasmas ancl magnetic fields throteh which the wave propagates. ic. to couduct the study of coronal seimnology.," An important extension to the coronal wave study is to develop diagnostic techniques of plasmas and magnetic fields through which the wave propagates, i.e., to conduct the study of coronal seismology."
 In our case. the period aud phase speed of the streamer wave which has beeu regarded as the propagating kiuk mode carried by the thin pasma sheet. if well resolved frou. observations. can be used to provide iuforiiation on magnetic properties of streamers.," In our case, the period and phase speed of the streamer wave which has been regarded as the propagating kink mode carried by the thin plasma sheet, if well resolved from observations, can be used to provide information on magnetic properties of streamers."
 Gonerallv speakiug. the phase speed for tje wave phenomenon investigated in this study is eiven by the suu of two compoucuts.," Generally speaking, the phase speed for the wave phenomenon investigated in this study is given by the sum of two components."
 The first one is the speed of the solar wind along the plasma sheet. the mediun carving the mode outwards.," The first one is the speed of the solar wind along the plasma sheet, the medium carrying the mode outwards."
 The other is of course the phase speed of the wave mode in the asma rest franc., The other is of course the phase speed of the wave mode in the plasma rest frame.
 The phase speed for the kink uxnde under thin plasma sheet econpetry cau be tentatively described with available ATID theory developed or a plasina-slab configuration In cartesia1 geonietzy (Re)berts. 1981: Edwin Roberts. 1982).," The phase speed for the kink mode under thin plasma sheet geometry can be tentatively described with available MHD theory developed for a plasma-slab configuration in cartesian geometry (Roberts, 1981; Edwin Roberts, 1982)."
" Substituting noniial parameters m the sIow-iud plasma sheet region above the streamer ctsp iuto the dispersion relation given x Edwin& Roberts (1982). we fiud that t16 phase spece of the relevaut fast kink body mode ej. Is sunaller than vet rather close to he external Alfvénn speed (y,=Boipni. where à is the proou number deusitv aid y, the protou lass."," Substituting nominal parameters in the slow-wind plasma sheet region above the streamer cusp into the dispersion relation given by Edwin Roberts (1982), we find that the phase speed of the relevant fast kink body mode $c_k$ is smaller than yet rather close to the external Alfvénn speed $v_{Ae}=B_e / \sqrt{\mu_0 n m_p}$, where $n$ is the proton number density and $m_p$ the proton mass."
 The differeice between the dediced Ch and cas is senucral voless tlm one third of ey)., The difference between the deduced $c_k$ and $v_{Ae}$ is generally less than one third of $v_{Ae}$.
" Therefore. to implement a preliminary seimnological stuv on the magnetic field sreneth D,. we take ey, to be equal to the kink mode phase spec Ch estimated frou οἱr observations."," Therefore, to implement a preliminary seismological study on the magnetic field strength $B_e$, we take $v_{Ae}$ to be equal to the kink mode phase speed $c_k$ estimated from our observations."
 Regarding the solar wind conditious in the coucerred reeion. the readers are referred. to relevant observational studies (Sheclev e al..," Regarding the solar wind conditions in the concerned region, the readers are referred to relevant observational studies (Sheeley et al.,"
 1997: Wang et al..," 1997; Wang et al.,"
 2000: Straclal, 2000; Strachan
While a more complete testing of our method will be presented in a subsequent paper. in (his section we present some resulis [rom applving our method to simulated catalogs (hat illustrate the effects of small.scale. nonlinear power and how thev are mitigated in our analvsis.,"While a more complete testing of our method will be presented in a subsequent paper, in this section we present some results from applying our method to simulated catalogs that illustrate the effects of small–scale, nonlinear power and how they are mitigated in our analysis."
 For our testing we have chosen simulated catalogs with 21000 galaxies designed to mimic the characteristics of the SET survey (claCostaefaf.1995)., For our testing we have chosen simulated catalogs with $\approx1000$ galaxies designed to mimic the characteristics of the SFI survey \citep{dacosta95}.
. The catalogs were drawn [rom a 256% Nbody PM (particle mesh) simulation with D=0.25 and 4=O'oy0.46.," The catalogs were drawn from a $256^3$ N–body PM (particle mesh) simulation with $\Gamma = 0.25$ and $\beta =
\Omega^{0.6}\sigma_8 = 0.46$."
 In these simulations. (he box size was taken to be 512 Mpc and the IInbble constant kms !Mpe.1=0.75: thus the box size in redshift space corresponds to a diameter of 38.400 kms +.," In these simulations, the box size was taken to be $512$ Mpc and the Hubble constant $h=H/100$ km $^{-1}{\rm
Mpc}^{-1}=0.75$; thus the box size in redshift space corresponds to a diameter of 38,400 km $^{-1}$."
 Galaxies were identified in (hese simulations and assigned physical properties., Galaxies were identified in these simulations and assigned physical properties.
" To duplicate the characteristics of the SEI survey. galaxies were ""observed"" bv applying the sanie selection criteria."," To duplicate the characteristics of the SFI survey, galaxies were “observed” by applying the same selection criteria."
 Realistic scatter was added (ο galaxy. properties Chat duplicates the relative error in the SFI inferred distances., Realistic scatter was added to galaxy properties that duplicates the relative error in the SFI inferred distances.
 Finally. following Freudling (1995) we applied an inhomogeneous Malhlmequist correction to our catalogs.," Finally, following Freudling (1995) we applied an inhomogeneous Malmquist correction to our catalogs."
 We perlormed the analvsis described in Sec., We performed the analysis described in Sec.
 G on these simulated catalogs., \ref{sec-anal} on these simulated catalogs.
 In Fig., In Fig.
 1 we show the window functions for selected moments calculated for a (vpical catalog in order of increasing eigenvalue. with the top plot showing the window functions for the moments with the five lowest eigenvalues. the middle showing five others associated with somewhat larger eigenvalues. aud the bottom. plot showing five more selected from the whole range ol eigenvalues.," \ref{winfun} we show the window functions for selected moments calculated for a typical catalog in order of increasing eigenvalue, with the top plot showing the window functions for the moments with the five lowest eigenvalues, the middle showing five others associated with somewhat larger eigenvalues, and the bottom plot showing five more selected from the whole range of eigenvalues."
 This demonstrates that selecting moments that are least sensitive to small scales does in [act generally result in moments that are most sensitive to large scales: window functions of moments with larger eigenvalues are successively larger on nonlinear scales as expected., This demonstrates that selecting moments that are least sensitive to small scales does in fact generally result in moments that are most sensitive to large scales; window functions of moments with larger eigenvalues are successively larger on nonlinear scales as expected.
 Thus the information contained in laree eigenvalue moments comes mostly from scales where fluctuations are nonlinear and should not be included in a linear analvsis., Thus the information contained in large eigenvalue moments comes mostly from scales where fluctuations are nonlinear and should not be included in a linear analysis.
 For our simulated catalogs. we know the “true” values of D and 2.," For our simulated catalogs, we know the “true” values of $\Gamma$ and $\beta$."
" If: we use these true values as our ""guess"" (see Sec. 6))", If we use these true values as our “guess” (see Sec. \ref{sec-anal}) )
 lo calculate the optimum moments. then the values of these moments calculated. [rom the velocities should have unit variance. since (he power spectrum model should be au excellent fit to the data.," to calculate the optimum moments, then the values of these moments calculated from the velocities should have unit variance, since the power spectrum model should be an excellent fit to the data."
 ILowever. nonlinear elfects can cause igher order moments to deviate from unit variance.," However, non–linear effects can cause higher order moments to deviate from unit variance."
 In Fig., In Fig.
 2. we show the sum of the first N moments versus moment number NV [or a (vpical catalog. where the moments are ranked in order of increasing eigenvalue.," \ref{ratio} we show the sum of the first $N$ moments versus moment number $N$ for a typical catalog, where the moments are ranked in order of increasing eigenvalue."
 Note that for small Αν the sum tracks a line with unit slope. whereas for large V the sum deviates from this line: (his is an indication that the ionlinear effects are causing the large No moments to deviate [rom unit variance.," Note that for small $N$, the sum tracks a line with unit slope, whereas for large $N$ the sum deviates from this line; this is an indication that the non–linear effects are causing the large $N$ moments to deviate from unit variance."
 In Fig., In Fig.
 3 we show the results of (he likelihood analvsis on a tvpical catalog Lor different vunber .V of moments kept., \ref{maxval} we show the results of the likelihood analysis on a typical catalog for different number $N^{\prime}$ of moments kept.
" For reference. we also give the value of A@, for each AN"" as discussed in Sec. 4.."," For reference, we also give the value of $\Delta\theta_q$ for each $N^{\prime}$ as discussed in Sec. \ref{sec-select}."
 Here the closed triangles correspond to the maximum likelihood values while the contours correspond (o 1/2. 1/10. and 1/100 of the maximum likelihood.," Here the closed triangles correspond to the maximum likelihood values while the contours correspond to $1/2$, $1/10$, and $1/100$ of the maximum likelihood."
" The asterisk svinbol corresponds to the input values used [or the simulation. the “true” values for D and ο,"," The asterisk symbol corresponds to the input values used for the simulation, the “true” values for $\Gamma$ and $\beta$."
 We see that in (his case. inclusion of all of the information leads to the location of the maximum likelihood being skewed away [from the true values (see the panel with NV’= N).," We see that in this case, inclusion of all of the information leads to the location of the maximum likelihood being skewed away from the true values (see the panel with $N^{\prime}=N$ )."
 However. when higher order moments are discarded. the location of (he maximum likelihood corresponds well with (he (rue values.," However, when higher order moments are discarded, the location of the maximum likelihood corresponds well with the true values."
" For this particular catalog. with σι,= 200km/s. the criterion of Eq. (21))"," For this particular catalog, with $\sigma_*=200$ km/s, the criterion of Eq. \ref{criterion}) )"
 would give A’~125 for the optimum number of moments to keep.," would give $N^{\prime}\simeq
125$ for the optimum number of moments to keep."
 The [act that the discarding of higher order moments leads to a much better agreement between the maximum likelihood location and the true values is a good indication that our analvsis method is effectively removing smallscale. nonlinear velocity information.," The fact that the discarding of higher order moments leads to a much better agreement between the maximum likelihood location and the true values is a good indication that our analysis method is effectively removing small–scale, nonlinear velocity information."
 thequark-cliquark cascade model with spin.We investigate thesofthadronic interactions,in the regions of $p_T < 2$ GeV/c. We analyze hyperon and anti-hyperon polarizations in
In general. (hermochemical equilibrium governs the composition of the deep atmospheres of giant planets and brown cdwarls because thev are warm enough lor chemical reactions to readily overcome energv barriers (o reaction kinetics.,"In general, thermochemical equilibrium governs the composition of the deep atmospheres of giant planets and brown dwarfs because they are warm enough for chemical reactions to readily overcome energy barriers to reaction kinetics."
 However. disequilibrium processes in substellar atmospheres are well known.," However, disequilibrium processes in substellar atmospheres are well known."
 In addition to photochemistry driven by ultraviolet irradiation. atmospheric mixing is one of the dominant mechanisms that drives the chemical composition out of equilibrium.," In addition to photochemistry driven by ultraviolet irradiation, atmospheric mixing is one of the dominant mechanisms that drives the chemical composition out of equilibrium."
 In this scenario. rapid vertical mixing may (ransport a parcel ol gas to higher. cooler altitudes before its chemical constituents have had sufficient. time to attain equilibrium via reaction chemistry a phenomenon that has been proposed to explain the overabundance of various “disequilibrium” species in the atmospheres of Jupiter. saturn. Uranus. ancl Neptune (e.g..Prinn&Owen1976:Barshavetal.2009:Visscher2010:Moses 2010).. brown cwarls 2010).. and extrasolar giant planets (e.g..Cooper&Showman2006:Fortneyetal.Mosesetal. 2011).," In this scenario, rapid vertical mixing may transport a parcel of gas to higher, cooler altitudes before its chemical constituents have had sufficient time to attain equilibrium via reaction chemistry — a phenomenon that has been proposed to explain the overabundance of various “disequilibrium” species in the atmospheres of Jupiter, Saturn, Uranus, and Neptune \citep[e.g.,][]{prinn1976,prinn1977,barshay1978,fegley1979,prinn1981jgr,prinn1984,lewis1984,fegley1985apj,fegley1986apj,fegley1988,fegley1991,fegley1994,lodders1994,lodders2002,bezard2002,taylor2004,visscher2005,fouchet2009,visscher2010icarus,moses2010}, brown dwarfs \citep[e.g.,][]{fegley1996,noll1997,griffith1999,griffith2000,saumon2000,lodders2002,golimowski2004,saumon2006,saumon2007,visscher2006,leggett2007,hubeny2007,geballe2009,king2010,yamamura2010}, and extrasolar giant planets \citep[e.g.,][]{cooper2006,fortney2006hd149,burrows2008,line2010,madhusudhan2011,stevenson2010,moses2011}."
. Prin&Barshav(1977). first developed an analvtical model to explain the observed overabundance of CO in Jupiters troposphere due (o strong vertical mixing., \citet{prinn1977} first developed an analytical model to explain the observed overabundance of CO in Jupiter's troposphere due to strong vertical mixing.
" In (his approach. a lime scale for convective mixing (7,,;.). based upon an estimated mixing length scale and vertical mixing rate. is compared to a (time scale for chemical kinetics (Τομ). based upon an assumption about which chemical pathways will be important for interconversion between atmospheric constituents."," In this approach, a time scale for convective mixing $\tau_{mix}$ ), based upon an estimated mixing length scale and vertical mixing rate, is compared to a time scale for chemical kinetics $\tau_{chem}$ ), based upon an assumption about which chemical pathways will be important for interconversion between atmospheric constituents."
 At high temperatures in the deep atmosphere. thermochemical equilibrium is maintained because reaction kinetics operate faster (han couveclive mixing Tehem< Tuis).," At high temperatures in the deep atmosphere, thermochemical equilibrium is maintained because reaction kinetics operate faster than convective mixing $\tau_{chem}<\tau_{mix}$ )."
" However. departures from. equilibrium can occur at colder. higher altitudes when convective mixing begins (o dominate over reaction kinetics(16. when Them> τε]. and (he abundance of a molecular constituent may become ""quenched al a value representative of the quench level (defined by 74,5,= τηε]."," However, departures from equilibrium can occur at colder, higher altitudes when convective mixing begins to dominate over reaction kinetics, when $\tau_{chem}>\tau_{mix}$ ), and the abundance of a molecular constituent may become “quenched” at a value representative of the quench level (defined by $\tau_{chem}=\tau_{mix}$ )."
 This level is different for each species (Feelev&Prinn1985) ancl. in principle. species that is subject to reaction chemistry and atmospheric transport will quench if the appropriate time scale for reaction kinetics becomes longer than the time scale for convective mixing (e.g... see Fig.," This level is different for each species \citep{fegley1985apj} and, in principle, species that is subject to reaction chemistry and atmospheric transport will quench if the appropriate time scale for reaction kinetics becomes longer than the time scale for convective mixing \citep[{e.g.}, , see Fig."
" 62010).Here we investigate the chemical interconversion between CO and CII,. which becomes quenched when ΤμΩω=CII;)>7,,, in a substellar atmosphere."," 6.Here we investigate the chemical interconversion between CO and $_{4}$, which becomes quenched when $\tau_{chem}(\textrm{CO}\rightleftarrows\textrm{CH}_{4})>\tau_{mix}$ in a substellar atmosphere."
Subelwarl B (scd) stars are helium core burning stars with very thin hvdrogen envelopes that [ie at the blue end. of the horizontal branch and hence are identified with the extreme horizontal branch (EIID) stars (seearecentre- 2009).,Subdwarf B (sdB) stars are helium core burning stars with very thin hydrogen envelopes that lie at the blue end of the horizontal branch and hence are identified with the extreme horizontal branch (EHB) stars \citep[see a recent review by][]{heb2009}.
. D'Cruzetal.(1996). showed. that a high mass-loss rate on the red. giant branch produces a thin hydrogen envelope and prevents the star from ascending the asymptotic giant branch., \citet{dcr1996} showed that a high mass-loss rate on the red giant branch produces a thin hydrogen envelope and prevents the star from ascending the asymptotic giant branch.
 The evolution of single 5L stars [rom zero-age to helium exhaustion may be followed on a series of tracks narrowly centred on 0.475 AZ. (Dormanetal. 1993)., The evolution of single EHB stars from zero-age to helium exhaustion may be followed on a series of tracks narrowly centred on 0.475 $M_\odot$ \citep{dor1993}.
.. After helium exhaustion these objects evolve directly onto the white chvarl sequence., After helium exhaustion these objects evolve directly onto the white dwarf sequence.
 On the other hand. following the original. proposal of Alengeletal.(1976). [or the formation of sdB stars through binary evolution. it has been found that a significant fraction of these stars reside in close binary svstenis2003).," On the other hand, following the original proposal of \citet{men1976} for the formation of sdB stars through binary evolution, it has been found that a significant fraction of these stars reside in close binary systems."
. Phe onset of a common envelope phase or Roche lobe overflow contributes to the removal of the hydrogen envelope and directs the star toward the 11112., The onset of a common envelope phase or Roche lobe overflow contributes to the removal of the hydrogen envelope and directs the star toward the EHB.
 Lanοἱal.(2002.2003). propose three formation channels for the formation of sdB stars through binary interaction. either involving common envelope (CLE) phases. episodes of Roche lobe overllow (191ΟΙ). or the merger of two helium white cwarfs.," \citet{han2002,han2003} propose three formation channels for the formation of sdB stars through binary interaction, either involving common envelope (CE) phases, episodes of Roche lobe overflow (RLOF), or the merger of two helium white dwarfs."
 Phe CI scenario involving primary stars that experience a helium Uash accompanied by a low-mass or white cwarl secondary star is expected to create short-period binaries (logP(d)=] to 1) and a linal primary mass distribution. narrowly centred: on 0.46.., The CE scenario involving primary stars that experience a helium flash accompanied by a low-mass or white dwarf secondary star is expected to create short-period binaries $\log{P(d)}\approx -1$ to 1) and a final primary mass distribution narrowly centred on $M_\odot$.
 The CLE scenario with primary stars massive enough to avoid a helium flash is expected. to achieve a much lower final mass for the primary (0.33-0.35. Alo)., The CE scenario with primary stars massive enough to avoid a helium flash is expected to achieve a much lower final mass for the primary (0.33-0.35 $M_\odot$ ).
 On the other hand. the RLOF scenario creates longer period binaries ancl a wider distribution of primary final masses.," On the other hand, the RLOF scenario creates longer period binaries and a wider distribution of primary final masses."
 Studies of the binary components. anc an estimate of the frequency of such systems are required to constrain these models ancl determine the relative contribution of these formation channels to the sdB population.," Studies of the binary components, and an estimate of the frequency of such systems are required to constrain these models and determine the relative contribution of these formation channels to the sdB population."
 In this context. we have initiated a program to identify," In this context, we have initiated a program to identify"
The final sample for cach chip was formed by using the magnitude from the long exposure image for all unsaturated stars and from the short image for those stars that are saturated on the long image.,The final sample for each chip was formed by using the magnitude from the long exposure image for all unsaturated stars and from the short image for those stars that are saturated on the long image.
 The change from long to short data is at. σετ., The change from long to short data is at $\approx$ 17.5.
 In NCC1805 we find that there is à colour shift between the chips of z0.04 mag., In NGC1805 we find that there is a colour shift between the chips of $\approx$ 0.04 mag.
 Similar shifts have been found by other groups in cluster colour magnitude diagrams (Johnsonetal.1999). and attributed to errors in CPE and aperture corrections and in zeropoints., Similar shifts have been found by other groups in cluster colour magnitude diagrams \cite{John99} and attributed to errors in CTE and aperture corrections and in zeropoints.
 Similar errors are likely causing the colour shift seen in our data., Similar errors are likely causing the colour shift seen in our data.
 The NICMOS data were combined using the WRAP task miscombine. which sums the data ancl performs cosmic rav rejection.," The NICMOS data were combined using the IRAF task mscombine, which sums the data and performs cosmic ray rejection."
 Stars were detected in the NICMOS image and aperture photometry was performed., Stars were detected in the NICMOS image and aperture photometry was performed.
 Detections near to bright stars and alone dillraction spikes were masked. as in the optical data.," Detections near to bright stars and along diffraction spikes were masked, as in the optical data."
 There are several cdillieulties with NICMOS data. fortunately none of these had a big elfect on this projecCt.," There are several difficulties with NICMOS data, fortunately none of these had a big effect on this project."
 The Ipedestal. a constant which remains after runningg the calibration pipeline and. causes an inverse Uatficld pattern to be imprinted on the image. is not a big problem for these data as we take a local background for each star.," The pedestal, a constant which remains after running the calibration pipeline and causes an inverse flatfield pattern to be imprinted on the image, is not a big problem for these data as we take a local background for each star."
 Ghosts. which appear at congruent. positions in the other quadrants when a bright star is present in one quadrant. do appear in our images. but it is possible to look carefully at. the positions where ghosts are expected to occur and eliminate false detections.," Ghosts, which appear at congruent positions in the other quadrants when a bright star is present in one quadrant, do appear in our images, but it is possible to look carefully at the positions where ghosts are expected to occur and eliminate false detections."
 An aperture correction to 0.5 aresec was calculated from bright stars in the image., An aperture correction to 0.5 arcsec was calculated from bright stars in the image.
" The cata were calibrated to magnitudes in an approximate Vega svstem Using wherePHOTENE and Pos are caleulated by STScl at the time of writing to be 2.337E-6 S/DN. and 1039.3 Jv respectively (NICMOS Data Handbook. v4. Table 5.1). and we assume ZPy4,20 (as in the CET. infrared photometry scale)."," The data were calibrated to magnitudes in an approximate Vega system using where and $_{\nu \rm{Vega}}$ are calculated by STScI at the time of writing to be 2.337E-6 $\times$ s/DN and 1039.3 Jy respectively (NICMOS Data Handbook, v4, Table 5.1), and we assume $\it{ZP}_{\rm{Vega}}$ =0 (as in the CIT infrared photometry scale)."
 The final detected. star. lists in. NICALOS FIGOW and WEPC2 F555\W were matched: using the positional information in the image headers., The final detected star lists in NICMOS F160W and WFPC2 F555W were matched using the positional information in the image headers.
" lo was found that there can be as much as 2"" offset between the RAs ancl Dees calculated for a star from the NICMOS image and. those calculated for the same star from the WEPC2 image.", It was found that there can be as much as $^{\prime\prime}$ offset between the RAs and Decs calculated for a star from the NICMOS image and those calculated for the same star from the WFPC2 image.
 According to SUScl this is due to the combined uncertainty of the guide star positions. the location of the fine guidance sensors relative to the telescope axis ancl the measured locations of the instrument apertures.," According to STScI this is due to the combined uncertainty of the guide star positions, the location of the fine guidance sensors relative to the telescope axis and the measured locations of the instrument apertures."
 Figure 3. shows the de-reddened V. vs V-L Clohnson-C'ousins magnitudes) colour-magnituce diagrams {CAIDs)) for all four chips of NGCISIS (top) and NGC1805 (bottom)., Figure \ref{fig:VIcolmag} shows the de-reddened V vs V-I (Johnson-Cousins magnitudes) colour-magnitude diagrams (CMDs) for all four chips of NGC1818 (top) and NGC1805 (bottom).
 The different chips have dillerent. svmibols., The different chips have different symbols.
 Stars marked. with bold squares are Be stars (see subsection 3.2))., Stars marked with bold squares are Be stars (see subsection \ref{subsec:Beid}) ).
 The data have been de-reddened: assuming I5(D-V)=0.075., The data have been de-reddened assuming E(B-V)=0.075.
 Tables 4 and 5 tabulate these data for Νις1515 and NGCISO5 respectively (full versions of these tables are available on the MNIUAS web site)., Tables \ref{tab:n1818dat} and \ref{tab:n1805dat} tabulate these data for NGC1818 and NGC1805 respectively (full versions of these tables are available on the MNRAS web site).
 lsochrones from Bertelli οἱ shortciteDert94.. with a range of age and metallicity values encompassing those found in the literature for these clusters. are plotted. on the CALDs.," Isochrones from Bertelli et \\shortcite{Bert94}, with a range of age and metallicity values encompassing those found in the literature for these clusters, are plotted on the CMDs."
 “Lo illustrate the effects. of age ancl metallicity the top plot in Figure 3. shows 25 40 Myr isochrones for two metallicities. Fe/1]. of -0.4 and 0 and he bottom plot shows solar metallicity isochrones for ages of 25. 40 and 63 Myr.," To illustrate the effects of age and metallicity the top plot in Figure \ref{fig:VIcolmag} shows 25 40 Myr isochrones for two metallicities, [Fe/H], of -0.4 and 0 and the bottom plot shows solar metallicity isochrones for ages of 25, 40 and 63 Myr."
 The two clusters have very similar CMDs. which are raceck well by the 25Mvyr solar metallicity. isochrone.," The two clusters have very similar CMDs, which are traced well by the 25Myr solar metallicity isochrone."
" The ages and metallicities of these clusters are. investigated ""urther by comparison with simulations in section 4..", The ages and metallicities of these clusters are investigated further by comparison with simulations in section \ref{sec:sim}.
 Note hat. even with the very short exposure times used here. the xieghtest stars are still saturated in the FSI4W images.," Note that, even with the very short exposure times used here, the brightest stars are still saturated in the F814W images."
 The red giant branch. of the field. population of the LAIC is apparent in the CMDs at V-Iz1. Vz 18.5.," The red giant branch of the field population of the LMC is apparent in the CMDs at $\approx$ 1, $\approx$ 18.5."
 We have not subtracted these background stars from our data as we are predominantly interested in the brighter stars x19) where the contribution from the field is negligible (see [unter et 11997. Figure)., We have not subtracted these background stars from our data as we are predominantly interested in the brighter stars $\le$ 19) where the contribution from the field is negligible (see Hunter et 1997 Figure4).
 Figure 4 show the V vs V-LI diagrams for both clusters., Figure \ref{fig:VHcolmag} show the V vs V-H diagrams for both clusters.
 Phe isochrones (Bertellietal.L994) are for 25 and 40 Myr and for solar moetallicitv., The isochrones \cite{Bert94} are for 25 and 40 Myr and for solar metallicity.
 ‘Tables G and 7 tabulate these data lor NGCTSIS and NGCISO5 respectively (full versions of these tables are available on the MNIUAS web site)., Tables \ref{tab:n1818vhdat} and \ref{tab:n1805vhdat} tabulate these data for NGC1818 and NGC1805 respectively (full versions of these tables are available on the MNRAS web site).
 Νάς1515 and NGCLSO5 have five and three red supergiants respectively., NGC1818 and NGC1805 have five and three red supergiants respectively.
 Although our having to use NIC? reduced 10 number of stars in the NICMOS data. these colour-magnitude ciagrams are still of some use as they show us that the red supergiants lic towards the red end of the isochrones.," Although our having to use NIC2 reduced the number of stars in the NICMOS data, these colour-magnitude diagrams are still of some use as they show us that the red supergiants lie towards the red end of the isochrones."
 Although the poisson errors on the magnitudes of the bright stars are small. there could well be svstematic errors of a few tenths of a magnitude due to uncertainty in the Ll calibration and the isochrones.," Although the poisson errors on the magnitudes of the bright stars are small, there could well be systematic errors of a few tenths of a magnitude due to uncertainty in the H calibration and the isochrones."
 In. λές1515 two of the red. supergiants are located. on the 25 Myr isochrone. and the other three are consistent with either the 25 or 40 Alvr isochrone.," In NGC1818 two of the red supergiants are located on the 25 Myr isochrone, and the other three are consistent with either the 25 or 40 Myr isochrone."
 ln ας1505 two of the red supergiants are located on the 40 Myr isochrone. and one which has lower limits in V is consistent with either isochrone.," In NGC1805 two of the red supergiants are located on the 40 Myr isochrone, and one which has lower limits in V is consistent with either isochrone."
 In Figure 3. there are many stars in both clusters with 15 VIT that are significantly redder than the isochrones., In Figure \ref{fig:VIcolmag} there are many stars in both clusters with $\le$ $\le$ 17 that are significantly redder than the isochrones.
 ]t was suspected that these are De stars., It was suspected that these are Be stars.
" An ellective wav to identify Be stars in clusters is to use the fact that these stars show Balmer emission anc hence will separate from non-Be stars in ία ""colour.", An effective way to identify Be stars in clusters is to use the fact that these stars show Balmer emission and hence will separate from non-Be stars in V-Ha `colour'.
 An image, An image
The signalOo component ofA is where the pixc-beam function F; is given bv Iu the flat sky limit. the theory covariance matrix is eiven bv where F;(1) is theFourier transform of F;itr).,"The signal component of$\Delta$ is where the pixel-beam function $F_i$ is given by In the flat sky limit, the theory covariance matrix is given by where ${\tilde F}_i({\bf l})$ is theFourier transform of $F_i({\bf r})$."
 To perform the iterative quadratic baud-power cstimation procedure. it is necessary to know the partial derivative of Cy with respect to cach of the band-powers q5. which according to equation(1)) is given by Note that this algorithm docs not assuue that the instrument beams stav constant duriug the observations.," To perform the iterative quadratic band-power estimation procedure, it is necessary to know the partial derivative of $C_T$ with respect to each of the band-powers $q_B$, which according to \ref{dl}) ) is given by Note that this algorithm does not assume that the instrument beams stay constant during the observations."
 As described in ?.. the ACBAR beam sizes are weak functions of the chopper position.," As described in \citet{runyan03a}, the ACBAR beam sizes are weak functions of the chopper position."
 NOL adopted a semianalytic expansion to correct for these effects to first order., K04 adopted a semi-analytic expansion to correct for these effects to first order.
 To verify that the effects due to non-uuiforii beams are sanall. we developed. two eud-to-cud pipelines.," To verify that the effects due to non-uniform beams are small, we developed two end-to-end pipelines."
" In the first pipeline. the pixebbeani functious F;(r) are calculated explicitly caving the co-adding process,"," In the first pipeline, the pixel-beam functions $F_i({\bf r})$ are calculated explicitly during the co-adding process."
 The baudpowers in Table 3. are analyzed with this aleoritlin., The bandpowers in Table \ref{tab:bands} are analyzed with this algorithm.
 Iu the secoud pipeline. au averaged beam is used for the cutire map.," In the second pipeline, an averaged beam is used for the entire map."
 The difference in the power spectra from the two pipelines is ucelieible., The difference in the power spectra from the two pipelines is negligible.
" Iu the aualvsis of οι, we assuued that the nolse is stationary in chopper position. after LAIT subtraction."," In the analysis of K04, we assumed that the noise is stationary in chopper position after LMT subtraction."
 Iu the cureut treatment. we relax this assumption and calculate the full two dimensional correlation matrix directly from the time stream data without using Fourier transforms.," In the current treatment, we relax this assumption and calculate the full two dimensional correlation matrix directly from the time stream data without using Fourier transforms."
 All the uunerical calculations are performed on the National Energy Research Scieutific Computing Center (NERSC) IDM SP RS/6000., All the numerical calculations are performed on the National Energy Research Scientific Computing Center (NERSC) IBM SP RS/6000.
 The evaluation of F;(1) and its Fourier transform are the most computationally demanding steps in this analysis., The evaluation of $F_i({\bf r})$ and its Fourier transform are the most computationally demanding steps in this analysis.
 After Cr. Ον and Crp are calculated. standard likelihood maximizing procedures are used to find the baud-powoers qp and uncertainties (?)..," After ${\bf C}_T$ , ${\bf C}_N$ and ${\bf C}_{T,B}$ are calculated, standard likelihood maximizing procedures are used to find the band-powers ${\bf q}_B$ and uncertainties \citep{bond98}."
 The results of this analysis are preseuted in Table 3 aud Figure 1.., The results of this analysis are presented in Table \ref{tab:bands} and Figure \ref{fig:acbar}.
 RCW38 is a compact UIT region in the Calactic plane at a declination simile to the ACBAR CAMB fields., RCW38 is a compact HII region in the Galactic plane at a declination similar to the ACBAR CMB fields.
 It has a large aud stable flux and serves as the primary calibrator for the ACDBAR observations., It has a large and stable flux and serves as the primary calibrator for the ACBAR observations.
 We determine the absolute flux of ΠΟΟδ usine maps from the 2003 flieht of BOOMERANG (?.. hereafter BO3). which are calibrated. relative to the WAIAP experiment with an absolute uncertaintv of LSU.," We determine the absolute flux of RCW38 using maps from the 2003 flight of BOOMERANG \cite{masi05}, hereafter B03), which are calibrated relative to the WMAP experiment with an absolute uncertainty of $1.8\%$."
 ROW3s docs not have a black-body spectrum. requiring spectral corrections for the calibration of CAIB anisotropies.," RCW38 does not have a black-body spectrum, requiring spectral corrections for the calibration of CMB anisotropies."
 However the similarity in the spectral respouses of the 150 GITz bauds iu the DOS and ACBAR experiments cusures these corrections to be sanall.," However the similarity in the spectral responses of the $150\,$ GHz bands in the B03 and ACBAR experiments ensures these corrections to be small."
 Were we outline the calibration procedure. leaving the details to Appeudix ?7..," Here we outline the calibration procedure, leaving the details to Appendix \ref{app:calib}."
 ACBAR typically observed. ROW3s before aud after cach CAIB observation., ACBAR typically observed RCW38 before and after each CMB observation.
 Comparisons between the Bos and ACBAR maps of RCW38 are used to determine the absolute calibration of the CMD fields to an unucertaiutv of6., Comparisons between the B03 and ACBAR maps of RCW38 are used to determine the absolute calibration of the CMB fields to an uncertainty of.
0%... For roughly of tlie 2002 season. we observed RCW358 with only half the 150 GIIz detectors CL out of SJ," For roughly of the 2002 season, we observed RCW38 with only half the $150\,$ GHz detectors (4 out of 8)."
 During these periods. the ROWS3s calibration was applied. to the remaiming detectors bv comparing CMD power spectra derived frou each half of the detectors.," During these periods, the RCW38 calibration was applied to the remaining detectors by comparing CMB power spectra derived from each half of the detectors."
 The calibration of the CAIB| field (observed im 2002) is extended to the overlapping CMD2 field (observed iu 2001) bv comparing power spectra frou cach field., The calibration of the CMB4 field (observed in 2002) is extended to the overlapping CMB2 field (observed in 2001) by comparing power spectra from each field.
" In the first ACBAR release. the 2001 and 2002 data sets were calibrated with au accuracy of using observations of Mags aud Veuus respectively,"," In the first ACBAR release, the 2001 and 2002 data sets were calibrated with an accuracy of using observations of Mars and Venus respectively."
 We determine the corrections o this planct-hased calibration to be 0.911-E0.072. for CAB? (2001) aud 1.1284x0.066 for CMD41-7. (2002) in CXMB temperature., We determine the corrections to this planet-based calibration to be $\pm$ 0.072 for CMB2 (2001) and $\pm$ 0.066 for CMB4-7 (2002) in CMB temperature.
 The 2002 observations dominate the final power spectra. and the final results have csseutially he same temperature calibration uucertainty as the 2002 data.," The 2002 observations dominate the final power spectra, and the final results have essentially the same temperature calibration uncertainty as the 2002 data."
 We performed a series of tests to constrain the amplitude of potential systematic errors in the power spectra. results., We performed a series of tests to constrain the amplitude of potential systematic errors in the power spectrum results.
 As described by I&01. the “first half iuinus secoud half jackkuite is a very powerful test for time clepeudcut errors. such as a chaugiug calibration. iuconsistenew in the beam or poiutiug reconstruction. and time varving sidelobe pickup.," As described by K04, the “first half minus second half"" jackknife is a very powerful test for time dependent errors, such as a changing calibration, inconsistency in the beam or pointing reconstruction, and time varying sidelobe pickup."
 Iu addition. high-f jackknife baud-powers constrain the uais-cstimation of noise.," In addition, $\ell$ jackknife band-powers constrain the mis-estimation of noise."
 We perform this test on the joint CAIB power spectrmu and find the baud-powers of the chronologically differenced maps are consistent with zero (Fig. 2))., We perform this test on the joint CMB power spectrum and find the band-powers of the chronologically differenced maps are consistent with zero (Fig. \ref{fig:sys}) ).
 Siuilarh. the data can be divided in two halves according to the direction of the chopper notion.," Similarly, the data can be divided in two halves according to the direction of the chopper motion."
 Microphouic vibrations due to the chopper turn-aromuds. erroneous trauster function corrections. or effects of wiud direction could produce a nonvanishing signal in the jackknife baud-powers.," Microphonic vibrations due to the chopper turn-arounds, erroneous transfer function corrections, or effects of wind direction could produce a nonvanishing signal in the jackknife band-powers."
 We fiud that the power spectruui of the left-right differenced maps is also consistent with ZOrO., We find that the power spectrum of the left-right differenced maps is also consistent with zero.
 To cusure that any residual chopper svuchronous offset is below the noise level. we have developed a new systematic test that the coutribution of auv such offsets relative to the CAIB.," To ensure that any residual chopper synchronous offset is below the noise level, we have developed a new systematic test that the contribution of any such offsets relative to the CMB."
 In this test. the baud-powers are derived. from an LAIT παρ. £|AL T. i which the residual svuchronous offsets (the sameiu each field) are enhanced relative to the (random) CAB fluctuations by a factor of 3 iu power Gicelecting the," In this test, the band-powers are derived from an LMT map, $L+M+T$ , in which the residual synchronous offsets (the samein each field) are enhanced relative to the (random) CMB fluctuations by a factor of 3 in power (neglecting the"
"CCs. and for the outward temperature and density declines in all clusters. is set by the underlying dominant DM distribution; specifically. such a seale is set by the peak of the DM velocity dispersion σ΄) at the position rj,2ro. that divides the ‘inner’ from the ‘outer’ cluster regions (see Figs.","CCs, and for the outward temperature and density declines in all clusters, is set by the underlying dominant DM distribution; specifically, such a scale is set by the peak of the DM velocity dispersion $\sigma^2(r)$ at the position $r_m\approx r_{-2}$, that divides the `inner' from the `outer' cluster regions (see Figs."
 | and 3 of CLFFO9)., 1 and 3 of CLFF09).
" In fact. the SM ensures that the approximation T—a7/d, [with J,=ΗμFy)kpT(ry). see CLFFO9] is to hold closely for all clusters around. |7j,. and over a wide radial range for theCCs*."," In fact, the SM ensures that the approximation $T\simeq\sigma^2/\beta_m$ [with $\beta_m\equiv \mu m_p\sigma^2(r_m)/k_BT(r_m)$, see CLFF09] is to hold closely for all clusters around $r_m$, and over a wide radial range for the."
. In the inner ICP regions. the temperature and density profiles are governed primarily by the central entropy level κ. set by events. 1.8.. energy discharged and blasts driven by major mergers or AGN outbursts: in à number of cases the latter imprint substructure on the small scale r; comparable to η.," In the inner ICP regions, the temperature and density profiles are governed primarily by the central entropy level $k_c$ set by events, i.e., energy discharged and blasts driven by major mergers or AGN outbursts; in a number of cases the latter imprint substructure on the small scale $r_f$ comparable to $r_m$."
 A novel feature ofthe SM (relative to handy isothermal or polytropie J-models where the entropy is assumed to be a functional &x11? of the density. see Cavaliere Fusco-Femiano 1976) is constituted by the radial entropy run entering Eq. (," A novel feature ofthe SM (relative to handy isothermal or polytropic $\beta$ -models where the entropy is assumed to be a functional $k\propto n^{\Gamma-5/3}$ of the density, see Cavaliere Fusco-Femiano 1976) is constituted by the radial entropy run entering Eq. ("
2).,2).
 We have collected in Table | the SM fitting parameters for all clusters in our set., We have collected in Table 1 the SM fitting parameters for all clusters in our set.
 By inspection it is apparent a correlation between the inner ICP profile type (marked by the CC/NCC tags) with the central entropy level κ... the outer entropy slope a. and the DM concentration c; high values of c and low values of a and κ. correspond to the CC class. while the opposite trend holds for NCC.," By inspection it is apparent a correlation between the inner ICP profile type (marked by the CC/NCC tags) with the central entropy level $k_c$, the outer entropy slope $a$ , and the DM concentration $c$; high values of $c$ and low values of $a$ and $k_c$ correspond to the CC class, while the opposite trend holds for NCC."
 We understand these trends in the framework of two-stage cluster formation (see 1) as follows., We understand these trends in the framework of two-stage cluster formation (see 1) as follows.
 For example. low values of & correspond to high values of bp (see Eq.," For example, low values of $a$ correspond to high values of $b_R$ (see Eq."
 4) owing to low values of Ao (see Eq., 4) owing to low values of $\Delta\phi$ (see Eq.
 3): these are related to high concentrations c=3.5(14z;) (see | and 2). which imply early transition redshifts ο. 1.e.. an old age.," 3); these are related to high concentrations $c=
3.5\,(1+z_t)$ (see 1 and 2), which imply early transition redshifts $z_t$ , i.e., an old age."
 We stress thatin the six clusters considered here the, We stress thatin the six clusters considered here the
(1998). vielding an overestimate of extinction in mocerate to high extinction regions.,"(1998), yielding an overestimate of extinction in moderate to high extinction regions."
 We estimate a calibration correction factor of for the ELI extinction values., We estimate a calibration correction factor of for the FIR extinction values.
 For the intermediate 17«|b]3° region. the relation between DIRBE/IRAS anc 2MASS extinction values departs more significantly from the identity line. as expected due to the larger contribution bv background dust.," For the intermediate $1^{\circ}<|{\it b}|<3^{\circ}$ region, the relation between DIRBE/IRAS and 2MASS extinction values departs more significantly from the identity line, as expected due to the larger contribution by background dust."
 In this region. the tvpical νομοςντι ratio is and could. be explained. by background dust. contribution and the calibration. factor allectine the Elli data.," In this region, the typical $A_{K,2MASS}/A_{K,FIR}$ ratio is and could be explained by background dust contribution and the calibration factor affecting the FIR data."
 An enhancement in the foreground. dust. with respect το the dust. model is observed in many cells in the northern strip., An enhancement in the foreground dust with respect to the dust model is observed in many cells in the northern strip.
 In the southern strip. several cells have cliourasscdpri smaller than expected. probably due to dense dust. clouds and temperature variations. currently not incorporated into our model.," In the southern strip, several cells have $A_{K,2MASS}/A_{K,FIR}$ smaller than expected, probably due to dense dust clouds and temperature variations, currently not incorporated into our model."
 For the regions very close to the Galactic Plane (|6 0.57). we have a typical value for the cliourτοςprin ratio of," For the regions very close to the Galactic Plane $|{\it b}| 
< 0.5^{\circ}$ ), we have a typical value for the $A_{K,2MASS}/A_{K,FIR}$ ratio of."
 Even considering the background. dust contribution and the calibration factor. this ratio is still smaller. than that predicted. by our. simple model. for the dust. distribution.," Even considering the background dust contribution and the calibration factor, this ratio is still smaller than that predicted by our simple model for the dust distribution."
 This fact is probably due to. the overestimation of the νι values by heated clust above that. obtained. from DIRBE temperature maps.," This fact is probably due to the overestimation of the $A_{K,FIR}$ values by heated dust above that obtained from DIRBE temperature maps."
 Another possible contribution to this dillerence is the existence of systematic clleets on the lisoaass values in high extinction regions (κοτος 2.5). where the PALASS extinction should be significantly. underestimated or even unreliable.," Another possible contribution to this difference is the existence of systematic effects on the $A_{K,2MASS}$ values in high extinction regions $A_{K,2MASS} > 2.5$ ), where the 2MASS extinction should be significantly underestimated or even unreliable."
 AX systematic asvmmetry in the values relative to the plane of the Galaxy is observed both in the 2ALASS and. DIRBE/IRAS data., A systematic asymmetry in the values relative to the plane of the Galaxy is observed both in the 2MASS and DIRBE/IRAS data.
 The behaviour and amplitude of this asymmetry with position on the sky suggest that the dominant role in creating this north-south asymmetry is a more ellective presence of foreground dust. clouds in the northern Galactic strips. such as the pe Nebula (Sect.," The behaviour and amplitude of this asymmetry with position on the sky suggest that the dominant role in creating this north-south asymmetry is a more effective presence of foreground dust clouds in the northern Galactic strips, such as the Pipe Nebula (Sect."
 2.2)., 2.2).
 X possible explanation is stellar winds and supernovae from nearby OD stellar associations xoducing dust cloud shells (Bhatt 2000)., A possible explanation is stellar winds and supernovae from nearby OB stellar associations producing dust cloud shells (Bhatt 2000).
 Phe nearby clouds »ojected towards the central parts ofthe Galaxy at positive atitudes belong to the Ophiuchus cust complex., The nearby clouds projected towards the central parts of the Galaxy at positive latitudes belong to the Ophiuchus dust complex.
 They. are oobablv related to the association Sco0D2. which is at à distance of 145 pe from the Sun (Bhatt 2000: Onishi ct al.," They are probably related to the association ScoOB2, which is at a distance of 145 pc from the Sun (Bhatt 2000; Onishi et al."
 1999)., 1999).
 ScoO132. in turn. belongs to Upper Scorpius. which is he easternmost part of the Sco-Cen Association. as studied » means of Hipparcos (de Zeeuw οἱ al.," ScoOB2, in turn, belongs to Upper Scorpius, which is the easternmost part of the Sco-Cen Association, as studied by means of Hipparcos (de Zeeuw et al."
 1999)., 1999).
 In all regions. significant substructure in. the Ancarisses. dope relation is seen. with loops and. armis stretching out from the main relation.," In all regions, significant substructure in the ${\it A_{K,2MASS}} vs.$ ${\it A_{K,FIR}}$ relation is seen, with loops and arms stretching out from the main relation."
 These structures are probably caused. by intervening dust. clouds. with cdilleren temperatures and densities for dilferent lines of sight.," These structures are probably caused by intervening dust clouds, with different temperatures and densities for different lines of sight."
 One extremely interesting. perspective is to model the dus distribution within the Galaxy. trying to reproduce as close as possible the details of the maps currently available.," One extremely interesting perspective is to model the dust distribution within the Galaxy, trying to reproduce as close as possible the details of the maps currently available."
 This effort. demancs models that. incorporate. on top of a smooth dust. distribution. the clleets of individual dus clouds. spiral arms. molecular rings and other structure. possibly with variable density contrasts ancl temperatures.," This effort demands models that incorporate, on top of a smooth dust distribution, the effects of individual dust clouds, spiral arms, molecular rings and other structure, possibly with variable density contrasts and temperatures."
 This ellort is currently under way for the central region of the Galaxy., This effort is currently under way for the central region of the Galaxy.
" This publication makes use of data products from the “Pwo Micron. All Sky Survey, which is a joint. project. of the University of Massachusetts and the Enfrared Processing and Analvsis Center/California Institute of Technology. funded by the National Acronautics and Space AXcdministration and the National Science Foundation."," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
 We also use the electronic form of the extinction maps provided by Schultheis et. al. (, We also use the electronic form of the extinction maps provided by Schultheis et al. (
1999) and Schlegel et al. (,1999) and Schlegel et al. (
1998).,1998).
 We thank the anonymous referee for his/her interesting comments and suggestions., We thank the anonymous referee for his/her interesting comments and suggestions.
 We acknowledge support from. the Brazilian institutions FAPESP and CNPq., We acknowledge support from the Brazilian institutions FAPESP and CNPq.
 CAID acknowledges FAPIESP for a post-doe fellowship (proc., CMD acknowledges FAPESP for a post-doc fellowship (proc.
 00/11S64-6)., 00/11864-6).
been adequately sampled over a wide chough rauge of huninosity to observe this pattern directly.,been adequately sampled over a wide enough range of luminosity to observe this pattern directly.
 Several more sources that vary by no more than a factor of LO in X- intensity appear to form portions of this pattern., Several more sources that vary by no more than a factor of 10 in X-ray intensity appear to form portions of this pattern.
 IU 303 aud IU 31 in Figure 1 exhibit oulv the bottom aud diagonal portions of this pattern., 4U $-$ 303 and 4U $-$ 34 in Figure \ref{cconly} exhibit only the bottom and diagonal portions of this pattern.
 Ser N-1 exhibits only the bottom portion. aud GS 238 ouly the top (Figure 1)).," Ser X-1 exhibits only the bottom portion, and GS $-$ 238 only the top (Figure \ref{cconly}) )."
 The tracks from the Z sources (νο νι CON 1712. GN D. aud CX 31010 (Figure 1)) ave uuuistakablv Z-shaped. but are somewhat differeut from those of the atoll sources.," The tracks from the Z sources Cyg X-2, GX 17+2, GX $-$ 1, and GX $+$ 0 (Figure \ref{cconly}) ) are unmistakably Z-shaped, but are somewhat different from those of the atoll sources."
 Iu general. Z sources are softer than atoll sources.," In general, Z sources are softer than atoll sources."
 Moreover. the traditional Z sources trace out a full track on time scales as short as a dav. while atoll sources fori. their color-color diagrams on much longer time scales of 30-100 days.," Moreover, the traditional Z sources trace out a full track on time scales as short as a day, while atoll sources form their color-color diagrams on much longer time scales of 30-100 days."
 We have highlehted using darker syiubols data spanning5 20- davs for the sources which trace their Z ou short time scales in Figure 1.. in oxder to allow an easy coniparison between the short aud long-term spectral variability.," We have highlighted using darker symbols data spanning 20 days for the sources which trace their Z on short time scales in Figure \ref{cconly}, in order to allow an easy comparison between the short and long-term spectral variability."
 As has been previously noted. the width of the track perpendicular to the direction of motion is mach sanaller in traditional Z sources than in atoll sources whe- observed ou time scales of davs (c.¢vanderIlis.1995).," As has been previously noted, the width of the track perpendicular to the direction of motion is much smaller in traditional Z sources than in atoll sources when observed on time scales of days \cite[e.g][]{vdk95}."
. Ou the five-vear time scales of Figure 1.. both the Z and atoll sources exhibit variations in the colors perpendicular to the direction of motion that traces the Z (o.c.Ixkuulkersetal.199L).," On the five-year time scales of Figure \ref{cconly}, both the Z and atoll sources exhibit variations in the colors perpendicular to the direction of motion that traces the Z \citep[e.g.][]{kul94}."
. It is interesting to note that the color-color diagram of CX 1 in Figure 1 resembles that of the Z source CX 1712 more than those of the other atoll sources.," It is interesting to note that the color-color diagram of GX $+$ 1 in Figure \ref{cconly}
 resembles that of the Z source GX $+$ 2 more than those of the other atoll sources."
 Although CX 1311 sas classified an atoll source because it lacked. strong QPOs (Ilasiueer&vanderWlis1989).. oobservatious have revealed weak QPOs similar to those of Z sources (Tomanetal.1998)... ancl several observations have revealed sinularities between the X-ray. (Schulzetal.1989) and IR (Bandyopadhyayctal.1999). spectra of GX 1311 and other Z sources.," Although GX 13+1 was classified an atoll source because it lacked strong QPOs \citep{hv89}, observations have revealed weak QPOs similar to those of Z sources \citep{hom98}, and several observations have revealed similarities between the X-ray \citep{sht89} and IR \citep{ban99} spectra of GX 13+1 and other Z sources."
 CX 13|L appears to exhibit both Z aud atoll properties. aud may prove important for uderstanding what causes the distinctions between the two classes of source.," GX 13+1 appears to exhibit both Z and atoll properties, and may prove important for understanding what causes the distinctions between the two classes of source."
 We next exanunune how these colors evolve versus oeiteusity and. for the atoll sources. versus time.," We next examine how these colors evolve versus intensity and, for the atoll sources, versus time."
 In Figure we have plotted the hard. aud soft color as a functiou of the 2I8 keV PCA count rates for the atoll source Aql N-1 aud the Z source GN 1., In Figure \ref{hid} we have plotted the hard and soft color as a function of the 2–18 keV PCA count rates for the atoll source Aql X-1 and the Z source GX $-$ 1.
 Iu Figure 3. we lave plotted the N-raw colors and the 2-12 keV. N-rav intensity from the AASNMI from the atoll source. IU 110 over a span of 200x davs., In Figure \ref{cctime} we have plotted the X-ray colors and the 2-12 keV X-ray intensity from the ASM from the atoll source 4U $-$ 440 over a span of 200 days.
 It is well-known that Z sources trace their colorcolor diagrams s1ioothlyv with time im less than a day (e.g.I&uulkersetal.199[:vanderWhs 1995).," It is well-known that Z sources trace their color-color diagrams smoothly with time in less than a day \citep[e.g.][]{kul94,vdk95}."
. These data are the best-sauipled examples from the LAINBs in Table 1.., These data are the best-sampled examples from the LMXBs in Table \ref{stats}.
 Figures 2 and 3 indicate that atoll sources also trace their Z-shaped track snoothly., Figures \ref{hid} and \ref{cctime} indicate that atoll sources also trace their Z-shaped track smoothly.
 The horizontal portion at the top of the color-color track of the atoll sources (with a hard color of about 0.5 in Figure 1)) is traced from left to reht as the intensity increases by iiore than a factor of 10., The horizontal portion at the top of the color-color track of the atoll sources (with a hard color of about 0.8 in Figure \ref{cconly}) ) is traced from left to right as the intensity increases by more than a factor of 10.
 For Αα X-1 and IU 1608-522. it is traced on time scales of davs to mouths as the source rises from (or decays to) below the PCA detection threshold. with a factor of 250 chanee in intensity.," For Aql X-1 and 4U 1608-522, it is traced on time scales of days to months as the source rises from (or decays to) below the PCA detection threshold, with a factor of 250 change in intensity."
 Ou the other hand. GS 18526. 31 has relmained in a smailur faint aud hard state for the sis vears of nuuonitormue.," On the other hand, GS $-$ 34 has remained in a similar faint and hard state for the six years of monitoring."
 The diagonal brauch is traced. ou fine scales of days. and only has been sampled well in time in a few instances (ee. Figure 3)).," The diagonal branch is traced on time scales of days, and only has been sampled well in time in a few instances (e.g. Figure \ref{cctime}) )."
 As an atoll source moves down alone the diagonal portion of the color-color track. the cout rate usually increases by a factor of 2 (aud vice versa: see Figure 3)).," As an atoll source moves down along the diagonal portion of the color-color track, the count rate usually increases by a factor of 2 (and vice versa; see Figure \ref{cctime}) )."
 However. during a recent uuusual outburst of the count rate decreased by a factor of 2 while," However, during a recent unusual outburst of \\citep{bai01} the count rate decreased by a factor of 2 while"
After imposing our selection criteria we are able to compile a list of candidate zS9 star forming galaxies in the IIUDE. UDE-P34. UDE-P12 and ERS fields.,"After imposing our selection criteria we are able to compile a list of candidate $z\approx 8-9$ star forming galaxies in the HUDF, UDF-P34, UDF-P12 and ERS fields."
 In. ‘Table 3 we list. positions and. photometry of these objects. while thumbnails of the bezzYJif images of these. candidates (where available) are presented in Figure 4..," In Table \ref{tab:objects} we list positions and photometry of these objects, while thumbnails of the $bvizYJH$ images of these candidates (where available) are presented in Figure \ref{fig:stamps}."
 In total we find 24 Y-drop candidates (LIUDE:6. UDE-DP34:7. UDE-P12:2. ERS:9) covering a range of apparent Jig magnitudes. of 21.0.28.5.," In total we find 24 $Y$ -drop candidates (HUDF:6, UDF-P34:7, UDF-P12:2, ERS:9) covering a range of apparent $J_{AB}$ magnitudes of $27.0-28.5$."
 In the three deep single WECS pointings. the number of cancdidates is fairly consistent from field to field. with 3. 4 and 2 Y-drops for the HUDE. UDE-I34 and fields. respectively. at Jigκ28.2.," In the three deep single WFC3 pointings, the number of candidates is fairly consistent from field to field, with 3, 4 and 2 $Y$ -drops for the HUDF, UDF-P34 and UDF-P12 fields, respectively, at $J_{AB}<28.2$."
 There are 9 fof the 24) objects in the Y-drop list (‘Table 3)) whieh we llag as being more marginal than the other candidates as they sit at the limits of our selection. although they are plausible z28ϐ galaxies (our effective volume calculation. already corrects for. galaxies. excluded as lving just outside the selection region).," There are 9 (of the 24) objects in the $Y$ -drop list (Table \ref{tab:objects}) ) which we flag as being more marginal than the other candidates as they sit at the limits of our selection, although they are plausible $z\approx 8-9$ galaxies (our effective volume calculation already corrects for galaxies excluded as lying just outside the selection region)."
 C'andidates ERS.YD7 and ERS.YDS in the ERS are llagged. as we only have a lower limit on the (YJ) colour (they are =lo in Y-baned).," Candidates ERS.YD7 and ERS.YD8 in the ERS are flagged, as we only have a lower limit on the $(Y-J)$ colour (they are $\lesssim\,1\,\sigma$ in $Y$ -band)."
" Adopting the Le lower limit on the (Y—J) colour places them in or above the ""contaminant triangular region of Figure 2.. fully consistent with entering our selection area."," Adopting the $1\,\sigma$ lower limit on the $(Y-J)$ colour places them in or above the `contaminant' triangular region of Figure \ref{fig:cc_1}, fully consistent with entering our selection area."
 Similarly. objects ERS.YD2. IS.YD5. ERS.YD9 and P34.YD7 are llagged: using a 10 lower limit on the (YJ) colour these candidates would fully meet our selection criteria (see Ligures 2. and 3)). while a more conservative 20 lower limit could. potentially locate them just. below our selection box. although with colours consistent with falling within the selection window.," Similarly, objects ERS.YD2, ERS.YD5, ERS.YD9 and P34.YD7 are flagged: using a $1\,\sigma$ lower limit on the $(Y-J)$ colour these candidates would fully meet our selection criteria (see Figures \ref{fig:cc_1} and \ref{fig:cc_2}) ), while a more conservative $2\,\sigma$ lower limit could potentially locate them just below our selection box, although with colours consistent with falling within the selection window."
 Deeper l-band imaging is required to show unambiguously that they are not in the ‘contaminant’ region of the colour:colour space., Deeper $Y$ -band imaging is required to show unambiguously that they are not in the `contaminant' region of the colour:colour space.
 Object P34.Y in P34 is also llagged. because it has a ~ 22a detection in D5the z-band.," Object P34.YD5 in P34 is also flagged, because it has a $\sim$ $\sigma$ detection in the $z$ -band."
 There are no detections in vc. 7 and Y-bands. though. so it is still a likely (2> 6) object he z-band Hux might be statistical Iluctuation or perhaps a high-equivalent-width emission line within the z-band.," There are no detections in $v$ -, $i$ - and $Y$ -bands, though, so it is still a likely high-redshift $z>6$ ) object – the $z$ -band flux might be statistical fluctuation or perhaps a high-equivalent-width emission line within the $z$ -band."
 We also llag as marginal two potential high redshift ealaxies in Ποια UDE-DP12. on the grounds that the short exposure time of the Z/-band image in this field. CFable 1)) made it impossible to measure the /Z/ig. magnitude.," We also flag as marginal two potential high redshift galaxies in field UDF-P12, on the grounds that the short exposure time of the $H$ -band image in this field (Table \ref{tab:exptimes}) ) made it impossible to measure the $H_{AB}$ magnitude."
 The upper limits on the (/I) colours. place them. away [rom the red contaminant region with (/fl)1.5 (Figure 2)). but we require the UV. luminosity in the //-flter (uncontaminated: by the elfeets o£. Lyman-a forest absorption) to infer the absolute UV. magnitude (as described in Section 4.1)).," The upper limits on the $(J-H)$ colours place them away from the red contaminant region with $(J-H)>1.5$ (Figure \ref{fig:cc_1}) ), but we require the UV luminosity in the $H$ -filter (uncontaminated by the effects of $\alpha$ forest absorption) to infer the absolute UV magnitude (as described in Section \ref{sec:lumfunc}) )."
 We now consider whether these single-band. cleteetions might be due to transients (such as was the case for the likely supernova in the WECS3 images of the HIUDE. object zD0 in Bunker ct 22010).," We now consider whether these single-band detections might be due to transients (such as was the case for the likely supernova in the WFC3 images of the HUDF, object zD0 in Bunker et 2010)."
 The P12 Ποιά was observed in J-band in two observing blocks. with S freies taken on 22009 November 02. and the other 16 frames taken over 2009 November 1015.," The P12 field was observed in $J$ -band in two observing blocks, with 8 frames taken on 2009 November 02, and the other 16 frames taken over 2009 November 10–15."
 As a check. we combined the two cillerent epochs separately with 7multidrizzle.," As a check, we combined the two different epochs separately with “multidrizzle""."
 The magnitude of P12.YDI is consistent between the two epochs. with J=28.07£0.25 (4.30) and J=27.9540.16 (6.80) respectively.," The magnitude of P12.YD1 is consistent between the two epochs, with $J=28.07\pm 0.25$ $4.3\,\sigma$ ) and $J=27.95\pm 0.16$ $6.8\,\sigma$ ) respectively."
 Lowever. D12.YD2 might show some variability in the J-band with J=27.36£013 (Ss.3sigma) For the first block of data and J=28.14E0.19 (5.8sigma) for the second.," However, P12.YD2 might show some variability in the $J$ -band with $J=27.36\pm 0.13$ $8.3\,sigma$ ) for the first block of data and $J=28.14\pm 0.19$ $5.8\,sigma$ ) for the second."
 Hence it is plausible that. P12.YD2 might. be a transient rather than a high-redshift Y-drop., Hence it is plausible that P12.YD2 might be a transient rather than a high-redshift $Y$ -drop.
 When this WECS3 program (GO-11563) is complete. the Z/-band will be much deeper on P12. allowing a further check on the robustness of the candidates in this field.," When this WFC3 program (GO-11563) is complete, the $H$ -band will be much deeper on P12, allowing a further check on the robustness of the candidates in this field."
 Llowever. the two candidates in P12 represent less than 10 per cent of our Y-drop sample. so will not quantitively affect our conclusions: for the moment. we exclude this field from our fitting of the UV. luminosity function.," However, the two candidates in P12 represent less than 10 per cent of our $Y$ -drop sample, so will not quantitively affect our conclusions; for the moment, we exclude this field from our fitting of the UV luminosity function."
 We now compare our new [ist of candidates within the IIUDE field with other groups! previous studies (Ocsch et 22010. Bouwens et 22010a. McLure et 22010. Yan et 22010 and Finkelstein ct al.," We now compare our new list of candidates within the HUDF field with other groups' previous studies (Oesch et 2010, Bouwens et 2010a, McLure et 2010, Yan et 2010 and Finkelstein et al."
 2010). and particularly with our previous paper (Bunker et 22010).," 2010), and particularly with our previous paper (Bunker et 2010)."
 A matched catalog between the Bunker et ((2010). AleLure ct ((2010) ancl Bouwens ct ((2010a) samples has already oen presented in Bunker Wilkins (2009).," A matched catalog between the Bunker et (2010), McLure et (2010) and Bouwens et (2010a) samples has already been presented in Bunker Wilkins (2009)."
 Our refined HUDE sample. based on à new reduction of rw HUDE data. has 6 Y-band drop-outs.," Our refined HUDF sample, based on a new reduction of the HUDF data, has 6 $Y$ -band drop-outs."
 In Bunker ct ((2010) we presented. a list of 7 Y-drop candidates within 1e HIUDE. field. the brightest. four (in. J-band) of which are reproduced with the new selection (LIUDE-YDI.2.3 4).," In Bunker et (2010) we presented a list of 7 $Y$ -drop candidates within the HUDF field, the brightest four (in $J$ -band) of which are reproduced with the new selection (HUDF-YD1,2,3 4)."
 Of the 3 other Y-drops from Bunker et ((2010). one (YD5) has a discrepant (Y1054Jan)=00.2 colour in the new data reduction. much bluer than our selection criteria of (VionJose)0.9.," Of the 3 other $Y$ -drops from Bunker et (2010), one (YD5) has a discrepant $(Y_{105w}-J_{125w})=0.2$ colour in the new data reduction, much bluer than our selection criteria of $(Y_{105w}-J_{125w})>0.9$."
 Phe faintest Y-drop in. Bunker et ((2010). YD7. is mareinally too faint (/=28.65) in our new reduction of the LUDE images to enter our new sample.," The faintest $Y$ -drop in Bunker et (2010), YD7, is marginally too faint $J=28.65$ ) in our new reduction of the HUDF images to enter our new sample."
 However. applying our new colour selection criteria to the old photometry (where J=28.44) would have resulted in the selection of YD.," However, applying our new colour selection criteria to the old photometry (where $J=28.44$ ) would have resulted in the selection of YD7."
 The remaining one (YD6) is only mareinally too blue for the Lyman-break selection in the newly-reduced data. with Oyios.Jpiesa)=0489. very close to the (Yriossγιου)20.9 cut.," The remaining one (YD6) is only marginally too blue for the Lyman-break selection in the newly-reduced data, with $(Y_{f105w}-J_{f125w})=0.89$, very close to the $(Y_{f105w}-J_{f125w})>0.9$ cut."
 This object has slight (~ 2260) detections in the ACS bands. too. and does not meet the selection. criterion. (Yos;οτι}0.73(JonaLlisow)| 0.9. so we did not include it in our list.," This object has slight $\sim$ $\sigma$ ) detections in the ACS bands, too, and does not meet the selection criterion $(Y_{105w}-J_{125w}) > 0.73\times (J_{125w}-H_{160w})+0.9$ , so we did not include it in our list."
 Moreover. no other group has found or listed. this object as a candidate.," Moreover, no other group has found or listed this object as a candidate."
 Two objects in our new catalog. (IIUDE.YDs. and IHIUDE.YDO) were not found in Bunker et ((2010): our previous study of Y-drops in the ICDL used. slightly cillerent magnitude and. colour cuts (Jag<28.5 and YoJig 1.0). ancl an older reduction ancl photometric zeropoints.," Two objects in our new catalog (HUDF.YD8 and HUDF.YD9) were not found in Bunker et (2010); our previous study of $Y$ -drops in the HUDF used slightly different magnitude and colour cuts $J_{AB}<28.5$ and $(Y-J)_{AB}>1.0$ ), and an older reduction and photometric zeropoints."
 These two objects were slightly too faint in the previous version of our HUDIE reductions (./=28.59 and J=2855. respectively) and slightly too blue (OYproseγιου.= O77. 0.92 respectively) to be selected with our original criteria in Bunker et ((2010).," These two objects were slightly too faint in the previous version of our HUDF reductions $J=28.59$ and $J=28.55$, respectively) and slightly too blue $(Y_{f105w}-J_{f125w})=0.77$ , $0.92$ respectively) to be selected with our original criteria in Bunker et (2010)."
 The new candidate HLUDE.YDs lies only Laaresee from the z-drop zD5 in Bunker et ((2010). and it is conceivable that both objects might be physically associated and might have similar redshifts at zὃν ," The new candidate HUDF.YD8 lies only arcsec from the $z$ -drop zD5 in Bunker et (2010), and it is conceivable that both objects might be physically associated and might have similar redshifts at $z\sim 8$."
We note that no other group has identified LLUDE.YD9 as a candidate., We note that no other group has identified HUDF.YD9 as a candidate.
 In Table 4 we show the Y -drop galaxy candidates from our LIUDE catalog which have been previously reported with their corresponding catalog names from other groups. while in Table 5. we show all the objects found by these groups," In Table \ref{tab:common} we show the $Y$ -drop galaxy candidates from our HUDF catalog which have been previously reported with their corresponding catalog names from other groups, while in Table \ref{tab:spare} we show all the objects found by these groups"
"available unipolar potential drop across the polar cap reads The αμα energv of the electrons accelerated iu tlis potential drop is s4a(ομως=7.6104B,oRep-(P/77min) 7.","available unipolar potential drop across the polar cap reads The maximum energy of the electrons accelerated in this potential drop is $\gamma_{e,M} = e \Phi_{\rm max} / m c^2 = 7.6 \times 10^4 B_{p,9}
{R^3_{\rm WD,8.7}} ({P}/{77~{\rm min}})^{-2}$ ."
 The potential drop is about 2 orders of magnitudes sinaller than that in radio pulsars., The potential drop is about 2 orders of magnitude smaller than that in radio pulsars.
 The surface laver of a white dwarf is composcel of a nou-degencrate clectron eas anc possibly an ionic lattice (Shapiro Teulolsky 1983)., The surface layer of a white dwarf is composed of a non-degenerate electron gas and possibly an ionic lattice (Shapiro Teukolsky 1983).
" The lattice imnelting temperature is T5,~NNονLeτς(p10?CresClu1δεσ 121015, where pis the density iu the surface laver. aud Z is the atoniüc umber (Alestal Ruderman 1967)."," The lattice melting temperature is $T_m \sim 8.8 \times 10^5
~{\rm K}~ (\rho / 10^2 ~{\rm ergs ~ cm^{-3}})^{1/3} (Z/12)^{5/3}$ , where $\rho$ is the density in the surface layer, and $Z$ is the atomic number (Mestal Ruderman 1967)."
 Given ai typical surface teiiperature 74~«10H K. we can see that eenerallv the surface is in the ionic lattice state.," Given a typical surface temperature $T_s \sim 3 \times 10^4$ K, we can see that generally the surface is in the ionic lattice state."
 For au anti-parallel rotator. ic. O0.B« where Q and B are the vectors for the rotational and maguetic axes. the polar cap region is populated with positively charged particles.," For an anti-parallel rotator, i.e. ${\bf \Omega \cdot B} < 0$, where ${\bf
\Omega}$ and ${\bf B}$ are the vectors for the rotational and magnetic axes, the polar cap region is populated with positively charged particles."
 Ta a co-rotating magnetic white dwarf magnetosphere. whether or not the surface can provide a free ionic flow iuto the polar cap region depends on two factors. ic. whether theriuionic enüssion could overcome the atomic cohesive energv in the surface laver. aud whether the free atoms are adequately ionized.," In a co-rotating magnetic white dwarf magnetosphere, whether or not the surface can provide a free ionic flow into the polar cap region depends on two factors, i.e, whether thermionic emission could overcome the atomic cohesive energy in the surface layer, and whether the free atoms are adequately ionized."
 In stroug magnetic fields. the ion cohesive cherev is eulianced aud depends on the streneth of the ficld. ic. xBY and is about several hundred eV. for B=10! © for iron (Jones 1986: Usov Melrose 1996).," In strong magnetic fields, the ion cohesive energy is enhanced and depends on the strength of the field, i.e. $\propto
B^{0.7}$ and is about several hundred eV for $B=10^{12}$ G for iron (Jones 1986; Usov Melrose 1996)."
 When B is lower than ~10° C. which is the general case of the white dwarf pulsar discussed here. however. the atoms are essentially uot influcucedby the magnetic field (Lai 2001). - that the atom cohesive energies are simular to the B=soM case.," When $B$ is lower than $\sim
10^9$ G, which is the general case of the white dwarf pulsar discussed here, however, the atoms are essentially not influenced by the magnetic field (Lai 2001), so that the atom cohesive energies are similar to the $B=0$ case."
" For carbon. which is Likely the composition in the WD surface laver. the cohesive cucrey is Ae,~8 eV/atoni and the first ionization euerev of carbon is Ae;~11.3 eV. The Coldveich-Julian (1969) charge umuber deusity at the maguetic pole is a,=GuPee)=1.5105«n3D,o9(PPTOdu) ldfor GCRT J1T15-3009. which is mich lower than the wumber deusitv of carbou atomis iu the surface laver. l6. Varuncpilin,=5.0«10?!cin. Ops."," For carbon, which is likely the composition in the WD surface layer, the cohesive energy is $\Delta\epsilon_c \sim 8$ eV/atom, and the first ionization energy of carbon is $\Delta\epsilon_i \sim 11.3$ eV. The Goldreich-Julian (1969) charge number density at the magnetic pole is $n_{_{\rm
GJ}}=(B_p/Pce)=1.5 \times 10^4 ~{\rm cm^{-3}}~ B_{p,9} (P/77~{\rm
min})^{-1}$ for GCRT J1745-3009, which is much lower than the number density of carbon atoms in the surface layer, i.e. $n_{\rm atom} \sim \rho/Am_p = 5.0 \times 10^{24} ~{\rm
cm^{-3}} \rho_2$ ."
 For photoionization and thermionic ejection of the ions. even if the ionization/cejection rate decreases exponentially with temperature. the critical enperatures should be still several teus siualler than the enperature defined by the colesiveionization cucrey.," For photoionization and thermionic ejection of the ions, even if the ionization/ejection rate decreases exponentially with temperature, the critical temperatures should be still several tens smaller than the temperature defined by the cohesive/ionization energy."
" Siuuilar to the troatinent for the neutron star surface (Ruclerman Sutherland 1975: Usov Melrose 1996). or the parameters of CCRT JT15-3009. this reductiou ‘actor ds ~oglZeplGT)?nin)Cn)B,2pDB|=logs 38."," Similar to the treatment for the neutron star surface (Ruderman Sutherland 1975; Usov Melrose 1996), for the parameters of GCRT J1745-3009, this reduction factor is $\sim \log [Z e
\rho (kT)^{1/2} (Am_p)^{-3/2} P/B]=\log [3.8 \times 10^{16} (kT/2.6
~{\rm eV})^{1/2} (P/77 ~{\rm min}) B_{p,9}^{-1} \rho_2 (Z/6)
(A/12)^{-3/2} \sim 38$ ."
 As a result. the critical temperatures for ionization of 1ο carbon atom aud for thernionic ejection of the carbou tonis are ~3.2«10? I and ~2.3«10? Is. respectively. voth are «OT.~«104 K. the typical temperature of je white chwarf surface.," As a result, the critical temperatures for ionization of the carbon atom and for thermionic ejection of the carbon atoms are $\sim 3.2\times 10^3$ K and $\sim 2.3\times 10^3$ K, respectively, both are $\ll T_s \sim 3\times 10^4$ K, the typical temperature of the white dwarf surface."
 The temperature at the magnetic vole could be cooler (simular to Sunspots). but as long as je temperature is higher than ~3«10? K. the white dwarf surface is able to provide copious ious to supply a Coldreich-Julian space charge linited flow from the xilar cap region (Arous Scharlemaun 1979: Tarding Aluslinov 1998).," The temperature at the magnetic pole could be cooler (similar to Sunspots), but as long as the temperature is higher than $\sim 3\times 10^3$ K, the white dwarf surface is able to provide copious ions to supply a Goldreich-Julian space charge limited flow from the polar cap region (Arons Scharlemann 1979; Harding Muslimov 1998)."
 Α similay conclusion applies for other surface compositions. as well as for the case of a parallel rotator (OQ.B0) in which case an electron free flow is pplied.," A similar conclusion applies for other surface compositions, as well as for the case of a parallel rotator ${\bf
\Omega \cdot B} > 0$ ) in which case an electron free flow is supplied."
 Iu a space-charec-linited flow. a charge depleted region is developed in the polar cap region due to the eeucral relativistic frame draseiug effect (Muslinov Tsvean 1992) and the curvature effect of the maguetic field mes (Arous Scharlemanun 1979).," In a space-charge-limited flow, a charge depleted region is developed in the polar cap region due to the general relativistic frame dragging effect (Muslimov Tsygan 1992) and the curvature effect of the magnetic field lines (Arons Scharlemann 1979)."
" The ratio between the potential drops developed for these two compoucuts is roushlv (ILudiug Miaslimov 1998) ~(8/0,etary. where \ ds the inclination angle between the magnetic aud the rotational axes. (,. is the opening augle of the polar cap region. &(RF). Ry is the Scluwarzschild radius. aud A. is the radius of the star (neutron star or white dwiutf)."," The ratio between the potential drops developed for these two components is roughly (Harding Muslimov 1998) $\sim (\kappa / \theta_{pc})~ 
{\rm ctan} \chi$, where $\chi$ is the inclination angle between the magnetic and the rotational axes, $\theta_{pc}$ is the opening angle of the polar cap region, $\kappa \sim (R_g/R_*)$, $R_g$ is the Schwarzschild radius, and $R_*$ is the radius of the star (neutron star or white dwarf)."
 For neutron star pulsars. the frame-drageine term dominates unless the inclination is near 90 degrees (Warding Abushinov 1998).," For neutron star pulsars, the frame-dragging term dominates unless the inclination is near 90 degrees (Harding Muslimov 1998)."
" For the white dwart pulsars. with the paramcters of CCRT J1715-3009. we find that αν0,.~10 5 which means that both contributions are comparable."," For the white dwarf pulsars, with the parameters of GCRT J1745-3009, we find that $\kappa \sim \theta_{pc} \sim
10^{-3}$ , which means that both contributions are comparable."
" The poteutial drop develops with height 7 in the form of (Ixidiug ADuslimov 1998) Oh)-—(2xD,/Pe(h/Rwpy which achieves the maxiuun potential Bas,I. essentially at Ro~Bap."," The potential drop develops with height $h$ in the form of (Harding Muslimov 1998) $\Phi (h) \sim (2 \pi B_p/Pc) R_{pc}^2 (h/R_{\rm
WD})^2$, which achieves the maximum potential $\Phi_{\rm max}$ essentially at $R \sim R_{\rm WD}$."
" Without pair production. clectrous iu the acceleration region could gain anu energv close to 54,44."," Without pair production, electrons in the acceleration region could gain an energy close to $\gamma_{e,M}$."
" Below we take a typical electron Lorentz factor 5,=slot.", Below we take a typical electron Lorentz factor $\gamma_e=5\times 10^4$ .
 Can the electron-positron. plasua needed by pulsar activity be eeuerated iu such white dwarts?, Can the electron-positron plasma needed by pulsar activity be generated in such white dwarfs?
 This would require generation of seed eauuna photons with energies well exceeding the electron rest enerev., This would require generation of seed gamma photons with energies well exceeding the electron rest energy.
 The dipolar field curvature radius in the wme dwarf maguectosphere is very large. Le. py=vanARR.Lbs103ein(RARwpyh?Pj 3 ," The dipolar field curvature radius in the white dwarf magnetosphere is very large, i.e. $\rho_{d}=(4/3)\sqrt{R R_{lc}}=1.4 \times 10^{11} ~{\rm
cm} ~ (R/R_{\rm WD})^{1/2} (P/77 ~{\rm min})^{1/2}$ ."
Near the surface. uou-dipolar magnetic fields could develop. as is tle case of the Sun (suuspots) and omeblv also the case of neutron stars (Ruderman Sutherland 1975: Arous Scharlemann 1979: Cal Altra 2001).," Near the surface, non-dipolar magnetic fields could develop, as is the case of the Sun (sunspots) and presumably also the case of neutron stars (Ruderman Sutherland 1975; Arons Scharlemann 1979; Gil Mitra 2001)."
" However. even if we choose a much xnaller curvature radius. ess p~10? cun the typical curvature radiation energv ds i too snl. Le. eeu,=(3/2)(hepyr?=3.7ονpo well below the pair production threshold."," However, even if we choose a much smaller curvature radius, e.g. $\rho \sim 10^9$ cm, the typical curvature radiation energy is still too small, i.e. $\epsilon_{\rm CR}=(3/2)
(\hbar c/\rho)\gamma_e^3 = 3.7 ~{\rm eV} ~ \gamma_{e,4.7}^3
\rho_9^{-1}$ , well below the pair production threshold."
" The typical thermal photou energy is ἐν~2.8&T7.375 ον for T,~3«104 K. Given the vpical electron energv and the streneth of the local maguetic field. the resonant inverse Compton(IC) scattering is uniniportant."," The typical thermal photon energy is $\epsilon_{\rm th} \sim 2.8 k T
= 7.3 T_{4.5}$ eV for $T_s \sim 3 \times 10^4$ K. Given the typical electron energy and the strength of the local magnetic field, the resonant inverse Compton (IC) scattering is unimportant."
 The typical resonant IC photon energy (Zhang ct al., The typical resonant IC photon energy (Zhang et al.
" 1997) is eft,~Se€p=D805,17Dug Κον. hich is slightly larecr than the clectrou vest eneregv ne?=511 keV. but does not mect the pair production threshold."," 1997) is $\epsilon_{\rm IC}^{R} \sim \gamma_{e} \epsilon_{\rm B} =
580 \gamma_{e,4.7} B_{p,9}$ keV, which is slightly larger than the electron rest energy $m c^2 = 511$ keV, but does not meet the pair production threshold."
 The most cfiicicut eamuna-ray production nechauisui ina white dwarf maguctosplere is non-resonant iuverse Compton (IC) scattering (Zhang et al., The most efficient gamma-ray production mechanism in a white dwarf magnetosphere is non-resonant inverse Compton (IC) scattering (Zhang et al.
 1997)., 1997).
 Thetypical eanuua-ray photon cucrey reads, Thetypical gamma-ray photon energy reads
We have studied the spin parameter of uniformly rotating compact stars in general relativity.,We have studied the spin parameter of uniformly rotating compact stars in general relativity.
 Our numerical results show that the behavior of the spin parameter of quark stars is quite different from that of neutron stars., Our numerical results show that the behavior of the spin parameter of quark stars is quite different from that of neutron stars.
" In particular, the spin parameter of neutron stars is bounded above by jmax£0.7, while quark stars can have a value larger than unity."," In particular, the spin parameter of neutron stars is bounded above by $ j_{\rm max} \approx 0.7$, while quark stars can have a value larger than unity."
" In this section, we shall discuss (in our view) the astrophysical implications of our results."," In this section, we shall discuss (in our view) the astrophysical implications of our results."
" First, how could the spin parameter of a compact star be measured?"," First, how could the spin parameter of a compact star be measured?"
" Unfortunately, so far there is no general technique to infer the spin parameter j of compact stars directly."," Unfortunately, so far there is no general technique to infer the spin parameter $j$ of compact stars directly."
" As far as we are aware, the spin parameter of a compact star could be potentially measured in disk-accreting compact-star systems."," As far as we are aware, the spin parameter of a compact star could be potentially measured in disk-accreting compact-star systems."
" In particular, the neutron stars (or quark stars) in low-mass X-ray binaries (LMXBs) provide the most natural cosmic laboratories for studying the spin parameter."," In particular, the neutron stars (or quark stars) in low-mass X-ray binaries (LMXBs) provide the most natural cosmic laboratories for studying the spin parameter."
" In order to understand how the spin parameter might be inferred in disk-accreting systems, it should be noted that the spin parameter of the central compact star directly affects the particle motion around the star."," In order to understand how the spin parameter might be inferred in disk-accreting systems, it should be noted that the spin parameter of the central compact star directly affects the particle motion around the star."
" For example, to first order in j, the orbital frequency of a point particle in a prograde orbit around a compact star is given by (see, e.g., vanderKlis (2006))) where is the orbital radius."," For example, to first order in $j$, the orbital frequency of a point particle in a prograde orbit around a compact star is given by (see, e.g., \citet{van_der_klis2006}) ) where $r$ is the orbital radius."
" For infinitesimally tilted and eccentricr orbits, the disk particles will have radial (v.) and vertical (vg) epicyclic frequencies which also depend on j vanderKlis(2006) for the expressions)."," For infinitesimally tilted and eccentric orbits, the disk particles will have radial $\nu_r$ ) and vertical $\nu_\theta$ ) epicyclic frequencies which also depend on $j$ (see \citet{van_der_klis2006} for the expressions)."
" Furthermore, (seethe combination vg—v, also gives rise to the periastron frequency of the orbit."," Furthermore, the combination $\nu_\theta - \nu_r$ also gives rise to the periastron frequency of the orbit."
" These frequencies in general depend on M and j, and hence their observations (possibly needed to be combined with the measurement of other stellar parameters such as the mass) would in principle provide useful information on the spin parameter."," These frequencies in general depend on $M$ and $j$, and hence their observations (possibly needed to be combined with the measurement of other stellar parameters such as the mass) would in principle provide useful information on the spin parameter."
" In fact, there are strong evidences that these frequencies have already been observed in LMXBs."," In fact, there are strong evidences that these frequencies have already been observed in LMXBs."
" It should be noted that existing algebraic relations, such as Equation (1)) which relate various orbital frequencies to stellar parameters are in general only valid for small spin rate j~0.1."," It should be noted that existing algebraic relations, such as Equation \ref{eq:orbital_freq}) ), which relate various orbital frequencies to stellar parameters are in general only valid for small spin rate $j \sim 0.1$."
 The difficulty of obtaining the corresponding algebraic relations for rapidly rotating stars (j~ 1) lies in the fact that there is no exact analytic representation of the vacuum spacetime outside a rapidly rotating compact star., The difficulty of obtaining the corresponding algebraic relations for rapidly rotating stars $j \sim 1$ ) lies in the fact that there is no exact analytic representation of the vacuum spacetime outside a rapidly rotating compact star.
 Having such an analytic representation of the spacetime metric will allow one to obtain the desired algebraic relations for a rapidly rotating compact star., Having such an analytic representation of the spacetime metric will allow one to obtain the desired algebraic relations for a rapidly rotating compact star.
 A starting point along this direction would be to take the closed-form asymptotically flat solution of the Einstein-Maxwell system obtained by Mankoetal. and study the geodesics in this spacetime., A starting point along this direction would be to take the closed-form asymptotically flat solution of the Einstein-Maxwell system obtained by \citet{manko2000} and study the geodesics in this spacetime.
" If (2000)there is no charge and magnetic moment, this analytic solution depends only on the mass, angular momentum, and quadrupole moment of the spacetime."," If there is no charge and magnetic moment, this analytic solution depends only on the mass, angular momentum, and quadrupole moment of the spacetime."
" Furthermore, this solution only involves rational functions, which helps to simplify the analytical study of geodesic motions."," Furthermore, this solution only involves rational functions, which helps to simplify the analytical study of geodesic motions."
" Berti&Stergioulas(2004) have demonstrated that this analytic solution can describe the exterior spacetime of a rapidly rotating neutron star (e.g., j> 0.5) very well."," \citet{berti2004} have demonstrated that this analytic solution can describe the exterior spacetime of a rapidly rotating neutron star (e.g., $j > 0.5$ ) very well."
" Nevertheless, further investigation is needed to check whether this analytic solution is also valid for rapidly rotating quark stars."," Nevertheless, further investigation is needed to check whether this analytic solution is also valid for rapidly rotating quark stars."
 One of the well-observed features of LMXBs is the high-frequency (~ kHz) QPOs., One of the well-observed features of LMXBs is the high-frequency $\sim$ kHz) QPOs.
" To date, QPOs have been observed in more than 20 LMXBs."," To date, QPOs have been observed in more than 20 LMXBs."
 These QPOs often come in pairs with frequencies νι and νι., These QPOs often come in pairs with frequencies $\nu_u$ and $\nu_l$ .
" In all systems in which the spin frequencies of the compact Stars Vstar have been measured, the frequency separation Av=v,—νι is approximately equal to vaa: OF Vstar/2."," In all systems in which the spin frequencies of the compact stars $\nu_{\rm star}$ have been measured, the frequency separation $\Delta \nu = \nu_u - \nu_l$ is approximately equal to $\nu_{\rm star}$ or $\nu_{\rm star}/2$."
 We refer the reader to vanderKlis(2006) and Lamb for recent reviews., We refer the reader to \citet{van_der_klis2006} and \citet{lamb2008} for recent reviews.
" While the physical mechanism responsible for producing the high-frequency QPOs is not known yet, most physical models involve orbital motion and disk oscillations."," While the physical mechanism responsible for producing the high-frequency QPOs is not known yet, most physical models involve orbital motion and disk oscillations."
" Hence, the frequencies ν,, vg, and various combinations of them are often invoked v; directly or indirectly) to explain high-frequency QPOs."," Hence, the frequencies $\nu_r$, $\nu_\theta$, $\nu_\phi$ and various combinations of them are often invoked (either directly or indirectly) to explain high-frequency QPOs."
 (eitherThis is the reason why one might hope to obtainuseful information on the spin parameter of the central compact star in an LMXB by observing, This is the reason why one might hope to obtainuseful information on the spin parameter of the central compact star in an LMXB by observing
lines are overplotted in red on the observed spectra in black in the bottom panels.,lines are overplotted in red on the observed spectra in black in the bottom panels.
" Our analysis has shown that the spectra of the B[e]SGs 773 and 112 display emission from the !?CO band heads (see reffig:fits)) which means that their circumstellar disk material is strongly enriched with 12Ο; hence, these stars are indeed evolved post-main sequence objects."," Our analysis has shown that the spectra of the B[e]SGs 73 and 12 display emission from the $^{13}$ CO band heads (see \\ref{fig:fits}) ) which means that their circumstellar disk material is strongly enriched with $^{13}$ C; hence, these stars are indeed evolved post-main sequence objects."
 Our model computations have also confirmed the relatively low temperature of the CO gas as indicated by the about equal strengths of the !*CO band heads., Our model computations have also confirmed the relatively low temperature of the CO gas as indicated by the about equal strengths of the $^{12}$ CO band heads.
" Such low temperatures are not expected within the canonical picture of B[e]SGs, which assumes that their disks are formed by enhanced mass flux from the stars’ equatorial regions."," Such low temperatures are not expected within the canonical picture of B[e]SGs, which assumes that their disks are formed by enhanced mass flux from the stars' equatorial regions."
" If the disks of the two B[e]SGs studied here were continuously supplied, then the CO in the gaseous parts of the disk should be visible from the hottest component having temperatures close to the CO dissociation temperature of KK. In this case the observable CO band spectrum, which always reflects the hottest available CO component, would be dominated by the peaks of the higher vibrational transitions occurring at longer wavelengths, while the lowest band head (at um) would be strongly suppressed (seeKraus2009)."," If the disks of the two B[e]SGs studied here were continuously supplied, then the CO in the gaseous parts of the disk should be visible from the hottest component having temperatures close to the CO dissociation temperature of K. In this case the observable CO band spectrum, which always reflects the hottest available CO component, would be dominated by the peaks of the higher vibrational transitions occurring at longer wavelengths, while the lowest band head (at $\mu$ m) would be strongly suppressed \citep[see][]{Kraus09}."
. The apparent deficiency in CO gas hotter that ~2800 KK therefore suggests that there is no CO gas closer to the star., The apparent deficiency in CO gas hotter that $\sim 2800$ K therefore suggests that there is no CO gas closer to the star.
" Hence, the disks seen around 112 and 773 cannot extend down to the stellar surface."," Hence, the disks seen around 12 and 73 cannot extend down to the stellar surface."
 Instead it would be more appropriate to assume that the stars are surrounded by a dense and coolring of material., Instead it would be more appropriate to assume that the stars are surrounded by a dense and cool of material.
 Support for such a scenario comes from the recent detection of a detached ring in quasi-Keplerian rotation around the Small Magellanic Cloud (SMC) B[e]SG star 665 (Krausetal.2010)., Support for such a scenario comes from the recent detection of a detached ring in quasi-Keplerian rotation around the Small Magellanic Cloud (SMC) B[e]SG star 65 \citep{Kraus2010}.
". Alternatively, gravitational darkening according to the von Zeipel theorem (vonZeipel1924) in rapidly rotating stars can lead to equatorial temperatures being sufficiently low (i.e., S3000 KK) for the hottest CO gas to not exceed ~2800 KK. In this picture, the disk still could extend from the stellar surface out to large distances, with the CO emitting gas located closest to the star within the innermost, hottest disk region."," Alternatively, gravitational darkening according to the von Zeipel theorem \citep{vonZeipel} in rapidly rotating stars can lead to equatorial temperatures being sufficiently low (i.e., $\la 3000$ K) for the hottest CO gas to not exceed $\sim 2800$ K. In this picture, the disk still could extend from the stellar surface out to large distances, with the CO emitting gas located closest to the star within the innermost, hottest disk region."
" While this scenario is somewhat speculative (and heating of the disk through light from hotter regions of the star would have to be considered as well), we can estimate the minimum rotation speed of the stars required to achieve an equatorial surface temperature of KK. Stellar rotation leads the poles to be hottest."," While this scenario is somewhat speculative (and heating of the disk through light from hotter regions of the star would have to be considered as well), we can estimate the minimum rotation speed of the stars required to achieve an equatorial surface temperature of K. Stellar rotation leads the poles to be hottest."
" Assuming that both stars are seen pole-on and that their effective temperatures, To, listed in reftab:parameters correspond to the polar values, requires a decrease in Τομ from polar to equatorial values by factors of 7.7 112) and 4 773)."," Assuming that both stars are seen pole-on and that their effective temperatures, $T_{\rm eff}$, listed in \\ref{tab:parameters} correspond to the polar values, requires a decrease in $T_{\rm eff}$ from polar to equatorial values by factors of 7.7 12) and 4 73)."
" This, in turn, requires the stars to rotate with >>99% of their critical velocity (seeFig.5ofKraus2006)."," This, in turn, requires the stars to rotate with $\gg
99\%$ of their critical velocity \citep[see Fig.\,5 of][]{Kraus06}."
". The assumed pole-on orientation only provides us withlimits of the real rotation velocity, but even in the unexpected case that both stars are seen perfectly pole-on, rotation rates close to or even at the critical limit are highly unlikely."," The assumed pole-on orientation only provides us with of the real rotation velocity, but even in the unexpected case that both stars are seen perfectly pole-on, rotation rates close to or even at the critical limit are highly unlikely."
" For comparison, the two most rapidly rotating B[e]SGs known, the SMC stars S223 and 665, rotate with =75% of their critical velocity (Zickgraf2000;Krausetal.2008,2010)."," For comparison, the two most rapidly rotating B[e]SGs known, the SMC stars 23 and 65, rotate with $\ga 75\%$ of their critical velocity \citep{Zickgraf00, Kraus08, Kraus2010}."
". If our program stars rotated at ~75% of their critical velocity, resulting equatorial temperatures would still be ~72% of the polar values (cf. reftab:parameters)),"," If our program stars rotated at $\sim 75\%$ of their critical velocity, resulting equatorial temperatures would still be $\sim 72\%$ of the polar values (cf. \\ref{tab:parameters}) ),"
" i.e. much hotter than KK. The deficiency of CO gas hotter than KK is thus real, and the most plausible scenario is hence that the molecular gas is located within a detached ring of material rather than in a disk extending down to the stellar surface."," i.e. much hotter than K. The deficiency of CO gas hotter than K is thus real, and the most plausible scenario is hence that the molecular gas is located within a detached ring of material rather than in a disk extending down to the stellar surface."
" In we include the sizes of the CO emitting regions (projected to the line of sight, Acocos i) as obtained from our"," In \\ref{tab:bestfit} we include the sizes of the CO emitting regions (projected to the line of sight, $A_{\rm CO} \cos{i}$ ) as obtained from our"
"Gamma-Ray Bursts (GRBs) without an optical counterpart of their X-ray afterglow are usually called ""dark bursts”: a more precise definition based on their X-ray to optical spectral energy distribution (SED) has been proposed (Jakobssonetal.2004:Rol2005:vanderHorstetal. 2009).","Gamma-Ray Bursts (GRBs) without an optical counterpart of their X-ray afterglow are usually called “dark bursts”; a more precise definition based on their X-ray to optical spectral energy distribution (SED) has been proposed \citep{jakobsson04,rol05,vanderhorst09}."
. It is clear now that in most cases extinction by dust in their host galaxies makes these afterglows optically dim (Greineretal.20101).. while cosmological Lyman drop out (high redshift) or intrinsic faintness are the exception rather than the rule (Cenkoetal.2009;Perley 2011).," It is clear now that in most cases extinction by dust in their host galaxies makes these afterglows optically dim \citep{greiner11}, while cosmological Lyman drop out (high redshift) or intrinsic faintness are the exception rather than the rule \citep[][]{cenko09,perley09,rossi11}."
 The long-duration lis one of these truly dark GRBs., The long-duration is one of these truly dark GRBs.
 It was detected by Swiftf/BAT on 2008 February 7 at 21:30:21 UT., It was detected by /BAT on 2008 February 7 at 21:30:21 UT.
 Swift/XRT began observing the field 124 s after the BAT trigger and found a bright X-ray afterglow. with a positional uncertainty radius of 1744 (Racusinetal. 2008a).," /XRT began observing the field 124 s after the BAT trigger and found a bright X-ray afterglow, with a positional uncertainty radius of 4 \citep{racusin08a}."
. No optical/near-infrared afterglow was detected for080207.. despite heroic efforts by several ground-based observatories.," No optical/near-infrared afterglow was detected for, despite heroic efforts by several ground-based observatories."
 The deepest limit relative to the fading X-ray afterglow was achieved by the Zeiss-600 telescope at Mt.Terskol observatory which did not detect the afterglow down to Rag>20.5 at 1.69 hr after the burst (Andreevetal, The deepest limit relative to the fading X-ray afterglow was achieved by the Zeiss-600 telescope at Mt.Terskol observatory which did not detect the afterglow down to $R_{\rm AB}>20.5$ at 1.69 hr after the burst \citep{andreev08}.
.2005). Racusinetal.(2008b.c) report that the XRT light curve between |.30hhr and hhr after the GRB start time declines monotonically with a power-law of index ay=1.85+0.10.," \citet{racusin08b,racusin08c} report that the XRT light curve between hr and hr after the GRB start time declines monotonically with a power-law of index $\alpha_X=1.85\pm0.10$."
" Therefore. we have assumed that the optical light curve in the same time interval is also decreasing monotonically. so that its behavior is sufficiently regular that the criterion of Jakobssonetal.(2004) to define GRB ""darkness"" — based on the comparison of the X-ray and optical flux levels at hhr — can be applied to the optical upper limit Rap>20.5 measured at 1.69 hr after the trigger."," Therefore, we have assumed that the optical light curve in the same time interval is also decreasing monotonically, so that its behavior is sufficiently regular that the criterion of \citet{jakobsson04} to define GRB “darkness” – based on the comparison of the X-ray and optical flux levels at hr – can be applied to the optical upper limit $R_{\rm AB}>20.5$ measured at 1.69 hr after the trigger."
 The X-ray flux at 1.69 hr is ~0.35 c/s (~3.1x100! eem). which. assuming the spectral index fitted by Racusin et al.," The X-ray flux at 1.69 hr is $\sim$ 0.35 c/s $\sim
3.1 \times 10^{-11}$ $^{-1}$ $^{2}$ ), which, assuming the spectral index fitted by Racusin et al."
 to the average XRT spectrum in the above time interval.By=1440.2. corresponds to a kkeV flux of 4.2 jJy.," to the average XRT spectrum in the above time interval, $\beta_{X}=1.4\pm0.2$, corresponds to a keV flux of 4.2 $\mu$ Jy."
 Together with thesimultaneous optical upper limit. this yields an. optical-to-X-ray index of Box<0.27. which leads to à dark GRB classification. according. to Jakobssonetal.(2004) and vanderHorstetal.(2009).," Together with thesimultaneous optical upper limit, this yields an optical-to-X-ray index of $\beta_{OX} < 0.27$, which leads to a dark GRB classification, according to \citet{jakobsson04} and \citet{vanderhorst09}."
 In the vicinity of the XRT error circle. Rossietal.(2011) found two very faint host candidates. one better visible in VIMOS/R-band and the second brighter in Κ.," In the vicinity of the XRT error circle, \citet{rossi11} found two very faint host candidates, one better visible in $R$ -band and the second brighter in $K$."
" However. the satellite observed the field of nnine days after the burst detection bySwiff (Racusinetal.2008a).. and was able to localize the GRB X-ray afterglow with very high accuracy of 07667 (Evans @== 113:50:02.97, 6 007:30:07.8 (J2000)."," However, the satellite observed the field of nine days after the burst detection by \citep{racusin08a}, and was able to localize the GRB X-ray afterglow with very high accuracy of 67 \citep{evans10}: $\alpha$ 13:50:02.97, $\delta$ 07:30:07.8 (J2000)."
" This position. coincides with one of the candidates found by Rossietal.(2011)— (source ""Bk a very faint (Raz- mmag). very red [R-K26. (R-Κα~4.7] galaxy."," This position coincides with one of the candidates found by \citet{rossi11} (source “B”): a very faint $R_{\rm AB}\sim$ mag), very red $R-K\ga$ 6, $(R-K)_{\rm AB}\sim4.7$ ] galaxy."
" The other candidate 1s bluer (dubbed""A""by.Rossietal.2011). with (R—K)ag.~2.1. and located just outside the XRT error circle. ~ 2733 northeast of the position."," The other candidate is bluer \citep[dubbed ``A'' by ][]{rossi11}, with $(R-K)_{\rm AB}\sim2.1$, and located just outside the XRT error circle, $\sim$ 3 northeast of the position."
 Hence we associate the first source (B) with080207., Hence we associate the first source (B) with.
. In this Letter. we present proprietary and archival optical and near-infrared (NIR) observations. combined with archival IIRAC and MIPS data of the host of080207.," In this Letter, we present proprietary and archival optical and near-infrared (NIR) observations, combined with archival IRAC and MIPS data of the host of."
. We estimate the photometric redshift oof the host by fitting the observed SED with a vast library of GRASIL models (Silvaetal.1998:[glestas-Páramoetal.2007;Michalowski2008. 2010).," We estimate the photometric redshift of the host by fitting the observed SED with a vast library of GRASIL models \citep{silva98,iglesias07,michal08,michal10}."
 Despite its optical faintness. because of the clear photospherie NIR bump at rest-frame .t=micron.. we are able to estimate the redshift of the host of wwith a precision of £0.3.," Despite its optical faintness, because of the clear photospheric NIR bump at rest-frame $\lambda$, we are able to estimate the redshift of the host of with a precision of $\pm$ 0.3."
 Section 2. describes the data reduction and the methods used to derive the photometry., Section \ref{sec:data} describes the data reduction and the methods used to derive the photometry.
 The SED models are presented in § 3.. and we discuss the results and their implications in 8 4..," The SED models are presented in $\S$ \ref{sec:sed}, and we discuss the results and their implications in $\S$ \ref{sec:discussion}."
" Throughout the paper. we assume a Q,, 00.3. Q4 2200.7 cosmology. with Hubble constant Hy ss! MMpe .In the course of an analysis of GRB host galaxy SEDs (Hunt et al."," Throughout the paper, we assume a $\Omega_m$ 0.3, $\Omega_\Lambda$ 0.7 cosmology, with Hubble constant $H_0$ $^{-1}$ $^{-1}$ .In the course of an analysis of GRB host galaxy SEDs (Hunt et al.,"
 in. preparation). we culled the aarchive for observations of080207.," in preparation), we culled the archive for observations of."
.. It was observed by A. Levan (PID 50562) more than one year after the burst in all four IRAC channels (Fazio and about five months after the burst at MIPS citepriekemips.., It was observed by A. Levan (PID 50562) more than one year after the burst in all four IRAC channels \citep{fazioirac} and about five months after the burst at MIPS \\citep{riekemips}.
 R- and K-band data were taken from Rossietal. (2011)..., $R$ - and $K$ -band data were taken from \citet{rossi11}. .
 B- and R-band images were acquired at the Large Binocular Telescope (LBT) in the course of our approved observing program: 7and, $B$ - and $R$ -band images were acquired at the Large Binocular Telescope (LBT) in the course of our approved observing program; $i^\prime$and
2000).,.
. When Ao7Apg. the perturbation growth becomes much slower.," When $\lambda_{\Omega} > \lambda_{BH}$, the perturbation growth becomes much slower."
 An additional constraint euters when Αι becomes comparable to the scale height fi: at this point. saturation of the MBI is thought to occur (Balbus&Hawley1998).," An additional constraint enters when $\lambda_{crit}$ becomes comparable to the scale height $h$; at this point, saturation of the MRI is thought to occur \citep{bal98}."
. Sanoetal.(2000) investigated where the MRI operates in protoplauetarv disks by including a detailed chemical reaction scheme aud the effect of the recombination of ious and clectrons on erain surfaces., \citet{san00} investigated where the MRI operates in protoplanetary disks by including a detailed chemical reaction scheme and the effect of the recombination of ions and electrons on grain surfaces.
 They showed that for the ΜΗπάσα solar nebula (fe= 1) the 23 AU reeio- of the disk would be stabilized by Olunic dissipation., They showed that for the minimum-mass solar nebula $f_{\Sigma}=1$ ) the 2–3 AU region of the disk would be stabilized by Ohmic dissipation.
 However. as the surface density paraueter. fe. is lowered. he unstable region expands (sce their Figure 8).," However, as the surface density parameter, $f_{\Sigma}$, is lowered, the unstable region expands (see their Figure 8)."
 Also. as the disk evolves. dust evains settle to a thin laver in he midplane. and so the erai abundance in the bulk of he nebula eradually decreases.," Also, as the disk evolves, dust grains settle to a thin layer in the midplane, and so the grain abundance in the bulk of the nebula gradually decreases."
 For this reason. Sano oet al. (," For this reason, Sano et al. ("
2000) introduced. a dust depletion factor. fy. which ucasures the abundance of dust erains relative to that in uolecular clouds.,"2000) introduced a dust depletion factor, $f_d$, which measures the abundance of dust grains relative to that in molecular clouds."
 Thus a low value of £; 1i. correspoud o an old. evolved. disk.," Thus a low value of $f_d$ may correspond to an old, evolved disk."
" They demonstrated that as fi, decreases for a given surface deusity. the stable region (the ""dead gone”) slau iu size. aud for fy<107. the entire disk becomes unstable for Ro>3 AU (see their Figure 11)."," They demonstrated that as $f_d$ decreases for a given surface density, the stable region (the “dead zone”) shrank in size, and for $f_d \lesssim 10^{-2}$, the entire disk becomes unstable for $R > 3$ AU (see their Figure 11)."
 This happens because the recombination rate on grain surfaces decreases., This happens because the recombination rate on grain surfaces decreases.
 They found a similar trend as erain erowthn occurs., They found a similar trend as grain growth occurs.
 From these results. we conclude that the region at Rw3 AU will become magnetically active at a sutiicicutly late stage of the evolution.," From these results, we conclude that the region at $R \approx 3$ AU will become magnetically active at a sufficiently late stage of the evolution."
 We use auc diseuss this result further in rofapplv.., We use and discuss this result further in \\ref{apply}.
 Figure compares the three relevant length scales at an carly (a) aud a late stage (b) of solar nebula evolution., Figure \ref{fig1} compares the three relevant length scales at an early (a) and a late stage (b) of solar nebula evolution.
 The region of interest where magnetic fields amplify due to the MRI extends from 2 AU to 1. AU in (aj). while with a reduced gas density. the active region moves inward to 12 AU (b).," The region of interest where magnetic fields amplify due to the MRI extends from 2 AU to 4 AU in (a), while with a reduced gas density, the active region moves inward to 1–2 AU (b)."
 Iu this aud the following sections. we cousider a simple. one-dimensional ecoetry illustrated iu Figure 2..," In this and the following sections, we consider a simple, one-dimensional geometry illustrated in Figure \ref{fig2}."
 Also shown are the directions of various terms in Equs. (33) , Also shown are the directions of various terms in Eqns. \ref{ve}) )
and CL)., and \ref{ind2}) ).
 Consider first a case in which a magnetic uull is present., Consider first a case in which a magnetic null is present.
 Near the uull a magnetic pressure eradieut. VGP/8z). acts ou ions aud pushes the feld lines into he null poiut iu both directious.," Near the null, a magnetic pressure gradient, $-\bnabla (B^2/8\pi)$, acts on ions and pushes the field lines into the null point in both directions."
 The ion-neutral diift is hen directed toward the null., The ion-neutral drift is then directed toward the null.
 Alone a gradient in the uaenetie feld. this nonliuear term produces faster field diffusion iu regions of stronger field. driving field iuto regions of weaker field aud steepenuiug the exadieut. rather han spreading it out as the linear Olumic term does.," Along a gradient in the magnetic field, this nonlinear term produces faster field diffusion in regions of stronger field, driving field into regions of weaker field and steepening the gradient, rather than spreading it out as the linear Ohmic term does."
 Ultimately. the eradicut sharpens towards a singularity. uarked by a sheet of hieh electric curreut.," Ultimately, the gradient sharpens towards a singularity, marked by a sheet of high electric current."
" This steepeuiug of magnetic field profile (BZ91) occurs within a timescale of Tap=I(2A4p)GO/6)5,707). caleulated o be z/12 of the orbital period when v;=v4."," This steepening of magnetic field profile (BZ94) occurs within a timescale of $\tau_{AD} \equiv L_{BH}^2/(2 \, \lambda_{AD}) = (\pi^2/\,6) \, (\nu_{in}/\Omega^2)$, calculated to be $\pi/12$ of the orbital period when $\chi_i=\chi_{crit}$."
 The actor of 1/2 in τι takes mto account the fact that the characteristic leugth decreases as steepeniug proceeds., The factor of 1/2 in $\tau_{AD}$ takes into account the fact that the characteristic length decreases as steepening proceeds.
Remarkably. this sharpening has been shown to occur even in the absence of uulls due to a variety of processes including a rotating eddy (201.Zweibel&burg 1997).. accretion disks sustainiug the MRI etal. 1995). lavee amplitude Alfven waves1996).. aud the nonlinear stage of Wardle C-shocks (MacLow&Sinith1997).,"Remarkably, this sharpening has been shown to occur even in the absence of nulls due to a variety of processes including a rotating eddy \citep[BZ94;][]{zwe97}, accretion disks sustaining the MRI \citep{mac95}, large amplitude Alfven waves, and the nonlinear stage of Wardle C-shocks \citep{mac97}."
. As previously mentioned. we ake the example of maguetolbydrodvuamic turbulence dviven by the MRI," As previously mentioned, we take the example of magnetohydrodynamic turbulence driven by the MRI."
 Iu this case. velocity shears produce naenetic nulls aud current siugulanties.," In this case, velocity shears produce magnetic nulls and current singularities."
 Stretched by shear flows. magnetic nulls develop where the magnetic field changes directions.," Stretched by shear flows, magnetic nulls develop where the magnetic field changes directions."
 We therefore crudely estimate hat the separation between two adjacent magnetic nulls Lau=Apu/?. We note that ambipolar diffusion allows the field to slip hrough the neutrals and tends to weaken the MBI., We therefore crudely estimate that the separation between two adjacent magnetic nulls $L_{BH}=\lambda_{BH}/2$ We note that ambipolar diffusion allows the field to slip through the neutrals and tends to weaken the MRI.
" For strong οποιο] diffusion. 1.0. when the iou-neutral collision vate (v7j,) is πιο ligher than the erowth rate of the MRI (—9). the iustabilitv is eutirelvprevented (MacLowetal.1995:Tlawley&Stone 1998).."," For strong enough diffusion, i.e. when the ion-neutral collision rate $\nu_{in}$ ) is much higher than the growth rate of the MRI $\sim\Omega$ ), the instability is entirelyprevented \citep{mac95,haw98}. ."
 On the other haud. ambipolar diffusion is required to generate curent sheets.," On the other hand, ambipolar diffusion is required to generate current sheets."
 For our model. we use parameters on the interface between the two reguues.," For our model, we use parameters on the interface between the two regimes."
" This choice is justified by dimensional simulations performed by MacLowetal.(1995)... which showed that for drag cocficieuts 5 a few times the miununni value that allows the instability. to develop. 5,5=Of(Gn;nyNeg) both the instability aud aüubipolar diffusion sharpening occur."," This choice is justified by two-dimensional simulations performed by \citet{mac95}, which showed that for drag coefficients $\gamma$ a few times the minimum value that allows the instability to develop, $\gamma_{min} \equiv \Omega /(m_i \, n_n \, \chi_{crit})$, both the instability and ambipolar diffusion sharpening occur."
" In other words. we require LS5/5,;4,100."," In other words, we require $1 \lesssim \gamma/\gamma_{min} < 100$."
 Alteruatively. magnetic field auplification could have preceeded the epoch of current sheet formation.," Alternatively, magnetic field amplification could have preceeded the epoch of current sheet formation."
 We now lake quantitative estimates of the propertics of current sheets;, We now make quantitative estimates of the properties of current sheets.
 BZOL showed that. in the absence of resistivity. the magnetic field profile relaxes to Box wand the current density JoxOD/Orss2 becomes singular at the origin.," BZ94 showed that, in the absence of resistivity, the magnetic field profile relaxes to $B \propto x^{1/3}$ , and the current density $J \propto \partial B/\partial x \propto x^{-2/3}$ becomes singular at the origin."
 In realitv. both the finite resistivity and the ion pressure act to remove this suguluitv (Drandeuburg&Zaweibel1995j.," In reality, both the finite resistivity and the ion pressure act to remove this singularity \citep{bra95}."
. Under the physical couditious of protoplanetary disks. om pressure can be neglected (Ieitsch&Zxveibel2003).," Under the physical conditions of protoplanetary disks, ion pressure can be neglected \citep{hei03}."
. The induction term in Equ. (1)). V«(eB).," The induction term in Eqn. \ref{ind2}) ), $\bnabla \times (\bov \times \bb)$,"
 becomes progressively less important as the field steepeus: Z/Ozzο) where I and O denote the magnitudes of the induction aud Olunic dissipation term respectively. aud is therefore ignored.," becomes progressively less important as the field steepens; $I/O \approx vL/\eta$ where $I$ and $O$ denote the magnitudes of the induction and Ohmic dissipation term respectively, and is therefore ignored."
" The quasi-steady-state maeuctic field profile (""Quasi since the magnetic field profile lasts only as lone as the velocity shear is present. for about a dynamical time) is calculated bv balancing the Olunic dissipation term against the züubipolar diffusion term iu Equ. (LI):"," The quasi-steady-state magnetic field profile (“Quasi” since the magnetic field profile lasts only as long as the velocity shear is present, for about a dynamical time) is calculated by balancing the Ohmic dissipation term against the ambipolar diffusion term in Eqn. \ref{ind2}) ):"
 Tere.for simplicity. we use the definition which ieasures the )relative importanceTjj of Olunic dissipation to ambipolar diffusion.," Here,for simplicity, we use the definition which measures the relative importance of Ohmic dissipation to ambipolar diffusion."
" We used 7415 and Ty; to denote »,/(1015coin?) and T,(10KR). respectively,"," We used $n_{g13}$ and $T_{g3}$ to denote $n_g/(10^{13} \; \rm{cm^{-3}})$ and $T_g/(10^3 \; \rm{K})$, respectively."
 Olunic dissipation dominateswhere B<By: ambipolar diffusion wins where P>By (see Equ. 7))., Ohmic dissipation dominateswhere $B < B_0$; ambipolar diffusion wins where $B > B_0$ (see Eqn. \ref{ratio}) ).
 Remarkably. Bo is indepeudent of the ionization fraction: so is the magnetic field profile.," Remarkably, $B_0$ is independent of the ionization fraction; so is the magnetic field profile."
Solving Equ. (10)).,"Solving Eqn. \ref{balance}) ),"
 we obtain asstunine that the field vanishes at the origin., we obtain assuming that the field vanishes at the origin.
" Au integration constant. D. is determined by the boundary condition B,(cLop)= Do. which leads to DP= 3)/Lpy."," An integration constant, $\Gamma$ , is determined by the boundary condition $B_y(x=L_{BH})=B_{max}$ which leads to $\Gamma = B_{max}(B_0^2+B_{max}^2/3)/L_{BH}$ ."
" The amplitude of the magnetic field. B, dis determined frou the roiinaut magneticfield"," The amplitude of the magnetic field, $B_{max}$ , is determined from the remnant magneticfield"
"was found to best match the Magorrian relation between bulge mass and black hole mass, Man~Mie; observed by Háring&Rix(2004).","was found to best match the Magorrian relation between bulge mass and black hole mass, $M_{\rm BH} \sim M_{\rm bulge}^{1.12}$, observed by \scite{Haring04}."
". In this study, the black hole accretion fraction is set to zero for the purposes of consitency with the simulation, where this aspect of evolution is not considered."," In this study, the black hole accretion fraction is set to zero for the purposes of consitency with the simulation, where this aspect of evolution is not considered."
" Being physically well motivated, the condition (81]) has been applied throughout this study."," Being physically well motivated, the condition \ref{StabilityLimit}) ) has been applied throughout this study."
" Conveniently, the precise formalism used need not cause too much concern here; the change that its application produces in the component masses (in Fig. ??,,"," Conveniently, the precise formalism used need not cause too much concern here; the change that its application produces in the component masses (in Fig. \ref{baryons},"
 for example) is barely enough to be noticed., for example) is barely enough to be noticed.
" This can be understood from Fig. ??,,"," This can be understood from Fig. \ref{stab},"
 which shows that only a very limited adjustment to the disk mass is required to prevent the limit in from being reached., which shows that only a very limited adjustment to the disk mass is required to prevent the limit in \ref{StabilityLimit}) ) from being reached.
" The mass aggregation of this particular (1)system, and the fact that the specific angular momentum of the gas is relatively (A>0.07 for the latter half of the halo’s history) leads high""to the stabilisation of the system, even if the redistribution of mass is not enforced."," The mass aggregation of this particular system, and the fact that the specific angular momentum of the gas is relatively $\lambda > 0.07$ for the latter half of the halo's history) leads to the stabilisation of the system, even if the redistribution of mass is not enforced."
 The size and rotation speed are shown for both realisations in Fig. ??.., The size and rotation speed are shown for both realisations in Fig. \ref{rotation}.
 The predicted value is based on conserving the initial angular momentum (19)) of in-falling gas; an assumption which is particularly effective in this case where the distributions for the initial and final mass were appropriate (Fig)., The predicted value is based on conserving the initial angular momentum \ref{haloAM}) ) of in-falling gas; an assumption which is particularly effective in this case where the distributions for the initial and final mass were appropriate \ref{halogas}) ).
" This correspondence is encouraging for the approach of calculating disk sizes from the spin parameter, A, of the halo from which the gas cooled."," This correspondence is encouraging for the approach of calculating disk sizes from the spin parameter, $\lambda$, of the halo from which the gas cooled."
" The total cooled mass (stars and cold gas), calculated using the methods of refAccretion,, can be seen from Fig. ??.."," The total cooled mass (stars and cold gas), calculated using the methods of \\ref{Accretion}, can be seen from Fig. \ref{baryons}."
" Agreement with the simulation is excellent, given the inherent difference in the methods of calculation, and bolsters confidence in the rather intricate calculation of the energy radiated by the system throughout its complex merger history (13}{T8))."," Agreement with the simulation is excellent, given the inherent difference in the methods of calculation, and bolsters confidence in the rather intricate calculation of the energy radiated by the system throughout its complex merger history \ref{e_th}- \ref{r_cool}) )."
 The division of the cooled mass into stars and gas are also shown in Fig. ??.., The division of the cooled mass into stars and gas are also shown in Fig. \ref{baryons}.
" The evident agreement can be understood by reference to Fig’s ?? and ?7,, which demonstrate that the considerable J,complexities of the simulation reduce to relatively simple relationships when integrated over the entire disk."," The evident agreement can be understood by reference to Fig's \ref{halogas}, \ref{SFtimescale} and \ref{feedback}, which demonstrate that the considerable complexities of the simulation reduce to relatively simple relationships when integrated over the entire disk."
coronal backeround.,coronal background.
 However. limitations exist when the cospatial cross-sectional background profile significantly deviates from a linear interpolation. or when substantial density and temperature eradieuts along the averaged loop segment exist. in which case a self&consistent DEM solution may be iulibited.," However, limitations exist when the cospatial cross-sectional background profile significantly deviates from a linear interpolation, or when substantial density and temperature gradients along the averaged loop segment exist, in which case a self-consistent DEM solution may be inhibited."
 In addition. the Poissonian photon noise aud calibration errors contribute to the uncertainty of DEM fits.," In addition, the Poissonian photon noise and calibration errors contribute to the uncertainty of DEM fits."
 Tow do we explain the result of predominant isothermal loops iu terms of plivsical models?, How do we explain the result of predominant isothermal loops in terms of physical models?
 Let us first discuss isotherinalitv iu terms of nauoflare models., Let us first discuss isothermality in terms of nanoflare models.
 Since nanofares occur ou unresolved spatial scales aid have no cross-ficld transport. every nanoflare model predicts a highly iuhomogeueous density and temperature structure of macroscopic loops (IKInuchuk 2006).," Since nanoflares occur on unresolved spatial scales and have no cross-field transport, every nanoflare model predicts a highly inhomogeneous density and temperature structure of macroscopic loops (Klimchuk 2006)."
 The only way to make a nanoflaring loop structure more isothermal is to svuchronize a storm of nanofiares and to streamline them to identical energv outputs. so that their release appears to be simultaneous and homogeneous across a loop cross-section. which iu the continuum limit is macroscopically iudistiguislhable from a monolithic loop.," The only way to make a nanoflaring loop structure more isothermal is to synchronize a storm of nanoflares and to streamline them to identical energy outputs, so that their release appears to be simultaneous and homogeneous across a loop cross-section, which in the continuum limit is macroscopically indistiguishable from a monolithic loop."
 Following Occam's razor. assumndue a coherent isothermal upflow in a single flux tube is a weaker assumption than svuchrouized uptlows with equal temperatures in a multithread structure.," Following Occam's razor, assuming a coherent isothermal upflow in a single flux tube is a weaker assumption than synchronized upflows with equal temperatures in a multi-thread structure."
 Towever. a strouger argument that currently supports a fragmented energv release. such as uanoflares. is the observed lifetime of some coronal loops that are more extended than expected for an πρι]νο heating phase with subsequent cooling (c.e.. Warren et al.," However, a stronger argument that currently supports a fragmented energy release, such as nanoflares, is the observed lifetime of some coronal loops that are more extended than expected for an impulsive heating phase with subsequent cooling (e.g., Warren et al."
 2008)., 2008).
 However. the time evolution of isolated loop strauds has first to be studied with high spatial resolution aud hieh time cadence in many temperature filters. such as ATA provides. before clear-cut cases can be established.," However, the time evolution of isolated loop strands has first to be studied with high spatial resolution and high time cadence in many temperature filters, such as AIA provides, before clear-cut cases can be established."
 At this point we can only conclude that the observed isothermality of coronal loops is uot consistent with standard nanoflare scenarios. nor do nanoflare mocels explain or predict the isothermal property.," At this point we can only conclude that the observed isothermality of coronal loops is not consistent with standard nanoflare scenarios, nor do nanoflare models explain or predict the isothermal property."
 If au isothermal loop cross-section cannot be produced by nanoflares. what other physical process cau account for it?," If an isothermal loop cross-section cannot be produced by nanoflares, what other physical process can account for it?"
 The most natural mechauisin seeuis to be that of chromosphierie evaporation as known in flares (oe. Antomucel et al.," The most natural mechanism seems to be that of chromospheric evaporation as known in flares (e.g., Antonucci et al."
 1999). where either coronal or chromospheric magnetic reconnection processes Cause a local heating of the chromosphere auc drive colercut upflows of heated plasina iuto a coroual loop conduit.," 1999), where either coronal or chromospheric magnetic reconnection processes cause a local heating of the chromosphere and drive coherent upflows of heated plasma into a coronal loop conduit."
 Although the details of the heating cross-section transverse to the maeuetic field is not fully understood. flare observations vield clear evidence that postflare loops are filled impulsively with heated plasima over a cross-section of several thousand kilometers.," Although the details of the heating cross-section transverse to the magnetic field is not fully understood, flare observations yield clear evidence that postflare loops are filled impulsively with heated plasma over a cross-section of several thousand kilometers."
" Distributions of hard X-ray. footpoiut sources in flare loops have typical EWIIM of z2”SZ, oy iz1.5.6.0 Mui. according to RITESSI iieasureiments with a spatial resolution of z2"" (Dennis aud Pernals 2009)."," Distributions of hard X-ray footpoint sources in flare loops have typical FWHM of $\approx 2\arcsec-8\arcsec$, or $w \approx 1.5-6.0$ Mm, according to RHESSI measurements with a spatial resolution of $\approx 2\arcsec$ (Dennis and Pernak 2009)."
 One of the few flare observations that resolve the bluc-shifted upflows from near-cospatial red-shifted dowuflows is described in Czavkowska et al. (, One of the few flare observations that resolve the blue-shifted upflows from near-cospatial red-shifted downflows is described in Czaykowska et al. (
1999). which vields evidence for coherent near-isothermial upflows over a cross-section of zzI (5:3 Maa).,"1999), which yields evidence for coherent near-isothermal upflows over a cross-section of $\approx 4\arcsec$ $\lapprox 3$ Mm)."
 Similarly. IBuode/EIS spectroscopic observations showed unear-isothermal upflows in active region loops (Tripathi et al.," Similarly, Hinode/EIS spectroscopic observations showed near-isothermal upflows in active region loops (Tripathi et al."
 2009)., 2009).
 The heating process at a cospatial location lasts iu the order of münutes for flare loops. while a flare can last for hours. with the energy release propagating along and perpendicular to the neutral line.," The heating process at a cospatial location lasts in the order of minutes for flare loops, while a flare can last for hours, with the energy release propagating along and perpendicular to the neutral line."
 Of course. tle uptlows in flare loops are more or less continuous during the heating phase. aud thus no thermal equilibria. is reached that obevs the theoretically expected couductive and radiative cooling. times. explaimine the discrepancies between observed loop lifetimes aud theoretically calculated cooling times (Warren et al.," Of course, the upflows in flare loops are more or less continuous during the heating phase, and thus no thermal equilibrium is reached that obeys the theoretically expected conductive and radiative cooling times, explaining the discrepancies between observed loop lifetimes and theoretically calculated cooling times (Warren et al."
 2008)., 2008).
 The temperature of an arbitrary loop cross-section can be nearly coustant during some part of the heating phase due to the continuous upflow of isothermal plasma. which explains the slow time evolution observed in coronal loops during time intervals of hours (see Fie.," The temperature of an arbitrary loop cross-section can be nearly constant during some part of the heating phase due to the continuous upflow of isothermal plasma, which explains the slow time evolution observed in coronal loops during time intervals of hours (see Fig."
 7)., 7).
 It secs natural to sugeestOO a flare-like heating mechanisi for active region loops. although there night be significaut differences.," It seems natural to suggest a flare-like heating mechanism for active region loops, although there might be significant differences."
 While the coronal height of, While the coronal height of
cpus) code. by imposing that all the particles are active.,"cpus) code, by imposing that all the particles are active."
 “Phis indeed represents the most non uniform situation., This indeed represents the most non uniform situation.
 It turns out that the deviation due to the interpolation amounts at maximun to 0S4., It turns out that the deviation due to the interpolation amounts at maximun to $0.8\%$.
 The maximum deviation is located at the time of the maximum compression (t. = 0.88), The maximum deviation is located at the time of the maximum compression (t $\approx$ 0.88).
 To evaluate the code performances. we use the acliabatic collapse just. described. and perform simulations αἱ increasing number of processors.," To evaluate the code performances, we use the adiabatic collapse just described, and perform simulations at increasing number of processors."
 We believe that this test is very stringent. and can give a lower limit of the code performances due to the high density contrast that is present at the time of maximum compression. when the particles are highly clustered.," We believe that this test is very stringent, and can give a lower limit of the code performances due to the high density contrast that is present at the time of maximum compression, when the particles are highly clustered."
 We are going to check the code timing. overall load-halance anc scalability.," We are going to check the code timing, overall load-balance and scalability."
 Moreover we shall analyze in details particular sections of the code. like the gravity computations.the SPII and the neighbor searching.," Moreover we shall analyze in details particular sections of the code, like the gravity computations,the SPH and the neighbor searching."
 An estimate of the parallel over-head will be given as well., An estimate of the parallel over-head will be given as well.
 We run the adiabatie collapse test up to the time of the maximum compression (t & 1.1) using 2lot particles on l. 2. 4. 8. 16. 32 and 64 processors. and. looked. at the performances in the following code sections (see also ‘Table 1): The results summarized. in Table 1 present the total walkclock time per processor over the last 50 lime-steps. together with the time spent in cach of the 5 subroutines (cata updating. neighbor searching. SPL computation. gravitational interaction ane parallel computation).," We run the adiabatic collapse test up to the time of the maximum compression (t $\simeq 1.1$ ) using $2 \times 10^{4}$ particles on 1, 2, 4, 8, 16, 32 and 64 processors, and looked at the performances in the following code sections (see also Table 1): The results summarized in Table 1 present the total wall-clock time per processor over the last 50 time-steps, together with the time spent in each of the 5 subroutines (data updating, neighbor searching, SPH computation, gravitational interaction and parallel computation)."
 Phe eravitation interaction takes about one-hird of the total time. while the search for neighbors takes roughly comparable time.," The gravitation interaction takes about one-third of the total time, while the search for neighbors takes roughly comparable time."
 Phe evaluation of hvdrodynamical quantities (see Section 3.5) takes about one-fourth of the ime. the remaining time being divided between L/O ancl data up-dating.," The evaluation of hydrodynamical quantities (see Section 3.5) takes about one-fourth of the time, the remaining time being divided between I/O and data up-dating."
 Phe parallel over-head does not appear to be a serious problem. being at maximum about 1% of the total ime.," The parallel over-head does not appear to be a serious problem, being at maximum about $1\%$ of the total time."
 Fhis timing refers. as indicated above. to simulations stopped at roughly the time of maximum compression.," This timing refers, as indicated above, to simulations stopped at roughly the time of maximum compression."
 A run with 8 processors up to /&2.5. the time at which the svsten is almost completely virialized. took 3800 secs.," A run with 8 processors up to $t \simeq 2.5$, the time at which the system is almost completely virialized, took 3800 secs."
 Global code performances are analvzed in the next sections., Global code performances are analyzed in the next sections.
 One of the most stringent requirements for a parallel code is the capability to distribute. the computational work equally between all processors., One of the most stringent requirements for a parallel code is the capability to distribute the computational work equally between all processors.
 This can be done defining a suitable work-loacl criterion. as cliscussecl in Section 3.2.," This can be done defining a suitable work-load criterion, as discussed in Section 3.2."
 This is far from being an easy task (Dave et al 1997). and in practice some processors stand. idlv for some time waiting that the processors with the heaviest computational load. accomplish. their work.," This is far from being an easy task (Davè et al 1997), and in practice some processors stand idly for some time waiting that the processors with the heaviest computational load accomplish their work."
 Ες is true. also when an asvnchronous communication scheme is adopted. as in our ‘TreeSPLL code.," This is true also when an asynchronous communication scheme is adopted, as in our TreeSPH code."
 Xs outlined. in Section 3> we are using individual work-Ioads. based on the time spent to evaluate he gravitational interaction on one particle with all the other ones.," As outlined in Section 3.2, we are using individual work-loads, based on the time spent to evaluate the gravitational interaction on one particle with all the other ones."
 A better choice would be to define the work-oac only for active particles. which are the particles evolving ally.," A better choice would be to define the work-load only for active particles, which are the particles evolving fatly."
 This possibility is currently under investigation. due o the memory problems that can arise. as discussed. in Section 3.5.," This possibility is currently under investigation, due to the memory problems that can arise, as discussed in Section 3.5."
 To evaluate the code Ioad-balance we adopted he same strategv of Dave et al (1997). measuring the pactional amount of time spent idle in a time-step while another processor performs Computation: Llere fines is the time spent bv the slowest processor. while /; is the time taken by the 7ff processors to perform computation.," To evaluate the code load-balance we adopted the same strategy of Davè et al (1997), measuring the fractional amount of time spent idle in a time-step while another processor performs computation: Here $t_{max}$ is the time spent by the slowest processor, while $t_i$ is the time taken by the $i-th$ processors to perform computation."
 The results are shown in Fig 5. where we plot the [oad-balance for simulations at increasing. number of processors. from 1 to 64.," The results are shown in Fig 5, where we plot the load-balance for simulations at increasing number of processors, from 1 to 64."
 The load balance maintains always, The load balance maintains always
"associated with Pj, are generally too large to make statistically significant statements about changes in the iron line flux. even though there are indications for this (e.g. Fig. 10):","associated with $F_{\rm K\alpha}$ are generally too large to make statistically significant statements about changes in the iron line flux, even though there are indications for this (e.g. Fig. \ref{fig11-ascaflareplt}) );"
 notice especially thedifference in Zi; between andΑ2., notice especially thedifference in $F_{\rm K\alpha}$ between and.
 We next investigate the time intervals surrounding the flare (Fig. 119: , We next investigate the time intervals surrounding the flare (Fig. \ref{fig12-xteflare}) );
due unfortunately to an absence of data during this time interval. we cannot make analogous Comparisons.," due unfortunately to an absence of data during this time interval, we cannot make analogous comparisons."
 However. we find that changes to the intrinsic power law slope during these intervals follow similar trends presented for the counterpart of the flare in the previous section. although events surrounding this bright— flare event appear to be much more erratic and complicated. (," However, we find that changes to the intrinsic power law slope during these intervals follow similar trends presented for the counterpart of the flare in the previous section, although events surrounding this bright flare event appear to be much more erratic and complicated. ("
We remind the reader that is ∡⋡also apparent in :hardness ratio comparisons: during: this. time ∣∢∣⋘⋨⋮⇂∥⊲∊,We remind the reader that this is also apparent in hardness ratio comparisons during this time interval.)
∥⊳⋅∣∖⋟⇀∖∕∣↳∣∣↖≟∣∏∟⋅⋯⋅∖∟⋅↳∣∟⋅∣⋯⋅∐∏↖≟⋔∪∣⋯∟⋅∣⋅∖⊲∙⇂∣∖⋅↭∣∣⋯∟⋅∣⋅∟⋅∖⋅⇂∣⇆∣⋅∪∟⋅∪↳∣∣∏≌. ⋡similarities with thes ASCA2. flare in observed followingtrends to changes in DL; io., The similarities with the flare lie in observed trends to changes in $\Gamma_{3-10}$.
" Inparticular, there is a noticeable flattening of LF; 0 inthe transition between pre-  and  post- ""flare states to X2. which continues through intervalm. before"," In particular, there is a noticeable flattening of $\Gamma_{3-10}$ in the transitionbetween ` pre- ' and ` post- ' flare states to , which continues through interval, before"
angular velocitv. the Coriolis force is stronger than in the solar case. so meridional flow does not be fast.,"angular velocity, the Coriolis force is stronger than in the solar case, so meridional flow does not be fast."
 In our model. meridional flow generates latitudinal entropy. eracdient in the subaciabatically stratified overshoot region.," In our model, meridional flow generates latitudinal entropy gradient in the subadiabatically stratified overshoot region."
 Since the meridional [low is not fast. the entropy eracdient is insufficient to move differential rotation far from the Tavlor-Proucdimaan state in rapidly rotating stars.," Since the meridional flow is not fast, the entropy gradient is insufficient to move differential rotation far from the Taylor-Proudman state in rapidly rotating stars."
 As a result. the differential rotation of stars with large stellar angular velocity is close to the Tavlor-Proudiman The temperature difference between latitudes is probably controlled by (wo important factors. i.e.. the subachabatic laver below the convection zone and anisotropic heat transport caused by turbulence aud rotation.," As a result, the differential rotation of stars with large stellar angular velocity is close to the Taylor-Proudman The temperature difference between latitudes is probably controlled by two important factors, i.e., the subadiabatic layer below the convection zone and anisotropic heat transport caused by turbulence and rotation."
 We suggest that the former is important in slow rotators like the sun. and the latter in rapid rotators.," We suggest that the former is important in slow rotators like the sun, and the latter in rapid rotators."
 The subadiabatic-laver effect is included in our model. while anisotropic heat transport is not.," The subadiabatic-layer effect is included in our model, while anisotropic heat transport is not."
 We found that the effect of the subacdiabatie laver can generate a temperature difference AT=10Ix in the solar case. which moderately increases with higher rotation speeds. and AT=30K in case O4=SQ...," We found that the effect of the subadiabatic layer can generate a temperature difference $\Delta T=10\ \mathrm{K}$ in the solar case, which moderately increases with higher rotation speeds, and $\Delta T=30\ \mathrm{K}$ in case $\Omega_0=8\Omega_\odot$."
 The ihree-dimensional simulations by Brownetal.(2008) include a self-consistent caleulation of anisotropy. of turbulent thermal transport but not the subaciabatic laver at the bottom boundary., The three-dimensional simulations by \cite{2008ApJ...689.1354B} include a self-consistent calculation of anisotropy of turbulent thermal transport but not the subadiabatic layer at the bottom boundary.
 In (heir caleulation AT is most likely smaller than LOIx in the solar case. since {μον cannot reproduce the solar differential rotation only with anisotropy. of thermal transport.," In their calculation $\Delta T$ is most likely smaller than $10\ \mathrm{K}$ in the solar case, since they cannot reproduce the solar differential rotation only with anisotropy of thermal transport."
 Also. AZ=LOOIx in case Oy=50... which is lareer than the case with the subacdiabatic laver.," Also, $\Delta T=100\ \mathrm{K}$ in case $\Omega_0=5\Omega_\odot$, which is larger than the case with the subadiabatic layer."
 We speculate that anisotropic heat (ransport becomes more significant in rapidly rotating stars., We speculate that anisotropic heat transport becomes more significant in rapidly rotating stars.
 There is also a possibilitv that our caleulated entropy gradient at the base of the convection zone can be used as a boundary condition for a self-consistent tliree dimensional simulation of stellar convection (Mieschetal.2006)..., There is also a possibility that our calculated entropy gradient at the base of the convection zone can be used as a boundary condition for a self-consistent three dimensional simulation of stellar convection \citep{2006ApJ...641..618M}.
 Note that differential rotation in rapidly rotating stars in IxXüker&Stix(2001) is not in the Tavlor-Proudiman state when anisotropy of turbulent thermal conductivity is included., Note that differential rotation in rapidly rotating stars in \cite{2001A&A...366..668K} is not in the Taylor-Proudman state when anisotropy of turbulent thermal conductivity is included.
 A future study of the simultaneous effects of the attached subacliabatic laver beneath convection zone ancl anisoloropy of the turbulent thermal conductivity on stellar differential rotation would, A future study of the simultaneous effects of the attached subadiabatic layer beneath convection zone and anisotoropy of the turbulent thermal conductivity on stellar differential rotation would
he most huniuous detected voung sstellar objec (YSO) with thermal jets.,the most luminous detected young stellar object (YSO) with thermal jets.
 Brooks et al. (, Brooks et al. (
2005) have reported he detection of a chain of IT». 2.12 jan endsson knots hat trace the collamated outflow that emanuates from the ceuter of the source.,2005) have reported the detection of a chain of $_{2}$ 2.12 $\mu$ m emission knots that trace the collimated outflow that emanates from the center of the source.
 The SED of the IR. cutission cau ο described by a imnodi&ed blackbody fiction with a yeas temperature of 30 EK. The total mass of the cloud is Ma29107 M. (Garay et al., The SED of the IR emission can be described by a modified blackbody function with a peak temperature of 30 K. The total mass of the cloud is $M_{\rm cl}=9\times 10^{2}$ $M_{\odot}$ (Garay et al.
 2003)., 2003).
 Molecular liue observations indicate that the size of the cloud is ~0.38 pe iu diameter (m1.1<1075 ein).," Molecular line observations indicate that the size of the cloud is $\sim0.38$ pc in diameter $\approx 1.1\times
10^{18}$ cm)."
 If we assume a spherical eeonietry. the averaged particle (atoms of IT) density of the cloud is Áom52.10 7.," If we assume a spherical geometry, the averaged particle (atoms of H) density of the cloud is $n_{\rm cl}\approx5.2 \times
10^{5}$ $^{-3}$."
" The euergev density of IR photons in the cloud. αππο an homogeneous distribution. is uy29loss109 ere 5,"," The energy density of IR photons in the cloud, assuming an homogeneous distribution, is $w_{\rm ph}\approx 1.8 \times
10^{-9}$ erg $^{-3}$."
 The radio observatious made by Caray et al. (, The radio observations made by Garay et al. (
2003) with he ATCA iux the deeper observations by Rodriguez et al. (,2003) with the ATCA and the deeper observations by guez et al. (
2005) with the VLA show the existence of a triple radio source inside the molecular cloud.,2005) with the VLA show the existence of a triple radio source inside the molecular cloud.
 The three components of the radio source are aligned iu the northwest-southeast direction. with the outer lobes separated frou the COYOC bv a projected distance of Oll pe.," The three components of the radio source are aligned in the northwest-southeast direction, with the outer lobes separated from the core by a projected distance of 0.14 pc."
 The ceutral source is elongated and las a spectral index. of 0.33zx0.05. consistent with Πουfree. emission. from a collimated jet (Bodríguez oet al.," The central source is elongated and has a spectral index of $0.33\pm0.05$, consistent with free-free emission from a collimated jet guez et al."
 2005)., 2005).
 The radio lobes have some substructure., The radio lobes have some substructure.
 The integrated cinission from the northern lobe has a spectral index of 0.32c0.29. of dubius hornal/uou-theniual nature.," The integrated emission from the northern lobe has a spectral index of $-0.32\pm
0.29$, of dubius thermal/non-thermal nature."
 The spectrum of the southern lobe radiation. instead. has an index a=0.59+0.15.," The spectrum of the southern lobe radiation, instead, has an index $\alpha=-0.59\pm 0.15$."
 Actually. the shock heated iaterial can radiate enough as to ionize the surroundingC» imediuu ancl free-free absorption could modify to some extent the radio spectrmu.," Actually, the shock heated material can radiate enough as to ionize the surrounding medium and free-free absorption could modify to some extent the radio spectrum."
" Ποπονα, we have no enough data to characterize the region in order to derive its thermal properties or the ionization deeree."," However, we have no enough data to characterize the region in order to derive its thermal properties or the ionization degree."
 Otherwise. diffusive shock acceleration by stroug nowrelativistic shocks naturally produces a power-law particle distribution with an index sinilku to what is interred from the observations.," Otherwise, diffusive shock acceleration by strong non-relativistic shocks naturally produces a power-law particle distribution with an index similar to what is inferred from the observations."
 Thus. we take as a first order approximation the observed radio spectrum as the original one. which would correspond to a nou-thermal particle distribution.," Thus, we take as a first order approximation the observed radio spectrum as the original one, which would correspond to a non-thermal particle distribution."
 The imferred linear size for this lobe is LdsP015 cm., The inferred linear size for this lobe is $\approx 1.1\times 10^{16}$ cm.
 Non-thermal radio emission has been associated in the past with the outflows of a few massive YSO (sce Rodriguez ct al., Non-thermal radio emission has been associated in the past with the outflows of a few massive YSO (see guez et al.
 2005. and references therein).," 2005, and references therein)."
 In any case the molecular cloucl is so massive aud. linens as in the case of IRAS 1217., In any case the molecular cloud is so massive and luminous as in the case of IRAS $-$ 4247.
 The flux deusity of the southern lobe is 2.54U0.] mJy at 8.16 GIIz., The flux density of the southern lobe is $2.8\pm 0.1$ mJy at 8.46 GHz.
 In Figure asketch of the scenario discussed im this paper is shown., In Figure \ref{fig_1} a sketch of the scenario discussed in this paper is shown.
 At the termination point of the thermal jet a stroug shock frout is expected to be formed., At the termination point of the thermal jet a strong shock front is expected to be formed.
 Chareccd particles can be accelerated at the shock. with an acceleration rate x2 (see below). where Bis the maguetic field aud the cficieney η depends on the acceleration mechauisi aud its details.," Charged particles can be accelerated at the shock, with an acceleration rate $\propto \eta B$ (see below), where $B$ is the magnetic field and the efficiency $\eta$ depends on the acceleration mechanism and its details."
 Tn the case of diffusive shock acceleration (fey fe. where ος is the shock velocity aud. fi. is the ratio offthe mean free path of particles to their evro-raclins.," In the case of diffusive shock acceleration $\eta\sim (v_{\rm s}/c)^2/f_{\rm sc}$ , where $v_{\rm s}$ is the shock velocity and $f_{\rm sc}$ is the ratio of the mean free path of particles to their gyro-radius."
 Close to the Bolin limit f.—1., Close to the Bohm limit $f_{\rm sc}\sim 1$.
 The terminal velocity of he collimated outflow in IRAS 1217 is uukuown. mtAarti. Rodríguez Reiputh (1995) have determined a velocity iu the range GOQO—1400 kins + for the ILI 80-81 thermal radio jet. which is powered also by a inassive ¥SO.," The terminal velocity of the collimated outflow in IRAS $-$ 4247 is unknown, but, guez Reipurth (1995) have determined a velocity in the range $600-1400$ km $^{-1}$ for the HH 80-81 thermal radio jet, which is powered also by a massive YSO."
 Adopting a value of ~1000 lan + for IRAS 1217. we have a reasonable officieucy of jj~10.7 for the southern lobe of his source.," Adopting a value of $\sim 1000$ km $^{-1}$ for IRAS $-$ 4247, we have a reasonable efficiency of $\eta\sim10^{-5}$ for the southern lobe of this source."
 The efficiency in the northern lobe would be ower. since the nou-thermal cussion does uot secu to be donuuaut there.," The efficiency in the northern lobe would be lower, since the non-thermal emission does not seem to be dominant there."
 The average matter deusitv of the cloud is 5.2«LO 7.," The average matter density of the cloud is $5.2 \times
10^{5}$ $^{-3}$."
 Tt is noted nevertheless that in the shocked material the deusitv should be ~1I times this value., It is noted nevertheless that in the shocked material the density should be $\sim 4$ times this value.
 Iu our calculations of the Bremsstrallune auc piou-decayv radiation we have adopted this higher density., In our calculations of the Bremsstrahlung and pion-decay radiation we have adopted this higher density.
 As incutioncd above. the rate of euergev eain for electrons at the acceleration region is: where 5 is the Lorentz factor of the particle. aud 0. is the electron rest mass.," As mentioned above, the rate of energy gain for electrons at the acceleration region is: where $\gamma$ is the Lorentz factor of the particle, and $m_e$ is the electron rest mass."
" Similarly. for protons. Iu order to get the imaxiumnuu enerey of the primary particles we have to balance the energwsain and loss rates,"," Similarly, for protons, In order to get the maximum energy of the primary particles we have to balance the energygain and loss rates."
 In the case of electrous. the relevant losses are svuchrotron. IC. aud relativistic Breisstrahhime losses.," In the case of electrons, the relevant losses are synchrotron, IC, and relativistic Bremsstrahlung losses."
then cascading will be unlikely to play a role.,then cascading will be unlikely to play a role.
 Finally. one-dimensional cascading should hold in the free pulsar wind as long as the pairs move strictly along the magnetic field.," Finally, one-dimensional cascading should hold in the free pulsar wind as long as the pairs move strictly along the magnetic field."
 In ? and ?.. the cascade radiation is computed up to the termination shock using a Monte Carlo approach.," In \citet{2005MNRAS.356..711S} and \citet{2008APh....30..239S}, the cascade radiation is computed up to the termination shock using a Monte Carlo approach."
 ?— also include a contribution from the region beyond the shock., \citet{2005MNRAS.356..711S} also include a contribution from the region beyond the shock.
 The cascade electrons in this region are assumed to follow the magnetic field lines (in contrast with the pulsar wind zone where the propagation is radial)., The cascade electrons in this region are assumed to follow the magnetic field lines (in contrast with the pulsar wind zone where the propagation is radial).
 There is no reacceleration at the shock and synchrotron losses are neglected., There is no reacceleration at the shock and synchrotron losses are neglected.
 In the method expounded here. the cascade radiation is calculated semi-analytically from a point-like gamma-ray source at the compact object location up to infinity. providing the maximum possible contribution of the one-dimensional cascade in gamma-ray binaries.," In the method expounded here, the cascade radiation is calculated semi-analytically from a point-like gamma-ray source at the compact object location up to infinity, providing the maximum possible contribution of the one-dimensional cascade in gamma-ray binaries."
 In order to compute the contribution from the cascade. the radiative transfer equation and the kinetic equation of the pairs have to be solved simultaneously.," In order to compute the contribution from the cascade, the radiative transfer equation and the kinetic equation of the pairs have to be solved simultaneously."
" The radiative transfer equation for the gamma-ray density ny=dNy/dtdedQ at a distance r from the source is where n,=dN,/drdE,dQ, is the electrons distribution. 7, the seed photon density from the massive star and dN/dtde, the Compton kernel."," The radiative transfer equation for the gamma-ray density $n_{\gamma}\equiv dN_{\gamma}/dt d\epsilon_1 d\Omega$ at a distance $r$ from the source is where $n_e\equiv dN_{e}/dr dE_e d\Omega_e$ is the electrons distribution, $n_{\star}$ the seed photon density from the massive star and $dN/dtd\epsilon_1$ the Compton kernel."
" The kernel is normalised to the soft photon density and depends on the energy E, of the electron and the angle between the photon and the direction of motion of the electron (2)..", The kernel is normalised to the soft photon density and depends on the energy $E_e$ of the electron and the angle between the photon and the direction of motion of the electron \citep{2008A&A...477..691D}.
" In the mono-energetic and point-like star approximation the stellar photon density can be estimated as [ινΑποκ”ει. where Ly is the stellar luminosity. &j=2.7k7T, the mean thermal photon energy and RA the distance to the massive star (see Fig. 1))."," In the mono-energetic and point-like star approximation the stellar photon density can be estimated as $L_{\star}/4\pi c R^2 \bar{\epsilon_0}$, where $L_{\star}$ is the stellar luminosity, $\bar{\epsilon_0}\approx 2.7 k T_{\star}$ the mean thermal photon energy and $R$ the distance to the massive star (see Fig. \ref{fig_bin}) )."
" The absorption rate d7,4./dr is given by Eq. (C8))."," The absorption rate $d\tau_{\gamma\gamma}/dr$ is given by Eq. \ref{taukern}) ),"
 convoluted to the soft photon density., convoluted to the soft photon density.
" The kinetic equation for thepairs ts given by thefollowing integro-differential equation for y,>>1 (2???) where P(E...E;) is the transition rate for an electron of energy E, down-scattered at an energy Ez<E, at r."," The kinetic equation for thepairs is given by thefollowing integro-differential equation for $\gamma_e\gg 1$ \citep{1970RvMP...42..237B,1988ApJ...335..786Z,2007A&A...469..857D}
 where $\mathcal{P}(E_e,E'_e)$ is the transition rate for an electron of energy $E_e$ down-scattered at an energy $E'_e\le E_e$ at $r$."
 The first two terms on the right side of the equation describe the inverse Compton cooling of pairs. taking into account catastrophic losses in the deep Klein-Nishina regime.," The first two terms on the right side of the equation describe the inverse Compton cooling of pairs, taking into account catastrophic losses in the deep Klein-Nishina regime."
" In this case. most of the electron energy is lost in the interaction and the scattered photon carries away most of its energy since ej=E,—E,Ey."," In this case, most of the electron energy is lost in the interaction and the scattered photon carries away most of its energy since $\epsilon_1=E_e-E'_e\approx E_e$."
 A continuous losses equation inadequately describes sizeable stochastic losses in a single interaction (??)..," A continuous losses equation inadequately describes sizeable stochastic losses in a single interaction \citep{1970RvMP...42..237B,1989ApJ...342.1108Z}."
" Since the inverse Compton kernel gives the probability per electron of energy E, to produce a gamma-ray of energy ej. the scattering rate can be rewritten as The expression of 4N/dtdE, is the same as the Compton kernel as described before but gives the spectrum of scattered electrons instead of the outcoming photon."," Since the inverse Compton kernel gives the probability per electron of energy $E_e$ to produce a gamma-ray of energy $\epsilon_1$, the scattering rate can be rewritten as The expression of $dN/dtdE'_e$ is the same as the Compton kernel as described before but gives the spectrum of scattered electrons instead of the outcoming photon."
 The first integral in Eq. (2)), The first integral in Eq. \ref{kin_eq}) )
" is the inverse Compton scattering rate and can be analytically expressed as where f, 1s the electron velocity in the observer frame and oj, 18 the total inverse Compton cross-section (for the full expression see ?.. Eq."," is the inverse Compton scattering rate and can be analytically expressed as where $\beta_e$ is the electron velocity in the observer frame and $\sigma_{ic}$ is the total inverse Compton cross-section (for the full expression see \citealt{1979rpa..book.....R}, , Eq."
 7.5)., 7.5).
 The last term in the kinetic equation is à source of pairs from yy-absorption coupled with, The last term in the kinetic equation is a source of pairs from $\gamma\gamma$ -absorption coupled with
Iu this section. we will show a stability result for the term A eiven byProposition 2.1 uuder a perturbation of (2.1)) of the form: where +20 d8 a small real ummber.,"In this section, we will show a stability result for the term $\l$ given byProposition \ref{ergo} under a perturbation of \ref{ergo_eq}) ) of the form: where $\g>0$ is a small real number."
" More precisely, we have the following proposition: We assiue. for thesake of simplicity. that g-=g|>."," More precisely, we have the following proposition: We assume, for thesake of simplicity, that $g_{\g}=g+\g$."
 In this case. it is easy to check that A-=Ay.," In this case, it is easy to check that $\l_{\g}\geq \l_{0}$."
 The other case with y-=y5 is treated similarly., The other case with $g_{\g}=g-\g$ is treated similarly.
 We first transform our ODE problem iuto a PDE oue by setting ( as the solution of the following equation: The proof is divided into three Using comparison principle aremmeuts for (3.3)). it is easily checked that ονι). is a non-decreasing fiction satisfying ο(foe|1)=etry1.," We first transform our ODE problem into a PDE one by setting $v^{\g}$ as the solution of the following equation: The proof is divided into three Using comparison principle arguments for \ref{ode-pde}) ), it is easily checked that $v^{\g}(t,.)$ is a non-decreasing function satisfying $v^{\g}(t,x+1)=v^{\g}(t,x)+1$."
 We want to control e2(f..) for any f.," We want to control $v^{\g}_{x}(t,.)$ for any $t$."
 For this reason. we proceed as follows.," For this reason, we proceed as follows."
" Define for 20: We compute which proves that jj is à sub-solution of (3.3)) with Οι)=er(ee|2)iοἱ (Qo). and therefore. by the conrparison principle. we obtain hence for auv f>0. we have 0xc(fc|:)etέν).«€sb then etter) ds Lipschitz continuous iu the variable 0, satisfviug: Iu a similar way. we can obtain a positive bound from belowou 02. and finally get"," Define for $z\geq 0$: We compute which proves that $\eta$ is a sub-solution of \ref{ode-pde}) ) with $\eta(0,x)=v^{\g}(0,x+z)-z=v^{\g}(0,x)$ , and therefore, by the comparison principle, we obtain hence for any $t\geq 0$, we have $0\leq v^{\g}(t,x+z)-v^{\g}(t,x)\leq ze^{Lt}$, then $v^{\g}(t,x)$ is Lipschitz continuous in the variable $x$ , satisfying: In a similar way, we can obtain a positive bound from belowon $v^{\g}_{x}$ , and finally get"
in Fig.,in Fig.
 7. although they were obtained in very different energy ranges.," 7, although they were obtained in very different energy ranges."
 Therefore one can be tempted to conclude that thev originate Irom (he same mechanism satisIving Belore discussing this point. we recall (hat in the collinear approximation. (he mechanism lo generata SSA is based on higher-twvist quark-gluon correlators (Elremov-Tervaev 1932. Qiu-Sterman 1991).," Therefore one can be tempted to conclude that they originate from the same mechanism satisfying Before discussing this point, we recall that in the collinear approximation, the mechanism to generata SSA is based on higher-twist quark-gluon correlators (Efremov-Teryaev 1982, Qiu-Sterman 1991)."
 LLowever. if one introduces transverse momentum dependence (ΤΟ). two QCD mechanisms have been - TAID parton distributions = Sivers ellect - TMD fragmentation distributions = Collins effect The eauge-invariance properties of the TAID PDF have been first clarified for DIS aud Drell-Yan processes in Ref. |20]..," However, if one introduces transverse momentum dependence (TMD), two QCD mechanisms have been - TMD parton distributions $\Rightarrow$ Sivers effect - TMD fragmentation distributions $\Rightarrow$ Collins effect The gauge-invariance properties of the TMD PDF have been first clarified for DIS and Drell-Yan processes in Ref. \cite{bhs}."
 In general both Sivers and. Collins effects contribute to a specific reaction. although there are some cases in which only one of them contributes.," In general both Sivers and Collins effects contribute to a specific reaction, although there are some cases in which only one of them contributes."
 For example in semi inclusive DIS. the Collins effect is the only mechanism that can lead lo asvimetries ορ and Apy.," For example in semi inclusive DIS, the Collins effect is the only mechanism that can lead to asymmetries $A_{UT}$ and $A_{UL}$."
 On the other hand. it does not appear in some electroweak interaction processes. where there is only the Sivers effect.," On the other hand, it does not appear in some electroweak interaction processes, where there is only the Sivers effect."
 In prompt photon production in pp collisions. which is dominated by ο—q5. the SSA is sensitive {ο either the quark or the gluon Sivers functions. according to the value of the photon wry |21]..," In prompt photon production in $pp$ collisions, which is dominated by $qG \to q \gamma$, the SSA is sensitive to either the quark or the gluon Sivers functions, according to the value of the photon $x_F$ \cite{ssy}."
 Now let us ask: do we understand the SSA displaved on Fig., Now let us ask: do we understand the SSA displayed on Fig.
 7. given the fact that STAR is at a very small angle 2.6 deg..," 7, given the fact that STAR is at a very small angle 2.6 deg.,"
 whereas E704 is ad à much larger angle. between 9 deg.," whereas E704 is at a much larger angle, between 9 deg."
 and 64 deg.?, and 64 deg.?
 A negative answer is partially obtained by looking at the eross section., A negative answer is partially obtained by looking at the cross section.
 The pQCD NLO calculation uuderestimates the cross section alfow energies aud angles. namely for the E704 kinematic region.," The pQCD NLO calculation underestimates the cross section at energies and angles, namely for the E704 kinematic region."
 This is shown on Fig., This is shown on Fig.
 7 and it means that one should not ignore other contributions., 7 and it means that one should not ignore other contributions.
 This is not the case al 90 deg., This is not the case at 90 deg.
 and at very small angles at high energv. which is the STAR kinematic range.," and at very small angles at high energy, which is the STAR kinematic range."
 To conclude. one should not try to “explain” the SSA. ignoring the unpolarized cross section [22]..," To conclude, one should not try to ""explain"" the SSA, ignoring the unpolarized cross section \cite{BS}."
 Of course one should not lorget resummation effects. which might help clarifving the situation.," Of course one should not forget resummation effects, which might help clarifying the situation."
oof ?..,of \citet{mar05}.
" Having such a low mass but still the luminosity and effective temperature of an O star imply the occurrence of processes not present in single-star evolution,", Having such a low mass but still the luminosity and effective temperature of an O star imply the occurrence of processes not present in single-star evolution.
 Another indication of interaction within the system comes from the analysis of the surface abundances., Another indication of interaction within the system comes from the analysis of the surface abundances.
 In reftig noweshowthenitrogenovercarbonratioasa functiono fFthenigamgenaeeeweseNi, In \\ref{fig_cno} we show the nitrogen over carbon ratio as a function of the nitrogen over oxygen.
-Aehiamaten starevolutionarvtracksineludingrotationalinixing fronPare shown," Again, single-star evolutionary tracks including rotational mixing from \citet{mm03} are shown."
iMthse ΕΚΕ vada cdonaaahatitnirecerved. —-leftcornerof fig. no..displayssurfacechemistryconsistentwithsinglestaret," We see that the primary, located at the bottom--left corner of \\ref{fig_cno}, displays surface chemistry consistent with single star evolution."
lsdusuwfakfeowwtuogehe ses n," However, the secondary features extreme nitrogen enrichment and carbon/oxygen depletion."
 Iunentand Ravetphaseinsinglestarevolution(theW ol RavetphaseisindicatedbyboldsvmbolsinFig. nop., Such high N/C and N/O ratios are only reached in the Wolf-Rayet phase in single star evolution (the Wolf-Rayet phase is indicated by bold symbols in \\ref{fig_cno}) ).
Interestinglv. thisisconsistentwiththehighheliumcontemrogehesarfaedcalyundance He/H=0.4. correspondingtoY 20.6. isonlvobservedinadvanced Sf itherisgle —starevolutionarvtrack swithinitialinassesofabout20," Interestingly, this is consistent with the high helium content of the secondary: $\rm{He/H}=0.4$, corresponding to $\rm{Y}=0.6$, is only observed in advanced states of evolution in single--star evolutionary tracks with initial masses of about 20."
 Thisraiht/thgupyetibwetini ebihasycexatutiónosdmthcodepen, This raises the question of whether the secondary is not in fact more evolved than the primary and is currently in a state of central helium burning.
e —huninosityrelationo f? forcoreH ιο παλια vsre ce os," Using the mass--luminosity relation of \citet{schaerer92} for core H--free stars, we find that a star with a current mass of about 6 (resp."
s tr, 5 ) should have a luminosity $\lL=4.85$ (resp.
, 4.69).
y log7L , This is remarkably consistent with the derived luminosity of the secondary $\lL=4.69$ ).
All indications thus favor the scenario in. which. the secondary was initially more massive. evolved faster. transferred mass to the companion and is now close to be a core He-burning object.," All indications thus favor the scenario in which the secondary was initially more massive, evolved faster, transferred mass to the companion and is now close to be a core He–burning object."
 This could also explain the rather low gravity and mass of the (if it is secondarytruly a main-sequence O star. sshould be much higher — see also above).," This could also explain the rather low gravity and mass of the secondary (if it is truly a main-sequence O star, should be much higher – see also above)."
 However. there are two caveats to this scenario.," However, there are two caveats to this scenario."
 First. the spectroscopic appearance of the secondary is not that of an evolved massive star.," First, the spectroscopic appearance of the secondary is not that of an evolved massive star."
 With a spectral type 99.7V. it looks like a normal main sequence star.," With a spectral type 9.7V, it looks like a normal main sequence star."
 Said differently. despite of all the indications. the absence of strong mass loss characterizing evolved massive stars (and affecting spectral types) is challenging.," Said differently, despite of all the indications, the absence of strong mass loss characterizing evolved massive stars (and affecting spectral types) is challenging."
" ? have shown that the mass loss rate of Wolf-Rayet stars depended on the Eddington factor L,.", \citet{gh08} have shown that the mass loss rate of Wolf-Rayet stars depended on the Eddington factor $\Gamma_{e}$.
 Estimating this factor from the parameters we determined. we obtain Εν~0.17.," Estimating this factor from the parameters we determined, we obtain $\Gamma_{e} \sim\ 0.17$."
" This is rather low. and much lower than the range of values for which ? provide mass loss prescriptions for Wolf-Rayet stars (their computations are restricted to D,> 0.4)."," This is rather low, and much lower than the range of values for which \citet{gh08} provide mass loss prescriptions for Wolf–Rayet stars (their computations are restricted to $\Gamma_{e} > 0.4$ )."
 Hence. the secondary may very well have a rather normal wind.," Hence, the secondary may very well have a rather normal wind."
 If we use the mass loss recipe of ?.. we obtain logM—7.0. which is much more typical ofO stars.," If we use the mass loss recipe of \citet{vink01}, we obtain $\log \dot{M} \sim -7.0$, which is much more typical of O stars."
 Consequently. it does not seem impossible that the secondary is an evolved object with a rather normal wind and thus the appearance of an O star.," Consequently, it does not seem impossible that the secondary is an evolved object with a rather normal wind and thus the appearance of an O star."
 The second caveat is the chemical composition of the primary., The second caveat is the chemical composition of the primary.
 If almost all the envelope of the secondary was dumped onto the primary. the chemical. pattern. should. be different.," If almost all the envelope of the secondary was dumped onto the primary, the chemical pattern should be different."
 The primary should show evidence for N enrichment and CO depletion since material from the secondary would be processed. the secondary being significantly evolved.," The primary should show evidence for N enrichment and CO depletion since material from the secondary would be processed, the secondary being significantly evolved."
 According to our determinations. there are only weak signatures of CNO processing.," According to our determinations, there are only weak signatures of CNO processing."
 This favors a weak aceretion of material from the donor., This favors a weak accretion of material from the donor.
 A way out of this puzzle is to invoke a very non-conservative mass transfer even though this kind of process was not observed in the two post Roche lobe overflow key systems: @pper (2) and SScuti (?).., A way out of this puzzle is to invoke a very non–conservative mass transfer even though this kind of process was not observed in the two post Roche lobe overflow key systems: $\phi$ per \citep{bozic1995} and Scuti \citep{grundstrom2007}.
 All in all. the scenario according to which the secondary was initially more massive and evolved faster to the core-He burning phase while filling its Roche lobe and transferring very inefficiently material onto the primary seems a valuable explanation of the system.," All in all, the scenario according to which the secondary was initially more massive and evolved faster to the core–He burning phase while filling its Roche lobe and transferring very inefficiently material onto the primary seems a valuable explanation of the system."
 But other explanations can be invoked., But other explanations can be invoked.
 ? and ? have argued that strong nitrogen enhancement can be obtained on interacting binaries., \citet{demink07} and \citet{langer08} have argued that strong nitrogen enhancement can be obtained on interacting binaries.
 Two effects are at work: the mass dM usually from the inner layers enc rte from the primary triggers rotational mixing which further increases παποπ μι onde roteen, Two effects are at work: the mass gainer receives N–rich material usually from the inner layers of the donor; and the angular momentum received from the primary triggers rotational mixing which further increases the surface nitrogen content of the secondary.
ien aic solar metallicities. the enrichment can be of a factor of ten compared to the initial N abundance.," In total, at solar metallicities, the enrichment can be of a factor of ten compared to the initial N abundance."
 According to Table 2.. the of the secondary star in CCep stake ena&st value.," According to Table \ref{tab_orb}, the nitrogen surface abundance of the secondary star in Cep is 25 times the initial value."
 This is qualitatively consistent cen odnedthantlie the orbital parameters and masses of the components., This is qualitatively consistent with the predictions of binary evolution which depend on the orbital parameters and masses of the components.
 In that case. contrary to the scenario described above. the primary is ms," In that case, contrary to the scenario described above, the primary is the mass donor."
i aaa transfer. one also expects helium enrichment in the secondary. which is what we observe in the CCep.," Since rotational mixing is triggered by mass transfer, one also expects helium enrichment in the secondary, which is what we observe in the Cep."
 However. according to ?.. to reach a mass fraction Y=0.6 requires very short periods and occurs only on stars more massive than ~ 40M.," However, according to \citet{dem09}, to reach a mass fraction $\rm{Y}=0.6$ requires very short periods and occurs only on stars more massive than $\sim$ 40."
.. Another concern ts that the rotational velocity of the secondary is not extreme., Another concern is that the rotational velocity of the secondary is not extreme.
 With == 80s. and assuming that the rotation axis is parallel to the orbital axis. this gives a rotation rate of ~ 100s7!.," With = 80, and assuming that the rotation axis is parallel to the orbital axis, this gives a rotation rate of $\sim$ 100."
. But as shown by ?.. the rotational velocity can reach very high values after mass transfer. and subsequently come back to more moderate values.," But as shown by \citet{langer08}, the rotational velocity can reach very high values after mass transfer, and subsequently come back to more moderate values."
 The question is whether fast mixing can be triggered during this short spin-up phase., The question is whether fast mixing can be triggered during this short spin–up phase.
 It seems unlikely. but further dedicated simulations should be run to tackle this problem.," It seems unlikely, but further dedicated simulations should be run to tackle this problem."
 With the inclination. projected rotational velocities and radius. we can estimate the rotational period of each star and check if synchronization has been achieved in the system.," With the inclination, projected rotational velocities and radius, we can estimate the rotational period of each star and check if synchronization has been achieved in the system."
" Assuming an inclination angle of 49.1° (the average of the two values from Table 3)). we obtain an orbital period of 3.0940.24 d using R = 10.50 R4 (average of the polar radii for the primary) and 3.7540.31 d using R = 12.75 R,. (equatorial radius)."," Assuming an inclination angle of $^{\circ}$ (the average of the two values from Table \ref{tab_lc}) ), we obtain an orbital period of $\pm$ 0.24 d using R = 10.50 $R_{\odot}$ (average of the polar radii for the primary) and $\pm$ 0.31 d using R = 12.75 $R_{\odot}$ (equatorial radius)."
 Similarly. the rotation period of the secondary. ts 3.1540.45 d (polar radius) and 4.21+0.60 d (equatorial radius).," Similarly, the rotation period of the secondary is $\pm$ 0.45 d (polar radius) and $\pm$ 0.60 d (equatorial radius)."
 Compared to the orbital period. (3.07 d). the true rotational periods of both components (likely in between the polar and equatorial cases) are thus roughly 10 to larger. implying that the synchronous co-rotation is not completely established in the system.," Compared to the orbital period (3.07 d), the true rotational periods of both components (likely in between the polar and equatorial cases) are thus roughly 10 to larger, implying that the synchronous co-rotation is not completely established in the system."
 However. given the small differences and the similarity of the primary and secondary rotational periods. the system is certainly on the verge of achieving it.," However, given the small differences and the similarity of the primary and secondary rotational periods, the system is certainly on the verge of achieving it."
 Allin all. we have thus gathered evidence for mass transfer and tidal interaction in the massive binary system CCep.," All in all, we have thus gathered evidence for mass transfer and tidal interaction in the massive binary system Cep."
 The scenario according to which the secondary was initially the more massive star appears to explain more observational constraints. although it is not free of challenges.," The scenario according to which the secondary was initially the more massive star appears to explain more observational constraints, although it is not free of challenges."
 What is clear is that single-star evolutionary tracks are certainly not relevant to explain the position of the two components in the HR diagram., What is clear is that single–star evolutionary tracks are certainly not relevant to explain the position of the two components in the HR diagram.
 This was highlighted by ? who computed evolutionary tracks for interacting binaries and showed that as, This was highlighted by \citet{wellstein01} who computed evolutionary tracks for interacting binaries and showed that as
function.,function.
 We also determine the optical depth. which is a function of the ratio between the SCO intensity and the excitation temperature. to apply a correction for saturation.," We also determine the optical depth, which is a function of the ratio between the $^{13}$ CO intensity and the excitation temperature, to apply a correction for saturation."
 Finally. we (transform the column density of CO to that of CO assuming an isotopic ratio of 65 (Densch 22001a) in the relatively well-shielded portions of the cloud defined by the “mask 2 region of Goldsmith ((2008).," Finally, we transform the column density of $^{13}$ CO to that of $^{12}$ CO assuming an isotopic ratio of 65 (Bensch 2001a) in the relatively well-shielded portions of the cloud defined by the “mask 2"" region of Goldsmith (2008)."
" The determination of [IN (CO), was improved relative to that presented by. Goldsmith ((2008) bv including an updated value of (he spontaneous decay rate. using an exact numerical rather than approximate analvtical ealeulation of the partition function."," The determination of $N$ $_{\rm gas}$ was improved relative to that presented by Goldsmith (2008) by including an updated value of the spontaneous decay rate, using an exact numerical rather than approximate analytical calculation of the partition function."
 Additionally. the data were corrected for error beam pick-up using the method presented by Beusch ((2001b).," Additionally, the data were corrected for error beam pick-up using the method presented by Bensch (2001b)."
" The resulting values of NV(CO).,. are about ~20% larger than those of Goldsmith ((2008).", The resulting values of $N$ $_{\rm gas}$ are about $\sim$ larger than those of Goldsmith (2008).
 Full details of these improvements are presented elsewhere (Pineda 22010)., Full details of these improvements are presented elsewhere (Pineda 2010).
 Motivated by observations of core-to-edge temperature differences in molecular clouds (Evans 22001) which can be found even in regions of only moderate radiation field intensity. Pineda (2010) studied the effects of temperature gradients on the determination of V(CO)....," Motivated by observations of core-to-edge temperature differences in molecular clouds (Evans 2001) which can be found even in regions of only moderate radiation field intensity, Pineda (2010) studied the effects of temperature gradients on the determination of $N$ $_{\rm gas}$."
 The radiative transfer code RATRAN (IIlogerheijde van der Tak 2000) was used to model the ?CO and CO emission emerging [rom a model cloud., The radiative transfer code RATRAN (Hogerheijde van der Tak 2000) was used to model the $^{12}$ CO and $^{13}$ CO emission emerging from a model cloud.
 Pineda found that using CO to determine the excitation temperature of the CO gas only traces the temperature at low column densities while (he excitation temperature is overestimated or larger column densities., Pineda found that using $^{12}$ CO to determine the excitation temperature of the CO gas only traces the temperature at low column densities while the excitation temperature is overestimated for larger column densities.
 This produces an underestimate of the “CO optical depth. and in consequence the opacity correction of NCCO). which is usually evaluated assuming an isothermal cloud.," This produces an underestimate of the $^{13}$ CO optical depth, and in consequence the opacity correction of $N(^{13}{\rm CO})$, which is usually evaluated assuming an isothermal cloud."
" The column densities presented here were corrected. by (this method including modest edge-center temperature gradients of ~4 Ix. This procedure tvpically increases the value of IN(CO)44, by ~LO4056... with larger increases occurring at higher column densities."," The column densities presented here were corrected by this method including modest edge-center temperature gradients of $\sim4$ K. This procedure typically increases the value of $N$ $_{\rm gas}$ by $\sim$, with larger increases occurring at higher column densities."
 Final values are presented in Table 1. together with ice ancl extinction data.," Final values are presented in Table 1, together with ice and extinction data."
dispersion of those 10000 realizations.,dispersion of those 10000 realizations.
 The results from this exercise are shown in Fig. 4..., The results from this exercise are shown in Fig. \ref{fig:evol}.
 We find an evolution in the number density of the massive end of the red sequence of a factor ~4 between z2| and z-0., We find an evolution in the number density of the massive end of the red sequence of a factor $\sim 4$ between $z=1$ and $z=0$.
 Our inference of a factor ~4 evolution in the number density of red sequence galaxies with M..>10!!M... is somewhat larger than a number of recent estimates., Our inference of a factor $\sim 4$ evolution in the number density of red sequence galaxies with $M_*>10^{11}M_\sun$ is somewhat larger than a number of recent estimates.
 The main reason for this ts our adoption of the recent results by vanDokkum&derMarel(2007) for the evolution of M/L determined through FP evolution. who find relatively rapid evolution of M/L leading to more rapid inferred fading of the stellar population than assumed by many previous works.," The main reason for this is our adoption of the recent results by \citet{vd} for the evolution of M/L determined through FP evolution, who find relatively rapid evolution of M/L leading to more rapid inferred fading of the stellar population than assumed by many previous works."
 For example. Cimattietal.(2006) found a result compatible with à constant number density of M. 10!! galaxies in the interval O<z«0.8.," For example, \citet{cimatti} found a result compatible with a constant number density of $M_*>10^{11}M_\sun$ galaxies in the interval $0<z<0.8$."
 In order to obtain that Mresult.. they applied a passive luminosity correction of Alog(M/L5)=0.46£0.04 to a sample drawn from COMBO-17 and DEEP2 (Davisetal.2003;Faberal. 2007).," In order to obtain that result they applied a passive luminosity correction of $\Delta \log(M/L_B)=0.46 \pm
0.04$ to a sample drawn from COMBO–17 and DEEP2 \citep{davis03, faber07}."
. This luminosity correction was obtained from the works by vanDokkum&Stanford(2003) and Treu(2005) and is lower than the one found a few years later by vanDokkum&derMarel(2007) with a larger sample.," This luminosity correction was obtained from the works by \citet{vdstanford} and \citet{treu} and is lower than the one found a few years later by \citet{vd}
 with a larger sample."
 This apparently small difference in Alog(M/Lg) has strong implications when studying the number density evolution of massive red galaxies at the level of ~1.5—2between z2| and z20., This apparently small difference in $\Delta \log(M/L_B)$ has strong implications when studying the number density evolution of massive red galaxies at the level of $\sim 1.5-2$between $z=1$ and $z=0$.
 Also. the number of galaxies in their COMBO-17 or DEEP? samples is relatively small when compared to the ~7 sq.," Also, the number of galaxies in their COMBO–17 or DEEP2 samples is relatively small when compared to the $\sim 7$ sq."
 deg., deg.
 of the NOAO Deep Wide-Field Survey Jannuzi&Dey1999) used by Brownetal.(2007).," of the NOAO Deep Wide–Field Survey \citep[NDWFS, ][]{jannuzi} used by \citet{brown}."
. Cooletal.(2008) used a sample of luminous red galaxies (LRG) from SDSS and spectroscopic follow-up from MMT to address a similar question., \citet{cool} used a sample of luminous red galaxies (LRG) from SDSS and spectroscopic follow–up from MMT to address a similar question.
 They found no evolution in the number density of LRGs since z= 0.9., They found no evolution in the number density of LRGs since $z=0.9$ .
 We believe that the main driver of the difference between Cool et al, We believe that the main driver of the difference between Cool et al.
s result and ours is. again. the correction for ML ratio evolution.,"'s result and ours is, again, the correction for $M/L$ ratio evolution."
 They used stellar population models from Bruzual/&Charlot(2003) and Maraston(2005). which gives slower evolution than FP zeropoint evolution constraints.," They used stellar population models from \citet{bc03} and \citet{maraston}, which gives slower evolution than FP zeropoint evolution constraints."
 In a recent work. Matsuoka&Kawara(2010). studied the number density evolution of massive galaxies since z| from a sample drawn from UKIRT Infrared Digital Sky Survey and SDSS.," In a recent work, \citet{mat} studied the number density evolution of massive galaxies since $z\sim 1$ from a sample drawn from UKIRT Infrared Digital Sky Survey and SDSS."
 Instead of applying an average M/L correction. they calculated stellar masses for individual galaxies at all redshifts. finding a factor ~4 evolution at M.>10!M... (when adding-up their two mass bins).," Instead of applying an average $M/L$ correction, they calculated stellar masses for individual galaxies at all redshifts, finding a factor $\sim
4$ evolution at $M_*>10^{11}M_\sun$ (when adding-up their two mass bins)."
 Unlike our present. work. they did not differenciate between red and blue galaxies. but given the predominance of red galaxies at such high stellar masses (they report a blue fraction «30% at all redshifts in the same mass range) we believe that our factor ~4 evolution agrees remarkably well with their result.," Unlike our present work, they did not differenciate between red and blue galaxies, but given the predominance of red galaxies at such high stellar masses (they report a blue fraction $<30\%$ at all redshifts in the same mass range) we believe that our factor $\sim 4$ evolution agrees remarkably well with their result."
 We now address the question of whether galaxy mergers between massive galaxies can drive the observed number density evolution of massive red galaxies., We now address the question of whether galaxy mergers between massive galaxies can drive the observed number density evolution of massive red galaxies.
 For that purpose. we use the result for the merger rate found in Section ?? and work under the assumption that mergers between massive galaxies will produce remnants with red optical colors.," For that purpose, we use the result for the merger rate found in Section \ref{sec:rate} and work under the assumption that mergers between massive galaxies will produce remnants with red optical colors."
 Dry (gas-free) mergers are likely to play an important role in the build-up of the massive end of the red sequence (vanDokkum2005:Belletal. 2006b).. and the descendant systems of such interactions will clearly be red systems as well.," Dry (gas–free) mergers are likely to play an important role in the build-up of the massive end of the red sequence \citep{vd_dry, bell06b}, and the descendant systems of such interactions will clearly be red systems as well."
 Moreover. mergers involving massive blue. disk-like galaxies are also expected to produce red ellipticals as remnants.," Moreover, mergers involving massive blue, disk–like galaxies are also expected to produce red ellipticals as remnants."
 Major interactions. independently of the color of the progenitors. are expected to perturb the orbits of the stars in the galaxies and form elliptical systems. but the observed fraction of spheroidal galaxies with blue optical colors is a negligible fraction of the total at masses above 10!M... (seeSchawinskietal.2009:KannappanHuertas-Companyetal. 2010).," Major interactions, independently of the color of the progenitors, are expected to perturb the orbits of the stars in the galaxies and form elliptical systems, but the observed fraction of spheroidal galaxies with blue optical colors is a negligible fraction of the total at masses above $10^{11}M_\sun$ \citep[see][]{schawinski, kan, huertas}."
. It is likely that even those few systems will passively evolve to the red sequence in =| Gyr. so. our assumption that merger remnants with masses above such stellar mass will become red sequence galaxies 1s justified.," It is likely that even those few systems will passively evolve to the red sequence in $\lesssim 1$ Gyr, so, our assumption that merger remnants with masses above such stellar mass will become red sequence galaxies is justified."
 We use the number density of massive systems at z~0.9 that we have obtained as the starting point and calculate the growth in the number of massive red galaxies implied by our measurement of the merger rate., We use the number density of massive systems at $z\sim 0.9$ that we have obtained as the starting point and calculate the growth in the number of massive red galaxies implied by our measurement of the merger rate.
 We show the result of this exercise as the solid line in Fig. 4+.., We show the result of this exercise as the solid line in Fig. \ref{fig:evol}.
 We stress that the observed evolution (filled circles) and the one implied by our measurement (solid line) areindependent except for the fact that we use the observed density at z~0.9 to anchor the evolution predicted by our close pair fractions., We stress that the observed evolution (filled circles) and the one implied by our measurement (solid line) are except for the fact that we use the observed density at $z\sim0.9$ to anchor the evolution predicted by our close pair fractions.
 We find that mergers of massive galaxies the evolution in the observed number density of massive red galaxies since z=|., We find that mergers of massive galaxies the evolution in the observed number density of massive red galaxies since $z=1$.
 In different words. the majority of galaxies observed today at the massive end of the red sequence have been assembled in the last 8 Gyrs through mergers of massive galaxies.," In different words, the majority of galaxies observed today at the massive end of the red sequence have been assembled in the last 8 Gyrs through mergers of massive galaxies."
 We have used 720.5 Gyrs. but using the 7~| Gyr timescale from Kitzbichler&White(2007) produces a somewhat slower evolution that is still compatible within the error bars.," We have used $\tau=0.5$ Gyrs, but using the $\tau \sim 1$ Gyr timescale from \citet{kw} produces a somewhat slower evolution that is still compatible within the error bars."
 There are two caveats we would like to mention., There are two caveats we would like to mention.
 Firstly. given the nature of our method. some of the progenitor galaxies have masses above 10!!M... so strictly speaking they are not massive galaxies.," Firstly, given the nature of our method, some of the progenitor galaxies have masses above $10^{11}M_\sun$, so strictly speaking they are not massive galaxies."
 Second. because we adopt a lower mass limit of 5«10M... we underestimate the number of major mergers that could lead to the formation of a massive galaxy.," Second, because we adopt a lower mass limit of $5\times 10^{10}M_\sun$, we underestimate the number of major mergers that could lead to the formation of a massive galaxy."
 For example. a pair with individual masses M.=6«10M απά M.24«10'?M would not make it into our pair sample but would produce a major.. merger remnant of 10!!M ...," For example, a pair with individual masses $M_*=6\times 10^{10}M_\sun$ and $M_*=4\times 10^{10}M_\sun$ would not make it into our pair sample but would produce a major merger remnant of $10^{11} M_{\sun}$ ."
 Assuming that the merging population has a composition similar to our overall sample. we estimate that this latter lower mass limit issue has an impact a factor ~2larger on the creation of >10!M galaxies than the overestimate caused by double counting already.. massive red galaxies (1.e.. for every 10 massive," Assuming that the merging population has a composition similar to our overall sample, we estimate that this latter lower mass limit issue has an impact a factor $\sim 2$larger on the creation of $>10^{11}M_\sun$ galaxies than the overestimate caused by double counting already massive red galaxies (i.e., for every 10 massive"
"extragalactic background falls as E*dN/dExE~°* in the GeV to 100 GeV range, the allowed Qppy increases with WIMP mass.","extragalactic background falls as $E^2 dN/dE \propto E^{-0.4}$ in the GeV to 100 GeV range, the allowed $\Omega_{\rm PBH}$ increases with WIMP mass."
" The Galactic signal constrains OpgyS107 for a WIMP mass of 100 GeV even if Br(y)=0.01, smaller than the cosmic bound (Eq. 9))."," The Galactic signal constrains $\Omega_{\rm PBH} \la 10^{-6}$ for a WIMP mass of 100 GeV even if $\Br (\gamma) = 0.01$, smaller than the cosmic bound (Eq. \ref{eqn:CosmicGammaLimits}) )."
" With these abundances, we can calculate the lower limits on the mean distance to the nearest PBH, [37/(46(nppy))]'/3, where the local 6=89000: for WIMP masses greater than 100 GeV. This implies a γ- flux of Unlike gamma rays, neutrinos do not cascade down in energy as they travel through the Universe, although they redshift."," With these abundances, we can calculate the lower limits on the mean distance to the nearest PBH, $\lambda_{\rm PBH} = [3 \pi / (4 \delta \mean{n_{\rm PBH}})]^{1/3}$ , where the local $\delta = 89000$: for WIMP masses greater than 100 GeV. This implies a $\gamma$ -ray flux of Unlike gamma rays, neutrinos do not cascade down in energy as they travel through the Universe, although they redshift."
" We the atmospheric neutrino spectrum (or diffuse neutrino integratebackground limits above 100 TeV) from mpMC?/e to mpmc? (Gaisser&Honda2002),, and require the neutrino flux from UCMHs be less than this; otherwise, they would have been detected (Beacometal.2007;seletal. 2007)."," We integrate the atmospheric neutrino spectrum (or diffuse neutrino background limits above 100 TeV) from $m_{\rm DM} c^2 / e$ to $m_{\rm DM} c^2$ \citep{Gaisser02}, and require the neutrino flux from UCMHs be less than this; otherwise, they would have been detected \citep{Beacom07,Yuksel07}."
". The measured data is reported in Ashieetal. (2005),, Gonzalez-Garciaetal. (2006),, Achterbergetal. (2007), Hoshinaetal. (2008), Abbasietal. (2009),, and DeYoungetal.(2009)."," The measured data is reported in \citet{Ashie05}, \citet{Gonzalez06}, \citet{Achterberg07}, \citet{Hoshina08}, \citet{Abbasi09}, and \citet{DeYoung09}."
". In Figure 1, we show the Galactic (solid)) and cosmic (dashed)) bounds on PBHs from neutrinos."," In Figure \ref{fig:BoundsWithMass}, we show the Galactic ) and cosmic ) bounds on PBHs from neutrinos."
" The neutrino also steeply falls with energy atmospheric(E?dN/dEοςE 13), so backgroundthat the bound on Qpgy is fairly energy-independent up to 100 TeV. In Figure 2,, we show that even with the uncertainties in our estimates, our constraints are powerful for PBHs inside UCMHs."," The atmospheric neutrino background also steeply falls with energy $E^2 dN/dE \propto E^{-1.3}$ ), so that the bound on $\Omega_{\rm PBH}$ is fairly energy-independent up to 100 TeV. In Figure \ref{fig:PBHFraction}, we show that even with the uncertainties in our estimates, our constraints are powerful for PBHs inside UCMHs."
" PBHs of many masses are ruled out to very small abundances, for a wide range in WIMP masses."," PBHs of many masses are ruled out to very small abundances, for a wide range in WIMP masses."
" At high PBH masses (Z1 Μο), CMB constraints become more powerful (Ricottietal.2008); at very low masses, radiation limits are more (e.g.,CarretHawkingal. 2010).."," At high PBH masses $\ga 1\ \Msun$ ), CMB constraints become more powerful \citep{Ricotti08}; at very low masses, Hawking radiation limits are more powerful \citep[e.g.,][]{Carr09}."
" Microlensing constraints are powerfulweaker, but make no assumptions about WIMPs (Alcocketal.2001;Tisserandal."," Microlensing constraints are weaker, but make no assumptions about WIMPs \citep{Alcock01,Tisserand07}."
" 2007).. For 107?Mc€Mpgu<10?Mo, ours are the constraints on PBHs aside from Qpm."," For $10^{-15}\ \Msun \la M_{\rm PBH} \la 10^{-9}\ \Msun$, ours are the constraints on PBHs aside from $\Omega_{\rm DM}$."
 Our conclusion depends on the standard assumption that most dark matter is a self-annihilating thermal relic., Our conclusion depends on the standard assumption that most dark matter is a self-annihilating thermal relic.
" Our analysis does not apply if all of the dark matter is made of PBHs (e.g.,Frampton2009),, because there will not be any WIMPs to annihilate."," Our analysis does not apply if all of the dark matter is made of PBHs \citep[e.g.,][]{Frampton09}, because there will not be any WIMPs to annihilate."
" For the smallest PBH and WIMP masses (MppuS10P miyMe), radial accretion may not hold."," For the smallest PBH and WIMP masses $M_{\rm PBH} \la 10^{-15} m_{100}^{-3} \Msun$ ), radial accretion may not hold."
" PBHs could remain important for other reasons (e.g.,asseedsofmassiveblackholes;seeMacketal.2007) which do not require them to have a substantial abundance."," PBHs could remain important for other reasons \citep[e.g., as seeds of massive black holes; see][]{Mack07} which do not require them to have a substantial abundance."
" The combination of our results with previous limits imply that PBHs either make up almost all of the dark matter, or almost none of it."," The combination of our results with previous limits imply that PBHs either make up almost all of the dark matter, or almost none of it."
 Our results could be improved by detailed simulations of how the UCMH evolves as its inner regions annihilate., Our results could be improved by detailed simulations of how the UCMH evolves as its inner regions annihilate.
" Annihilation could lower the luminosity of the inner regions of the halo over time, weakening our bounds."," Annihilation could lower the luminosity of the inner regions of the halo over time, weakening our bounds."
" Steeper density profiles than η0/2 may increase the annihilation luminosity, but shorten the lifetime of WIMPs so that the halo annihilates away before z=0."," Steeper density profiles than $r^{-3/2}$ may increase the annihilation luminosity, but shorten the lifetime of WIMPs so that the halo annihilates away before $z = 0$."
" We expect that some combination of constraints on gamma-ray and neutrino backgrounds, reionization history, and the CMB energy spectrum will generally require a small Opgy."," We expect that some combination of constraints on gamma-ray and neutrino backgrounds, reionization history, and the CMB energy spectrum will generally require a small $\Omega_{\rm PBH}$."
 Further improvements could come from more detailed studies of the annihilation products and their detectability., Further improvements could come from more detailed studies of the annihilation products and their detectability.
" When considering annihilation into charged particles, we considered only gamma rays from internal bremsstrahlung, but can themselves radiate, such as through Inverse chargedCompton particlesscattering (e.g.,Cirelli&Panci2009)."," When considering annihilation into charged particles, we considered only gamma rays from internal bremsstrahlung, but charged particles can themselves radiate, such as through Inverse Compton scattering \citep[e.g.,][]{Cirelli09}."
. We are confident the limits for WIMP masses above 100 GeV can be improved by new observations., We are confident the limits for WIMP masses above 100 GeV can be improved by new observations.
" Finally, Ricotti&Gould(2009) that some UCMHs may exist without PBHs, formed from suggestweaker initial perturbations in the early Universe."," Finally, \citet{Ricotti09} suggest that some UCMHs may exist without PBHs, formed from weaker initial perturbations in the early Universe."
 Limits on dark matter annihilation in these UCMHs may strongly constrain these weaker perturbations., Limits on dark matter annihilation in these UCMHs may strongly constrain these weaker perturbations.
" We thank A. Gould, J. Rich, M. Ricotti, K. Stanek, G.Steigman, and T. Thompson for helpful discussions."," We thank A. Gould, J. Rich, M. Ricotti, K. Stanek, G.Steigman, and T. Thompson for helpful discussions."
 This work was supported by NSF CAREER Grant PHY-0547102 to J.E.B. and an Alfred P. Sloan Fellowship to T. Thompson., This work was supported by NSF CAREER Grant PHY-0547102 to J.F.B. and an Alfred P. Sloan Fellowship to T. Thompson.
 among others and need not be repeated here.,\nocite{fer98} among others and need not be repeated here.
" In light of these discussions, we feel — perhaps optimistically — that the true external uncertainty in the Virgo distance is close to 20.2 mag."," In light of these discussions, we feel – perhaps optimistically – that the true external uncertainty in the Virgo distance is close to $\pm 0.2$ mag."
 The distance to Fornax is not as well determined: fewer galaxies have been studied: the differences among the PNLF. Cepheid and SBF distance estimates are larger than their internal error margins: and a TRGB calibration is not yet in hand.," The distance to Fornax is not as well determined: fewer galaxies have been studied; the differences among the PNLF, Cepheid and SBF distance estimates are larger than their internal error margins; and a TRGB calibration is not yet in hand."
" Of the three published Cepheid galaxies (again, with the incompleteness-corrected values from Ferrareseetal. 1999b)), two (NGC 1326A, 1365) have published moduli ~0.2 mag larger than the SBF or PNLF values, while the third (NGC 1425) is ~0.6 mag larger."," Of the three published Cepheid galaxies (again, with the incompleteness-corrected values from \cite{fer99b}) ), two (NGC 1326A, 1365) have published moduli $\sim 0.2$ mag larger than the SBF or PNLF values, while the third (NGC 1425) is $\sim 0.6$ mag larger."
" This discrepancy reinforees the serious concern that the outlying Cepheid spirals may either be much more widely spread than the central ellipticals which are most relevant for GCLF studies, or at a different mean In principle. a better way to approach the comparison is to use methods relating to the ellipticals alone."," This discrepancy reinforces the serious concern that the outlying Cepheid spirals may either be much more widely spread than the central ellipticals which are most relevant for GCLF studies, or at a different mean In principle, a better way to approach the comparison is to use methods relating to the ellipticals alone."
" The measured GCLF turnovers of Virgo and Fornax E galaxies can be used to compute a relative Fornax-Virgo distance, with the that they have fundamentally the same luminosity."," The measured GCLF turnovers of Virgo and Fornax E galaxies can be used to compute a relative Fornax-Virgo distance, with the that they have fundamentally the same luminosity."
" Six ellipticals in cach cluster have GCLFs measured to <1 mag beyond the turnover point, with results as summarized in Table 7.."," Six ellipticals in each cluster have GCLFs measured to $\gtsim 1$ mag beyond the turnover point, with results as summarized in Table \ref{tab:vlist}."
" The difference between the weighted averages of V"" gives Aj(Fornax-Virgo) =0.14+0.05.", The difference between the weighted averages of $V^0$ gives $\Delta \mu$ (Fornax-Virgo) $= 0.14 \pm 0.05$.
" Adding this distance offset to the adopted Virgo modulus then suggests ji(Fornax,GCLP) =31.13+ 0.08."," Adding this distance offset to the adopted Virgo modulus then suggests $\mu_0$ (Fornax,GCLF) $=31.13\pm0.08$ ."
" This value, in turn, is in entirely reasonable agreement with the SBF (31.23= 0.06) and PNLF (31.204 0.14) determinations of the Fornax distance."," This value, in turn, is in entirely reasonable agreement with the SBF $31.23\pm0.06$ ) and PNLF $31.20\pm0.14$ ) determinations of the Fornax distance."
" However, it disagrees strongly with the mean Cepheid distance or with Ferrarese al.’ss recalibrated SBF For the purposes of this discussion, we will defer further use of the Fornax system and employ only the Virgo ellipticals as our calibrators for the GCLF turnover luminosity."," However, it disagrees strongly with the mean Cepheid distance or with Ferrarese s recalibrated SBF For the purposes of this discussion, we will defer further use of the Fornax system and employ only the Virgo ellipticals as our calibrators for the GCLF turnover luminosity."
 There are two obvious ways to obtain the Coma/Virgo distance ratio through the GCLF data: (a) compare the GCLF turnovers of just the two central cD galaxies. M87 and NGC 4874: or (b) compare the mean GCLF turnovers of all Virgo giant ellipticals with the mean for our two Coma ellipticals.," There are two obvious ways to obtain the Coma/Virgo distance ratio through the GCLF data: (a) compare the GCLF turnovers of just the two central cD galaxies, M87 and NGC 4874; or (b) compare the mean GCLF turnovers of all Virgo giant ellipticals with the mean for our two Coma ellipticals."
" The latter route has the advantage of greater statistical weight over many galaxies, while the former has the advantage of matching similar types of galaxies as strictly as possible, albeit at the cost of greater internal uncertainty."," The latter route has the advantage of greater statistical weight over many galaxies, while the former has the advantage of matching similar types of galaxies as strictly as possible, albeit at the cost of greater internal uncertainty."
" As we will see below, the two approaches turn out to agree quite well."," As we will see below, the two approaches turn out to agree quite well."
" At various times in the literature. concerns have been raised regarding possible dependences of the GCLF shape and peak on environment, galaxy type, and cluster metallicity."," At various times in the literature, concerns have been raised regarding possible dependences of the GCLF shape and peak on environment, galaxy type, and cluster metallicity."
" By restricting the comparison to M87 and NGC 4874 alone, we can minimize these worries."," By restricting the comparison to M87 and NGC 4874 alone, we can minimize these worries."
" Both are centrally placed giant ellipticals with cD envelopes, and though they are not identical in size or luminosity (NGC 4874 is a magnitude more luminous, and Coma is distinctly richer than Virgo), the comparison 1s certainly closer than with an average E galaxy."," Both are centrally placed giant ellipticals with cD envelopes, and though they are not identical in size or luminosity (NGC 4874 is a magnitude more luminous, and Coma is distinctly richer than Virgo), the comparison is certainly closer than with an average E galaxy."
 M87 has one ofthe best-studied of GCS luminosity functions: repeated observations to ever-increasing depth and radial coverage have extended the GCLF far past the turnover and have now made it possible (for example) to define GCLFs separately for the two parts of its clear bimodal color distribution or as a function of galactocentric distance., M87 has one of the best-studied of GCS luminosity functions: repeated observations to ever-increasing depth and radial coverage have extended the GCLF far past the turnover and have now made it possible (for example) to define GCLFs separately for the two parts of its clear bimodal color distribution or as a function of galactocentric distance.
" For M87, the best determination of the GCLF turnover is by Kundu ((1999).. who find V""(M87) =23.67+0.06 from their deep WFPC? photometry, after subtraction of an adopted foreground extinction Ay=0.067+0.04 mag."," For M87, the best determination of the GCLF turnover is by Kundu \nocite{kun99}, , who find $V^0$ (M87) $= 23.67 \pm 0.06$ from their deep WFPC2 photometry, after subtraction of an adopted foreground extinction $A_V = 0.067 \pm 0.04$ mag."
" Subtracting V°(M87) from our NGC 4874 determination, we immediately derive Ni (Coma-Virgo) =4.15x0.14 (with the plausible assumption that both galaxies are at or near the physical centers of their clusters)."," Subtracting $V^0$ (M87) from our NGC 4874 determination, we immediately derive $\Delta\mu_0$ (Coma-Virgo) $=4.15
\pm 0.14$ (with the plausible assumption that both galaxies are at or near the physical centers of their clusters)."
 A further sophistication on this comparison can be taken if we look more narrowly at the metallicity and radial dependences., A further sophistication on this comparison can be taken if we look more narrowly at the metallicity and radial dependences.
 Several deep photometric studies (Grillmairetal.1986;; Lauer&Kormendy 1986;; McLaughlinetal. 1994:; Harrisetal. 1998b:: Kunduetal. 1999)) show radial trends in the GCLE turnover can be ignored as the level of variation is at the £0.2—maglevel or less., Several deep photometric studies \cite{gri86}; \cite{lau86}; \cite{mcl94}; \cite{har98b}; ; \cite{kun99}) ) show radial trends in the GCLF turnover can be ignored as the level of variation is at the $\pm 0.2-$ maglevel or less.
" Forthe metallicity issue, Whitmore and Kundu ((1999) find that thepeak of the M87 GCLF has à"," Forthe metallicity issue, Whitmore \nocite{whi95} and Kundu (1999) \nocite{kun99} find that thepeak of the M87 GCLF has a"
{wo stars we used instead the ASI26A Li E line.,two stars we used instead the $\lambda 8126$ Li I line.
 In any case. all the Li abundances were derived by spectrum svnthesis ancl corrected by N-LTE effects according to Laverny(1999) (see this paper for details).," In any case, all the Li abundances were derived by spectrum synthesis and corrected by N-LTE effects according to \citet{abi99} (see this paper for details)."
 ILowever. one has to be very careful when interpreting the Li abundances derived rom the resonance Li I line.," However, one has to be very careful when interpreting the Li abundances derived from the resonance Li I line."
 The presence of a circumstellar component in the AGTOS Li absorption would lead us (ο overestimate the Li abundance when derived from this line., The presence of a circumstellar component in the $\lambda 6708$ Li absorption would lead us to overestimate the Li abundance when derived from this line.
 In fact. VN And. DM Gem ane V614 Mon show weak blueshifted Na D Iine absorptions. probably indicating the presence of a circumstellar envelope in these stars.," In fact, VX And, BM Gem and V614 Mon show weak blueshifted Na D line absorptions, probably indicating the presence of a circumstellar envelope in these stars."
 This circumstellar absorption is. however. not observed in (he strong AT698 Ix I resonance line which should form at about the same depth in the atmosphere (han the Li resonance line (see Darnbaum 1992).," This circumstellar absorption is, however, not observed in the strong $\lambda 7698$ K I resonance line which should form at about the same depth in the atmosphere than the Li resonance line (see Barnbaum 1992)."
 The lack of Ix I circumstellar absorption probably rules oul significant contamination of the photospheric feature., The lack of K I circumstellar absorption probably rules out significant contamination of the photospheric feature.
 Final Li abundances are shown in Table 4 in the scale log e(Lij= 12+ log(Li/Il). where li/llis the abundance of Li by number.," Final Li abundances are shown in Table 4 in the scale log $\epsilon$ $=12 +$ log(Li/H), where Li/H is the abundance of Li by number."
 From Table 4 it is clear Chat all the stars have unusual Li abundances (log e(Li)zZ 1). larger than the twpical value found in normal C-stars (log e(Lij~ 0.0). but significantly smaller than those found in the so-called super Li-vich.," From Table 4 it is clear that all the stars have unusual Li abundances (log $\epsilon$ $\gtrsim 1$ ), larger than the typical value found in normal C-stars (log $\epsilon$ $\sim 0.0$ ), but significantly smaller than those found in the so-called super Li-rich."
 WX Cve and WZ Cas are certainly super Livich stars. although these stars may not be J-tvpe stars (see above).," WX Cyg and WZ Cas are certainly super Li-rich stars, although these stars may not be J-type stars (see above)."
 The formal uncertainty in Li abundances of Table 4 ranges from 0.3-0.4 dex., The formal uncertainty in Li abundances of Table 4 ranges from 0.3-0.4 dex.
 Figure 2 shows the correlation of Li abundances versus. PC/PC ratios [ound in J- (this work) ancl N-twpe carbon stars (Abia&Isern1997)., Figure 2 shows the correlation of Li abundances versus $^{12}$ $^{13}$ C ratios found in J- (this work) and N-type carbon stars \citep{abi97}.
. J-tvpe stars are all Li-rich and note that there are also some Li-rich N-type stars., J-type stars are all Li-rich and note that there are also some Li-rich N-type stars.
 The formal error in the PC/PC ratios in Figure 2 is &G (see Abia Isern 1997)., The formal error in the $^{12}$ $^{13}$ C ratios in Figure 2 is $\pm 6$ (see Abia Isern 1997).
neither to be isothermal nor polvtropic.,neither to be isothermal nor polytropic.
 As a result our local temperature and. pressure are not simple functions of the density but arise from the evolution of the thermal energy., As a result our local temperature and pressure are not simple functions of the density but arise from the evolution of the thermal energy.
 Jecause we consider processes (c.g. radiative cooling. selferavity. star formation) whose ellectiveness depends on the density. the hydrodynamic equations are no longer scaleinvariant.," Because we consider processes (e.g. radiative cooling, self--gravity, star formation) whose effectiveness depends on the density, the hydrodynamic equations are no longer scale--invariant."
 Therefore the condition of randomness between subsequent density Uuctuations is violated and one cannot expect a log.normal density PDE (e.g. Vázzquez-Semacdeni's (19943)., Therefore the condition of randomness between subsequent density fluctuations is violated and one cannot expect a log–normal density PDF (e.g. Vázzquez-Semadeni's (1994)).
 In our series of experiments of increasing complexity (see Table 13). the simplest simulation we performed was of nonisothermal supersonic turbulence (run A).," In our series of experiments of increasing complexity (see Table \ref{sims}) ), the simplest simulation we performed was of non–isothermal supersonic turbulence (run A)."
 Despite the inclusion of density.dependent cooling processes. we found that the structure of the eas quickly evolved. to a density PDE consistent with a lognormal.," Despite the inclusion of density–dependent cooling processes, we found that the structure of the gas quickly evolved to a density PDF consistent with a log–normal."
 This is not à surprising result since without a heat source the majority of the gas quickly cools to a nearly isothermal state (see bottom row of figure 9)) with an average temperature corresponding to the minimum of the cooling curve (horizontal dashed. line in bottom row of figure 9))., This is not a surprising result since without a heat source the majority of the gas quickly cools to a nearly isothermal state (see bottom row of figure \ref{pdf_nofbknosg}) ) with an average temperature corresponding to the minimum of the cooling curve (horizontal dashed line in bottom row of figure \ref{pdf_nofbknosg}) ).
 Furthermore. the scaling [or the dispersion of the PDE given by Padoan. Nordlund Jones (1997) continued to hold.," Furthermore, the scaling for the dispersion of the PDF given by Padoan, Nordlund Jones (1997) continued to hold."
" In fact. rather than fit lognormal functions to our density. PDEs. we measured the average ol the logarithm of the gas density. <logyp and the A, ofthe gas at dillerent time instances and then overplotted Padoan. Norcluned Jones’ (1997). prediction for the lognormal distribution."," In fact, rather than fit log--normal functions to our density PDFs, we measured the average of the logarithm of the gas density, $<{\mathrm {log}}_{10} \rho>$, and the $M_{\mathrm {rms}}$ of the gas at different time instances and then overplotted Padoan, Nordlund Jones' (1997) prediction for the log–normal distribution."
" For the runs where we formecl stars (without selfgravity or feedback) in addition to having radiative cooling (runs DI and €1). the gas density PDF continued to have the same behavior: the Ma, of the system progressively declined with time. while the density PDP remained consistent with a lognormal clistribution (lie. 10))."," For the runs where we formed stars (without self–gravity or feedback) in addition to having radiative cooling (runs B1 and C1), the gas density PDF continued to have the same behavior: the $M_{\mathrm {rms}}$ of the system progressively declined with time, while the density PDF remained consistent with a log–normal distribution (fig. \ref{pdf_nofbknosgstars}) )."
 The runs which showed the first departure from normal density PDEs were the runs which incluced: selt-eravity (runs D2 and €2) but still no feedback (figure 11))., The runs which showed the first departure from log--normal density PDFs were the runs which included self-gravity (runs B2 and C2) but still no feedback (figure \ref{pdf_nofbksg}) ).
" Repeating the exercise of measuring the average of the logarithm of the gas density. <logyp>. and the Mj, of the gas at dillerent times. we found two cdillerences: (a) the Aus initially declined but then stabilized at a value higher than that seen in the runs without selfgravity. and (b) the lognormal PDI predicted by Padoan. Nordlund Jones (1997) consistenthy uncerpredicted the distribution at high eas densitv."," Repeating the exercise of measuring the average of the logarithm of the gas density, $<{\mathrm {log}}_{10} \rho>$, and the $M_{\mathrm {rms}}$ of the gas at different times, we found two differences: (a) the $M_{\mathrm {rms}}$ initially declined but then stabilized at a value higher than that seen in the runs without self–gravity, and (b) the log–normal PDF predicted by Padoan, Nordlund Jones (1997) consistently underpredicted the distribution at high gas density."
 A powerlaw fit the high density tail well., A power–law fit the high density tail well.
 In one-dimensional simulations of Burgers Lows. ic. infinitely compressible Hows. powerlaws were also found to be good fits to the density PDFs (Gotoh Ixraichnan. 1993).," In one-dimensional simulations of Burgers flows, i.e. infinitely compressible flows, power–laws were also found to be good fits to the density PDFs (Gotoh Kraichnan 1993)."
 We therefore interpret the powerlaw behavior for the run with ποgravity. as rellecting the added: possibility. of the gas. once it has a high density. to compress to even higher density. reminiscent of the behavior in Burgers Lows.," We therefore interpret the power–law behavior for the run with self–gravity, as reflecting the added possibility of the gas, once it has a high density, to compress to even higher density, reminiscent of the behavior in Burgers flows."
 IxIessen (2000) also explored the form of the density PDE for the cases of decaving and driven scll-eravitating turbulence., Klessen (2000) also explored the form of the density PDF for the cases of decaying and driven self-gravitating turbulence.
 Although he found a departure from log-normal at high densities. the departure could not be characterized by a power law.," Although he found a departure from log-normal at high densities, the departure could not be characterized by a power law."
 When we add feedback to the list of simulated processes. either with self.e&ravitv (runs B4 and €4) or without (runs BS and €3). the density PDE becomes markedly bimocal (figure 12)). illustrating that the majority of the simulation volume is occupied by low density gas.," When we add feedback to the list of simulated processes, either with self–gravity (runs B4 and C4) or without (runs B3 and C3), the density PDF becomes markedly bimodal (figure \ref{pdfcomp}) ), illustrating that the majority of the simulation volume is occupied by low density gas."
 | bimocdal density distribution is also a sign of a thermal instability (Vázzquez-Semacdeni. Gazol Scalo (2000)) the consequences of which we will discuss in a future paper (Slvz. Devriendt. Brvan Silk. preparation).," A bimodal density distribution is also a sign of a thermal instability (V\'{a}zzquez-Semadeni, Gazol Scalo (2000)) the consequences of which we will discuss in a future paper (Slyz, Devriendt, Bryan Silk, )."
 For the runs with self.gravity. the high density powerlaw tail disappears.," For the runs with self–gravity, the high density power–law tail disappears."
 Perhaps it can be argued that the high density part of the density PDI may be fit with a lognormal distribution (figure 13))., Perhaps it can be argued that the high density part of the density PDF may be fit with a log–normal distribution (figure \ref{pdf_b4_xlnfit}) ).
 The exercise of overplotting the lognormal given by Padoan. Nordlund Jones (1997) is not possible because the Ada; measured for the entire simulation box does not correspond to the Aus of the high density gas for which the lognormal function nw be a good description.," The exercise of overplotting the log–normal given by Padoan, Nordlund Jones (1997) is not possible because the $M_{\mathrm {rms}}$ measured for the entire simulation box does not correspond to the $M_{\mathrm {rms}}$ of the high density gas for which the log–normal function may be a good description."
 Llence we can only lognormals to the high density gas. similar to what others. e.g. Wada Norman (2001). Ixravstov (2003). do in their global simulations of the ISM.," Hence we can only log--normals to the high density gas, similar to what others, e.g. Wada Norman (2001), Kravstov (2003), do in their global simulations of the ISM."
 The interest. of describing the density structure of the ISAL with a single function. such as the lognormal. lies in finding a link between the σας density averaged: over kiloparsec sized regions and the high density regions which might form stars.," The interest of describing the density structure of the ISM with a single function, such as the log–normal, lies in finding a link between the gas density averaged over kiloparsec sized regions and the high density regions which might form stars."
 This is precisely the link required. for an explanation of the Schmidt law., This is precisely the link required for an explanation of the Schmidt law.
 Rewriting the Schmidt law in a form where the star formation rate is equal to some constants multiplied by the fraction of gas in. high density regions and by the gas density averaged over large scales (his equation 7). Elmegreen (2002) emphasized. that star formation rates depend on the geometry of the density field. i.c. the PDP.," Rewriting the Schmidt law in a form where the star formation rate is equal to some constants multiplied by the fraction of gas in high density regions and by the gas density averaged over large scales (his equation 7), Elmegreen (2002) emphasized that star formation rates depend on the geometry of the density field, i.e. the PDF."
 Lo the shape of the density. PDE is universal. then the fraction of gas in high density regions is known.," If the shape of the density PDF is universal, then the fraction of gas in high density regions is known."
 Consequently. if the high. density. regions are also sell&ravitating. then the fraction of gas available for star formation is also known.," Consequently, if the high density regions are also self–gravitating, then the fraction of gas available for star formation is also known."
 AXdmittedlv. the density PDE contains no spatial information. hence there is no reason for which the high density regions should find themselves to be spatially contiguous. so that they comprise regions of mass ereater than the Jeans mass.," Admittedly, the density PDF contains no spatial information, hence there is no reason for which the high density regions should find themselves to be spatially contiguous, so that they comprise regions of mass greater than the Jeans mass."
 In fact. figure Ll clearly shows that at least some of the dense gas regions are not contiguous because if they were they would simply not persist as all the eas would be converted to stars on a dynamical timescale since these regions are well above pig and cold.," In fact, figure \ref{pdf_nofbksg} clearly shows that at least some of the dense gas regions are not contiguous because if they were they would simply not persist as all the gas would be converted to stars on a dynamical timescale since these regions are well above $\rho_{\mathrm {crit}}$ and cold."
 We therefore have to identify these regions with divergent gas Lows., We therefore have to identify these regions with divergent gas flows.
 AX two-dimensional study of the ISM in a galactic disk by Wada Norman (2001) has claimed that the log.normal distribution is à robust description. of the ISM density distribution. over many orders of magnitude in censity. regardless of the simulated physics.," A two-dimensional study of the ISM in a galactic disk by Wada Norman (2001) has claimed that the log–normal distribution is a robust description of the ISM density distribution over many orders of magnitude in density, regardless of the simulated physics."
 More. specifically. in their simulations the presence of stellar feedback does not change the shape of the PDE but increases the dispersion," More specifically, in their simulations the presence of stellar feedback does not change the shape of the PDF but increases the dispersion"
 (Widrow2002).. (Caasso&Ru," \citep{Widrow2002}, \citep{Grasso2001}."
binstein2O01).. Kroubere1991:Blasietal.1999) (Barrowctal.1997: ~100° (Aharonianetal.1991) (Plaga—1995).. (Neronoy& (Neronov&Senikoz2009).," \citep{Kronberg1982,
	Kronberg1994,Blasi1999} \citep{Barrow1997,Durrer2000} $\sim10^{-9}$ \citep{Aharonian1994} \citep{Plaga1995}, \citep{Neronov2007,Elyiv2009}. \citep{Neronov2009,Dolag2009}."
. Ez100 zE300 measured fiux or upper luit iu the WE band (Ahwaseetal.2008:Neronov&Vovk 2010).," $E\gtrsim100$ $\lesssim E\lesssim 300$ measured flux or upper limit in the HE band \citep{Murase2008,Neronov2010}."
. Analytic cascade models assmunuime a simple relationship between the cascade flux and ECAIF streneth/ have. put the lower lianit at LO196 to 10.1 Gauss when the source livetine is uulimited (Tavecchioetal.2010.2011:Neronov&Vovk 2010).. similar to the results of Monte Carlo siuulatious (Dolagetal.2011:Taxlor 2011)..," Analytic cascade models assuming a simple 	relationship between the cascade flux and EGMF strength have put the lower limit at $10^{-16}$ to $10^{-15}$ Gauss when the source livetime is 	unlimited \citep{Tavecchio2010,Tavecchio2011,Neronov2010}, similar to the results of Monte Carlo simulations \citep{Dolag2011,Taylor2011}. ."
 If the cascade time delay is limited to the ~3 vears of simultaneous ITE and VITE observations. the lower μπιτ becomes 10.1? to 10.15 οπως according to the simple cascade models (Deineretal.2011). or 10. to 10.17 Gauss according to the simulations (Tayloretal.2011).," If the cascade time delay is limited to the $\sim3$ years of simultaneous HE and VHE observations, the lower limit becomes $10^{-19}$ to $10^{-18}$ Gauss according to the simple cascade models \citep{Dermer2011}, or $10^{-18}$ to $10^{-17}$ Gauss according to the simulations \citep{Taylor2011}."
. Iu this Letter. we preseut a new seui-uialvtie model of the electromagnetic cascade.," 	In this Letter, we present a new semi-analytic model of the electromagnetic cascade."
 Iu coutrast to previous analytic cascade models (Denieretal.2011:Nerouov&Senukoz2009:Tavecchioetal. 2011)... our cascade model considers the full track of the primary photon without the assmuption of interacting exclusively at the mean free path.," In contrast to previous analytic cascade models \citep{Dermer2011,Neronov2009,Tavecchio2011}, our cascade model considers 	the full track of the primary photon without the assumption of interacting exclusively at the mean free path."
 This simultancously accounts for both angular and temporal constraimts iu a natural wav., This simultaneously accounts for both angular and temporal constraints 	in a natural way.
 In addition. we model the radiation backeromnds and source emission iu greater detail.," In addition, we model the radiation backgrounds and source emission in greater detail."
 As a complementary approach to full Moute Carlo simulations (Dolagetal.2011:Tavlor2011:Arlenetal. 2011)... our model serves as a tool for clarifving the cascade picture. rapidly searching through the parameter space. and interpreting simulation results.," As a complementary approach to full Monte Carlo simulations \citep{Dolag2011,Taylor2011,preparation}, 	our model serves as a tool for clarifying the cascade picture, rapidly searching through the parameter space, and interpreting simulation results."
 We also develop a systematicframeworkforapplying our cascade models predictionsto derive lower linitsonthe ECAIF streneth at specific confidence levels., 	We also develop a systematicframeworkforapplying our cascade model's predictionsto derive lower limitsonthe EGMF strength at specific confidence levels.
 This framework is applicable to theresults of Monte Carlo sinmlations as well., This framework is applicable to theresults of Monte Carlo simulations as well.
concentrated. on the observational charcteristics related to ongoing star formation. e.g. the presence of O stars. CO emission.,"concentrated on the observational charcteristics related to ongoing star formation, e.g. the presence of O stars, CO emission."
 Furthermore. the birth sites of their clusters were assumed to lie along an imposecl spiral. rather than derived from numerical moclels.," Furthermore, the birth sites of their clusters were assumed to lie along an imposed spiral, rather than derived from numerical models."
 In this paper we carry out numerical simulations of gas How in mocel galaxies in which the spiral arms are excited bv the various theoretical mechanisms and then locate the regions in which ‘star clusters’. formed. in the dense arm regions. can be expected to be found at later times.," In this paper we carry out numerical simulations of gas flow in model galaxies in which the spiral arms are excited by the various theoretical mechanisms and then locate the regions in which `star clusters', formed in the dense arm regions, can be expected to be found at later times."
 We consider four canonical galaxy models. corresponding to the four basic theoretical moclels for spiral arm and/or bar formation.," We consider four canonical galaxy models, corresponding to the four basic theoretical models for spiral arm and/or bar formation."
" These are: (i) a galaxy with an imposed. spiral potential with fixed pattern speed. (ii) a galaxy which is bar unstable. (ii) a ""occulent galaxy in which the arms are intermittent and driven by local gravitational instabilities. and (iv) à galaxy which is subject to a strong external tidal interaction."," These are: (i) a galaxy with an imposed spiral potential with fixed pattern speed, (ii) a galaxy which is bar unstable, (iii) a `flocculent' galaxy in which the arms are intermittent and driven by local gravitational instabilities, and (iv) a galaxy which is subject to a strong external tidal interaction."
 We then make predictions for the distribution of stellar clusters of dilferent. ages., We then make predictions for the distribution of stellar clusters of different ages.
 We find that dilferent spatial clistributions arise at dilferent cluster ages depending on the underlving dynamics of the galaxy. and on the spiral excitation mechanism.," We find that different spatial distributions arise at different cluster ages depending on the underlying dynamics of the galaxy, and on the spiral excitation mechanism."
 We finally suggest. that by using methods for age-dating clusters. e.g. in the recent work by 7.. it may be possible to identify what is the underlving mechanism producing spiral structure in individual nearby galaxies.," We finally suggest that by using methods for age-dating clusters, e.g. in the recent work by \citet{Fall2009}, it may be possible to identify what is the underlying mechanism producing spiral structure in individual nearby galaxies."
 In Section 2.. we outline our basic numerical mocel. roth for the galaxy dwnamies. and for the identification of ocations for age-dated star clusters.," In Section \ref{models}, we outline our basic numerical model, both for the galaxy dynamics, and for the identification of locations for age-dated star clusters."
 We then apply. these echniques to the four excitation mechanisms mentioned above: ‘density wave theory! in Section 3.. bar-driven waves in Section 4.. Hlocculent behaviour in Section 5 and external ides in Section 6..," We then apply these techniques to the four excitation mechanisms mentioned above: `density wave theory' in Section \ref{linshu}, bar-driven waves in Section \ref{bar}, flocculent behaviour in Section \ref{flocculent} and external tides in Section \ref{m51}."
 In Section 7. we illustrate our findings by plotting the distributions of cluster ages across spiral armis or the dillerent excitation mechanisms., In Section \ref{xsection} we illustrate our findings by plotting the distributions of cluster ages across spiral arms for the different excitation mechanisms.
" We summarise our indings ancl provide discussion in Section δι,", We summarise our findings and provide discussion in Section \ref{discussion}.
 All the calculations presented here use an SPL code. developed originally by ? and. substantially modified. to include sink particles (?).. individual timesteps (?) and individual smoothing lengths (7)..," All the calculations presented here use an SPH code, developed originally by \citet{Benz1990} and substantially modified to include sink particles \citep{Batesph1995}, individual timesteps \citep{Bate1995} and individual smoothing lengths \citep{Price2004}."
 For all the calculations. except the. fixed. spiral (Section 3)). we fully model the galaxy. with particles for stars. gas. and the dark matter halo.," For all the calculations except the fixed spiral (Section \ref{linshu}) ), we fully model the galaxy with particles for stars, gas, and the dark matter halo."
 Since we are primarilv interested. in the gas flow. and in particular the σας [ow downstream of any density maxima which might be taken to give rise to star cluster formation. we use relatively low eas surface densities.," Since we are primarily interested in the gas flow, and in particular the gas flow downstream of any density maxima which might be taken to give rise to star cluster formation, we use relatively low gas surface densities."
 In the case of the fixed spiral we do not include the self-eravity of the gas., In the case of the fixed spiral we do not include the self-gravity of the gas.
 In the moclels which clo include gas selt-gravity (Sections 4.. 5.. and 6)) we choose a relatively high gas temperature (103 Ix). in order to avoid widespread gravitational collapse.," In the models which do include gas self-gravity (Sections \ref{bar}, \ref{flocculent}, and \ref{m51}) ) we choose a relatively high gas temperature $10^4$ K), in order to avoid widespread gravitational collapse."
 La all models the gas is taken to be isothermal., In all models the gas is taken to be isothermal.
 Thus we make no pretence of modelling star-formation. feedback. radiative processes in the ISM. and so on. in any detail.," Thus we make no pretence of modelling star-formation, feedback, radiative processes in the ISM, and so on, in any detail."
 Rather we are trying to identily those regions in which the gas tends to have higher than average density. and will then identify those regions as the ones in which star clusters are most likely to form.," Rather we are trying to identify those regions in which the gas tends to have higher than average density, and will then identify those regions as the ones in which star clusters are most likely to form."
 Even so. in the runs with self-gravitv. we have still had to insert a few sink particles to replace the highest density regions in order to ensure that the simulations did not take too long.," Even so, in the runs with self-gravity, we have still had to insert a few sink particles to replace the highest density regions in order to ensure that the simulations did not take too long."
 The number of sink particles used is low (see Sections 4.. 5 and ?)).," The number of sink particles used is low (see Sections \ref{bar}, \ref{flocculent} and \citealt{Dobbs2010}) )."
 Since in our simulations we cannot resolve the formation of star clusters we instead locate dense gas. which would. be where stars are more likely to form. given a more realistic surface density. (or self-gravitv. for the fixed spiral) and sullicient resolution.," Since in our simulations we cannot resolve the formation of star clusters we instead locate dense gas, which would be where stars are more likely to form, given a more realistic surface density (or self-gravity, for the fixed spiral) and sufficient resolution."
 For cach simulation. we locate gas at high clensity at times between 2 Myr and 130 Myr before a given time frame.," For each simulation, we locate gas at high density at times between 2 Myr and 130 Myr before a given time frame."
 So for example. for the calculation with a fixed spiral potential (Section 3)). we take a time frame of 255 Myr to represent the present.," So for example, for the calculation with a fixed spiral potential (Section \ref{linshu}) ), we take a time frame of 255 Myr to represent the present."
 We then locate dense eas al times of 125 Myr. 155 Myr. 205 Myr and 253 Myr.," We then locate dense gas at times of 125 Myr, 155 Myr, 205 Myr and 253 Myr."
 We assume that the dense gas represents the location of stellar clusters forming at those times. and then plot the location of that gas at the time of 255 Myr.," We assume that the dense gas represents the location of stellar clusters forming at those times, and then plot the location of that gas at the time of 255 Myr."
 Thus we obtain estimates for the location of star clusters of ages of 2 Myr. 50 Myr. 100 Myr and 130 Myr.," Thus we obtain estimates for the location of star clusters of ages of 2 Myr, 50 Myr, 100 Myr and 130 Myr."
 The definition of ‘dense gas? requires some care. since in general the mean eas density decreases with radius.," The definition of `dense gas' requires some care, since in general the mean gas density decreases with radius."
 In. practice. we found for the models except the fixed spiral. it was sullicient to select gas which has a density more than 3 times the average (mass weighted) surface density for a given racial bin. as being regions where star formation would be Likely to take place.," In practice, we found for the models except the fixed spiral, it was sufficient to select gas which has a density more than 3 times the average (mass weighted) surface density for a given radial bin, as being regions where star formation would be likely to take place."
 For the fixed spiral (Section 3). we used colder gas. which shocks to higher densities. and thus chose a density of 10 times the average.," For the fixed spiral (Section 3), we used colder gas, which shocks to higher densities, and thus chose a density of 10 times the average."
 ltadial bins were chosen to have width NAIs0.05., Radial bins were chosen to have width $\Delta R/R \approx 0.08$.
 As we do not explicitly include star formation in these simulations. we cannot. distinguish. between SPILL particles which represent star clusters and those which represent gas.," As we do not explicitly include star formation in these simulations, we cannot distinguish between SPH particles which represent star clusters and those which represent gas."
 Thus. we assume that the trajectories of the gas particles ave not dissimülar to stars.," Thus, we assume that the trajectories of the gas particles are not dissimilar to stars."
 In reality stellar clusters. lose their gas over a time period of a few Myr. after. which they are not subject to gas pressure.," In reality stellar clusters lose their gas over a time period of a few Myr, after which they are not subject to gas pressure."
 However. the role of gas pressure is most important as eas passes through a spiral shock.," However, the role of gas pressure is most important as gas passes through a spiral shock."
 In the simulations. the dense gas that we assume represents the locations of star cluster formation has alreacly passed through the shock.," In the simulations, the dense gas that we assume represents the locations of star cluster formation has already passed through the shock."
 The velocity of the gas is highly supersonic and thus the elfect of gas pressure is in eeneral small. except in shocks.," The velocity of the gas is highly supersonic and thus the effect of gas pressure is in general small, except in shocks."
 We therefore expect those stars formed in an arm. together with the eas from which they form. to emerge from a spiral arm with similar space velocities. and to continue on neighbouring trajectories until jev pass through the next spiral shock.," We therefore expect those stars formed in an arm, together with the gas from which they form, to emerge from a spiral arm with similar space velocities, and to continue on neighbouring trajectories until they pass through the next spiral shock."
 In our calculations. the majority of the cluster locations we identify occur before 10 gas has gone through the next spiral arm.," In our calculations, the majority of the cluster locations we identify occur before the gas has gone through the next spiral arm."
 We test this assumption that gas pressure is. [or the most part. negligible. explicitly in Section 3.," We test this assumption that gas pressure is, for the most part, negligible, explicitly in Section 3."
 The model galaxy discussed in this Section is subjected to a fixed spiral potential rotating with a fixed global pattern speed., The model galaxy discussed in this Section is subjected to a fixed spiral potential rotating with a fixed global pattern speed.
 Lt is intended to represent the ? model of quasi- spiral structure. which predicts the presence of a global spiral mode.," It is intended to represent the \citet{Lin1964} model of quasi-steady spiral structure, which predicts the presence of a global spiral mode."
 We note that a galaxy which exhibits a, We note that a galaxy which exhibits a
lu order to formulate the ecquatious of trausler in a gaseous medium the most convenient representaion of polarized radiation is by a set of four parameters called Stokes parameters.,In order to formulate the equations of transfer in a gaseous medium the most convenient representation of polarized radiation is by a set of four parameters called Stokes parameters.
 Chandrasekharκ(1960). first introduced the Stokes parameters in the equation of radiative transfer with a slightM modification of Stoke’s representation., \cite{chandra60} first introduced the Stokes parameters in the equation of radiative transfer with a slight modification of Stoke's representation.
 In au elliptically polarized beam. the vibrations ol the electricI aud the maguetic vectors in the plane trausverse to the direction of propagation are such tIHuo tlie ratio of the amplitucles aud cdiference in phases of the components in any two directions at right angles to each otlier are absolute constants.," In an elliptically polarized beam, the vibrations of the electric and the magnetic vectors in the plane transverse to the direction of propagation are such that the ratio of the amplitudes and difference in phases of the components in any two directions at right angles to each other are absolute constants."
 A regular vibration of this character can be represented by where £j aud & are the compouents of the vibration aoug two directious { aud rat right angles to each other. w the circular [requency of the vibration. and £j(0). £&(0). ej and e; are constants.," A regular vibration of this character can be represented by where $\xi_{l}$ and $\xi_{r}$ are the components of the vibration along two directions $l$ and $r$ at right angles to each other, $\omega$ the circular frequency of the vibration, and $\xi^{(0)}_{l}$, $\xi^{(0)}_{r}$, $\epsilon_{l}$ and $\epsilon_{r}$ are constants."
 If the principal axes of the ellipse clescribec by (£j£.) are in directions making angles y and \+dx to the direction £. the equations represenine the vibration take the simplified [orum: where 9 denotes au augle whose taigent is the ratio of the axes of the ellipse traced by the eud point of the electric for magnetic) veclor. aud the utumerical values of it lies between 0 aud 5.1 ," If the principal axes of the ellipse described by $(\xi_{l},\xi_{r})$ are in directions making angles $\chi$ and $\chi+\frac{1}{2}\pi$ to the direction $l$, the equations representing the vibration take the simplified forms: where $\beta$ denotes an angle whose tangent is the ratio of the axes of the ellipse traced by the end point of the electric (or magnetic) vector, and the numerical values of it lies between 0 and $\frac{1} 
{2}\pi$."
The sigu of ij is positive or negative accordiug as the polarization is riglit-haudecd or left-bauclect., The sign of $\beta$ is positive or negative according as the polarization is right-handed or left-handed.
" Iu equation (3)). &£c(0)"" denotes a quautity proportional to the mean amplitude of the electric vector aid whose square is equal to the 1uteusity of the beam: Following the representation given i1 equation (2)) one obtains for the vibrations in the / aud r directions aud Tlje iuteusities1ties { aud LiJ, iu the directionsclirect /{ audl x cau |be writtentt as"," In equation \ref{2}) ), $\xi^{(0)}$ denotes a quantity proportional to the mean amplitude of the electric vector and whose square is equal to the intensity of the beam: Following the representation given in equation \ref{2}) ) one obtains for the vibrations in the $l$ and $r$ directions and The intensities $I_{l}$ and $I_{r}$ in the directions $l$ and $r$ can be written as"
We compiled radio observations of a sample of IFRS. and added dedicated high-frequency. high-resolution observations to investigate the nature of these objects.,"We compiled radio observations of a sample of IFRS, and added dedicated high-frequency, high-resolution observations to investigate the nature of these objects."
 Using the IR detection limits we computed their ratio of GGHz to 4m flux density and compared them with the general radio source population and a sample of high-redshift radio galaxies., Using the IR detection limits we computed their ratio of GHz to $\mu$ m flux density and compared them with the general radio source population and a sample of high-redshift radio galaxies.
 Our conclusions are as follows., Our conclusions are as follows.
 We have presented further evidence that IFRS are a distinct group of radio sources which are principally detected and studied via radio observations., We have presented further evidence that IFRS are a distinct group of radio sources which are principally detected and studied via radio observations.
 Although the brighter IFRS are likely to be similar to HZRG (Seymouretal.2007.. Jarvisetal. 2001)). they are. on average. probably much less luminous.," Although the brighter IFRS are likely to be similar to HzRG \citealt{Seymour2007}, \citealt{Jarvis2001}) ), they are, on average, probably much less luminous."
 So the IFRS probably represent obscured. radio-loud AGN which have not previously been studied.," So the IFRS probably represent obscured, radio-loud AGN which have not previously been studied."
" It is expected that future radio surveys of the sensitivity of the ATLAS survey. such as ASKAP-EMU. combined with infrared observations. will uncover IFRS in the thousands. allowing further examination of their characteristics,"," It is expected that future radio surveys of the sensitivity of the ATLAS survey, such as ASKAP-EMU, combined with infrared observations, will uncover IFRS in the thousands, allowing further examination of their characteristics."
papers in (his series we adopt the exünction law of Schlegel οἱ al. (,papers in this series we adopt the extinction law of Schlegel et al. (
"1998) and fit a straight line to the relation (1—M),=(nAY)Ay(mAL),E(B—V)sRy.",1998) and fit a straight line to the relation $(m-M)_{0} = (m-M)_{\lambda} - A_{\lambda} = (m-M)_{\lambda} - E(B-V) * R_{\lambda}$.
" Using the corrected. reddened distance moduli in the V ancl I1 photometric bands as given above. together with the values for the J and IX bands obtained in (his paper. we obtain the following values for the reddening aud (he true distance modulus of WEM trom the multiwavelength analysis: E(B-VW)=0.0820.020 (n—M),24.924+0.042. This corresponds to a distance of WLM of 0.97 + 0.02 Alpe."," Using the corrected, reddened distance moduli in the V and I photometric bands as given above, together with the values for the J and K bands obtained in this paper, we obtain the following values for the reddening and the true distance modulus of WLM from the multiwavelength analysis: $ E(B-V) = 0.082 \pm 0.020$ $(m-M)_{0} = 24.924 \pm 0.042$, This corresponds to a distance of WLM of 0.97 $\pm$ 0.02 Mpc."
 In Table 4. we give the adopted values of 2) ancl the unreddened. true distance moduli in each band which are obtained with the reddening value determined in our multi-wavelength solution.," In Table 4, we give the adopted values of $R_{\lambda}$ and the unreddened, true distance moduli in each band which are obtained with the reddening value determined in our multi-wavelength solution."
 The agreement between (he unreddened distance moduli obtained in each band is verv good., The agreement between the unreddened distance moduli obtained in each band is very good.
 In Fig., In Fig.
 6. we plot the apparent Cepheid distance moduli for WLM in VIJI as a [function of Ay. and the best fitting straight line to the data.," 6, we plot the apparent Cepheid distance moduli for WLM in VIJK as a function of $R_{\lambda}$, and the best fitting straight line to the data."
 It is appreciated that the total reddening. and the true distance modulus of WLM ave indeed. very. well determined from (his fit.," It is appreciated that the total reddening, and the true distance modulus of WLM are indeed very well determined from this fit."
 By comparison with the foreground reddening of 0.02 mag (Schlegel et al., By comparison with the foreground reddening of 0.02 mag (Schlegel et al.
" 1998) it is seen (hat most of the total average reddening of the WLAI Cepheids is produced inside their host galaxy,", 1998) it is seen that most of the total average reddening of the WLM Cepheids is produced inside their host galaxy.
 As in the previous distance determinations of Local ancl Seulptor Group irregular and spiral galaxies [rom combined Cepheid near-infrared and optical photometry (see references cited in the Introduction) obtained in the Aranearia Project. our combined near-inlrared and optical data for a sample of 31 Cepheids in WLM have led to a very. robust. distance determination for this relatively distant Local Group galaxy.," As in the previous distance determinations of Local and Sculptor Group irregular and spiral galaxies from combined Cepheid near-infrared and optical photometry (see references cited in the Introduction) obtained in the Araucaria Project, our combined near-infrared and optical data for a sample of 31 Cepheids in WLM have led to a very robust distance determination for this relatively distant Local Group galaxy."
 Again. as for the other galaxies we have studied so far. we find that the slopes of the Cepheid PL relations in WLM in the near-inlrared J and Ix bands are statistically in agreement with the one defined by the LAIC! Cepheids (Persson et al.," Again, as for the other galaxies we have studied so far, we find that the slopes of the Cepheid PL relations in WLM in the near-infrared J and K bands are statistically in agreement with the one defined by the LMC Cepheids (Persson et al."
 2004)., 2004).
 Indeed. the [ree fits to the data discussed in section 3 of this paper vield slopes in both J and Ix which agree within 1 σ with the adopted slope of Persson et al. (," Indeed, the free fits to the data discussed in section 3 of this paper yield slopes in both J and K which agree within 1 $\sigma$ with the adopted slope of Persson et al. ("
2004) from the LAIC Cepheids.,2004) from the LMC Cepheids.
 If we exclude the long-period Cepheid in the fits which carries a strong weight in the linear regression. the resulting values for the slopes of the PL relation in J and Ix become -2.368 + 0.426. and -2.756 + 0.368. respectively.," If we exclude the long-period Cepheid in the fits which carries a strong weight in the linear regression, the resulting values for the slopes of the PL relation in J and K become -2.868 $\pm$ 0.426, and -2.756 $\pm$ 0.368, respectively."
 Due to their large uncertainties. which is a consequence of (he now very restricted. period range of the remaining 23 Cepheicl sample. (hese values still agree within 1.5 σ with the," Due to their large uncertainties, which is a consequence of the now very restricted period range of the remaining 23 Cepheid sample, these values still agree within 1.5 $\sigma$ with the"
coverage.,coverage.
 Depending on the geometries and locations of the continuum source and broad emission line regions (and empirically on their positions relative to the emission lines) the absorber can partially cover either or both 1999)., Depending on the geometries and locations of the continuum source and broad emission line regions (and empirically on their positions relative to the emission lines) the absorber can partially cover either or both \citep{gan99}.
. There may be overlap between the aand ccategories as some absorption systems can possess both strong aand lines., There may be overlap between the and categories as some absorption systems can possess both strong and lines.
 Further observations covering both of these lines are needed to investigate this issue., Further observations covering both of these lines are needed to investigate this issue.
 These different types of NALs are of interest because they may allow us to probe different regions of the outflow., These different types of NALs are of interest because they may allow us to probe different regions of the outflow.
 These three families of NALs may represent different lines of sight through the outflow. which is also suggested by the relations between the properties of UV NALs and the X-ray properties of the quasars that display them (e.g..Chartasetal.2009).," These three families of NALs may represent different lines of sight through the outflow, which is also suggested by the relations between the properties of UV NALs and the X-ray properties of the quasars that display them \citep[e.g.,][]{chartas09}."
. The question we address in this paper is the origin of the intrinsic NALs., The question we address in this paper is the origin of the intrinsic NALs.
 To this end we construct photoionization models for the strong absorption systems in the spectra of three radio-quiet quasars from the HIRES/Keck sample studied by Misawaetal. (2007).., To this end we construct photoionization models for the strong absorption systems in the spectra of three radio-quiet quasars from the HIRES/Keck sample studied by \citet{mis07}. .
 This sample contains 37 optically bright quasars at z=2.4., This sample contains $37$ optically bright quasars at $z=2-4$.
" We choose these particular three quasars because their spectra exhibit characteristics of the ""strong V7 family."," We choose these particular three quasars because their spectra exhibit characteristics of the “strong "" family."
 In addition. these three systems offer many observational constraints because many ions are covered in their spectra. and the aabsorption lines have simple profiles.," In addition, these three systems offer many observational constraints because many ions are covered in their spectra, and the absorption lines have simple profiles."
 Similar systems. containing multiple absorption components spread over thousands of hhave been reported in the spectra of RX | OLS (Gangulyetal.2003) and 3C 351 (Yuanetal.2002).," Similar systems, containing multiple absorption components spread over thousands of have been reported in the spectra of RX $+$ 0115 \citep{gan03} and 3C 351 \citep{yua02}."
. In Section 2.. we describe the absorption profiles of the three intrinsic NAL systems.," In Section \ref{data}, we describe the absorption profiles of the three intrinsic NAL systems."
 In Section 3.. we introduce our method for modelling the three systems using the Cloudy photoionization code (Ferland2006).," In Section \ref{method}, we introduce our method for modelling the three systems using the Cloudy photoionization code \citep{fer06}."
. The modelling results. in the form of constraints on metallicity. ionization parameter. and volume number density are presented in Section 4..," The modelling results, in the form of constraints on metallicity, ionization parameter, and volume number density are presented in Section \ref{results}."
 In Section 6.. we discuss possible interpretations of our results on the location of the absorbers in the quasar winds.," In Section \ref{discussion}, we discuss possible interpretations of our results on the location of the absorbers in the quasar winds."
 We present a summary and conclusion in Section 7.., We present a summary and conclusion in Section \ref{conclusion}.
 The cosmology we use in this paper is Q4=0.7. ον=0.3. fi=0.7. which leads to the luminosity distances listed in Table I..," The cosmology we use in this paper is $\Omega_\Lambda=0.7$, $\Omega_{\rm
 M}=0.3$, $h=0.7$, which leads to the luminosity distances listed in Table \ref{tab-sumsys}."
 The Keck/HIRES spectra of our three intrinsic systems are described in Misawaetal.(2007)., The Keck/HIRES spectra of our three intrinsic systems are described in \citet{mis07}.
. The spectral resolution is R=37.500. or ~ 7kms7!.," The spectral resolution is $R=37,500$, or $\sim 7$ km $^{-1}$."
 Table 1 summarizes the basic properties of the three absorption systems and of the quasars that host them., Table \ref{tab-sumsys} summarizes the basic properties of the three absorption systems and of the quasars that host them.
 The transitions listed are detected at a 56 confidence level at the NAL redshift., The transitions listed are detected at a $5\sigma$ confidence level at the NAL redshift.
 The sign of the velocity of a system is taken to be positive. if the line ts blueshifted. Le.. if the gas appears to be outflowing relative to the quasar: this is the opposite convention from that adopted in etal. (2007).," The sign of the velocity of a system is taken to be positive, if the line is blueshifted, i.e., if the gas appears to be outflowing relative to the quasar; this is the opposite convention from that adopted in \citet{mis07}."
. In the same table. we also list the 4400 fflux densities (takenfromMisawaetal.2007).. as well as the bolometric luminosities and ionizing photon rates (obtained as described in refmethod)).," In the same table, we also list the 4400 flux densities \citep[taken
from][]{mis07}, as well as the bolometric luminosities and ionizing photon rates (obtained as described in \\ref{method}) )."
 Figures |.. 2.. and 3. present. for each system. absorption profiles. for transitions which are used as. modelling constraints.," Figures \ref{fig-q0130-sys}, \ref{fig-q1009-sys}, and \ref{fig-q1700-sys} present, for each system, absorption profiles for transitions which are used as modelling constraints."
 An example of the best models we have found (discussed in latersections) is also shown in each figure., An example of the best models we have found (discussed in latersections) is also shown in each figure.
 The velocity of an entire system is defined by the optical depth- center of the strongest member of the doublet., The velocity of an entire system is defined by the optical depth-weighted center of the strongest member of the doublet.
We describe each system below.,We describe each system below.
this aremment that if 2006 SQszo is from the Oort Cloud. it is very likely from the inner 104 AU of the Oort Cloud.,"this argument that if 2006 $_{372}$ is from the Oort Cloud, it is very likely from the inner $^4$ AU of the Oort Cloud."
 Comet C/2003 A2 (2?) is the only previously shown LPC with perihchou bevoud Saturn.," Comet C/2003 A2 \citep{gleas03,green03} is the only previously known LPC with perihelion beyond Saturn."
 Because its serihelion is still within 2 AU of Saturn (q=11.1 AU). its dynamics differ significantly from 2006 SQzr2.," Because its perihelion is still within 2 AU of Saturn $q = 11.4$ AU), its dynamics differ significantly from 2006 $_{372}$."
 This object's proximity to the strong energy kicks of Saturn nake it just as likely to come from the outer Oort Cloud as the inner Oort Cloud., This object's proximity to the strong energy kicks of Saturn make it just as likely to come from the outer Oort Cloud as the inner Oort Cloud.
 This is because the timescale or Saturn to modifv the seninajor axis of a comet is uuch shorter than Uranus or Neptune., This is because the timescale for Saturn to modify the semimajor axis of a comet is much shorter than Uranus or Neptune.
 As a result. our areuineut for the inner Oort Cloud origin of 2006 50)” cannot be applied to the orbit of Comet C/2003 À2. aud we cannot coustrain the region of the Oort Cloud where this comet previously resided.," As a result, our argument for the inner Oort Cloud origin of 2006 $_{372}$ cannot be applied to the orbit of Comet C/2003 A2, and we cannot constrain the region of the Oort Cloud where this comet previously resided."
 Civeu that we believe 2006 SQazo is from the inner Oort Cloud. a comparison with Sedna. another purported inner Oort Cloud body (?).. is justified.," Given that we believe 2006 $_{372}$ is from the inner Oort Cloud, a comparison with Sedna, another purported inner Oort Cloud body \citep{brown04}, is justified."
 Sedua is thought to have been scattered ou to its current orbit (q=76 AU. «=ls? AU) by strong external perturbations of the Sus birthplace cuviromment carly in the solar systems history (2??)..," Sedna is thought to have been scattered on to its current orbit $q = 76$ AU, $a = 487$ AU) by strong external perturbations of the Sun's birthplace environment early in the solar system's history \citep{morblev04,bras06,kaibquinn08}."
 Because external forces of this magnitude no longer act in the Suus vicinity. Sedua's orbit essenutiallv ceased evolving Corrs ago.," Because external forces of this magnitude no longer act in the Sun's vicinity, Sedna's orbit essentially ceased evolving Gyrs ago."
 In coutrast. because of a larger original ΠΠ axis. 2006 SQsr> has receutlv had its perilelion pushed closer to Earth by the Calactic tide and passing stars.," In contrast, because of a larger original semimajor axis, 2006 $_{372}$ has recently had its perihelion pushed closer to Earth by the Galactic tide and passing stars."
 Thus. 2006 ο represents a more proximate (q¢2 15 AU) population of inner Oort Cloud bodies whose orbits are still coutinnally evolving.," Thus, 2006 $_{372}$ represents a more proximate $q \gtrsim$ 15 AU) population of inner Oort Cloud bodies whose orbits are still continually evolving."
 Qut of all known TNOs. 2000 OO (?2:: =20.7 AU. a=552 AU. 7= 20°) has an orbitgs most similar o 2006 SQs-2.," Out of all known TNOs, 2000 $_{67}$ \cite{mil02,veil01}; $q = 20.7$ AU, $a = 552$ AU, $i = 20^\circ$ ) has an orbit most similar to 2006 $_{372}$."
 When we use Fieure 2 to determine he probability that it is from the Oort Cloud. we find hat the Oort Cloud should produce orbits similar to 2000 ος at a rate that is at least Ll times that of the scattered disk. assuming the comet population inbers quoted in section 12.," When we use Figure 2 to determine the probability that it is from the Oort Cloud, we find that the Oort Cloud should produce orbits similar to 2000 $_{67}$ at a rate that is at least 14 times that of the scattered disk, assuming the comet population numbers quoted in section 4.2."
 Furthermore. because its perihelion is far bevoud the orbit of Saturn. similar inescale arguments can be used to show that 2000 ος ust also come from the inner —10! AU of the Oort Cloud.," Furthermore, because its perihelion is far beyond the orbit of Saturn, similar timescale arguments can be used to show that 2000 $_{67}$ must also come from the inner $\sim$ $^4$ AU of the Oort Cloud."
 Thus. as suggested in 7.. it appears that 2000 OQOgz is also a comet from the inner Oort Cloud.," Thus, as suggested in \citet{emel05}, it appears that 2000 $_{67}$ is also a comet from the inner Oort Cloud."
 We have discovered an object (2006 SQ3-2) on an unusual orbit with a perihelion of 21.2 AU and a sclnimajor axis of 796 AU., We have discovered an object (2006 $_{372}$ ) on an unusual orbit with a perihelion of 24.2 AU and a semimajor axis of 796 AU.
 Because ofthe ability of the scattered disk and Oort Cloud to produce plauct-crossing orbits in the outer Solar System we believe this object previously resided m one of these reeious., Because of the ability of the scattered disk and Oort Cloud to produce planet-crossing orbits in the outer Solar System we believe this object previously resided in one of these regions.
 To cetermine which scenario is more probable we have simulated the production of similar orbits from: both the scattered disk and the Oort Cloud., To determine which scenario is more probable we have simulated the production of similar orbits from both the scattered disk and the Oort Cloud.
 The results of our simmlations indicate that similar objects are produced from the Oort Cloud anywhere from 2.3 to 1100 times more often than the scattered disk. depending ou the the populations of both reeious.," The results of our simulations indicate that similar objects are produced from the Oort Cloud anywhere from 2.3 to 1100 times more often than the scattered disk, depending on the the populations of both regions."
 We argue that a production rate ratio of 16 or higher is 1uost likely aud that an Oort Cloud origin for 2006 SOsz is therefore most probable., We argue that a production rate ratio of 16 or higher is most likely and that an Oort Cloud origin for 2006 $_{372}$ is therefore most probable.
 Iutrieuiugly. an Oort Cloud origin for 2006 SOsc πράος that is very likely from the inner LO! AU of the Oort Cloud.," Intriguingly, an Oort Cloud origin for 2006 $_{372}$ implies that is very likely from the inner $^4$ AU of the Oort Cloud."
 Otherwise. its perihelion would have beeu removed fou he plauetary region before its ποιαΊου axis could be drawn down to LO? AU.," Otherwise, its perihelion would have been removed from the planetary region before its semimajor axis could be drawn down to $10^3$ AU."
 This line of reasoning can also he applied to show that 2000 OOgr must have had a similar jstorv., This line of reasoning can also be applied to show that 2000 $_{67}$ must have had a similar history.
 An inner Oort Cloud origin makes 2006 SQ3-> and 2000 OOgs unique with respect to all known LPCs., An inner Oort Cloud origin makes 2006 $_{372}$ and 2000 $_{67}$ unique with respect to all known LPCs.
 Furthermore. these bodies are members of an evolving. nore proximate population of iuner Oort Cloud bodies hat are easier to detect because of their αμα] perilelia conrpared to Sedua aud other ucmbers of the extended scattered disk.," Furthermore, these bodies are members of an evolving, more proximate population of inner Oort Cloud bodies that are easier to detect because of their small perihelia compared to Sedna and other members of the extended scattered disk."
" Although we have shown 2006 πο is a probable LPC. this object cannot provide much new information about the structure of the Oort Cloud bv itself,"," Although we have shown 2006 $_{372}$ is a probable LPC, this object cannot provide much new information about the structure of the Oort Cloud by itself."
 The colmpilation of a lavee sample of similar Oort Cloud objects passing through the outer planetary region can. however. place better coustraints ou both the size and structure of the Oort Cloud.," The compilation of a large sample of similar Oort Cloud objects passing through the outer planetary region can, however, place better constraints on both the size and structure of the Oort Cloud."
 The prospects for compiling such a sample in the near future are eucouraeius., The prospects for compiling such a sample in the near future are encouraging.
 2006 SQaco was discovered by imaging the same fields of sky over iultiple uielits separated bw periods of mouths during the SDSS-II SN survey., 2006 $_{372}$ was discovered by imaging the same fields of sky over multiple nights separated by periods of months during the SDSS-II SN survey.
 In terms of both Inuitiug magnitude and skv coverage. this survey is inodest compared to the large synoptic surveys such as Pau-STARRS (?7) and LSST (?) πο planned for the conine decade.," In terms of both limiting magnitude and sky coverage, this survey is modest compared to the large synoptic surveys such as Pan-STARRS \citep{kais02,jew03} and LSST \citep{ivez08} being planned for the coming decade."
 Thus. we can expect iiv objects similar to 2006 SQaco to be discovered in the near future;," Thus, we can expect many objects similar to 2006 $_{372}$ to be discovered in the near future."
 Used In conjuction with dynamical modoeliug. these discoveries will provide mauyv clues about the current size aud structure of the Oort Cloud as well as the dynamical history of the solar svsteni.," Used in conjuction with dynamical modeling, these discoveries will provide many clues about the current size and structure of the Oort Cloud as well as the dynamical history of the solar system."
 We would like to thank the reviewer. Alessandro ALorbidelli. for insightful comments aud sugeestious that ercatlv improved the quality of this work.," We would like to thank the reviewer, Alessandro Morbidelli, for insightful comments and suggestions that greatly improved the quality of this work."
 This research was partially fuuded by a NASA Earth and Space Science. Fellowship., This research was partially funded by a NASA Earth and Space Science Fellowship.
 Most. of our computing work was performed using the Purdue Teragricd coluputing facilities managed with Condor scheduling software (see lttp:/Avwow.cswisc.edu/condor)., Most of our computing work was performed using the Purdue Teragrid computing facilities managed with Condor scheduling software (see http://www.cs.wisc.edu/condor).
 Funding for the SDSS aud SDSS-II has been provided bv the Alfred P. Sloan Foundation. the Participatiug Tustitutions. the National Science Foundation. the U.S. Department of Enerev. the National Acronautics aud Space Acuninistration. the Japanese \loubukagalasho. the Max Planck Society. and the IHigher. Education Funding Council for Enelaud.," Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England."
 The SDSS Web Site is http:/Awww.scdss.ore/ The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions., The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions.
 The Participating Tustitutious are the American Mauseuia of Natural Historv. Astrophysical Institute Potsdam. University of Basel. University of Cambridge. Case Western. Reserve University. Uuiversity of Chicago. Drexel Universitv. Fermilab. the Institute for Advauced Study. the Japan Participation Group. Jolus Hopkins University. the Joint Institute for Nuclear Astrophysics.," The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics,"
Iu section 2 we have pointed out that ZPI is possibly a mixture oPa general ZP structure anda OPP. and the QPP is a very short-period pulsation (WSP).,"In section 2 we have pointed out that ZP1 is possibly a mixture of a general ZP structure and a QPP, and the QPP is a very short-period pulsation (VSP)."
 From the work of Tan. et al (2007) and Tau. et al (2010). we may suppose that the OPP is possibly a result of modulations of tle resistive teariuganuodoe oscillations in tle cureut-curviug flare plasma loops with lieh temperature.," From the work of Tan, et al (2007) and Tan, et al (2010), we may suppose that the QPP is possibly a result of modulations of the resistive tearing-mode oscillations in the current-carrying flare plasma loops with high temperature."
 The pane (9) of Fie.l indicates that teniperature around ZP1 22 MIS) is very close to the maxima of the profile. it is possible that the modulations of the resistive tearine-imocde oscillations take place.," The panel (9) of Fig.1 indicates that temperature around ZP1 22 MK) is very close to the maximum of the profile, it is possible that the modulations of the resistive tearing-mode oscillations take place."
 The eurreut-cairviug ]asina loop can drive the tearing-mode oscillation aud modulate the microwave emission to foxiu VSP., The current-carrying plasma loop can drive the tearing-mode oscillation and modulate the microwave emission to form VSP.
 Ou his VSP backeround. some mechanisi generate the ZP structure.," On this VSP background, some mechanism generate the ZP structure."
 Ou 15 February 2011. there erupts an N2.2 flare event on the solar disk. which was the first N-class flare occurred since the solar evele 21.," On 15 February 2011, there erupts an X2.2 flare event on the solar disk, which was the first X-class flare occurred since the solar cycle 24."
 Associated with this flare eveut. Έπος microwave ZP structures at different frequencies are registered in ciffercut phases of the fare: the first is registered from SBRS/IInairou at frequency of 6:10 ~ 7.00 GITz. which is very unusual at such lieh frequency baud aud in the early rising phase of the flare: the second is also registered from SBRS/ITuairou. at frequency of 2.60 2.73 Cz in the decay phase of the flare. possibly it may extend to the frequency loweY than 2.60 CUz: the third is registered from SBRS/Yunnan at frequency of LOL 1.13 GIIz. in the decay phase after far from thle flare peak.," Associated with this flare event, three microwave ZP structures at different frequencies are registered in different phases of the flare: the first is registered from SBRS/Huairou at frequency of 6.40 $\sim$ 7.00 GHz, which is very unusual at such high frequency band and in the early rising phase of the flare; the second is also registered from SBRS/Huairou, at frequency of 2.60 – 2.73 GHz in the decay phase of the flare, possibly it may extend to the frequency lower than 2.60 GHz; the third is registered from SBRS/Yunnan at frequency of 1.04 – 1.13 GHz, in the decay phase after far from the flare peak."
 By scrutinizing the current prevalent theoretical wwodels of ZP structures Gucluding Derusteim model. whistler wave model. DPR τος]. and the Ledeney model). comparing their estimated maeuetic field streneths in the correspondiug source regions. we find that the DPR model is muc[um more possible for explaining the generation of microwave ZP structures.," By scrutinizing the current prevalent theoretical models of ZP structures (including Bernstein model, whistler wave model, DPR model, and the Ledenev model), comparing their estimated magnetic field strengths in the corresponding source regions, we find that the DPR model is much more possible for explaining the generation of microwave ZP structures."
" It derived the maeuetic field streneths as about 230 315 €. 126. 117 C. and 23. 26 Gin the source regions of ZP1. ZP2. aud ZP2. respectively,"," It derived the magnetic field strengths as about 230 – 345 G, 126 – 147 G, and 23 – 26 G in the source regions of ZP1, ZP2, and ZP3, respectively."
 Comparison with the diagnostics of fiber bursts aud the previous empirical 110dc ‘Lowe sugeest that such estimations are acceptable.," Comparison with the diagnostics of fiber bursts and the previous empirical model, we suggest that such estimations are acceptable."
 It should be noted that DPR model is not self-contained when we adopt it to diagnose the magnetic ficlds in ZP source regions., It should be noted that DPR model is not self-contained when we adopt it to diagnose the magnetic fields in ZP source regions.
 It needs the supplement of the imhomoecucity model of plasima density aie uaenetic feld., It needs the supplement of the inhomogeneity model of plasma density and magnetic field.
 ITowever. it is not easy to eet the exact inhomogencity models.," However, it is not easy to get the exact inhomogeneity models."
 So far. all the existing models (c.g. Dulk McLean model. Takakura model. ete ) of plasina deusity and maeguetic field are propose hat the plasina density and maeuetic field chanee with the height. aud. expressed as functious of height Gr or hj.," So far, all the existing models (e.g. Dulk McLean model, Takakura model, etc ) of plasma density and magnetic field are proposed that the plasma density and magnetic field change with the height, and expressed as functions of height $r$ or $h$ )."
 Such method implies that the inhomogencous scale heights of inagnuetic field is also a function of naenetic field. and this is the origin of the sclfcoutracictious.," Such method implies that the inhomogeneous scale heights of magnetic field is also a function of magnetic field, and this is the origin of the self-contradictions."
 We necd a iore perfect model which can xovide the inhomogencous scale leneths of plasima clensity aud maguctic field. aud they are independent of he maguitude of maenetic field streneth.," We need a more perfect model which can provide the inhomogeneous scale lengths of plasma density and magnetic field, and they are independent of the magnitude of magnetic field strength."
 The magnetic field diaguties of BAL model and WW. inodel are independent of the inhomogeneity models. they secu to be the perfect models to diagnose the magnetic ficl in the source region by ZP structures.," The magnetic field diagnostics of BM model and WW model are independent of the inhomogeneity models, they seem to be the perfect models to diagnose the magnetic field in the source region by ZP structures."
 But they are not likely to agree with the ZP structures observed iu his work., But they are not likely to agree with the ZP structures observed in this work.
 However. we should be noted that either the Dulk MeLean model. or the Talkaluras model is oulv a simplified model.," However, we should be noted that either the Dulk McLean model, or the Takakura's model is only a simplified model."
" The actual regime divine the microwave burst with ZP structures should be extremely dynamic processes, the real maeuetic topology is also uch mere conplex and changeable than the depiction of the above models."," The actual regime during the microwave burst with ZP structures should be extremely dynamic processes, the real magnetic topology is also much more complex and changeable than the depiction of the above models."
 So far. the ouly thing we can do is to obtain the approximated estimations.," So far, the only thing we can do is to obtain the approximated estimations."
 The relative laree difference of the magnetic estimations iun ZP3 source regiou between DPR model aud the Takakura’s model just reflect that its source region should be located at cdiffererit niaguetic loop with differeut distance, The relative large difference of the magnetic estimations in ZP3 source region between DPR model and the Takakura's model just reflect that its source region should be located at different magnetic loop with different distance
(c.g... Noguchi 1988)s,"(e.g., Noguchi 1988)."
patial Therefore. the bars in disk. galaxies can change the distribution of gas such that the distribution can be much more centrally concentrated.," Therefore, the bars in disk galaxies can change the spatial distribution of gas such that the distribution can be much more centrally concentrated."
 hus it could well also be possible that the origin of the proposed truncated: gas disks has something to do with dynamical action of stellar bars in disk galaxies., Thus it could well also be possible that the origin of the proposed truncated gas disks has something to do with dynamical action of stellar bars in disk galaxies.
 An important and. possibly testable prediction of the scenario presented hereis that any emission associated with the star formation in opticallypassive red spirals should be much more compact that that associated with star formation in blue spiral galaxies., An important and possibly testable prediction of the scenario presented here is that any emission associated with the star formation in optically–passive red spirals should be much more compact that that associated with star formation in blue spiral galaxies.
 On the other hand. the star formation rate per unit area (ie. star formation density measured. in units of M. | [or these two cdilferent. tvpes of spirals should not be so different. because the star formation densities within the inner regions of passive spirals are expected be as high as those in blue spirals.," On the other hand, the star formation rate per unit area (i.e., star formation density measured in units of ${\rm M}_{\odot}$ $^{-1}$ $^{-2}$ ) for these two different types of spirals should not be so different, because the star formation densities within the inner regions of passive spirals are expected to be as high as those in blue spirals."
 What is needed to test this is an emission line that traces star formation (and its rate) and is not heavily allected by dust extinction., What is needed to test this is an emission line that traces star formation (and its rate) and is not heavily affected by dust extinction.
 Here. the Hao line probably holds the most promise in terms of being the optical line least allected by dust extinction and one that can be readily niapped spatially and out to high redshifts via high resolution imaging and integral field unit spectroscopy.," Here, the $\alpha$ line probably holds the most promise in terms of being the optical line least affected by dust extinction and one that can be readily mapped spatially and out to high redshifts via high resolution imaging and integral field unit spectroscopy."
 As commented on above. a shortcoming of this study is its inability to show that disk galaxies with centrally-concentrated star formation exhibit passive k-twpe spectra.," As commented on above, a shortcoming of this study is its inability to show that disk galaxies with centrally-concentrated star formation exhibit passive k-type spectra."
 Although previous theoretical studies based. on one-zone models ancl numerical simulations showed. that e(a) and alk/kla spectra can be formed. in. clusty star-forming ealaxies (o... Shiova Bekki 200€: Bekki et al.," Although previous theoretical studies based on one-zone models and numerical simulations showed that e(a) and a+k/k+a spectra can be formed in dusty star-forming galaxies (e.g., Shioya Bekki 2000; Bekki et al."
 2001: Shiova et al., 2001; Shioya et al.
 2002). they did not clearly show k-tvpe spectra can be Formed from clusty star-forming galaxies.," 2002), they did not clearly show k-type spectra can be formed from dusty star-forming galaxies."
 μις it is crucially important that as a next step we conduct further numerical simulations that include spectrophotometric modeling which allow us to predict the spectroscopic signature associated with star formation that proceeds in the inner regions of clisk ealaxies., Thus it is crucially important that as a next step we conduct further numerical simulations that include spectrophotometric modeling which allow us to predict the spectroscopic signature associated with star formation that proceeds in the inner regions of disk galaxies.
 If the star-lorminge regionso of red. passive spirals are as spatially extended as those in blue star-forming spirals. then it will be necessary to consider alternative scenarios to the one presented here.," If the star-forming regions of red passive spirals are as spatially extended as those in blue star-forming spirals, then it will be necessary to consider alternative scenarios to the one presented here."
 One such possibility worthy of brief mention hereis that the upper-mass cutoll Cni) of the initial mass function (IATL) is significantly smaller(e.g.. < 2OAL.) in passive spirals.," One such possibility worthy of brief mention here is that the upper-mass cutoff $m_{\rm upp}$ ) of the initial mass function (IMF) is significantly smaller (e.g., $<20 {\rm M}_{\odot}$ ) in passive spirals."
 In this truncated IME scenario. there are few or no massive O stars that can ionize the ISA (Le... 220M. ). such that ooptical emission lines are very weak. ancl dedust in the ISM can obscure star formation quite efficiently. due to there being little destruction. of dust by ionizing photons.," In this truncated IMF scenario, there are few or no massive O stars that can ionize the ISM (i.e., $>20 {\rm M}_{\odot}$ ), such that optical emission lines are very weak, and dust in the ISM can obscure star formation quite efficiently due to there being little destruction of dust by ionizing photons."
 Εις truncated IME scenario has observational support through being able to explain the UV and Le properties of low surface. brightness. galaxies with low star formation rates (Aleurer et al., This truncated IMF scenario has observational support through being able to explain the UV and $\alpha$ properties of low surface brightness galaxies with low star formation rates (Meurer et al.
 2009)., 2009).
 More quantitative investigation based on numerical simulations of disk galaxy evolution with non-universal LMEs are require to test the viability of this scenario., More quantitative investigation based on numerical simulations of disk galaxy evolution with non-universal IMFs are required to test the viability of this scenario.
 Recent observational studies of distant. galaxies base onSpilzer 245m photometry and. optical imaging by the have revealed the clusty nature of red galaxies and have provided: new clues to the possible eraclual truncation of galactic star formation in dillercn environments (e.g. Callazzi ct al.," Recent observational studies of distant galaxies based on $\mu$ m photometry and optical imaging by the have revealed the dusty nature of red galaxies and have provided new clues to the possible gradual truncation of galactic star formation in different environments (e.g., Gallazzi et al."
 2009: Wolf et al., 2009; Wolf et al.
 2009), 2009).
 The present study suggests that truncation of star formation can occur more dramatically in the outer parts of clisk galaxies. where environmental processes (e.g. tidal and rani pressure stripping) can be more ellective.," The present study suggests that truncation of star formation can occur more dramatically in the outer parts of disk galaxies, where environmental processes (e.g., tidal and ram pressure stripping) can be more effective."
" Lt also sugeests that the possible inner dusty. star-forming regions in passive spirals would be due to 7outside-in truncation of star formation"" in the course of disk galaxy evolution. in particular. in groups and. clusters of ealaxies."," It also suggests that the possible inner dusty star-forming regions in passive spirals would be due to “outside-in truncation of star formation” in the course of disk galaxy evolution, in particular, in groups and clusters of galaxies."
 We are grateful to the referee Christian Wolf lor valuable comments. which contribute to improve the present. paper.," We are grateful to the referee Christian Wolf for valuable comments, which contribute to improve the present paper."
 KB and. WJC all acknowledge the financial support of the Australian Research Council throughout the course of this work., KB and WJC all acknowledge the financial support of the Australian Research Council throughout the course of this work.
 Numerical computations reported. here were carried out both on the GRAPE system at the. University of Western Australia and on those kindly made available by the Center for computational astrophysies (CICA) of the National Astronomical Observatory of Japan., Numerical computations reported here were carried out both on the GRAPE system at the University of Western Australia and on those kindly made available by the Center for computational astrophysics (CfCA) of the National Astronomical Observatory of Japan.
afinite limit. So fheintroduction ofthe cutoff proves,"procedure for the gauge-Higgs sector, the application to the Faddeev-Popov sector is obvious."
 to becrucial: idallows for thecancellation the fundamental ," In the Green' s function method the expression for $\calj(0,\Lambda^2)$, Eq. \ref{seffGFunsub}) )"
svstem A7.(7:22) the ratiocan be madeequal tounity independent summation over waves. Even thoughnow contribution infinite This willbe manifestin," is replaced by _n ^2) ] where of course dr r ^2) The diagonal components of the first order Green's function are given by ^2)+ _ir) and the first order part of the functions $h^{i\pm}_{n,i}$ is obtained by solving Eq. \ref{h1equation}) )"
the numerical results and willbe displaved in section ,"which for the diagonal elements reduces to _i ^2) }(r) The boundary condition is $h^{(1)i\pm}_{n,i}(\infty,\nu^2)\to 0$."
Theim =1 partial wavehas a boundstate alv7 = 0.Asthe partial wave is degenerate withm =—1 thebound state hasa d," There is one important point, however: there cannot be divergences with external gauge legs, i.e. proportional to $a^\mu A_\mu$, as in a gauge theory the is no mass counter term for the gauge field."
ett0.NoiThe123(5.1)[actor.v7Eberemovedbytakingthederivative-hasazeroat77>= computingthe determinant at twosullicientlvsmall values of 72. Sowehave," Indeed these tadpole contributions are compensated by second order terms, in scalar QED in $4$ dimensions this corresponds the cancellation of quadratic divergences proportional to $A^\mu A_\mu$ ."
in the aceretion dise; see above). we would be dealing with intermittency in. the generation of high-energy asymmetric fluctuations (within or close to the jet).,"in the accretion disc; see above), we would be dealing with intermittency in the generation of high-energy asymmetric fluctuations (within or close to the jet)."
 As we commented here above. the structure function. fypo(2 AA)) agrees with the presence of very short-timescale shots. which are not detected from the new LQLM I records.," As we commented here above, the structure function $f_{APO}$ ) agrees with the presence of very short-timescale shots, which are not detected from the new LQLM I records."
 These fluctuations may be caused by some kind of observational (systematic) noise or alternatively. they might be related to an episode of very short-timescale activity inside the disc. the Jet or other region.," These fluctuations may be caused by some kind of observational (systematic) noise or alternatively, they might be related to an episode of very short-timescale activity inside the disc, the jet or other region."
 In Fig. 4..," In Fig. \ref{sffits},"
" the adopted solution (STF+STF model with parameters wy) = 0.27. r, = I] d. and r» = 105 d) fits FApo(2100 AA)) reasonably well."," the adopted solution (STF+STF model with parameters $w_1$ = 0.27, $\tau_1$ = 11 d, and $\tau_2$ = 105 d) fits $f_{APO}$ ) reasonably well."
 We present a novel and rigorous analysis of the structure function of the UV variability of the gravitationally lensed quasarQ0957+561., We present a novel and rigorous analysis of the structure function of the UV variability of the gravitationally lensed quasar.
. New Liverpool telescope data (2005-2007seasons:Shalyapinetal.2008) allow us to construct normalized structure functions of the quasar luminosity at two restframe wavelengths: 4|~ 2100 ((g band data) and vt~ 2600 (r band data).," New Liverpool telescope data \citep[2005$-$2007 
seasons;][]{Sha08} allow us to construct normalized structure functions of the quasar luminosity at two restframe wavelengths: $\lambda \sim$ 2100 $g$ band data) and $\lambda \sim$ 2600 $r$ band data)."
 Old Apache Point Observatory records in the £ band (1995-1996seasons;Kundiéetal.1997) are also used to check the possible time evolution of the variability at ~ 2100Α.," Old Apache Point Observatory records in the $g$ band \citep[1995$-$1996 
seasons;][]{Kun97} are also used to check the possible time evolution of the variability at $\sim$ 2100."
. The observed shapes of the structure functions are compared to predictions of a large set of Poissonian models., The observed shapes of the structure functions are compared to predictions of a large set of Poissonian models.
 This set of models includes the simplest and well-known ones. consisting of only one variable component (CidFernandesetal.2000.andreferencestherein ).. as well as hybrid models incorporating two independent variable components.," This set of models includes the simplest and well-known ones, consisting of only one variable component \citep[][and references therein]{Cid00}, as well as hybrid models incorporating two independent variable components."
 Severalhybrid (or level-2) models are able to account for both Liverpool telescope structure functions (see Table 1))., Severalhybrid (or level-2) models are able to account for both Liverpool telescope structure functions (see Table \ref{fits}) ).
 Some of them contain flares with unrealistic profile (square flares) and lead to solutions that are difficult to interpret., Some of them contain flares with unrealistic profile (square flares) and lead to solutions that are difficult to interpret.
 Fortunately. we also find reasonable solutions in. which ~ 100-d time-symmetric and ~ 170-d. time-asymmetric flares are produced at both restframe wavelengths.," Fortunately, we also find reasonable solutions in which $\sim$ 100-d time-symmetric and $\sim$ 170-d time-asymmetric flares are produced at both restframe wavelengths."
 Exponentially decaying and starburst flares (and very. probably other time-asymmetric shots) work in a similar way., Exponentially decaying and starburst flares (and very probably other time-asymmetric shots) work in a similar way.
 Therefore. the good behaviour of the starburst ingredient does not necessarily implies the existence of supernova explosions. but the production of highly asymmetric shots.," Therefore, the good behaviour of the starburst ingredient does not necessarily implies the existence of supernova explosions, but the production of highly asymmetric shots."
 What about the mechanism of intrinsic variability?., What about the mechanism of intrinsic variability?.
 The old and very recent gr light curves of led to time delays between quasar components and between optical bands that mainly support a reverberation scenario (Collier.2001:Shalyapinet 2008).," The old and very recent $gr$ light curves of led to time delays between quasar components and between optical bands that mainly support a reverberation scenario \citep{Col01b,Sha08}."
. Thus. reverberation would be the main mechanism of variability.," Thus, reverberation would be the main mechanism of variability."
 The presence of an EUV/radiojet (e.g..Garrettal.1994;Hutchings2003) and a bright X-ray source (Chartas2000) also suggests the viability of this mechanism: two types of EUV/X-ray fluctuations that are generated within or close to the Jet. and later reprocessed by two rings of the disc (each ring corresponds to a different restframe wavelength).," The presence of an EUV/radio jet \citep[e.g.,][]{Gar94,Hut03} and a bright X-ray source \citep{Cha00} also suggests the viability of this mechanism: two types of EUV/X-ray fluctuations that are generated within or close to the jet, and later reprocessed by two rings of the disc (each ring corresponds to a different restframe wavelength)."
 On the other hand. one can also justify both kinds of flare profile.," On the other hand, one can also justify both kinds of flare profile."
 For example. the cellular-automaton model produces asymmetric shots (e.g.Kawaguchietal.1998).. and hydrodynamical simulations lead to symmetric flares (Manmotoetal.1996).," For example, the cellular-automaton model produces asymmetric shots \citep[e.g.,][]{Kaw98}, and hydrodynamical simulations lead to symmetric flares \citep{Man96}."
. The - 100-d time-symmetric shots seem to be also responsible for most of the ~ 2100 vvariability detected in. the Apache Point Observatory experiment. but there is no evidence of asymmetric shots in the old UV variability.," The $\sim$ 100-d time-symmetric shots seem to be also responsible for most of the $\sim$ 2100 variability detected in the Apache Point Observatory experiment, but there is no evidence of asymmetric shots in the old UV variability."
 This absence of asymmetric flares may be due to gaps in the light curves. a relatively short monitoring period. etc.," This absence of asymmetric flares may be due to gaps in the light curves, a relatively short monitoring period, etc."
 Alternatively. it could means an evolution of the central engine. L.e.. intermittent production of high-energy asymmetric. fluctuations.," Alternatively, it could means an evolution of the central engine, i.e., intermittent production of high-energy asymmetric fluctuations."
 The Apache Point Observatory structure function is also consistent with the presence of very short-lifetime (~ 10 d) symmetric flares., The Apache Point Observatory structure function is also consistent with the presence of very short-lifetime $\sim$ 10 d) symmetric flares.
 This kind of flare might be caused by observational systematic noise. or perhaps. represent additional evidence for time evolution.," This kind of flare might be caused by observational systematic noise, or perhaps, represent additional evidence for time evolution."
 Our results do not support a previous claim for the possible starburst origin of some events in the old. g- band light curves (Ullánetal.2003)., Our results do not support a previous claim for the possible starburst origin of some events in the old $g$ -band light curves \citep{Ull03}.
. Despite the presence of two twin events (one in each quasar component) with an anomalous delay (Goicoechea2002).. the associated. shot probably occurred at the base of the Jet or in the circumnuclear region. but it was not originated by a supernova explosion.," Despite the presence of two twin events (one in each quasar component) with an anomalous delay \citep{Goi02}, the associated shot probably occurred at the base of the jet or in the circumnuclear region, but it was not originated by a supernova explosion."
 Very recently. several studies have showed evidence that optical/UV variability of quasars on restframe timescales > 100 d is mainly driven by variations in accretion rate (e.g..hiteetal. 2008).," Very recently, several studies have showed evidence that optical/UV variability of quasars on restframe timescales $>$ 100 d is mainly driven by variations in accretion rate \citep[e.g.,][]{Wol07,Are08,LiC08,Wil08}."
. Here we are discussing the UV variability of at restframe lags < 100 d. is a very bright and massive object. and this population could not be studied by Wilhiteetal.(2008).," Here we are discussing the UV variability of at restframe lags $\leq$ 100 d. is a very bright and massive object, and this population could not be studied by \citet{Wil08}."
. The first lensed quasar has also an EUV/radio jet (and important X-ray activity: see above). so high-energy variations in the surroundings of the disc axis and their reverberation are possible.," The first lensed quasar has also an EUV/radio jet (and important X-ray activity; see above), so high-energy variations in the surroundings of the disc axis and their reverberation are possible."
 For example. on timescales below 100 days. the optical variations of the local Seyfert galaxy are related to its X-ray variations (Czernyetal.1999).," For example, on timescales below 100 days, the optical variations of the local Seyfert galaxy are related to its X-ray variations \citep{Cze99}."
 Hence. part of the optical variability of this AGN (timescales « 100 days) could be explained by X-ray reprocessing.," Hence, part of the optical variability of this AGN (timescales $<$ 100 days) could be explained by X-ray reprocessing."
 Chromatic delays for the local Seyfert galaxy also reveal a reverberation scenario (Collieretal. 1999).., Chromatic delays for the local Seyfert galaxy also reveal a reverberation scenario \citep{Col99}. .
 Arévaloetal.(2008) reported on an illustrative example of mixed variability., \citet{Are08} reported on an illustrative example of mixed variability.
 The local quasar has been monitored simultaneously in X-rays and optical bands., The local quasar has been monitored simultaneously in X-rays and optical bands.
 All spectral regions were significantly variable. and the fluctuations were clearly correlated.," All spectral regions were significantly variable, and the fluctuations were clearly correlated."
 Arévaloetal.(2008) indicated that pure reprocessing of X-rays cannot account for both ~ 100-d and ~ 500-d timescale optical variability., \citet{Are08} indicated that pure reprocessing of X-rays cannot account for both $\sim$ 100-d and $\sim$ 500-d timescale optical variability.
 They claimed that two distinet mechanisms produce the variability: aceretion rate variations plus reverberation. and the shortest timescale optical events are due to reverberation.," They claimed that two distinct mechanisms produce the variability: accretion rate variations plus reverberation, and the shortest timescale optical events are due to reverberation."
 For an irradiated disc. in general. we obtain a temperature profile shallower than the standard one Tx77?(Shakura& 1973)..," For an irradiated disc, in general, we obtain a temperature profile shallower than the standard one $T \propto r^{-3/4}$\citep{Sha73}. ."
 The reverberation hypothesis assumes that the optical/UV dise regions are irradiated by EUV/X-ray photons from the vicinity of the disc axis., The reverberation hypothesis assumes that the optical/UV disc regions are irradiated by EUV/X-ray photons from the vicinity of the disc axis.
 If the high-energy source is, If the high-energy source is
at 319 MIIz. with their RAZs indicated.,"at 349 MHz, with their $RM$ s indicated."
 The diameter of the zxaubol aud its shading indicate the value of RAL., The diameter of the symbol and its shading indicate the value of $RM$.
 Sources 1l aud 2 could have a sieuificant contzibution of iustrumental polarization. because they are observed im ouly oue pointing center aud are located about tto aaway frou he potting center.," Sources 1 and 2 could have a significant contribution of instrumental polarization, because they are observed in only one pointing center and are located about to away from the pointing center."
 Nevertheless. they show lineara of{A2)-relatious. which indicates that iustrunoental polarization is not imiportant.," Nevertheless, they show linear $\phi(\lambda^2)$ -relations, which indicates that instrumental polarization is not important."
 These sources are therefore iucluded in the analysis., These sources are therefore included in the analysis.
 The stroug correlation of RAL across the sky indicates a Galactic component to the RÀAZs of the extragalactic point sources., The strong correlation of $RM$ across the sky indicates a Galactic component to the $RM$ s of the extragalactic point sources.
 The best-fitting lincar gradieut to the As of the sources has a stecpes slope of 3.62 pper degree in position angle ie. roughly in the direction of increasing Calactic latitude (position augle 667)).," The best-fitting linear gradient to the $RM$ s of the sources has a steepest slope of 3.62 per degree in position angle, i.e. roughly in the direction of increasing Galactic latitude (position angle )."
 The staudux deviation of ΠΛ: around this eradient is 053;293.672, The standard deviation of $RM$ s around this gradient is $\sigma_{RM} \approx 3.6$.
 This is consistent with the result of Lealw C987). who finds a typical internal’ RAF coutribution from the extragalactic source or a halo around it of ~ 5?.," This is consistent with the result of Leahy (1987), who finds a typical 'internal' $RM$ contribution from the extragalactic source or a halo around it of $\sim$ 5."
. The change of sign of RAL of the extragalactic sources within the observed region indicates a reversal iu the Galactic magnetic field. (, The change of sign of $RM$ of the extragalactic sources within the observed region indicates a reversal in the Galactic magnetic field. (
Unlike the diffuse emission. sigu changes in RASS of extragalactic sources cannot be due to depolarization effects.),"Unlike the diffuse emission, sign changes in $RM$ s of extragalactic sources cannot be due to depolarization effects.)"
 The reversal exists on scales larger than our field of view. as is shown in Fie. 1ll..," The reversal exists on scales larger than our field of view, as is shown in Fig. \ref{f4:snk},"
 where we combine the RASS of our sources with those from the literature., where we combine the $RM$ s of our sources with those from the literature.
 The circles are our sources. the squares indicate sources that were detected by Simard-Normandin et ((1981) and/or by Tahara Inoue (1980). ancl the oue triangle indicates the ouly pulsar ucarby (ILuuiltou Lyne 1987).," The circles are our sources, the squares indicate sources that were detected by Simard-Normandin et (1981) and/or by Tabara Inoue (1980), and the one triangle indicates the only pulsar nearby (Hamilton Lyne 1987)."
 The umuubers next to the squares and triangle eive the magnitude of RAL of the source (the diameter of the viubol is only proportional to RAL up to RAL — 15 73)., The numbers next to the squares and triangle give the magnitude of $RM$ of the source (the diameter of the symbol is only proportional to $RM$ up to $RM$ = 15 ).
 One source at (0.80)=(91.257.[87) was onutted. because Simard-Normancdin et seive RAL —30b aand Tabara Inoue give RAL =61.9ον," One source at $(\alpha,\delta) = (91.25\dg,
48\dg)$ was omitted, because Simard-Normandin et give $RM$ $=
34$ and Tabara Inoue give $RM$ $=-64.9$."
", The source distribution shows a clear maguetic Ποιά reversal. which is uot ou Galactic scale but on smaller scales. as is shown in Figs."," The source distribution shows a clear magnetic field reversal, which is not on Galactic scale but on smaller scales, as is shown in Figs."
 1 and 2 of Simard-Normancdin Wroubere (1980). which show RASS of extragalactic sources over the whole sky.," 1 and 2 of Simard-Normandin Kronberg (1980), which show $RM$ s of extragalactic sources over the whole sky."
 The eradieut in RAL ucasured in the extragalactic »)ut sources is completely unrelated ando aliost rpenudieular to the structure in RAL frou he diffusechussion., The gradient in $RM$ measured in the extragalactic point sources is completely unrelated and almost perpendicular to the structure in $RM$ from the diffuse emission.
 This cau be explained from the widely differeut oath tha Ieneths enission extragalactic sources and diffuse probe., This can be explained from the widely different path lengths that extragalactic sources and diffuse emission probe.
 Diffuse emission cau only be observe out oa distance of a few hundred pc to a kpc. as radiation roni further away will be mostly depolarized.," Diffuse emission can only be observed out to a distance of a few hundred pc to a kpc, as radiation from further away will be mostly depolarized."
 On the contrary. extragalactic sources are Faraday-rotated over he complete path leneth through the Milky Way of many," On the contrary, extragalactic sources are Faraday-rotated over the complete path length through the Milky Way of many"
 [Fe/II]<2.5 [HES| [Fe/II]<—2.5 [FeT<3 (Beersetal.2005)... (I&oiivactal.2010.hereafterPaperD..," $\feoh<-2.5$ $\feoh\leq -2.5$ $\feoh < -3$ \citep{Beers05}. \citep[][hereafter Paper I]{Komiya10},"
 |Fe/II|]=2.5 (Ikonmüvactal.2009a).. [Fe/T]]<—2.5 ~10M. |Fe/TI]<2.5 [Fe/II]«—5 hyper iietal-poor (ITMDP) stars.," $\feoh\lesssim-2.5$ \citep{Komiya07,Komiya09}, $\feoh\leq -2.5$ $\sim 10 \msun$ $\feoh\leq -2.5$ $\feoh<-5$ hyper metal-poor (HMP) stars."
 2 IIMP stars detected iu the Milkv Wax halo are the most metal deficicut objects observed vet (ITEI327-2326. Frebel et 22005: HTEOT07-5210. Clhiristlieb et 22002).," 2 HMP stars detected in the Milky Way halo are the most metal deficient objects observed yet (HE1327-2326, Frebel et 2005: HE0107-5240, Christlieb et 2002)."
 Formation cuviromment of the EXIP population stars are thought to differ from present galaxies., Formation environment of the EMP population stars are thought to differ from present galaxies.
 EXIP stars are formed iu the process of galaxy formation in the carly universe., EMP stars are formed in the process of galaxy formation in the early universe.
 In the A cold dark matter (C'DAL) universe. aree galaxies like the Milkv Wavy are formed through nerecr of smaller galaxies as building blocks.," In the $\Lambda$ cold dark matter (CDM) universe, large galaxies like the Milky Way are formed through merger of smaller galaxies as building blocks."
 According o the licrarclical structure formation scenario. a stellar ido is an aeerceation of stars formed iu the niu snall galaxies (Searle&ZinnLOTS:Holiii2008).," According to the hierarchical structure formation scenario, a stellar halo is an aggregation of stars formed in the many small galaxies \citep{Searle78, Helmi08}."
. Earlier heoretical studies show that the first stars are formed. in very low mass halos with ~LO°AL.. (Teemarkctal.1997:Nishi&Susa1999). and host galaxies of the second gcucration of stars are also sinall (Ricottictal.2002:Wise&Abel 2008).," Earlier theoretical studies show that the first stars are formed in very low mass halos with $\sim 10^6\msun$ \citep{Tegmark97, Nishi99} and host galaxies of the second generation of stars are also small \citep{Ricotti02, Wise08}."
. Chemical abundances of these building blocks cau differ from cach other., Chemical abundances of these building blocks can differ from each other.
 Uulike meta rich stars. metals in EXIP stars are svuthesized by oulv oue or a few precursorv SN(e).," Unlike metal rich stars, metals in EMP stars are synthesized by only one or a few precursory SN(e)."
 Element abundances of the EXP stars cau reflect individual characteristics of the precursory SN(e) and their host galaxies., Element abundances of the EMP stars can reflect individual characteristics of the precursory SN(e) and their host galaxies.
 À soiui-iwalytic hierarchical approach can provide a framework witlin which to study the earliest phases of the chemica evolution aud the formation historv of the EMP stars., A semi-analytic hierarchical approach can provide a framework within which to study the earliest phases of the chemical evolution and the formation history of the EMP stars.
 One important point at issue for the earliest phases of chemical evolution is a possible difference iu the IME of EMP stars (c.g.Abiaetal.2001:IKomiva2007).," One important point at issue for the earliest phases of chemical evolution is a possible difference in the IMF of EMP stars \citep[e.g.][]{Abia01, Komiya07}."
. Theoretically. typical mass of stars is large in extremely mietal-poor cuvironment and/or in the carly universe.," Theoretically, typical mass of stars is large in extremely metal-poor environment and/or in the early universe."
 NMuevicalsimmlatiounsshow thatPopulationIII stars withoutmetalare very massive (ee.Brounctal.1999: 2002)..," Numericalsimulationsshow thatPopulationIII stars withoutmetalare very massive \citep[e.g.][]{Bromm99, Abel02}. ."
 For EXIP stars with a little metal. the," For EMP stars with a little metal, the"
Another possible argument for a slowly changing BI mass with time is the [act that the velocity widths. κοἳmU? of the broad lines ave similar for high luminosity QSOs at high z as they are for low-Iuminosity Sevfert Esa ow ον,"Another possible argument for a slowly changing BH mass with time is the fact that the velocity widths, $<~v^2>^{1/2}$, of the broad lines are similar for high luminosity QSOs at high $z$ as they are for low-luminosity Seyfert I's at low $z$."
 Now in the standard. view where QSOs radiate a approximately a fixed fraction of the IEddington luminosity. his is just a coincidence where the increase in DII mass w a factor of z30 at high z is approximately offse » a comparable increase in the BL radius.," Now in the standard view where QSOs radiate at approximately a fixed fraction of the Eddington luminosity, this is just a coincidence where the increase in BH mass by a factor of $\approx30$ at high $z$ is approximately offset by a comparable increase in the BLR radius."
 Since the BLE. radius-Iuminosity relation takes the form krzzL77 (Petersonetal.2004:Ixaspi2005:Bentz2009).. if his relation is unevolving with redshift then <(7ExLfrxLUS and so «ceTetoLOR17 perhaps slow enough o explain the similarity in line widths.," Since the BLR radius-luminosity relation takes the form $r \approx L^{0.5-0.7}$ \citep{bmp04,kaspi05,bentz09}, if this relation is unevolving with redshift then $<v^2> \propto L/r \propto L^{0.5-0.3}$ and so $<v^2>^{1/2} \propto L^{0.25-0.15}$, perhaps slow enough to explain the similarity in line widths."
 In this fixed fraction of Lear case. there is therefore some empirical evidence to icl explain this coincidence.," In this fixed fraction of $L_{Edd}$ case, there is therefore some empirical evidence to help explain this coincidence."
 Nevertheless. the explanation is stillp vulnerable if the r.—L relation actually proves to evolve with recshift.," Nevertheless, the explanation is still vulnerable if the $r-L$ relation actually proves to evolve with redshift."
 Alternatively. to maintain the single DII mass model. we could pursue the idea motivated by the earlv-tvpe galaxies that nothing evolves except the QSO Iuminosity.," Alternatively, to maintain the single BH mass model, we could pursue the idea motivated by the early-type galaxies that nothing evolves except the QSO luminosity."
" On this view. the BL1t radius might then remain constant with luminosity and hence redshilt ie reLU""."," On this view, the BLR radius might then remain constant with luminosity and hence redshift ie $r\approx L^0$."
 A fixed broad-line velocity width. <r2L7. and radius. r. gives a fixed. DL mass with luminosity/redshift.," A fixed broad-line velocity width, $<v^2>^{1/2}$, and radius, $r$, gives a fixed BH mass with luminosity/redshift."
 Phe question then is whether it is feasible. for example. that the factor of z30 increase in L over 0«z2.2 could leave the BLB radius unchanged if à balance has to be maintained between photolonisation and eravity and. this seems physically unlikely.," The question then is whether it is feasible, for example, that the factor of $\approx30$ increase in $L^*$ over $0<z<2.2$ could leave the BLR radius unchanged if a balance has to be maintained between photoionisation and gravity and this seems physically unlikely."
 Lt also seems at οὐος with the above evidence for the rL (and Mea £) correlation found from reverberation mapping., It also seems at odds with the above evidence for the $r-L$ (and $M_{BH}-L$ ) correlation found from reverberation mapping.
 Thus if these empirical results are correct then the constancy of QSO velocity width with Lane z must remain simply a coincidence. rather than providing further evidence for the single DII mass mocel.," Thus if these empirical results are correct then the constancy of QSO velocity width with $L$ and $z$ must remain simply a coincidence, rather than providing further evidence for the single BH mass model."
 Indeed. any single BI mass/flickering mocel motivated by OSO clustering clearly has a severe problem with the QSO Alen L relation found from reverberation mapping at low redshift.," Indeed, any single BH mass/flickering model motivated by QSO clustering clearly has a severe problem with the QSO $M_{BH}-L$ relation found from reverberation mapping at low redshift."
 Equa (9) of Petersonetal.(2004). rejects no correlation ie a slope of zero in the ο—L relation at the S.S. level!, Equn (9) of \citet{bmp04} rejects no correlation ie a slope of zero in the $M_{BH}-L$ relation at the $8.8\sigma$ level!
 So here we have potentially the most serious and direct conlliet with any simple single black hole mass interpretation of the QSO clustering results., So here we have potentially the most serious and direct conflict with any simple single black hole mass interpretation of the QSO clustering results.
 As Petersonetal.(2004) point out. the simple fixed fraction of IEddington model fits the reverberation data very well but this moclel most clearly contradicts the clustering results.," As \citet{bmp04} point out, the simple fixed fraction of Eddington model fits the reverberation data very well but this model most clearly contradicts the clustering results."
 H£ we are to reconcile the clustering cata with the reverberation data. we must conclude that there must be some decoupling of Alex and Adpary and this immeciately takes away the evidence for the Uiekering and any other single BIL mass moclel.," If we are to reconcile the clustering data with the reverberation data, we must conclude that there must be some decoupling of $M_{BH}$ and $M_{DMH}$ and this immediately takes away the evidence for the flickering and any other single BH mass model."
 Aleanwhile. we postpone our discussion of reverberation mapping and the PLE model to Section S and Fig.," Meanwhile, we postpone our discussion of reverberation mapping and the PLE model to Section 8 and Fig."
 4. where we shall directly compare BLE masses predicted by the PLE model to masses derived from reverberation mapping., \ref{fig:ple} where we shall directly compare BH masses predicted by the PLE model to masses derived from reverberation mapping.
 The space number censity of QSOs compared to. the numbers of remnant nuclear black holes in. early-tvpe galaxies has been used to argue for short (7=10' vr) QSO lifetimes Richstoneetal.(1998)., The space number density of QSOs compared to the numbers of remnant nuclear black holes in early-type galaxies has been used to argue for short $\tau\approx10^{7}$ yr) QSO lifetimes \cite{richstone}.
. However. if the QSOs were restricted to a small range of host galaxy/halo/BL masses as implied by the similarity of the QSO clustering results then this argument would start to push QSO Lifetimes higher.," However, if the QSOs were restricted to a small range of host galaxy/halo/BH masses as implied by the similarity of the QSO clustering results then this argument would start to push QSO lifetimes higher."
 Generally. these estimates depend on an assumption that QSO luminosity is proportional to host BIL mass ancl this assumption is also challenged by the lack of a luminosity dependence of QSO clustering.," Generally, these estimates depend on an assumption that QSO luminosity is proportional to host BH mass and this assumption is also challenged by the lack of a luminosity dependence of QSO clustering."
 Richstonectal.(1998) suggest that the bright QSO space density at zg2 is O.1% of the galaxy density at 2=0., \citet{richstone} suggest that the bright QSO space density at $z\approx2$ is $\approx0.1$ of the galaxy density at $z=0$.
 Thus. for example. if the QSOs are restricted to zz1% of galaxies then this lifetime would increase to zzLO’ vrs.," Thus, for example, if the QSOs are restricted to $\approx1$ of galaxies then this lifetime would increase to $\approx10^9$ yrs."
 Clearly the more restricted the galaxy. luminosity range. the longer these duty evele arguments allow the QSO lifetime to be.," Clearly the more restricted the galaxy luminosity range, the longer these duty cycle arguments allow the QSO lifetime to be."
 Alore detailed. arguments for short QSO Lifetimes have been similarly made on the basis of high peaks models of biased QSO clustering., More detailed arguments for short QSO lifetimes have been similarly made on the basis of high peaks models of biased QSO clustering.
 Martini&Weinbere(2001) made predictions for the QSO correlation function. based on the halo number densities inferred from OSO lifetimes estimated from QSO ancl model halo densities., \citet{martini} made predictions for the QSO correlation function based on the halo number densities inferred from QSO lifetimes estimated from QSO and model halo densities.
 Using the ACDAL results from these authors and the 2QZ clustering amplitude. Croometal.(2005). derived a QSO lifetime of 4-50 Myr at z2.," Using the $\Lambda$ CDM results from these authors and the 2QZ clustering amplitude, \citet{croom05} derived a QSO lifetime of 4-50 Myr at $z\approx2$."
 We note that a monotonic increase of halo mass with QSO luminosity is an implicit. assumption here. otherwise it is impossible to establish a minimum halo mass at a fixed QSO Luminosity limit.," We note that a monotonic increase of halo mass with QSO luminosity is an implicit assumption here, otherwise it is impossible to establish a minimum halo mass at a fixed QSO luminosity limit."
 In. what follows we shall also make this assumption., In what follows we shall also make this assumption.
 Hf we take the change in 5 as z9% [rom 28LACQ (ry=5.7h Mpe) to SDSS QSOs (ru.=63h !Mpe) then from the approximate relation. boxANS. fitted to Pig.," If we take the change in $b$ as $\approx9$ from 2SLAQ $r_0=5.7$ $^{-1}$ Mpc) to SDSS QSOs $r_0=6.3$ $^{-1}$ Mpc) then from the approximate relation, $b\propto M_{min}^{0.2}$, fitted to Fig."
 6 of Martini&Weinberg(2001).. his corresponds to an z1.5 οσο halo mass for the ACDAL model.," 6 of \citet{martini}, this corresponds to an $\approx1.5\times$ bigger halo mass for the $\Lambda$ CDM model."
" Now assuming an approximate CDM halo mass function with N(cAM)xMUT for M«I0HhTAL. (Jenkinsetal.2001). we find that the number of haloes in he23.107h!M. mass range is 104 of the total with Al2.10""h4AL..", Now assuming an approximate $\Lambda$ CDM halo mass function with $N(<M)\propto M^{0.15}$ for $M<10^{14}\rm{h}^{-1}M_\odot$ \citep{jenkins01} we find that the number of haloes in the $2-3\times10^{12}\rm{h}^{-1}M_\odot$ mass range is $\approx10$ of the total with $M>2\times10^{12}\rm{h}^{-1}M_\odot$.
 This means that the estimates of QSO ifetime would now lie in the range z40500 Myr., This means that the estimates of QSO lifetime would now lie in the range $\approx40-500$ Myr.
 Since he clustering of the 28LAQ. 20Z and SDSS samples are statistically consistent. the increase in QSO Lifetime could even be larger than this estimate.," Since the clustering of the 2SLAQ, 2QZ and SDSS samples are statistically consistent, the increase in QSO lifetime could even be larger than this estimate."
 This conclusion [rom the clustering is stronger than from the direct. QSO imagine results., This conclusion from the clustering is stronger than from the direct QSO imaging results.
" Llere there is an 5, 2mae range in host luminosity or a [actor of & 6. corresponding to zz404 of the haloes with M2LOMAL. giving an zz2.3. increase in QSO lifetime."," Here there is an $\approx2$ mag range in host luminosity or a factor of $\approx6$ , corresponding to $\approx40$ of the haloes with $M>2\times10^{12}M_\odot$ giving an $\approx2-3\times$ increase in QSO lifetime."
 Again. since the range for the host luminosities is clearly an upper limit due to no account being taken of galaxy magnitude errors. there is no necessary disagreement with the clustering results and indeed the QSO lifetimes could. be much higher if as expected the galaxy magnitucle errors dominate the = 2mag scatter.," Again, since the range for the host luminosities is clearly an upper limit due to no account being taken of galaxy magnitude errors, there is no necessary disagreement with the clustering results and indeed the QSO lifetimes could be much higher if as expected the galaxy magnitude errors dominate the $\approx2$ mag scatter."
 We therefore suggest that the short QSO lifetimes previously derived on the basis of the Martini&Weinberg arguments are at best lower limits if the QSOs are restricted to a range of halo masses and even on the basis of the present clustering data could imply that this lower limit is r2LO ve.," We therefore suggest that the short QSO lifetimes previously derived on the basis of the \citet{martini}
 arguments are at best lower limits if the QSOs are restricted to a range of halo masses and even on the basis of the present clustering data could imply that this lower limit is $\tau>10^8$ yr."
 IE QSO lifetimesbecome longer then there is a clear prediction for QSO clustering vie the long-lived model of, If QSO lifetimesbecome longer then there is a clear prediction for QSO clustering via the long-lived model of
the observational tests we have described will serve to distinguish them from ofl-axis orphliaus.,the observational tests we have described will serve to distinguish them from off-axis orphans.
 Conusion between conventioual GRBs aud dirty. fireball populations remaius possible., Confusion between conventional GRBs and dirty fireball populations remains possible.
 To remove this confusiou. the expected ‘ate of conveutional afterglows iu an orphan survey could be calculated using the known 5-r ay alxl afterglow properties of conventional GRBs.," To remove this confusion, the expected rate of conventional afterglows in an orphan survey could be calculated using the known $\gamma$ -ray and afterglow properties of conventional GRBs."
 An accurate correction might be difficult at preseαπ. because the ratio of afterglow to 5-ray flux shows a wide dispersion and because the data óncu ‘rently available samples are very lhieterogeueous.," An accurate correction might be difficult at present, because the ratio of afterglow to $\gamma$ -ray flux shows a wide dispersion and because the data in currently available samples are very heterogeneous."
 Fortunately. more horjogeneous aud statistically tractable samples should become available with the launch of the Swift mission.," Fortunately, more homogeneous and statistically tractable samples should become available with the launch of the Swift mission."
 Swill offers another potential way to mitigate coufusion by conventional GRBs: Orphan searches could be conducted preferentially in regions monitored by the Swift Burst Alert Telescope (BAT)., Swift offers another potential way to mitigate confusion by conventional GRBs: Orphan searches could be conducted preferentially in regions monitored by the Swift Burst Alert Telescope (BAT).
 With a set of  10-20 Ποια» eveuly spaced ou the sky. at least oue would always [all within the 2 steradian field of the BAT.," With a set of $\sim 10$ $20$ fields evenly spaced on the sky, at least one would always fall within the 2 steradian field of the BAT."
 TIe BAT duty eycle at any point on the sky is limited by Earth occultation to Z5 50%. with eac1 continuous observation <50 minutes.," The BAT duty cycle at any point on the sky is limited by Earth occultation to $\la 50\%$ , with each continuous observation $\la 50$ minutes."
 For each orphan search field. the nuuber of afterglows with Swift counterparts could be counted. aud tje total utunber of conventional alterglows in the sauple could then be inferred by a simple correction for the Swit duty cycle.," For each orphan search field, the number of afterglows with Swift counterparts could be counted, and the total number of conventional afterglows in the sample could then be inferred by a simple correction for the Swift duty cycle."
 This correction coul ye fairly precisely derived using the kuown Swi1 poiuting history. and would a factor between 2 at| 10.," This correction could be fairly precisely derived using the known Swift pointing history, and would a factor between $2$ and $10$."
 On-axls orpliau searches coLd be performed with a much wider field. brigtter Πιx limit. and faster observing cadence than the off-axis searches we have primarily considered hee (Nakar Pirau 2002).," On-axis orphan searches could be performed with a much wider field, brighter flux limit, and faster observing cadence than the off-axis searches we have primarily considered here (Nakar Piran 2002)."
 Such surveys are 1alural successors LO curreit robotic teephoto lens projects (e.g.. ]|xehoe et al 2002).," Such surveys are natural successors to current robotic telephoto lens projects (e.g., Kehoe et al 2002)."
 Iu this limit. it would be possible to ionior fields in tle BAT fielcl ol view. aud thereby know wheher any particular optica flasl Was Or was Hot accompanied by a GRB.," In this limit, it would be possible to monitor fields in the BAT field of view, and thereby know whether any particular optical flash was or was not accompanied by a GRB."
" luplementiug this strateey would require switchii fields appxinately twice per Swill orbit (assuming Swift changes argets to avoid Earth oc""ultatiou). aud would require real-time kuowledege of the Swilt pointiug. but neither of these shoud be an insumountable difficulty."," Implementing this strategy would require switching fields approximately twice per Swift orbit (assuming Swift changes targets to avoid Earth occultation), and would require real-time knowledge of the Swift pointing, but neither of these should be an insurmountable difficulty."
 A general and robust orpha1 searcli strategy is to requve sufficient areal aucl time coverage to ensure the detection of soije conventional CRB afterglows. and lenceaso of ou-axis dirty fireballs rour any population with au eveit rate comparable to tle GRB ‘ate.," A general and robust orphan search strategy is to require sufficient areal and time coverage to ensure the detection of some conventional GRB afterglows, and hence also of on-axis dirty fireballs from any population with an event rate comparable to the GRB rate."
" The search cadence srould be sept al dfZ515,4 (suggesting nieltly observatious) in order to catch the ou-anxis orphaus before their jet break aud so recoguize hem.", Tbe search cadence should be kept at $\delta t \la \tjet$ (suggesting nightly observations) in order to catch the on-axis orphans before their jet break and so recognize them.
 I such coverage [ails to iud on axis Orphans. it will mean tiat clirty ireballs do not couribute substantially to he overall oryj»hau afterglow rate.," If such coverage fails to find on axis orphans, it will mean that dirty fireballs do not contribute substantially to the overall orphan afterglow rate."
 More probably. we will iu Oll-axls orphatrs.," More probably, we will find on-axis orphans."
 Compariug their rate to the GRB ‘ate will tell us the fraction of cosiYological irealls tha achieves [Ty>2100.," Comparing their rate to the GRB rate will tell us the fraction of cosmological fireballs that achieves $\Gamma_0
\ga 100$."
 A large class o£ dirty |ireballs night. be fouud. and i it is. would ell us that GRBs are a minority population in a mucl largere class of cosmologicale fireballs.," A large class of dirty fireballs might be found, and if it is, would tell us that GRBs are a minority population in a much larger class of cosmological fireballs."
 The cotubinecd kiowledge of the on-axis dirty fireball rate aud the off-axis orphan rate would tLen allow correction o “the orphan statisties for cirv fireballs. aid lead to the desired constraint ou. CRBcollimation [rom orphan counts.," The combined knowledge of the on-axis dirty fireball rate and the off-axis orphan rate would then allow correction of the orphan statistics for dirty fireballs, and lead to the desired constraint on GRBcollimation from orphan counts."
on fairly elliptical orbits anc with apo-galactic distances around 20kkpc.,on fairly elliptical orbits and with apo-galactic distances around kpc.
 The globular clusters merged. within the supercluster eiving rise to an object with values for surface brightness and. absolute bolometric luminosity comparable to UCDs., The globular clusters merged within the supercluster giving rise to an object with values for surface brightness and absolute bolometric luminosity comparable to UCDs.
 The last formation scenario has been discussed in Oh. Lin Aarseth (1995) as well as Bekki et al. (," The last formation scenario has been discussed in Oh, Lin Aarseth (1995) as well as Bekki et al. ("
2001. 2003).,"2001, 2003)."
 The latter authors performed. numerical simulations of the dynamical evolution of nucleatecl dwarl galaxies. orbiting inside NGC 1390 and the Fornax cluster., The latter authors performed numerical simulations of the dynamical evolution of nucleated dwarf galaxies orbiting inside NGC 1399 and the Fornax cluster.
 Adopting a plausible scaling relation for dwarf galaxies. Bekki et al.," Adopting a plausible scaling relation for dwarf galaxies, Bekki et al."
 found that the outer stellar envelope of a nucleated dwarf was totally removed by tidally stripping over the course of several passages [rom the central region of their host., found that the outer stellar envelope of a nucleated dwarf was totally removed by tidally stripping over the course of several passages from the central region of their host.
 The nucleus is so dense that it will abwavs survive the tidal fiel ofa group or cluster potential., The nucleus is so dense that it will always survive the tidal field of a group or cluster potential.
 By construction in the initia conditions. the size and luminosity of the remnant were similar to those observed for UCDs and the host galaxy. was a Plummer mocel halo which has a constant density core au therefore easy. to tidally disrupt leaving behind the centra nucleus which would be associated with a UCD galaxy.," By construction in the initial conditions, the size and luminosity of the remnant were similar to those observed for UCDs and the host galaxy was a Plummer model halo which has a constant density core and therefore easy to tidally disrupt leaving behind the central nucleus which would be associated with a UCD galaxy."
 The main goal of the present study is to shed. ligh into the formation of UCDs investigating the last two formation scenarios in more detail and with more realistic initial conditions., The main goal of the present study is to shed light into the formation of UCDs investigating the last two formation scenarios in more detail and with more realistic initial conditions.
 First. we adopt the Fellhauer and Ixroupa (2002) model and combine i with the cold. dark. matter model (CDM) paradigm.," First, we adopt the Fellhauer and Kroupa (2002) model and combine it with the cold dark matter model (CDM) paradigm."
 We assume that globular clusters orm and subsequently: orbit around dark matter. haloes jwing masses comparable to that of he LFornax dwarl spheroidal and containing no other barvonic matter., We assume that globular clusters form and subsequently orbit around dark matter haloes having masses comparable to that of the Fornax dwarf spheroidal and containing no other baryonic matter.
 Due o dynamical friction. these globular cluster would spiral in owards the centre of the halo.," Due to dynamical friction, these globular cluster would spiral in towards the centre of the halo."
 We perform collisionless j/N-»ody simulations of this process and show that the object resulting from the coalescence of the globular clusters at the 1alo centre resembles a UCL., We perform collisionless $N$ -body simulations of this process and show that the object resulting from the coalescence of the globular clusters at the halo centre resembles a UCD.
 As we demonstrate below. the above model sullers from two major drawbacks which rule out its applicability for UCD formation.," As we demonstrate below, the above model suffers from two major drawbacks which rule out its applicability for UCD formation."
 The first issue is related to the observed high ML ratio (between 6 and 9 in solar units) recently reported for CDs (Llageganetal., The first issue is related to the observed high M/L ratio (between 6 and 9 in solar units) recently reported for UCDs \cite{ACS}.
2005).. One has to note here that the M/L ratio of UCDs is still debated (Micheal Drinkwater. private communication. compare Evstigneeva et al.," One has to note here that the M/L ratio of UCDs is still debated (Micheal Drinkwater, private communication, compare Evstigneeva et al."
 2007 who quote a M/L for UCDs 3 and 5 in solar units)., 2007 who quote a M/L for UCDs 3 and 5 in solar units).
 Such high values can probably not be achieved by the above mechanism., Such high values can probably not be achieved by the above mechanism.
 This is because sinking globular clusters will expel most of the dark matter particles from the halo centre (El-Zant.al. 2006)..," This is because sinking globular clusters will expel most of the dark matter particles from the halo centre \cite{amr,merritt,goerdt}."
 I£ there is no dark matter in UCDs. Fellbauer and Ixroupa (2006) describe a possible way to enhance the niass-to-light ratios of UCDs through tidal interactions.," If there is no dark matter in UCDs, Fellhauer and Kroupa (2006) describe a possible way to enhance the mass-to-light ratios of UCDs through tidal interactions."
 The second. problem with this scenario is related to the total luminosity of an UCD., The second problem with this scenario is related to the total luminosity of an UCD.
 At least today. dark. matter halos with a sullicient number of globular clusters to. produce such bright objects are very rare (Sharina.Puzia&Malsrov1996:Bokeretal.," At least today, dark matter halos with a sufficient number of globular clusters to produce such bright objects are very rare \cite{sharina,boeker1}."
 2002).. Second. we examine the scenario of Bekki et al. (," Second, we examine the scenario of Bekki et al. ("
2001. 2003) which is based on the hypothesis that UCDs are remnants of stripped. nucleatecl galaxies.,"2001, 2003) which is based on the hypothesis that UCDs are remnants of stripped nucleated galaxies."
 This mechanism js also been discussed by Ixazantzidis et al. (, This mechanism has also been discussed by Kazantzidis et al. (
2003).,2003).
 We est this model using SPILL simulations of a low mass galaxy which forms a nucleus via gas inflow to the inner zLOO xwsees., We test this model using SPH simulations of a low mass galaxy which forms a nucleus via gas inflow to the inner $\approx 100$ parsecs.
 Once this galaxy is placed. on a critical disrupting orbit within a cluster potential. the surviving nucleus is in excellent agreement with the latest observational constraints or UCDs. including their dark matter content and spatial distribution.," Once this galaxy is placed on a critical disrupting orbit within a cluster potential, the surviving nucleus is in excellent agreement with the latest observational constraints for UCDs, including their dark matter content and spatial distribution."
 The paper is organised as follows., The paper is organised as follows.
 In section 2 we describe the globular cluster numerical simulations anc compare the propertics of the resulting object with those of UCDs., In section 2 we describe the globular cluster numerical simulations and compare the properties of the resulting object with those of UCDs.
 Section 3 contains results from the simulations of tidal stripping of disk galaxies inside a host cluster potential., Section 3 contains results from the simulations of tidal stripping of disk galaxies inside a host cluster potential.
 Finally. in Section 4 we summarise our main conclusions.," Finally, in Section 4 we summarise our main conclusions."
 The globular cluster simulations were performed with a multi stepping. parallel N-bocky tree code 2001)..," The globular cluster simulations were performed with a multi stepping, parallel $N$ -body tree code \cite{stadel}."
" We create an NEW (Navarro.Frenk&White1996) halo. emploving the technique developed: by Ixazantzicis. Magorrian Aloore (2004). which has the following density. profile: In our case py=242 MAL. /pc? and kr,=L5kkpe."," We create an NFW \cite{nfw} halo, employing the technique developed by Kazantzidis, Magorrian Moore (2004), which has the following density profile: In our case $\rho_0 = 242$ $_\odot$ $^3$ and $r_{\rm s} = 1.5$ kpc."
 The halo has a virial mass of 15LO MAL., The halo has a virial mass of $1.5 \times 10^9$ $_{\odot}$.
 The concentration parameter is 20 which is the tvpical value for halos of this initial mass., The concentration parameter is 20 which is the typical value for halos of this initial mass.
 To increase mass resolution in the region of interest. we divide the halo into three shells (Zempetal.2007) cach of which contains 10° particles.," To increase mass resolution in the region of interest, we divide the halo into three shells \cite{zemp} each of which contains $10^5$ particles."
 The innermos shell has ppe radius., The innermost shell has pc radius.
 Phe second shell is between 100 and ppe while the third. shell contains the rest of the halo., The second shell is between 100 and pc while the third shell contains the rest of the halo.
 The softening lengths for these shells are 1. LO ane ppe respectively.," The softening lengths for these shells are 1, 10 and pc respectively."
 This shell model allows us to resolve the detailed kinematies within the central few parsecs whils retaining the global structure of the halo out to its viria radius of 29.39 προ., This shell model allows us to resolve the detailed kinematics within the central few parsecs whilst retaining the global structure of the halo out to its virial radius of 29.39 kpc.
 We use ten globular clusters consisting of LO” particles and being represented by the Ixing mocel (xing1966:Michie1963:Michie&Bodenheimer1963) Each elobular cluster is constructed. with a Wy=WO)far parameter of 6. a total mass of 4.2lo’ MIALT.. a central velocity. dispersion of kkm/s. and an absolute magnitudes of -8.5. assuming a mass to light ratio of 2 (Bokerctal.2004:Walcher2005).," We use ten globular clusters consisting of $10^5$ particles and being represented by the King model \cite{king,michie,bodenheimer}
 Each globular cluster is constructed with a $_0 = \Psi(0) / \sigma^2$ parameter of 6, a total mass of $4.2 \times 10^5$ $_{\odot}$, a central velocity dispersion of km/s, and an absolute magnitudes of -8.5, assuming a mass to light ratio of 2 \cite{boeker2,walcher}."
. We use Lppe for its gravitational softening length., We use pc for its gravitational softening length.
 The ten globular clusters are initially clüstributed within the halo between ppc and ppe., The ten globular clusters are initially distributed within the halo between pc and pc.
" They are spatially distributed according to pli)κ rl, which is in agreement with observations of globular"," They are spatially distributed according to $\rho (r) \propto r^{-4}$ , which is in agreement with observations of globular"
"In this section, we compare our results to previous work, especially those based on the Spitzer data.","In this section, we compare our results to previous work, especially those based on the Spitzer data."
" Comparisons are best done in the same wavelengths, since the conversion from either La,,,, or L124m to Lrrr involves the largest uncertainty."," Comparisons are best done in the same wavelengths, since the conversion from either $L_{8\mu m}$ or $L_{12\mu m}$ to $L_{TIR}$ involves the largest uncertainty."
 Hubble parameters in the previous work are converted to 0.7 for comparison., Hubble parameters in the previous work are converted to $h=0.7$ for comparison.
" Recently, using the Spitzer space telescope, restframe ""μπι LFs of z~1 galaxies have been computed in detail by (2007) in the GOODS fields and by Babbedgeet in the SWIRE field."," Recently, using the Spitzer space telescope, restframe $\mu$ m LFs of $z\sim$ 1 galaxies have been computed in detail by \citet{2007ApJ...660...97C} in the GOODS fields and by \citet{2006MNRAS.370.1159B} in the SWIRE field."
" In this section, we compare our restframe 81m LFs (Fig.4)) to these and discuss possible differences."," In this section, we compare our restframe $\mu$ m LFs \ref{fig:8umlf}) ) to these and discuss possible differences."
" In Fig.4,, we overplot Caputietal.(2007)’’s LFs at z=1 and z=2 in the cyan dash-dotted lines."," In \ref{fig:8umlf}, we overplot \citet{2007ApJ...660...97C}' 's LFs at $z$ =1 and $z$ =2 in the cyan dash-dotted lines."
 Their z=2 LF is in good agreement with our LF at 1.8<z<2.2., Their $z$ =2 LF is in good agreement with our LF at $<z<$ 2.2.
" However, their z=1 LF is larger than ours by a factor of 3-5 at logL>11.2."," However, their $z$ =1 LF is larger than ours by a factor of 3-5 at $\log L>11.2$."
 Note that the brightest ends (logL~11.4) are consistent with each other to within 1c., Note that the brightest ends $\log L\sim 11.4$ ) are consistent with each other to within $\sigma$.
" They have excluded AGN using optical-to-X-ray flux ratio, and we also have excluded AGN through the optical SED fit."," They have excluded AGN using optical-to-X-ray flux ratio, and we also have excluded AGN through the optical SED fit."
" Therefore, especially at the faint-end, the contamination from AGN is not likely to be the main cause of differences."," Therefore, especially at the faint-end, the contamination from AGN is not likely to be the main cause of differences."
" Since Caputietal.(2007) uses GOODS fields, cosmic variance may play a role here."," Since \citet{2007ApJ...660...97C} uses GOODS fields, cosmic variance may play a role here."
" The exact reason for the difference is unknown, but we point out that their ()7p estimate at z=1 is also higher than other estimates by a factor of a few (see their Fig.15)."," The exact reason for the difference is unknown, but we point out that their $\Omega_{IR}$ estimate at $z$ =1 is also higher than other estimates by a factor of a few (see their Fig.15)."
" Once converted into Lr;g, Magnellietal.(2009) also reported Caputietal.(2007)""'s z=1 LF is higher than their estimate based on 70m by a factor of several (see their Fig.12)."," Once converted into $L_{TIR}$, \citet{2009A&A...496...57M} also reported \citet{2007ApJ...660...97C}' 's $z$ =1 LF is higher than their estimate based on $\mu$ m by a factor of several (see their Fig.12)."
" They concluded the difference is from different SED models used, since their LF matched with that of Caputi once the same SED models were used."," They concluded the difference is from different SED models used, since their LF matched with that of \citet{2007ApJ...660...97C}' 's once the same SED models were used."
 We will our total LFs to those in the literature below., We will compare our total LFs to those in the literature below.
 compare Babbedgeetal.(2006) also computed restframe 8/:m LFs ising the Spitzer/SWIRE data., \citet{2006MNRAS.370.1159B} also computed restframe $\mu$ m LFs using the Spitzer/SWIRE data.
 We overplot their results at 70.25«z0.5 and 0.5<z« Lin Fig.4 with the pink dot-dashed lines., We overplot their results at $0.25<z<0.5$ and $0.5<z<1$ in \ref{fig:8umlf} with the pink dot-dashed lines.
" In both redshift ranges, good agreement is found zat higher luminosity bins (Lgjm>10195 1,9)."," In both redshift ranges, good agreement is found at higher luminosity bins $L_{8\mu m}>10^{10.5} L_{\odot}$ )."
" However, at all redshift ranges including the ones not shown here, Babbedgeetal.(2006) tends to show a flatter faint-end tail than ours, and a smaller $ by a factor of ~3."," However, at all redshift ranges including the ones not shown here, \citet{2006MNRAS.370.1159B} tends to show a flatter faint-end tail than ours, and a smaller $\phi$ by a factor of $\sim$ 3."
" Although the exact reason is unknown, the deviation starts toward the fainter end,where both works approach the flux limits of the surveys."," Although the exact reason is unknown, the deviation starts toward the fainter end,where both works approach the flux limits of the surveys."
" Therefore, possibly incomplete sampling may be one of the reasons."," Therefore, possibly incomplete sampling may be one of the reasons."
" It is also reported that the faint-end of IR LFs depends on the environment, in the sense that higher-density environment has steeper faint-end tail (Gotoetal.,2010).."," It is also reported that the faint-end of IR LFs depends on the environment, in the sense that higher-density environment has steeper faint-end tail \citep{cluster_LF}."
" Note that at z=1, Babbedgeetal. (2006)’’s LF (pink) deviates from that by Caputietal.(2007) (cyan) by almost a magnitude."," Note that at $z$ =1, \citet{2006MNRAS.370.1159B}' 's LF (pink) deviates from that by \citet{2007ApJ...660...97C} (cyan) by almost a magnitude."
" Our ""μπι LFs are between these works.", Our $\mu$ m LFs are between these works.
" These comparisons suggest that even with the current generation of satellites and state-of-the-art SED models, factor-of-several uncertainties still remain in estimating the 8jum LFs at z~1."," These comparisons suggest that even with the current generation of satellites and state-of-the-art SED models, factor-of-several uncertainties still remain in estimating the $\mu$ m LFs at $\sim$ 1."
 More accurate determination has to await a larger and deeper survey by the next generation IR satellites such as Herschel and WISE., More accurate determination has to await a larger and deeper survey by the next generation IR satellites such as Herschel and WISE.
" To summarise, our 84m LFs are between those by Babbedgeetal.(2006) and Caputietal. (2007), and discrepancy is by a factor of several at most."," To summarise, our $\mu$ m LFs are between those by \citet{2006MNRAS.370.1159B} and \citet{2007ApJ...660...97C}, and discrepancy is by a factor of several at most."
" We note that both of the previous works had to rely on SED models to estimate Lgym from the Spitzer $54,,,, in the MIR wavelengths where", We note that both of the previous works had to rely on SED models to estimate $L_{8\mu m}$ from the Spitzer $S_{24\mu m}$ in the MIR wavelengths where
our sink particles.,our sink particles.
 This likely overestimates the efficiency that would result were feedback from massive stars included., This likely overestimates the efficiency that would result were feedback from massive stars included.
 It is worth noting that our gas densities and core sizes are similar to the continuum surveys (Andre ..) that would require a 100 conversion in order to obtain a mapping from the core mass function to the stellar IMF’., It is worth noting that our gas densities and core sizes are similar to the continuum surveys (Andre ..) that would require a 100 conversion in order to obtain a mapping from the core mass function to the stellar IMF.
" Furthermore, Previous simulations including ionisation and winds (792) do not find a large change in the resultant masses or mass spectra."," Furthermore, Previous simulations including ionisation and winds \citep{Daleetal2005, DalBon2008} do not find a large change in the resultant masses or mass spectra."
 The simulation was followed for 1.02 free-fall times or &6.6x10° years and zz3.9x10° years after the first stars formed (see Fig. , The simulation was followed for $1.02$ free-fall times or $\approx 6.6 \times 10^5$ years and $\approx 3.9\times 10^5$ years after the first stars formed (see Fig. \ref{OAevol}) ).
"During this time, 2542 stars were formed with masses [).between 0.017 and 30Mo."," During this time, 2542 stars were formed with masses between $0.017$ and 30."
". The majority of these stars form in the upper gravitationally bound part of the cloud while some 7 per cent form in the lower, gravitationally unbound regions."," The majority of these stars form in the upper gravitationally bound part of the cloud while some 7 per cent form in the lower, gravitationally unbound regions."
 Figure shows the developing initial mass function during the star formation process., Figure \ref{imf9} shows the developing initial mass function during the star formation process.
 The stars form with masses comparable to the Jeans mass of the local gas., The stars form with masses comparable to the Jeans mass of the local gas.
 These initial masses are initially of the order of several tenths of, These initial masses are initially of the order of several tenths of
Orbital evolution of dust particles in the Solar System is investigated. partially iu a plivsical way. since the time of Povuting (1903).,"Orbital evolution of dust particles in the Solar System is investigated, partially in a physical way, since the time of Poynting (1903)."
 Besides action of solar gravity aud solar electromagnetic radiation. iu the form of the Povutine-Robertson effect (Povutiug 1903. Robertson 1937. Klackka 20041. 20082. 2008b. IKlackka et al.," Besides action of solar gravity and solar electromagnetic radiation, in the form of the Poynting-Robertson effect (Poynting 1903, Robertson 1937, Klačkka 2004, 2008a, 2008b, Klačkka et al."
 20092). also the effect of solar corpuscular radiation is taken iuto account (Whipple 1955. 1967. Dolnauvi 1978 and others).," 2009a), also the effect of solar corpuscular radiation is taken into account (Whipple 1955, 1967, Dohnanyi 1978 and others)."
 It is eencerallv considered that the eect of solar wind is simular to the effect of solar clectromaguetic radiation and the solar wind ds about (0.2-0.3)-times less nuportaut than the Povuting-Robertson eect AVipple 1955. 1967. Dohnauvi 1978. aud. e.g.. Mukai and Yamamoto 1982. Jacksou aud Zook 1989. Leiuert aud Coiun 1990. Dermott et al.," It is generally considered that the effect of solar wind is similar to the effect of solar electromagnetic radiation and the solar wind is about (0.2-0.3)-times less important than the Poynting-Robertson effect (Whipple 1955, 1967, Dohnanyi 1978, and, e.g., Mukai and Yamamoto 1982, Jackson and Zook 1989, Leinert and Grünn 1990, Dermott et al."
 1001. Ciustafson 1991. Reach et al.," 1994, Gustafson 1994, Reach et al."
 1995. Devmott et al.," 1995, Dermott et al."
 2001. Alfiuato et al.," 2001, Minato et al."
 2001. Abe 2009. Mann 2009).," 2004, Abe 2009, Mann 2009)."
 However. this conventional idea is in variance with reality. iu general (Klackka and Saniga 1993. Klackka 1991. IKkocifaj aud Wlackka 2008. IKlackka et al.," However, this conventional idea is in variance with reality, in general (Klačkka and Saniga 1993, Klačkka 1994, Kocifaj and Klačkka 2008, Klačkka et al."
 2009b. Klackka et al.," 2009b, Klačkka et al."
 2009c. Passtor et al.," 2009c, Pásstor et al."
 2010)., 2010).
 It turns out that orientation ou the correct plivsica treating of the solar wind effect (IKIlackka et al., It turns out that orientation on the correct physical treating of the solar wind effect (Klačkka et al.
 2009c) can siguificautle nuprove our knowledge of ie action of the solar wind on the orbital evolution of dust erains iu the Solar System., 2009c) can significantly improve our knowledge of the action of the solar wind on the orbital evolution of dust grains in the Solar System.
 This paper is based ou je theory preseuted by IKlackka et al. (, This paper is based on the theory presented by Klačkka et al. (
2009€).,2009c).
 The aiu of this paper is to apply solar wind action πι orbital evolution of dust particle moving in the zone of 16 Edeeworth-Ixuiper belt. which cau be defined by semi-najor axes 30-50 AU. approximately.," The aim of this paper is to apply solar wind action on orbital evolution of dust particle moving in the zone of the Edgeworth-Kuiper belt, which can be defined by semi-major axes 30-50 AU, approximately."
 We are interested im je time of the particle stav in the zone., We are interested in the time of the particle stay in the zone.
 More plysica uodoel of the solar wind action (Nlackka e al., More physical model of the solar wind action (Klačkka et al.
 20000) wi )6 Colmpared with the conventional access., 2009c) will be compared with the conventional access.
 The particles of several micrometers to tens of micrometers are considered. so the Loreutz force Guterplauctary magnetic feld) aux collisions among particles are ucelected (Doluzusi 1975. Leinert and Caüuu 1990. Dermott et al.," The particles of several micrometers to tens of micrometers are considered, so the Lorentz force (interplanetary magnetic field) and collisions among particles are neglected (Dohnanyi 1978, Leinert and Grünn 1990, Dermott et al."
 2001 aud Crüuu et al, 2001 and Grünn et al.
 1985)., 1985).
 παν of the Sun. Povutine-Robertson effect. action of the solar wind and fast interstellar gas flow are forces which influcnee motion of the particles.," Gravity of the Sun, Poynting-Robertson effect, action of the solar wind and fast interstellar gas flow are forces which influence motion of the particles."
 Siuultancous action of all these effects can shed a light 1 the relevance of the solar wind on the orbital evolution dust particles., Simultaneous action of all these effects can shed a light on the relevance of the solar wind on the orbital evolution of dust particles.
 The inportauce of the solar wind action id the P-R effect can be compared., The importance of the solar wind action and the P-R effect can be compared.
 The results can help uot ouly iu better understanding of dust evolution bevond planets in the Solar System. but also in understanding of dust belts in other stellar svstenis (sec. c.g.. Marley 2008).," The results can help not only in better understanding of dust evolution beyond planets in the Solar System, but also in understanding of dust belts in other stellar systems (see, e.g., Marley 2008)."
 Section 2 preseuts relevant equation of motion., Section 2 presents relevant equation of motion.
 Sec., Sec.
 3 offers the most importan results of nuimerical solution of the equation of motion., 3 offers the most important results of numerical solution of the equation of motion.
 The section compares the results for the couveutional time-independent radial solar wind with those obtained or nore real solar wind model., The section compares the results for the conventional time-independent radial solar wind with those obtained for more real solar wind model.
 Sec., Sec.
 L presents a short discussion on current situation iu 1uodeling of dust evolution in the Solar Systein., 4 presents a short discussion on current situation in modeling of dust evolution in the Solar System.
 Dealing with orbital evolution of dust particles i space bevoud plancts. eravitv of the Sun aud noneravitational forces acting ou the particles plav the most relevant role.," Dealing with orbital evolution of dust particles in space beyond planets, gravity of the Sun and nongravitational forces acting on the particles play the most relevant role."
 Equation of motion of the dust erain under the action of solar gravity. solar clectromaguetic aud corpuscular radiation. fast interstellar eas flow. aud eravity of other bodies. is where r is the evain’s position vector with respect to the Sun. velocity vector e = drdt. C is the eravitational constant. AZ. is mass of the Sum. r= [9].," Equation of motion of the dust grain under the action of solar gravity, solar electromagnetic and corpuscular radiation, fast interstellar gas flow, and gravity of other bodies, is where $\vec{r}$ is the grain's position vector with respect to the Sun, velocity vector $\vec{v}$ $=$ $d\vec{r}/dt$, $G$ is the gravitational constant, $M_{\odot}$ is mass of the Sun, $r$ $=$ $| \vec{r} |$."
 The first term on the right-haud side of Eq. (1)), The first term on the right-hand side of Eq. \ref{eqm}) )
 corresponds to eravitational acceleration of the Sun., corresponds to gravitational acceleration of the Sun.
 The second tem on the right-hand side of Eq. (13) , The second term on the right-hand side of Eq. \ref{eqm}) )
is the action of the solar electromagnetic radiation. on a spherical particle., is the action of the solar electromagnetic radiation on a spherical particle.
 It is represented by the, It is represented by the
number of emission lines redshifts in the range 0<z1 raises the question whether our model - based on a sample of redshifts measured from the afterglows - is consistent.,number of emission lines redshifts in the range $0 < z < 1$ raises the question whether our model - based on a sample of redshifts measured from the afterglows - is consistent.
" We calculate, using our model, the expected redshift distribution of the entire bursts population and compare it to the number of observed bursts with a known redshift (measured using all methods)."," We calculate, using our model, the expected redshift distribution of the entire bursts population and compare it to the number of observed bursts with a known redshift (measured using all methods)."
 Clearly for any range of redshifts the number of known bursts with a given redshift range should not exceed the number predicted by the model., Clearly for any range of redshifts the number of known bursts with a given redshift range should not exceed the number predicted by the model.
 Fig 7 depicts this comparison for the cumulative number and for the counts number in redshift bins., Fig \ref{fig:NzvsNtot} depicts this comparison for the cumulative number and for the counts number in redshift bins.
 Our model is consistent for all redshifts =0.4., Our model is consistent for all redshifts $\gtrsim 0.4$.
 However there is an excess of low redshift bursts not predicted by our model., However there is an excess of low redshift bursts not predicted by our model.
 The discrepancy arise due to three bursts with z<0.1 while less than one burst is predicted by the model., The discrepancy arise due to three bursts with $z < 0.1$ while less than one burst is predicted by the model.
 The redshift of the three bursts was determined using, The redshift of the three bursts was determined using
"density =106 cm, we get abundances of 4x107? (CH) and 9x1071! (CH*).","density $\gtrsim 10^6$ $^{-3}$, we get abundances of $4\times10^{-9}$ (CH) and $9\times10^{-11}$ $^+$ )."
 Why are the hydrides mostly in the ground-state?, Why are the hydrides mostly in the ground-state?
" The excitation temperature Τεκ of a species, assumed to have only two levels, is given by with the kineticAE)- Intemperature(1 of the collision partner Τμ, the energy difference between the levels ΔΕ, the escape probability for a line photon f, the critical density mej, and the density of a collision partner n(H2)."," The excitation temperature $T_{\rm ex}$ of a species, assumed to have only two levels, is given by with the kinetic temperature of the collision partner $T_{\rm kin}$ , the energy difference between the levels $\Delta E$ , the escape probability for a line photon $\beta$, the critical density $n_{\rm crit}$ and the density of a collision partner $n({\rm H}_2)$."
 The observed excitation temperature of CH*(J=1— 0) of 38.4 K can be reproduced by Eq., The observed excitation temperature of $^+$ $J=1-0$ ) of 38.4 K can be reproduced by Eq.
" 2 with B~0.04, assuming a kinetic temperature of 500 K, necessary for the formation of the molecule, and a H5 density of order 106 cm, typical at scales of the beam."," \ref{eq:exci} with $\beta \sim 0.04$, assuming a kinetic temperature of 500 K, necessary for the formation of the molecule, and a $_2$ density of order $10^6$ $^{-3}$, typical at scales of the beam."
" This B corresponds to an optical depth of a few at the line center, depending on the expression used for B=fir)."," This $\beta$ corresponds to an optical depth of a few at the line center, depending on the expression used for $\beta=\beta(\tau)$."
 Far infrared radiation by the dust continuum also pumps the molecule., Far infrared radiation by the dust continuum also pumps the molecule.
" Critical densities of other hydrides are not well known, but likely of similar order."," Critical densities of other hydrides are not well known, but likely of similar order."
" Thus, the observed low excitation temperature found for OH* and H5O* can be well explained by the origin of the absorption in tenuous gas."," Thus, the observed low excitation temperature found for $^+$ and $_2$ $^+$ can be well explained by the origin of the absorption in tenuous gas."
 Predictions for CH and CH* abundances and line fluxes can be made utilizing the photochemical and outflow model of Brudereretal.(2009b)., Predictions for CH and $^+$ abundances and line fluxes can be made utilizing the photochemical and outflow model of \citet{Bruderer09b}.
". In a spherical envelope model, Brudereretal.(2010) find volume averagedabundancesorders of magnitude lower than observed (~6x10:11 for CH and ~5x10776 for CH*)."," In a spherical envelope model, \citet{Bruderer10} find volume averagedabundancesorders of magnitude lower than observed $\sim6 \times 10^{-11}$ for CH and $\sim5 \times 10^{-16}$ for $^+$ )."
 They present also a model including, They present also a model including
realism.,realism.
" Physically, the random motions gained by gas clouds are dissipated in collisions so that they keep moving on near circular orbits."," Physically, the random motions gained by gas clouds are dissipated in collisions so that they keep moving on near circular orbits."
" Stars that form from the clouds therefore have similar motions, and a continuous supply of fresh stars on near-circular orbits maintains a responsive stellar distribution that allows spiral activity to continue."," Stars that form from the clouds therefore have similar motions, and a continuous supply of fresh stars on near-circular orbits maintains a responsive stellar distribution that allows spiral activity to continue."
" Quantitatively, fresh stars added to a disc at the rate of a few every year is sufficient for this purpose."," Quantitatively, fresh stars added to a disc at the rate of a few every year is sufficient for this purpose."
 Thus the transient spiral picture offers a natural explanation for the absence of spiral patterns in SO disc galaxies that have little or no gas and the long-noted correspondence that disc galaxies with significant gas components also manifest spiral patterns., Thus the transient spiral picture offers a natural explanation for the absence of spiral patterns in S0 disc galaxies that have little or no gas and the long-noted correspondence that disc galaxies with significant gas components also manifest spiral patterns.
" Note that this appealing argument again does not uniquely favour short-lived waves, since also invoke a gas component both to limit the amplitude of their mildly growing modes and to maintain the dynamically cool outer disc."," Note that this appealing argument again does not uniquely favour short-lived waves, since \cite{BL96} also invoke a gas component both to limit the amplitude of their mildly growing modes and to maintain the dynamically cool outer disc."
" It has been clear for some time that the velocity dispersion of disc stars in the solar neighbourhood rises with age 2004) and also, for main sequence stars, with colour ⊓ which is a surrogate for mean age."," It has been clear for some time that the velocity dispersion of disc stars in the solar neighbourhood rises with age \citep{Wiel77,Nord04}
 and also, for main sequence stars, with colour \citep{AB09}, which is a surrogate for mean age."
" The highest velocities cannot be produced by cloud scattering and some other accelerating agent, such as transient spirals, seems to be required."," The highest velocities cannot be produced by cloud scattering \citep{Lace91,HF02} and some other accelerating agent, such as transient spirals, seems to be required."
" Long-lived spirals, which do not heat the disc nearly as effectively, could not achieve the requisite high velocities."," Long-lived spirals, which do not heat the disc nearly as effectively, could not achieve the requisite high velocities."
" Studies of the metallicities and ages of nearby stars ∢↴olmberg,Nordstróm&Andersen||2007)find that older stars tend to have lower metallicities on average."," Studies of the metallicities and ages of nearby stars \citep{Edva93,Nord04,Reid07,Holm07} find that older stars tend to have lower metallicities on average."
" As it is difficult to estimate the ages of individual stars, the precise form of the relation is still disputed."," As it is difficult to estimate the ages of individual stars, the precise form of the relation is still disputed."
" However, there seems to be general agreement that there is a spread of metallicities at each age, which is also supported by other studies (Chen ef⋅⊔∙∆"," However, there seems to be general agreement that there is a spread of metallicities at each age, which is also supported by other studies \citep{CHW03,Hayw08,SH10}."
" metallicity spread amongst coeval stars is inconsistent with a simple chemical evolution model in which the metallicity of the disc rises monotonically in each annular bin, without mixing in the radial direction."," A metallicity spread amongst coeval stars is inconsistent with a simple chemical evolution model in which the metallicity of the disc rises monotonically in each annular bin, without mixing in the radial direction."
" Again (2002) and Ro&kkar et showed that when the disc supports recurrent transient spirals, the needed radial mixing arises naturally through angular momentum changes at corotation that do not heat the disc."," Again \cite{SB02} and Ros̆kkar (2008ab) showed that when the disc supports recurrent transient spirals, the needed radial mixing arises naturally through angular momentum changes at corotation that do not heat the disc."
 developed the first chemical evolution model for the Milky Way disc to include this radial churning., \cite{SB09} developed the first chemical evolution model for the Milky Way disc to include this radial churning.
 The transient nature of the patterns is an essential aspect to produce efficient churning., The transient nature of the patterns is an essential aspect to produce efficient churning.
" The horse-shoe orbits near co-rotation of a large-amplitude spiral cause a single change in the angular momentum if the spiral grows and decays in less than half the (long) orbit period Were the spiral to persist for longer than this, thenframe[] the exchanges a star experiences at one arm would be undone when the star encounters the next arm, and the star would merely take periodic, equal inward and outward steps in radius leading to no lasting change."," The horse-shoe orbits near co-rotation of a large-amplitude spiral cause a single change in the angular momentum if the spiral grows and decays in less than half the (long) orbit period Were the spiral to persist for longer than this, then the exchanges a star experiences at one arm would be undone when the star encounters the next arm, and the star would merely take periodic, equal inward and outward steps in radius leading to no lasting change."
" report substantial mixing due to combined influence of a bar and spirals, but their simulations were of short duration and it is likely the effect is large only as the bar forms."," \cite{Minc10} report substantial mixing due to combined influence of a bar and spirals, but their simulations were of short duration and it is likely the effect is large only as the bar forms."
" Only recurrent transient patterns, with corotation radii spread at random locations, can cause the home radii of stars to “diffuse” across the disc throughout its life."," Only recurrent transient patterns, with corotation radii spread at random locations, can cause the home radii of stars to “diffuse” across the disc throughout its life."
 Tremaine (private communication) pointed out that transient spirals could co-exist with long-lived patterns., Tremaine (private communication) pointed out that transient spirals could co-exist with long-lived patterns.
" This possibility could be consistent, for example, with the change in the appearance of spiral patterns between the blue and near-IR passbands (e.g.[al[1994) ,, although these data contain no information about the lifetimes of features of any colour."," This possibility could be consistent, for example, with the change in the appearance of spiral patterns between the blue and near-IR passbands \citep[\eg][]{BW91,Bloc94}, although these data contain no information about the lifetimes of features of any colour."
" The co-existence of short-lived spirals seems inconsistent with the specific mode theory for long-lived patterns proposed by (1996)., since transient waves would heat the outer disc to an extent that must render their proposed mechanism unworkable, as shown in reftheory."," The co-existence of short-lived spirals seems inconsistent with the specific mode theory for long-lived patterns proposed by \cite{BL96}, since transient waves would heat the outer disc to an extent that must render their proposed mechanism unworkable, as shown in \\ref{theory}."
.4., .4.
" However, some other possible (as yet unknown) mechanism for long-lived spirals might remain viable."," However, some other possible (as yet unknown) mechanism for long-lived spirals might remain viable."
 An observational test for the coexistence of both short- and long-lived patterns would be a daunting challenge., An observational test for the coexistence of both short- and long-lived patterns would be a daunting challenge.
" For a well-constructed sample of galaxies, one would have to determine the mean amplitude needed to produce the desired heating and churning by the short-lived patterns, and then show that the mean observed spiral amplitudes, estimated somehow from the passband dependent photometry, would or would not allow an additional significant long-lived spiral component."," For a well-constructed sample of galaxies, one would have to determine the mean amplitude needed to produce the desired heating and churning by the short-lived patterns, and then show that the mean observed spiral amplitudes, estimated somehow from the passband dependent photometry, would or would not allow an additional significant long-lived spiral component."
 The question of the lifetimes of spiral patterns is important because galaxy discs evolve to a much greater extent if spirals are transient than if they are quasi-steady refevolv))., The question of the lifetimes of spiral patterns is important because galaxy discs evolve to a much greater extent if spirals are transient than if they are quasi-steady \\ref{evolv}) ).
" In this paper, I have summarized the evidence that bears on the question; were any one piece decisive, there would have been no need to write this paper."," In this paper, I have summarized the evidence that bears on the question; were any one piece decisive, there would have been no need to write this paper."
" However, I find that the transient picture is favoured by the velocity structure in the solar neighbourhood refobserv)) and by the behaviour of simulations 9 reftheory."," However, I find that the transient picture is favoured by the velocity structure in the solar neighbourhood \\ref{observ}) ) and by the behaviour of simulations \\ref{theory}."
.4 , .4 \ref{sims}) ).
"Direct observationalE]). evidence from external galaxies refobserv)) is frustratingly inconclusive, largely because our single snapshot view can tell us little about the lifetimes of the patterns or, if self-excited, the mechanism"," Direct observational evidence from external galaxies \\ref{observ}) ) is frustratingly inconclusive, largely because our single snapshot view can tell us little about the lifetimes of the patterns or, if self-excited, the mechanism"
 forced boundary and fades towards theinterior of the computationalbox. The magneticfieldis,"normalized to have the same rms, but the spatial pattern was different for every simulation as determined by the random amplitudes."
" overall weakly correlated withthe forcing velocity,and atmost reaches a mild correlationclose tothe forcing boundary, never clo"," In the simulations presented here the forcing is exactly the same for each simulation, as it is simply a single Fourier mode [eq. \ref{eq:f0}) )]."
se[]toa strong correlation Cor— In Figure we showthe 2D averagesinthe x-y planes of the magneticj?/Rand velocity," As the forcing is the same for each simulation the overlap is more evident, and makes stronger the claim that total dissipation is."
" functionofzat differenttimes. Thebehaviourisvery similarto our previous simulations with −−different (vortical)) forcing patterns These macroscopic quantitiesaresmooth andpresent almostno variation alongthe axial direction. The velocitymust approachits boundary valuesat z=0 and 10,and"," We conjectured this hypothesis \citep{rved08} because the energy flux entering the system (Poynting flux) and the energy transport from large to small scales due to the development of a turbulent dynamics are both independent from the Reynolds number, where the turbulent transport exhibits this property only for a of the Reynolds number."
 in the volume higher thanthe boundaryvelocity fact alsothevelocity the volume, This unfortunately imply at all that the thermodynamical and radiative outcome is independent of the Reynolds number.
" is nota mappingofthe boundary forcingbut developsself- consistently,in particular the plasma jets at", In fact how field-lines are heated strongly depends on the dynamics and properties around the single current sheet.
 reconnection locations contributetoo to its average. The rea, These become thinner and thinner at higher Reynolds numbers and the overall dynamics more chaotic.
sonof the overall smooth behaviorof these quantities isthat every disturbanceor gradient alongthe axial directionis smoothed outby thefast propagation o, Hence the thermodynamical properties of the field-lines that cross the current sheets (elongated along the axial direction) and get so heated are all to be explored.
"fAlfvénnwaves alongthis direction;their propagation time7,is infact thefastest timescale present(in particular fasterthan nonlinear timescale 74«74;), tends tobe homogeneous alongthis direction. Dissipation"," An open question is whether the dissipation is independent of the Reynolds number also in the single current sheets, or their number and properties change to attain independence for the total dissipation."
 and 5.2. versus Reynolds Number Transition to Turbulence Simulations A-E differonlyforthevalueof dissipationfor the5 simulationsas[|a function oftime. The simulation describedin the previous," Research in this area is active \citep{lsc07,lap08,ser09, slu09,csd09,bhyr09,lus09,hb10,ser10} and has already shown that at relatively high Reynolds numbers reconnection departs the classic Sweet-Parker scalings in the MHD regime."
"paragraph hadRe = 800,after the firstbig dissipative peakits signal displaysa complex temporal structure thatarises fromthe underlying turbulent dynamics. Decreasing the valueof the Reynolds numberto"," Furthermore for a plasma in coronal conditions kinetic effects cannot be excluded a priori, in fact they might play an important role in dissipating energy through particle acceleration."
" Re—400the structure of the signal hasa simpler structure, withan almost sinusoi", Also the high value of the magnetic Prandtl number could have a bearing \citep{she04}.
"dalformofdifferent amplitudeand period ~874 diffusion,,At— i.e.200wwhendiffusion diffusionaffects dominatesthe system thesubstantiall"," Nonetheless the total dissipation, integrated over the whole volume, is likely to not depend on the detailed small-scale dissipative mechanism."
"y,dynamics.. instabilityRe takes longer t", In fact the energy injection rate [eq. \ref{eq:pf}) )]
"odevelop, and theactual signal and linear saturation curve differ afterward little. 'Thereis no first bigdissipative peak, for a time ~ 55074 the signal followsthe linear saturationlong upcurvetot(19)).it slightly displaying a sinusoidalbehaviorAfterwards ofvery smalldeparts amplitude. At lower Reynolds (Re —100 10are shown) diffusion dominatesnumberthe dynamics, and nosmall scalesis"," depends only on the large scale fields and the transfer of this energy toward the small scales appears to be local \citep{rv10}, it is determined by the fields at neighboring large scales, thus making also the transfer energy rate toward the small scales independent of the Reynolds number (provided it is beyond a minimal threshold to have scale separation)."
" formed, whileohmic follow curve (18))-(19)) describingdissipation and energy the is formed,as the termsdiffusive are equilibriumthat nonlinear Itis completelyvery dep"," In Figure \ref{fig9} we show a close-up of total dissipation for an equal interval of time $\Delta\, t = 600\, \tau_A$ on the same scale in order to highlight their temporal structures."
"leted.interesting noticethat dissipation for B andtoCwith Retotal—800, 400and simulationssubstantiallyA, overlap each other. As"," At higher Reynolds number smaller time frequencies are present, a clear indication of a transition to turbulence \citep{fr95}."
" typical with spectral200 numericalcodes, which use Fourier transforms to compute derivatives, eqs. (4))-(@)),"," For Reynolds numbers lower than $100$, the signal is completely flat following exactly the linear saturation curve eq. \ref{eq:diff3}) )."
 is numerically verified withina verysmall error:itis around 2%for the lower resolution simulations and decreasesbelow 0.5% athigher resolutions (see[Rap-] Figure 5.5).. presented here the dissipationrates , In order to study the spectral properties of the system and compare them with those of previous simulations we have performed a new simulation where the sheared forcing is applied on (run F) with reversed direction [eqs. \ref{eq:f0}) )
are substantiall, and \ref{eq:f2}) )].
y due to theexplicit dissipative terms presentin equations (4))- (8).while the diffusiondue to the numerical schemesis negligible. Wehad alreadyseen this behavior inour pre," We compare these results with those obtained from a previous simulation (run G), that has all the same parameters, except that on the two boundary planes a large-scale “vortex-like” velocity pattern with the same rms $\langle \mathbf{u^2} \rangle = 1/2$ ) is applied."
"vious simulations with vortex forcing al.][2008).., butin that casethe boundary (Rappazzovelocityetwasslightly differentin each simulation.In factit wasbuiltbya linear combination o"," Both simulations have been performed with a numerical resolution of $512\times 512\times 200$ grid points, and hyperdiffusion with dissipativaty $n=4$ and $R_4 = 10^{19}$ ."
"vious simulations with vortex forcing al.][2008).., butin that casethe boundary (Rappazzovelocityetwasslightly differentin each simulation.In factit wasbuiltbya linear combination of"," Both simulations have been performed with a numerical resolution of $512\times 512\times 200$ grid points, and hyperdiffusion with dissipativaty $n=4$ and $R_4 = 10^{19}$ ."
About of radio quiet luminous quasars (QSOs) exhibit Droad Absorption Lines (BAL) in the highh ionized resonant UV lines. such as NV. SiIV. and CIV. (Purnshek 1984. Wevmann. 1995).,"About of radio quiet luminous quasars (QSOs) exhibit Broad Absorption Lines (BAL) in the highly ionized resonant UV lines, such as NV, SiIV and CIV (Turnshek 1984, Weymann, 1995)."
 Approximately of these also show low ionization BAL such as Mgll. ALLE (Mevymann οἱ al.," Approximately of these also show low ionization BAL such as MgII, AlIII (Weymann et al."
 1991. Lartig Baldwin 1990. Boroson Alevers 1992. Voit. Wevman Ixorista 1993).," 1991, Hartig Baldwin 1990, Boroson Meyers 1992, Voit, Weyman Korista 1993)."
 The BALs are displaced to shorter wavelengths with velocities of up to 0.1c relative to their corresponding emission line peaks., The BALs are displaced to shorter wavelengths with velocities of up to 0.1c relative to their corresponding emission line peaks.
" ""Fhis indicates the presence of partially tonizecl gas outflows at enormous", This indicates the presence of partially ionized gas outflows at enormous
We begin with the spectrum from the source SWIFT J1753.5-0127.,We begin with the spectrum from the source SWIFT J1753.5-0127.
 ? took a 42 ksec observation and estimated an X-ray luminosity (0.5-10keV) of Lx/Lgaa=2.6x10-3(d/8.5 kpc)?(M/10 Μο).," \cite{2006ApJ...652L.113M} took a 42 ksec observation and estimated an X-ray luminosity (0.5-10keV) of $L_{\rm{X}}/L_{\rm{Edd}} = 2.6 \times 10^{-3} (d/8.5\,
$ $)^2 ($ M/10 $_\odot)$."
" They fit the spectrum to a power law with a photon index Γ=1.67 (2-10 keV), and interstellar absorption of Ny=2.3x10?! cm”."," They fit the spectrum to a power law with a photon index $\Gamma = 1.67$ (2-10 keV), and interstellar absorption of $N_{\rm{H}} = 2.3 \times 10^{21}$ $^{-2}$."
" Fitting the spectrum with an absorped power law component alone reveals a small soft excess below 2 keV, which they fit to a disk with ΚΤη«0.22 keV and Rj,« Rs (M/10Mo)(d/8.5kpc)/ cos!?j,"," Fitting the spectrum with an absorped power law component alone reveals a small soft excess below 2 keV, which they fit to a disk with $kT_{in} \simeq 0.22$ keV and $_{in}\simeq$ $_{\rm{S}}$ $_\odot)\;
(d/8.5$ $)/\cos^{1/2}i$."
 A very truncated disk will be too cool to be observable in X- while an untruncated disk will have a higher temperature than is observed (because of the effects of heating from the corona).," A very truncated disk will be too cool to be observable in X-rays, while an untruncated disk will have a higher temperature than is observed (because of the effects of heating from the corona)."
" To compare with the observed spectra we assume a moderate truncation radius of 15Rs, which is qualitatively very different from an untruncated disk (which will have an inner radius between 0.5— 3Rs depending on the spin of the black hole).We make the same assumptions for mass (M = 10Mo) and distance (d= 8.5kpc) as in ? and assume an inclination of µ=0.5."," To compare with the observed spectra we assume a moderate truncation radius of $_{\rm S}$, which is qualitatively very different from an untruncated disk (which will have an inner radius between $0.5-3$ $_{\rm S}$ depending on the spin of the black hole).We make the same assumptions for mass (M = $_\odot$ ) and distance $d =
8.5$ kpc) as in \cite{2006ApJ...652L.113M} and assume an inclination of $\mu = 0.5$."
 The simulation then predicts an accretion rate of M/Mgaa=1.5x10? for an efficiency of10%., The simulation then predicts an accretion rate of $\dot{M} / \dot{M}_{\rm Edd } = 1.5 \times 10^{-3}$ for an efficiency of.
". In order to match the measured flux in the power law component, we increase the flux in each component by a factor C=12, to give an accretion rate M/Mggag=1.7x1072."," In order to match the measured flux in the power law component, we increase the flux in each component by a factor $C =
12$, to give an accretion rate $\dot{M} / \dot{M}_{\rm{Edd}} =
1.7 \times 10^{-2}$."
" In the hot layer, the predicted surface density profile gives the optical depth (r)=0.87, and the energy balance between the cool disk and hot layer (see sect. 2.3))"," In the hot layer, the predicted surface density profile gives the optical depth $\langle\tau\rangle = 0.87$, and the energy balance between the cool disk and hot layer (see sect. \ref{sec:endist}) )"
" determines a temperature of kT,= 75 keV, which gives a photon index of Γ=1.87 (3-10 The density of the disk and flux in the hot layer also allows us to estimate the ionization parameter, ἕ=50 erg cm s, from which we get an albedo and find that the Comptonized flux incident on the disk heats it to KT,=0.11 keV, which looks like kT,=0.19 keV when spectral hardening is taken into account."," determines a temperature of $kT_e = $ 75 keV, which gives a photon index of $\Gamma = 1.87$ (3-10 The density of the disk and flux in the hot layer also allows us to estimate the ionization parameter, $\xi = 50$ erg cm $^{-1}$, from which we get an albedo and find that the Comptonized flux incident on the disk heats it to $kT_{e} = 0.11$ keV, which looks like $kT_{e} =
0.19$ keV when spectral hardening is taken into account."
" The predicted surface density in the hot inner ring also allows us to calculate the optical depth of this layer, which we find to be (r)=0.7."," The predicted surface density in the hot inner ring also allows us to calculate the optical depth of this layer, which we find to be $\langle\tau\rangle = 0.7$."
" Since we do not have a detailed model for the radiative transfer in this component, we take a plausible value of kT=200 keV, which gives a photon index of Γ=1.39 between 2 and 10 keV. The resulting spectrum is still too soft, so we set 4=0.89, meaning that of the disk and reflection photons are used to seed the hot layer, while the rest seed the hot inner ring."," Since we do not have a detailed model for the radiative transfer in this component, we take a plausible value of $kT = 200$ keV, which gives a photon index of $\Gamma = 1.39$ between 2 and 10 keV. The resulting spectrum is still too soft, so we set $\zeta =0.89$, meaning that of the disk and reflection photons are used to seed the hot layer, while the rest seed the hot inner ring."
 Figure 5 shows the resulting spectral energy distribution for SWIFT J1753.5-0127., Figure \ref{fig:SWIFTspec} shows the resulting spectral energy distribution for SWIFT J1753.5-0127.
 The different components shown in the figure are the same as in fig. 4.., The different components shown in the figure are the same as in fig. \ref{fig:TESTspec}.
 Figure 6 (taken from ?)) shows the spectrum divided by an absorped power law with an absorption column density of Ny=23x10?! cm?. Overlaid we show their best fit using the XSPEC “diskbb+pow” model (in red) and the soft excess predicted by our model (in green)., Figure \ref{fig:SWIFTcomp} (taken from \cite{2006ApJ...652L.113M}) ) shows the spectrum divided by an absorped power law with an absorption column density of $N_{\rm{H}} = 2.3\times 10^{21} $ $^{-2}.$ Overlaid we show their best fit using the XSPEC “diskbb+pow” model (in red) and the soft excess predicted by our model (in green).
" To obtain a better fit we change the absorption column density to Ny=2.35x10?! cm? to render the two excesses effectively indistinguishable, even though the temperature and peak fluxes of both disks are very different."," To obtain a better fit we change the absorption column density to $N_{\rm{H}} =
2.35\times 10^{21}$ $^{-2}$ to render the two excesses effectively indistinguishable, even though the temperature and peak fluxes of both disks are very different."
 We discuss the reasons for this at the beginning of sect. ??.., We discuss the reasons for this at the beginning of sect. \ref{sec:otherwork}.
" Figure 6 also shows that the deviation from a simple power law in the range 2-100 keV is less than10%,, which is consistent with observation."," Figure \ref{fig:SWIFTcomp} also shows that the deviation from a simple power law in the range 2-100 keV is less than, which is consistent with observation."
 Our second source for comparison is the relatively better constrained X-ray source GX 339-4., Our second source for comparison is the relatively better constrained X-ray source GX 339-4.
" This source has an estimated mass of 6Mo, distance of around 8 kpc and inclination yz= 0.9."," This source has an estimated mass of $M_\odot$, distance of around 8 kpc and inclination $\mu =
0.9$ ."
" ? took a 280ks observation of this source, which they observed at a flux of about Lx 0.05(M/10M.)(D/8kpc)."," \cite{2006ApJ...653..525M} took a 280ks observation of this source, which they observed at a flux of about $L_{\rm{X}}/L_{\rm{Edd}} \simeq 0.05
($ $/10$ $_\odot)(D/8$ $)$."
" They fit the data with a moderate absorption column density (Ny=3.72x10?! cm?) and very hard power law (T= 1.44) and find a soft excess below 3 keV which they fit to a disk with R;,= 0.6Rs and KT;,=0.38 keV. Additionally, they observe a broad asymmetric Fe-K line with a maximum at 6.9 keV, from which they measure a reflection fraction of about 0.2-0.3 and ionization parameter £~10° erg cm s."," They fit the data with a moderate absorption column density $N_{\rm{H}} = 3.72\times 10^{21}$ $^{-2}$ ) and very hard power law $\Gamma = 1.44$ ) and find a soft excess below 3 keV which they fit to a disk with $_{in} = 0.6$ $_{\rm{S}}$ and $kT_{in} = 0.38$ keV. Additionally, they observe a broad asymmetric Fe-K line with a maximum at 6.9 keV, from which they measure a reflection fraction of about 0.2-0.3 and ionization parameter $\xi
\sim 10^3 $ erg cm $^{-1}$."
 Fitting the Fe-K line using relativistic broadening suggests an inner radius of 0.7Rs., Fitting the Fe-K line using relativistic broadening suggests an inner radius of $_{\rm{S}}$ .
" For an assumed radius of the inner edge of the disk 19Rs, a good fit for this source is obtained with the following set of parameter values: C277 (or M/Meaa= 0.15), €=0.63, and the temperature of the hot layer andhot ring 75 keV, and 200 keV respectively."," For an assumed radius of the inner edge of the disk $R_{\rm in}
=19R_{\rm S}$ , a good fit for this source is obtained with the following set of parameter values: C=77 (or $\dot{M}/\dot{M}_{\rm Edd}
= 0.15$ ), $\zeta$ =0.63, and the temperature of the hot layer andhot ring 75 keV, and 200 keV respectively."
 The resulting spectral energy distribution is shown in fig. 7.., The resulting spectral energy distribution is shown in fig. \ref{fig:GX3394spec}. .
 Figure 8 shows the soft excess observed in, Figure \ref{fig:GX3394comp} shows the soft excess observed in
"We fit both the quadratic fnction of In(dt) aud (he transit curve to the data simultaneously using à Markov. Chain Monte Carlo method as described in relshort,orm.",We fit both the quadratic function of ln(dt) and the transit curve to the data simultaneously using a Markov Chain Monte Carlo method as described in \\ref{short_norm}.
.Asbefore.thedistribulionofvaluesiwascergeloselosimmelricinalleases.andtheredidnotap [ite," As before, the distribution of values was very close to symmetric in all cases, and there did not appear to be any strong correlations between variables."
clipsedeptlhsandlimesfromthesefilsaregiveninTablel..andlhetimeseriesbothbeforeandaflHerc," Best-fit eclipse depths and times from these fits are given in Table \ref{eclipse_depths}, and the time series both before and after correcting for detector effects are shown in Figures \ref{norm_plots} and \ref{four_eclipses}, respectively."
orreci from our quoted value.," As a check we repeated these fits using both a linear function in time and a linear function of ln(dt), and found that the eclipse depth in both cases varied by less than $\sigma$ from our quoted value."
 We ultimately. achieve noise levels in each of the four bandpasses that are 1.5. 1.9. 2.0. and 1.6 (mes higher than the predicted photon noise al 3.6. 4.5. 5.8. and 8.0jm. respectively.," We ultimately achieve noise levels in each of the four bandpasses that are 1.5, 1.9, 2.0, and 1.6 times higher than the predicted photon noise at 3.6, 4.5, 5.8, and 8.0, respectively."
 The additional noise is most likely the result of the jitter introduced by the need to repoint the telescope every 3.5 minutes in order (to switeli between bandpasses., The additional noise is most likely the result of the jitter introduced by the need to repoint the telescope every 3.5 minutes in order to switch between bandpasses.
 Although the eclipse is detected to a high degree of significance in all four bandpasses. it is worth noting that we ullimately achieve a precision al 4.5 and 8.0 comparable to that of the measured secondary eclipse depths for TrES-1 (Charbonneauetal. 2005).. even though this star is significantlv fainter (A&=9.82 vs A=6.32 for ID 209453).," Although the eclipse is detected to a high degree of significance in all four bandpasses, it is worth noting that we ultimately achieve a precision at 4.5 and 8.0 comparable to that of the measured secondary eclipse depths for TrES-1 \citep{char05}, , even though this star is significantly fainter $K=9.82$ vs $K=6.32$ for HD 209458)."
 This is partially explained by the reduced. cadence of the HD 209458 observations. which had of the total effective integration time per band for the TYES-1 observations. but the frequent repointings appear to have contributed additional noise as well.," This is partially explained by the reduced cadence of the HD 209458 observations, which had of the total effective integration time per band for the TrES-1 observations, but the frequent repointings appear to have contributed additional noise as well."
 We also note that Charbonneanetal.(2005) did not include the uncertainties contributed bv their fits to the trends in the out-of-transit data in their error estimates., We also note that \citet{char05} did not include the uncertainties contributed by their fits to the trends in the out-of-transit data in their error estimates.
 If we combine our estimates of the eclipse timing in each of the four bandpasses. we [ind that the center of the eclipse occurs at 2453702.525120.0012 ILJD.," If we combine our estimates of the eclipse timing in each of the four bandpasses, we find that the center of the eclipse occurs at $2453702.5251\pm0.0012$ HJD."
 This is 4941.7 minutes or 2.98 earlier than predicted. Gxuutsonetal.2007a).. and this is without. accounting for the additional 50 s delay in the predicted time due to the light travel (ime in (he svstem (Loeb2005).," This is $4.9\pm1.7$ minutes or $\sigma$ earlier than predicted \citep{knut07a}, and this is without accounting for the additional 50 s delay in the predicted time due to the light travel time in the system \citep{loeb05}."
. In (his case the uncertainty in the predicted time is negligible compared to the uncertainiw in our measurement., In this case the uncertainty in the predicted time is negligible compared to the uncertainty in our measurement.
 It is possible (hat we have under-estimated (he uncertainties in our estimates for the timing of the eclipse., It is possible that we have under-estimated the uncertainties in our estimates for the timing of the eclipse.
 Because the ingress and egress occur on relatively short (nme scales. fitted values for the eclipse times are particularly sensitive to the presence of correlated noise in (he (ime series. which is not included in (he error bars from (he Markov fits described above.," Because the ingress and egress occur on relatively short time scales, fitted values for the eclipse times are particularly sensitive to the presence of correlated noise in the time series, which is not included in the error bars from the Markov fits described above."
 To estimate (he effect of correlated noise on our best-fit eclipse depths and times. we use (he “praver-beacl” method as described in Gillonetal.(2007).," To estimate the effect of correlated noise on our best-fit eclipse depths and times, we use the “prayer-bead” method as described in \citet{gill07}."
. In Chis method we take (he time series of the residuals [rom our best-fit solution for each eclipse. shift the residuals forward in time with the points at the end of (he lime series wrapping back around (o the beginning. addthe best-fit solution back in. and fit the new timeseries for the full set of," In this method we take the time series of the residuals from our best-fit solution for each eclipse, shift the residuals forward in time with the points at the end of the time series wrapping back around to the beginning, addthe best-fit solution back in, and fit the new timeseries for the full set of"
In the meantime. a wealth of data has been accumulated.,"In the meantime, a wealth of data has been accumulated."
 It makes worthwhile the temporal and spatial features of soft. N-rav. (SAR) flaves on an extensive statistical basis. (1998), It makes worthwhile re-investigating the temporal and spatial features of soft X-ray (SXR) flares on an extensive statistical basis. )
) studied the distribution of tlie X-ray flares (M21) from 1987 to 1992 with respect to helio longitude., studied the distribution of the X-ray flares $\geq$ 1) from 1987 to 1992 with respect to helio longitude.
 They have shown that the fares were not uniformly distributed in longitude., They have shown that the flares were not uniformly distributed in longitude.
 The temporal analvsis of X-[lare statistics concerns basically the wailing-lime clis(vibution (see. for example. 1999:: 2001::13872001::likvinarkmainBoclyCilalionStart13902002:: and so on).," The temporal analysis of X-flare statistics concerns basically the waiting-time distribution (see, for example, ; ; and so on)."
 In the present analvsis we make use of SNR. flares observed by GOES cluring 1976-2006., In the present analysis we make use of SXR flares observed by GOES during 1976-2006.
 Our consideration will be devoted only to the energy statistics of soft. X-ray solar flares in time., Our consideration will be devoted only to the energy statistics of soft X-ray solar flares in time.
 Our analvsis is based on the properties of [racional autoregressive integrated moving average (FARIMA) processes 1994))., Our analysis is based on the properties of fractional autoregressive integrated moving average (FARIMA) processes ).
 Thev are widely used in modeling of various complex physical svstems., They are widely used in modeling of various complex physical systems.
" The FARIALA(p.d.qg) process is delined as the solution of the equation (B).NX(n)= O(B)e,. n€Z. where D is the shift operator BN(1)=X(n—1) and A is the difference operator. ee. NX(1)Αη)—Xin1)."," The $p$ $d$ $q$ ) process is defined as the solution of the equation $\Phi(B)\Delta^dX(n)=\Theta(B)\epsilon_n$ , $n\in{\bf Z}$, where $B$ is the shift operator $BX(n)=X(n-1)$ and $\Delta$ is the difference operator, e. $\Delta X(n)= X(n)-X(n-1)$."
 Here ® and © are the polynomials of degree p and q respectively. d takes Ilvactional values. either positive or negalive. and “innovations” e; are independent and identically distributed (1.0.)," Here $\Phi$ and $\Theta$ are the polynomials of degree $p$ and $q$ respectively, $d$ takes fractional values, either positive or negative, and “innovations” $\epsilon_j$ are independent and identically distributed (i.i.d.)"
 random variables., random variables.
 The polynomials ® and © correspond to autoregressive (AR) and moving average (ATA) parts. respectively.," The polynomials $\Phi$ and $\Theta$ correspond to autoregressive (AR) and moving average (MA) parts, respectively."
 The linear representation of FARIALA processes takes the form for details see 1994))., The linear representation of FARIMA processes takes the form for details see ).
 The innovations may be either Gaussian. non-Gaussian with finite variance or they may have infinite variance.," The innovations may be either Gaussian, non-Gaussian with finite variance or they may have infinite variance."
 For infinite variance innovations e.one may consider. for example. svaumnetric ancl skewed stable distributions. as well as Pareto distributions.," For infinite variance innovations $\epsilon$,one may consider, for example, symmetric and skewed stable distributions, as well as Pareto distributions."
 Both are characterized by the parameter a and (heir tails P(e>7) satisfy where F(7) denotes the corresponding distribution function aud ~ denotes that the ratio ol theleft-hand side to the right-hand one tends (ο 1. as à— o.," Both are characterized by the parameter $\alpha$ and their tails $P(\epsilon > x)$ satisfy where $F(x)$ denotes the corresponding distribution function and $\sim$ denotes that the ratio of theleft-hand side to the right-hand one tends to 1, as $x\to\infty$ ."
 H should be noted that the, It should be noted that the
"As a further development of the equations reported in the previous section for unobstructed mirrors, we hereafter provide the inverse formula of Eq. (14)).","As a further development of the equations reported in the previous section for unobstructed mirrors, we hereafter provide the inverse formula of Eq. \ref{eq:Aeff_alpha}) )."
" For a given A, from the effective area variation with 0 in [0,09], for 6=0, we derive the product of the reflectivities of the two segments (o1)where ri()ri(2),@=2a—a, (Eq. (5))."," For a given $\lambda$, from the effective area variation with $\theta$ in $[0, \alpha_0]$, for $\delta =0$, we derive the product of the reflectivities of the two segments where $\alpha_2 = 2\alpha_0-\alpha_1$ (Eq. \ref{eq:anglesum}) ))."
" Owing to the symmetry ofαι and o» with respect to the y-axis when 6= 0, it is"," Owing to the symmetry of$\alpha_1$ and $\alpha_2$ with respect to the $y$ -axis when $\delta =0$ , it is"
The Galactic microquasar SS 433 ts very luminous and unique in its continual ejection of plasma it two opposite jets at approximately one quarter the speed of light.,The Galactic microquasar SS 433 is very luminous and unique in its continual ejection of plasma in two opposite jets at approximately one quarter the speed of light.
 The system is a 13 day binary and probably powered by'! supercritcal accretion by the conpact member from its compauon., The system is a 13 day binary and probably powered by supercritcal accretion by the compact member from its companion.
 The orbital speed of the conpact object is fairly well established but in order to determine its mass either à measurement of the orbital velocity of the conpanion is needed or à meastre of the total mass of the system., The orbital speed of the compact object is fairly well established but in order to determine its mass either a measurement of the orbital velocity of the companion is needed or a measure of the total mass of the system.
 Stationary emission lines in the spectra of SS 433 display à persistent two horn structure of just the kind expected for emission from an orbiting ring. or disk. seen more or less edge on.," Stationary emission lines in the spectra of SS 433 display a persistent two horn structure of just the kind expected for emission from an orbiting ring, or disk, seen more or less edge on."
 The horn separation corresponds to a rotation speed in excess of 200 km s! and attributed to material orbiting the centre of mass of the binary system implies a system mass in excess of 40 M..., The horn separation corresponds to a rotation speed in excess of 200 km $^{-1}$ and attributed to material orbiting the centre of mass of the binary system implies a system mass in excess of 40 $M_\odot$ .
 The Ha spectra were originally discussed in Blundell et al (2008)., The $\alpha$ spectra were originally discussed in Blundell et al (2008).
 Departures from the pattern expected for a uniformly radiating ring are present and are more pronounced in He I emission lines., Departures from the pattern expected for a uniformly radiating ring are present and are more pronounced in He I emission lines.
 These departures have been explained in terms of emission stimulated by some kind of spotlight rotating with the binary (Bowler 2010): propinquity of the intense radiation source in the accretion disk is sufficient (Bowler 2011).Thus these observations fix rather well the mass of the SS 433 system (hence of the compact object). provided that the two horned structure in Hc and He [is indeed produced in an orbiting cireumbinary ring rather than 1n ejecta on escape trajectories or some radially expanding structure.," These departures have been explained in terms of emission stimulated by some kind of spotlight rotating with the binary (Bowler 2010); propinquity of the intense radiation source in the accretion disk is sufficient (Bowler 2011).Thus these observations fix rather well the mass of the SS 433 system (hence of the compact object), provided that the two horned structure in $\alpha$ and He I is indeed produced in an orbiting circumbinary ring rather than in ejecta on escape trajectories or some radially expanding structure."
 Origin in the inner rim of a circumbinary disk seems more likely to produce persistent and stable phenomena than the accretion disk blowing bubbles., Origin in the inner rim of a circumbinary disk seems more likely to produce persistent and stable phenomena than the accretion disk blowing bubbles.
 The stability and persistence of the two horned structure is therefore of some significance., The stability and persistence of the two horned structure is therefore of some significance.
 In this note I compare the SS 433 stationary Πα spectra shown in Fig., In this note I compare the SS 433 stationary $\alpha$ spectra shown in Fig.
 2 of Li Yan (2010) with the spectra in Fig., 2 of Li Yan (2010) with the spectra in Fig.
 2 of Schmidtobreick Blundell (2006a). as analysed in. Bluncell et al (2008) and Bowler (2010).," 2 of Schmidtobreick Blundell (2006a), as analysed in Blundell et al (2008) and Bowler (2010)."
 These two sets of spectra were taken under very ditferent conditions but are so similar in structure that it is my optnion that they show the same processes and that these are stable over a long period., These two sets of spectra were taken under very different conditions but are so similar in structure that it is my opinion that they show the same processes and that these are stable over a long period.
 The data of Li Yan (Purple Mourtal Observatory) were taken in the years 2004. 2007 ard 2008 and there are 17 spectra in all: labelled in that paper only accordiσ to the orbital phase.," The data of Li Yan (Purple Mountain Observatory) were taken in the years 2004, 2007 and 2008 and there are 17 spectra in all; labelled in that paper only according to the orbital phase."
 There is no other information provided except that the precessional phases lay in the range 0.155 and 0.296.where precessional," There is no other information provided except that the precessional phases lay in the range 0.155 and 0.296,where precessional"
"are given by à aud ve. respectively,","are given by $u$ and $v$, respectively."
" The (third term in the equation gives (he contribution from the planets rotation. where A, is the planets radius at 1: bar (=1.32 A, for all models). zis the heieht in (he atmosphere above the 1-bar level. and Q is the planets rotational speed in radians |."," The third term in the equation gives the contribution from the planet's rotation, where $R_p$ is the planet's radius at 1 bar $=1.32$ $R_{Jup}$ for all models), z is the height in the atmosphere above the 1-bar level, and $\Omega$ is the planet's rotational speed in radians $^{-1}$."
 The final term in Equation 4. gives the contribution from the orbital motion of the planet., The final term in Equation \ref{vlos} gives the contribution from the orbital motion of the planet.
 The orbital speed is denoted bv Όρμν. and 45 is the phase angle of the orbit. defined as ο=0 at the center of transit.," The orbital speed is denoted by $v_{orb}$, and $\varphi$ is the phase angle of the orbit, defined as $\varphi = 0$ at the center of transit."
 For the results presented in this paper we asstune a circular orbit resulting in a constant (4., For the results presented in this paper we assume a circular orbit resulting in a constant $v_{orb}$.
" To calculate the rotation ancl orbital speeds we assume a lidally locked planet (D,=D,,) wilh an orbital period of 3.53 clays.", To calculate the rotation and orbital speeds we assume a tidally locked planet $P_{orb} = P_{rot}$ ) with an orbital period of 3.53 days.
 Figure 4. shows the Doppler-shiftecl wind speeds along the planets terminator that result from each of the four atinospherie dynamic models described in Section 2.1.., Figure \ref{fig:dshift} shows the Doppler-shifted wind speeds along the planet's terminator that result from each of the four atmospheric dynamic models described in Section \ref{model}.
 The opacity & from Equation 2 is finally evaluated at wavelength where Ay is (he unshifted wavelength. to produce the properly Doppler-shifted absorption.," The opacity $\kappa$ from Equation \ref{tau} is finally evaluated at wavelength where $\lambda_0$ is the unshifted wavelength, to produce the properly Doppler-shifted absorption."
 Ligh resolution spectroscopy al a spectral resolution of ~10° along with sufficiently high signal-(o-noise is necessary to measure Doppler shifts in transmission spectra at the kins | level., High resolution spectroscopy at a spectral resolution of $\sim$ $^5$ along with sufficiently high signal-to-noise is necessary to measure Doppler shifts in transmission spectra at the km $^{-1}$ level.
 [ere we caleulate all of our transmission spectra αἱ a spectral resolution of 10°. which is higher than the resolution of most currently available high-resolution spectrographs by al least an order of magnitude. (," Here we calculate all of our transmission spectra at a spectral resolution of $10^6$, which is higher than the resolution of most currently available high-resolution spectrographs by at least an order of magnitude. ("
For example. the CRIRES spectrograph used by ? has a working resolution of ~10°).,"For example, the CRIRES spectrograph used by \citet{sne10} has a working resolution of $\sim$ $^5$ )."
 We compute spectra at such high resolution to clearly show the effects of Doppler shifts on our transmission spectra and also to show the power of very high spectral resolution for future instrumentation., We compute spectra at such high resolution to clearly show the effects of Doppler shifts on our transmission spectra and also to show the power of very high spectral resolution for future instrumentation.
 All of our spectra can be easily degraded to lower spectral resolution by convolution wilh a Gaussian of the appropriate width., All of our spectra can be easily degraded to lower spectral resolution by convolution with a Gaussian of the appropriate width.
 We calculate all of our spectra from 2291 (o 2350 nm as a representative wavelength range over which high resolution (ransnussion spectra can be obtained from the ground., We calculate all of our spectra from 2291 to 2350 nm as a representative wavelength range over which high resolution transmission spectra can be obtained from the ground.
 This wavelength range covers the 2.3 san first overtone (Av= 2) band of CO., This wavelength range covers the 2.3 $\mu$ m first overtone $\Delta \nu = 2$ ) band of CO.
 This is also the wavelength: coverage of the observations by ?.. which facilitates comparisons between our model spectra and (heir results.," This is also the wavelength coverage of the observations by \citet{sne10}, which facilitates comparisons between our model spectra and their results."
 It is important (o note that the stellar spectrum experiences velocity shifts of its own during transit., It is important to note that the stellar spectrum experiences velocity shifts of its own during transit.
 These shifts result from both the induced motion from the planets orbit (stellar radial velocity) along with the Bossiter-McLauehlin effect by which the planet blocks out a portion of the blue- or red-shifted limb of the star during transit., These shifts result from both the induced motion from the planet's orbit (stellar radial velocity) along with the Rossiter-McLaughlin effect by which the planet blocks out a portion of the blue- or red-shifted limb of the star during transit.
 Both effects produce a zero net Doppler shift at the center of transit. provided that the planets orbit is circular and (he stellar spin axis is aligned with the normal axis of the planet's orbit.," Both effects produce a zero net Doppler shift at the center of transit, provided that the planet's orbit is circular and the stellar spin axis is aligned with the normal axis of the planet's orbit."
 However. when either of these effects becomes non-zero. il is necessary (o divide out the appropriately stellar spectrum from the in-transit spectrum to obtain a transmission spectrum that," However, when either of these effects becomes non-zero, it is necessary to divide out the appropriately Doppler-shifted stellar spectrum from the in-transit spectrum to obtain a transmission spectrum that"
cluster.,cluster.
" It is encouraging that numerous different methods of studying the contributions of LSBs to (he optical huninosity density and O,, vield consistent results. both for the local Universe and for clusters."," It is encouraging that numerous different methods of studying the contributions of LSBs to the optical luminosity density and $\Omega_m$ yield consistent results, both for the local Universe and for clusters."
 These results all point to the conclusion that LSBs do not account [or a significant amount of matter., These results all point to the conclusion that LSBs do not account for a significant amount of matter.
 We would like to thank the ROTSE team. with special thanks to Tim Mcelxay. and Carl AkerloL. for allowing us access to their skvpatrol data.," We would like to thank the ROTSE team, with special thanks to Tim McKay and Carl Akerlof, for allowing us access to their skypatrol data."
 ROTSE is a collaboration of Lawrence Livermore National Lab. Los Alamos National Lab. and the University of Michigan umichliedu/-rotse).," ROTSE is a collaboration of Lawrence Livermore National Lab, Los Alamos National Lab, and the University of Michigan $\sim$ rotse)."
" This research has made use of the SIAIBAD database. operated at CDS. Strasbourg. France. the $TSclI Digitized Sky Survey (DSS) form). the International Astronomical Union's ""List of Supernovae"" (http://efa-harvarcedu/iat/lists/Supernovae.htiml). operated by the Central Dureau for Astronomical Telegrams (CBAT). and the NASA/IPAC Extragalactic Database (NED). which is operated bv the Jet. Propulsion Laboratory. California Institute of Technology. under contract with ihe National Aeronauties and Space Achuinistration."," This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France, the STScI Digitized Sky Survey (DSS) form), the International Astronomical Union's “List of Supernovae” (http://cfa-www.harvard.edu/iau/lists/Supernovae.html), operated by the Central Bureau for Astronomical Telegrams (CBAT), and the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
istance iuto the tenuous trausitional dust disk. making Roo>Ray.,"distance into the tenuous transitional dust disk, making $R_\mathrm{CO}>R_\mathrm{sub}$."
 A possible quantitative. test for this hypothesis would be to calculate a UV. penetration -epth anc compare to the radius discrepaucy., A possible quantitative test for this hypothesis would be to calculate a UV penetration depth and compare to the radius discrepancy.
" However. --lis is difficult to test im practice. since the local dust -eusitv at AJ, is not known. ("," However, this is difficult to test in practice, since the local dust density at $R_\mathrm{sub}$ is not known. ("
Although SED models -ui place some coustraints on the location aud surface enusitv of the iuner disk dust. they do not do so to the precision required here.),"Although SED models can place some constraints on the location and surface density of the inner disk dust, they do not do so to the precision required here.)"
" Yet. with sufficieutle low dust deusities. this explanation is plausible: for example. witl x~.10!cn,ο and a dust density p=10οσα? (equivalent to a vertical dust surface density of ~106 [m]s cm? for a scae height of 0.1. AU). the penetration depth (to 7= 1) is ~7 AU."," Yet, with sufficiently low dust densities, this explanation is plausible; for example, with $\kappa\sim10^4\ \mathrm{cm}^2\ \mathrm{g}^{-1}$ and a dust density $\rho=10^{-18}\ \mathrm{g\ cm}^{-3}$ (equivalent to a vertical dust surface density of $\sim10^{-6}\ $ g $^{-2}$ for a scale height of 0.1 AU), the penetration depth (to $\tau=1$ ) is $\sim7$ AU."
 For comparison. Eisner.Cliang.&Illebran(2006) measure a dust surface density of 6.3.«LO ogo» - at the inner edee of the TW Iva disk.," For comparison, \citet{Eisner06} measure a dust surface density of $6.3\times 10^{-7}$ g $^{-2}$ at the inner edge of the TW Hya disk."
 A final possibility we consider is that the CO, A final possibility we consider is that the CO
"set to 100x the ambient pressure, and with a uniform magnetic field at 45? to the grid axes.","set to $100\times$ the ambient pressure, and with a uniform magnetic field at $45^{\circ}$ to the grid axes."
 Very similar results were found for the case where the initial plasma beta parameter 8=0.2., Very similar results were found for the case where the initial plasma beta parameter $\beta=0.2$.
" Because of the robustness of the ? algorithm, and presumably also because the integration algorithm used is slightly more diffusive than the algorithm, our code can also simulate this test problem with a field 10x stronger (8= 0.002), shown in Fig. A4.."," Because of the robustness of the \citet{DedKemKroEA02} algorithm, and presumably also because the integration algorithm used is slightly more diffusive than the algorithm, our code can also simulate this test problem with a field $10\times$ stronger $\beta=0.002$ ), shown in Fig. \ref{fig:SG09BW_Beta0p002}."
 This extra robustness is important for the strong field simulations run here., This extra robustness is important for the strong field simulations run here.
" ? suggest that for zero-gradient boundary conditions (BCs), using zero-gradient for the unphysical scalar field ᾧ is sufficient."," \citet{DedKemKroEA02} suggest that for zero-gradient boundary conditions (BCs), using zero-gradient for the unphysical scalar field $\psi$ is sufficient."
" We have found, however, that more care is needed for magnetically dominated regions (plasma parameter ϐ=8xpg/B?« 1) when gas is flowing subsonically near the boundary."," We have found, however, that more care is needed for magnetically dominated regions (plasma parameter $\beta\equiv8\pi
p_g/B^2\ll1$ ) when gas is flowing subsonically near the boundary."
" In what follows X is the normal vector to the boundary, the computational domain is ‘left’ of the boundary, and off the domain is ‘right’."," In what follows $\hat{\mathbf{x}}$ is the normal vector to the boundary, the computational domain is `left' of the boundary, and off the domain is `right'."
 Note that « is not a physical variable so its gradient is not constrained., Note that $\psi$ is not a physical variable so its gradient is not constrained.
" In fact if the field is constant across the boundary then the formal solution to the 1D Riemann problem is that the flux, F[B.]=0 (as it is for all 1D Riemann problems)."," In fact if the field is constant across the boundary then the formal solution to the 1D Riemann problem is that the flux, $F[B_x]=0$ (as it is for all 1D Riemann problems)."
" So ideally a boundary condition for w should be selected which ensures F[B;,|=0.", So ideally a boundary condition for $\psi$ should be selected which ensures $F[B_x]=0$.
" From ?,eqs.41,42 it is easy to calculate this: the zero-gradient condition states that Br—Bl=0 so if additionally y=—4/! then F[B,]=0 is obtained."," From \citet[][eqs.~41,42]{DedKemKroEA02} it is easy to calculate this: the zero-gradient condition states that $B_x^r-B_x^l=0$ so if additionally $\psi^r=-\psi^l$ then $F[B_x]=0$ is obtained."
 A sequence of figures from a 2D slab-symmetric test simulation (this time using uunits) are shown in Fig., A sequence of figures from a 2D slab-symmetric test simulation (this time using units) are shown in Fig.
 Bl to demonstrate the effectiveness of the new boundary condition., \ref{fig:TestBC_PG_c15} to demonstrate the effectiveness of the new boundary condition.
" The simulation domains used have sizes [0.2,0.1] pc (large) and [0.1,0.1] pc (small) and the radiation source is at infinity in the —x direction."," The simulation domains used have sizes $[0.2,0.1]\,$ pc (large) and $[0.1,0.1]\,$ pc (small) and the radiation source is at infinity in the $-\hat{\mathbf{x}}$ direction."
 A group of randomly located dense clumps are placed between z—0.02 pc and z—0.08 pc on a uniform background density of ng=100cmὃ.," A group of randomly located dense clumps are placed between $x=0.02\,$ pc and $x=0.08\,$ pc on a uniform background density of $n_{\mathrm{H}}=100\,\mathrm{cm}^{-3}$."
" The initial magnetic field is B= μα, and the initially constant gas pressure is set to give ϐ=0.017 (T'= 1500K in the lowest density gas)."," The initial magnetic field is $\mathbf{B} = [50,1,0]\times \sqrt{4\pi}\,\mu$ G, and the initially constant gas pressure is set to give $\beta=0.017$ $T=1500$ K in the lowest density gas)."
" Zero gradient BCs are enforced in the +x directions, and periodic in xy."," Zero gradient BCs are enforced in the $\pm\hat{\mathbf{x}}$ directions, and periodic in $\pm\hat{\mathbf{y}}$."
 Fig., Fig.
 Bl shows the gas pressure at time t=12kyr using cooling model C2.," \ref{fig:TestBC_PG_c15} shows the gas pressure at time $t=12\,$ kyr using cooling model C2."
" The problem with the zero-gradient BC for v is very apparent, as is the dramatic improvement on switching to the V=—31/ BC."," The problem with the zero-gradient BC for $\psi$ is very apparent, as is the dramatic improvement on switching to the $\psi^r=-\psi^l$ BC."
 The black regions are hiding cells which obtained negative pressures and were set to an artificial pressure floor value in order to allow the simulation to continue., The black regions are hiding cells which obtained negative pressures and were set to an artificial pressure floor value in order to allow the simulation to continue.
" Regarding the other boundaries, the x:= 0.2pc boundary only has waves moving perpendicular to it and so there is negligible net flow across the boundary."," Regarding the other boundaries, the $x=0.2\,$ pc boundary only has waves moving perpendicular to it and so there is negligible net flow across the boundary."
" The z—0 pc boundary has strong supersonic outflows, so the details of the BC for w have little effect on it. 2)). C1.. C2.tab:clump,rops."," The $x=0\,$ pc boundary has strong supersonic outflows, so the details of the BC for $\psi$ have little effect on it. \ref{tab:rmhd_Fields}) \ref{tab:M17mhd_BoundaryTest}. \ref{fig:MHDboundary_comp}."
Solar eruptions. including flares. filament eruptions. and coronal mass ejections (CMEs) have been understood as the result of magnetic reconnection in the solar corona (e.g.. Kopp Pneuman 1976; Antiochos et al.,"Solar eruptions, including flares, filament eruptions, and coronal mass ejections (CMEs) have been understood as the result of magnetic reconnection in the solar corona (e.g., Kopp Pneuman 1976; Antiochos et al."
" 1999),", 1999).
 Although surface magnetic field evolution (such as new flux emergence and shear motion) play important roles in building energy and triggering eruption. most models of flares and CMEs have the implication that photospheric magnetic fields do not have rapid. irreversible changes associated with the eruptions.," Although surface magnetic field evolution (such as new flux emergence and shear motion) play important roles in building energy and triggering eruption, most models of flares and CMEs have the implication that photospheric magnetic fields do not have rapid, irreversible changes associated with the eruptions."
 The key reason behind this assumption is that the solar surface. where the coronal magnetic fields are anchored. has much higher density and gas pressure than the corona.," The key reason behind this assumption is that the solar surface, where the coronal magnetic fields are anchored, has much higher density and gas pressure than the corona."
 Recently. we note the work by Hudson. Fisher Welsch (2008. hereafter HFW08). who quantitatively assessed the back reaction on the solar surface and interior resulting from the coronal field evolution required to release energy. and made the prediction that after flares. the photospheric magnetic fields become more horizontal.," Recently, we note the work by Hudson, Fisher Welsch (2008, hereafter HFW08), who quantitatively assessed the back reaction on the solar surface and interior resulting from the coronal field evolution required to release energy, and made the prediction that after flares, the photospheric magnetic fields become more horizontal."
 Their analysis is based on the simple principle that any change of magnetic field energy must lead to a corresponding change in magnetic pressure., Their analysis is based on the simple principle that any change of magnetic field energy must lead to a corresponding change in magnetic pressure.
 This 1s one of the very few models that specifically predict that flares can be accompanied by rapid and irreversible changes of photospheric magnetic fields., This is one of the very few models that specifically predict that flares can be accompanied by rapid and irreversible changes of photospheric magnetic fields.
 Over a decade ago. Melrose (1997) used the concept of reconnection between two current-carrying systems to provide explanation for the enhancement of magnetic shear at the flaring magnetic polarity inverstor line (PIL). which 15 sometimes observed (see below).," Over a decade ago, Melrose (1997) used the concept of reconnection between two current-carrying systems to provide explanation for the enhancement of magnetic shear at the flaring magnetic polarity inversion line (PIL), which is sometimes observed (see below)."
 Perhaps this is in the same line as the tether-cutting model for sigmoids. which was elaborated by Moore et al. (," Perhaps this is in the same line as the tether-cutting model for sigmoids, which was elaborated by Moore et al. ("
2001) anc involves a two-stage reconnection processes.,2001) and involves a two-stage reconnection processes.
 At the eruptior onset. the near-surface reconnection between the two sigmoic elbows produces a low-lying shorter loop across the PIL and a larger twisted flux rope connecting the two far ends of the sigmoid.," At the eruption onset, the near-surface reconnection between the two sigmoid elbows produces a low-lying shorter loop across the PIL and a larger twisted flux rope connecting the two far ends of the sigmoid."
 The second stage reconnection occurs whe the large-scale loop cuts through the areade fields causing erupting plasmoid to become CME and precipitation of electrons to form flare ribbons., The second stage reconnection occurs when the large-scale loop cuts through the arcade fields causing erupting plasmoid to become CME and precipitation of electrons to form flare ribbons.
 If scrutinizing the magnetic topology close to the surface. one would find a permanent change of magnetic fields that conforms to the scenario of HFW208: the magnetic fields turn more horizontal near the flaring PIL due to the newly formed short loops there.," If scrutinizing the magnetic topology close to the surface, one would find a permanent change of magnetic fields that conforms to the scenario of HFW08: the magnetic fields turn more horizontal near the flaring PIL due to the newly formed short loops there."
 On the observational side. earlier studies were inconclusive on the flare-related changes of photospheric magnetic field topology.," On the observational side, earlier studies were inconclusive on the flare-related changes of photospheric magnetic field topology."
 Wang (1992) and Wang et al. (, Wang (1992) and Wang et al. (
1994) showed impulsive changes of vector fields after flares including some unexpected patterns such as increase of magnetic shear along the PIL. while mixed results were also reported (Ambastha et al.,"1994) showed impulsive changes of vector fields after flares including some unexpected patterns such as increase of magnetic shear along the PIL, while mixed results were also reported (Ambastha et al."
 1993; Hagyard et al., 1993; Hagyard et al.
 1999: Chen et al., 1999; Chen et al.
 1994. Li et al.," 1994, Li et al."
 2000a. 2000b).," 2000a, 2000b)."
 It is not until recently that rapid and permanent changes of photospheric magnetic fields. mainly the line-of-sight component. are observed to be consistently appear in major flares and considered as indicative of flare energy release (Kosovichev Zharkova 2001; Sudol Harvey 2005).," It is not until recently that rapid and permanent changes of photospheric magnetic fields, mainly the line-of-sight component, are observed to be consistently appear in major flares and considered as indicative of flare energy release (Kosovichev Zharkova 2001; Sudol Harvey 2005)."
 In particular. a number of papers of Big Bear Solar Observatory (BBSO)/New Jersey Institute of Technology group have been devoted to the finding of sudden unbalanced magnetic flux change (Spirock et al.," In particular, a number of papers of Big Bear Solar Observatory (BBSO)/New Jersey Institute of Technology group have been devoted to the finding of sudden unbalanced magnetic flux change (Spirock et al."
 2002; Wang et al., 2002; Wang et al.
 2002: Yurchyshyn et al., 2002; Yurchyshyn et al.
 2004: Wang et al., 2004; Wang et al.
 2004a: Wang 2006) and a new phenomenon of sunspot white-light structure change (Wang et al., 2004a; Wang 2006) and a new phenomenon of sunspot white-light structure change (Wang et al.
 2004b: Deng et al., 2004b; Deng et al.
 2005: Liu et al., 2005; Liu et al.
 2005: Chen et al., 2005; Chen et al.
 2006) associated with flares., 2006) associated with flares.
 By evaluating the mean relative motions between two magnetic polarities of five flaring ó spots. Wang (2006) revealed a sudden release of the overall magnetic shear and found a correspondence between converging/diverging motion and merease/decrease of magnetic gradient at the PIL. which suggests magnetic reconnection at or close to the photosphere.," By evaluating the mean relative motions between two magnetic polarities of five flaring $\delta$ spots, Wang (2006) revealed a sudden release of the overall magnetic shear and found a correspondence between converging/diverging motion and increase/decrease of magnetic gradient at the PIL, which suggests magnetic reconnection at or close to the photosphere."
 Liu et al. (, Liu et al. (
2005) discussed. outer penumbral decay and central penumbral darkening of 6 spot seen in seven X-class flares. and proposed,"2005) discussed outer penumbral decay and central penumbral darkening of $\delta$ spot seen in seven X-class flares, and proposed"
In this section we rederive the conservation equations for cosmic ταν mocified shocks in (heir eeneral Form. in order to emphasize the mathematical origin of (he escape flux.,"In this section we rederive the conservation equations for cosmic ray modified shocks in their general form, in order to emphasize the mathematical origin of the escape flux."
 The time dependent conservation equations in (he presence of accelerated particles at a shock ean be written in the following form: Here 2. DL. and Dy ave respectively the gas pressure. the cosmic ray pressure and the pressure in (he form of waves.," The time dependent conservation equations in the presence of accelerated particles at a shock can be written in the following form: Here $P_g$ , $P_c$ and $P_W$ are respectively the gas pressure, the cosmic ray pressure and the pressure in the form of waves."
 £y: is the energy density in (he form of waves and D is the rate al which the background. plasma is heated due to the damping of waves onto the plasma., $E_W$ is the energy density in the form of waves and $\Gamma$ is the rate at which the background plasma is heated due to the damping of waves onto the plasma.
 The rate of change of (he gas temperature is related to P£y: through: The cosmic ray. pressure can be calculated from the transport equation: where we put (Cr)=uC)—eyCr) and eyCr) is Che wave velocity., The rate of change of the gas temperature is related to $\Gamma E_W$ through: The cosmic ray pressure can be calculated from the transport equation: where we put $\tilde u(x) = u(x) - v_W(x)$ and $v_W(x)$ is the wave velocity.
 For our purposes here we are neglecting the injection term., For our purposes here we are neglecting the injection term.
 Multiplying this equation by the kinetic energy T(p)=mjc(5—1). where 5 is the Lorentz [actor of a particle with momentum p. and integrating (he transport equation in nmonmentunm. one has: where are (he enerey density and pressure in (he form of accelerated. particles.," Multiplying this equation by the kinetic energy $T(p)=m_p c^2
(\gamma-1)$, where $\gamma$ is the Lorentz factor of a particle with momentum $p$, and integrating the transport equation in momentum, one has: where are the energy density and pressure in the form of accelerated particles."
 Moreover we introduced the mean diffusion coefficient: The only assumption that we made here is that /(p)—0 for p—x., Moreover we introduced the mean diffusion coefficient: The only assumption that we made here is that $f(p)\to 0$ for $p\to \infty$.
" We shall conmnent later (84)) on whatwould happen il the spectrum were (truncated al some fixed p,,,,. instead."," We shall comment later\ref{sec:escape}) ) on whatwould happen if the spectrum were truncated at some fixed $p_{max}$ , instead."
But there is no analytical form for two-ldy motion in General Relativity. so we will use Newtonian otvsies to describe the motion of the interacting back es. (,"But there is no analytical form for two-body motion in General Relativity, so we will use Newtonian physics to describe the motion of the interacting black holes. ("
One effect of this choice is to confine hie motion to the initial orbital plane. =0.),"One effect of this choice is to confine the motion to the initial orbital plane, $z=0$ .)"
 ThIs OL 'esults will be accurate but. uninteresting Oo| Newonian motions (slow velocities. la‘oe impacl ‘ameters or orbital separations).," Thus our results will be accurate but uninteresting for Newtonian motions (slow velocities, large impact parameters or orbital separations)."
 For he iuerestii case of relativistic interactOLS (velociles near c. stall impact. parameters or orbital seyaratious) we will obtain qualitative resuts. which jetheless lead to estimates of kick ratios Or clilfe‘ent configurations aud allow ali analytical 1erstaudiug of the kick process.," For the interesting case of relativistic interactions (velocities near $c$, small impact parameters or orbital separations) we will obtain qualitative results, which nonetheless lead to estimates of kick ratios for different configurations and allow an analytical understanding of the kick process."
 For he aasicircular Case Our work is completuentary to that o. Schuittinanetal.(2008).. which is au extensive study of non-edlal miass nuergers with. eiher zero or equal but opposite spins: we sttονmass mergers with arbitrary raticys ancl directions of spins on the two black holes.," For the quasicircular case our work is complementary to that of \cite{2008PhRvD..77d4031}, which is an extensive study of non-equal mass mergers with either zero or equal but opposite spins; we study mergers with arbitrary ratios and directions of spins on the two black holes."
 The staudard expressions for the mass quadrupoles ancl octupoles are: wlere the stun is over both black 1jioles., The standard expressions for the mass quadrupoles and octupoles are: where the sum is over both black holes.
 We 1se Latiu letters lor spatial iudices. i.3.z. and do uot distinguish between covariaut and «'ontravariaut indices.," We use Latin letters for spatial indices, $x,y,z$, and do not distinguish between covariant and contravariant indices."
 We treat one black ile at a time., We treat one black hole at a time.
 The two-hole result is obtained by adding tie contribution from each hole., The two-hole result is obtained by adding the contribution from each hole.
 Because in our examples the black laoles have equal mass. tlie iiass «x'tupole vanishes on this addition. so that the third term in the foruula Eq (1)) can be ignored regardless of spin coufiguratiou.," Because in our examples the black holes have equal mass, the mass octupole vanishes on this addition, so that the third term in the formula Eq \ref{kickEQ}) ) can be ignored regardless of spin configuration."
" Since we —eed the third time derivative of the mass quadrupole. we evaluate or,(/) ancl its first tliree time «erlvatives."," Since we need the third time derivative of the mass quadrupole, we evaluate $x_k(t)$ and its first three time derivatives."
 The spin multipoles require the b:dIliug coucept of spin deusity. but since we are dealing with ]xerr black holes. which have an iutriusic spin dipole moment. we use a trick to evaluate the spin quadrupole.," The spin multipoles require the baffling concept of spin density, but since we are dealing with Kerr black holes, which have an intrinsic spin dipole moment, we use a trick to evaluate the spin quadrupole."
 We replace tle spin dipole by a fictitious pair of spin charges. of value g=ze separated by a distance m. centered at the actuil location of the black hole (Herrmannetal.2007a)..," We replace the spin dipole by a fictitious pair of spin charges, of value $q=\pm a$ separated by a distance m, centered at the actual location of the black hole \citep{2007ApJ...661..430H}."
 This reproduces the dipole augular momentun (ein). aud allows us to compute the quadrupole directly: Now the suia is over the two blacς holes aud over the two spin charges lor each hole. aud the zr; appearing in these formula are ollse in the clirection of the spin.," This reproduces the dipole angular momentum $am$ ), and allows us to compute the quadrupole directly: Now the sum is over the two black holes and over the two spin charges for each hole, and the $x_i $ appearing in these formula are offset in the direction of the spin."
 Strictly oue should take a limit, Strictly one should take a limit
holes are likely to be intimately related (Normencdy&Richstone1995:Magorrianοἱal.1993:Gebhardtetal.2000:Ferrarese&Merritt 2000).. the most powerful active nuclei will almost certainly be found in massive galaxies.,"holes are likely to be intimately related \citep{kor95,mag98,geb00,fer00}, the most powerful active nuclei will almost certainly be found in massive galaxies."
 At high redshifts. such galaxies are expected to have formed. almost exclusively in regions of high overdensitv. in which the processes οἱ ealaxv formation will have proceeded most rapidly. and which will also be likely seedbecs for protoclusters.," At high redshifts, such galaxies are expected to have formed almost exclusively in regions of high overdensity, in which the processes of galaxy formation will have proceeded most rapidly, and which will also be likely seedbeds for protoclusters."
 Recent observational evidence supports this picture (see Best2000. [or a review οἱ earlier results)., Recent observational evidence supports this picture (see \citealt{bes00} for a review of earlier results).
 In particular. although high-redshift radio sources are found in a variety of environments. evidence for clustering around many of (iem has been found in both optical/Ih (e.g..Bestetal.2003) and. X-ray (Pentericcietal.2002). survevs.," In particular, although high-redshift radio sources are found in a variety of environments, evidence for clustering around many of them has been found in both optical/IR \citep[\eg][]{bes03} and X-ray \citep{pen02} surveys."
 In some cases. (here are indications that the radio source is the central galaxy in the cluster. either [rom an overdensitv of galaxies al small scales around it (6.9..Bestοἱal.2003).. or [vomthe surface- profile of (he galaxy itself (Best.Longair.&ROUeering1998).," In some cases, there are indications that the radio source is the central galaxy in the cluster, either from an overdensity of galaxies at small scales around it \citep[\eg][]{bes03}, or fromthe surface-brightness profile of the galaxy itself \citep{bes98}."
. This means that the study of tbe morphologies and other properties of high-redshift radio galaxies Chat appear to be the dominant members of clusters or protoclusters can give us insights into the formation of the most massive galaxies in the present-day universe., This means that the study of the morphologies and other properties of high-redshift radio galaxies that appear to be the dominant members of clusters or protoclusters can give us insights into the formation of the most massive galaxies in the present-day universe.
 One such case is the 2=1.1736 radio galaxy 2294. which Τοetal.(2003). have found to be surrounded. by an overclensity of faint red. galaxies.," One such case is the $z=1.786$ radio galaxy 294, which \citet{tof03} have found to be surrounded by an overdensity of faint red galaxies."
" This object has been a favorite target for adaptive oplics (AQ) svstems on large telescopes both because it is one of the most powerful radio galaxies in (he observable universe and because ils optical center lies «10"" [from a 12th mag star.", This object has been a favorite target for adaptive optics (AO) systems on large telescopes both because it is one of the most powerful radio galaxies in the observable universe and because its optical center lies $<10$ from a 12th mag star.
" The first AO imaging of this object. using the University of Hawaii eurvature-sensing AO svstem on the Canacla-France-Lawaii Telescope (CFIUT). was reported by Stockton.Canalizo.&Rideway(1999).. who found a chlumpy structure in the A"" band and suggested that the various clamps might be clisty subunits in the process of merging and illuminated by a hidden nucleus to the south of most ofthe observed structure."," The first AO imaging of this object, using the University of Hawaii curvature-sensing AO system on the Canada-France-Hawaii Telescope (CFHT), was reported by \citet{sto99}, who found a clumpy structure in the $K'$ band and suggested that the various clumps might be dusty subunits in the process of merging and illuminated by a hidden nucleus to the south of most ofthe observed structure."
 Thev also emphasized (he uncertainty in the position in the nucleus., They also emphasized the uncertainty in the position in the nucleus.
" Quirrenbachetal.(2001) used the AO svstem on the Keck II telescope to observe 2294 in the HW and A"" bands.", \citet{qui01} used the AO system on the Keck II telescope to observe 294 in the $H$ and $K'$ bands.
" They found an essentially stellar profile (<50 mas FWHAD for the eastern component. separated by ~1"" from the diffuse western component."," They found an essentially stellar profile $<50$ mas FWHM) for the eastern component, separated by $\sim1\arcsec$ from the diffuse western component."
 Taking ihe USNO-A2.0 position for the AO guide star. they concluded Chat the active nucleus is coincident with the stellar eastern component.," Taking the USNO-A2.0 position for the AO guide star, they concluded that the active nucleus is coincident with the stellar eastern component."
 They interpret (he structure as an ongoing merger of two galaxies. with the active nucleus associated with the fainter galaxy.," They interpret the structure as an ongoing merger of two galaxies, with the active nucleus associated with the fainter galaxy."
 steinbring.Craampton.&Hutchings(2002) found a position for the radio source similar to that of Stockton.Canalizo.&Ricdeway(1999).. but they did not state how they calculated it.," \citet{ste02} found a position for the radio source similar to that of \citet{sto99}, but they did not state how they calculated it."
 Their AO imaging at // and A. obtained with on the CFIUT. showed 3 main," Their AO imaging at $H$ and $K$ , obtained with on the CFHT, showed 3 main"
"differs from the finding of ?) who found a weak trend of increasing Lx toward earlier spectral types, though their nearby sample could differ intrinsically from our more distant sample, and may include fewer highly active stars.","differs from the finding of \citet{zick05} who found a weak trend of increasing $L_X$ toward earlier spectral types, though their nearby sample could differ intrinsically from our more distant sample, and may include fewer highly active stars."
" The X-ray luminosities are typically higher than that of the contemporary Sun (?,adjustedtomatch Chandra""s spectralbands), with a small number of G- stars exhibiting X-ray luminosities similar to the active Sun."," The X-ray luminosities are typically higher than that of the contemporary Sun \citep[][adjusted to match {\it Chandra}' 's spectral, with a small number of G-type stars exhibiting X-ray luminosities similar to the active Sun."
" Figure 3 also shows the X-ray to bolometric luminosity ratio as a function of spectral type, using bolometric luminosities determined from the table of main-sequence bolometric magnitudes presented by ?).."," Figure \ref{lx} also shows the X-ray to bolometric luminosity ratio as a function of spectral type, using bolometric luminosities determined from the table of main-sequence bolometric magnitudes presented by \citet{krau07}."
 Our sample and that of are in agreement with the observed trend of increasing?)| Lx/Lpo; toward later spectral types (e.g.?).., Our sample and that of \citet{cove08} are in agreement with the observed trend of increasing $L_X / L_{bol}$ toward later spectral types \citep[e.g.][]{flem95}.
 Our sample includes a number of sources with distances >1 kpc that are likely to be members of the Galactic halo., Our sample includes a number of sources with distances $> 1$ kpc that are likely to be members of the Galactic halo.
" Partly because of our sensitivity limits, the majority of these are highly luminous with log Lx>31 erg s!."," Partly because of our sensitivity limits, the majority of these are highly luminous with log $L_X > 31$ erg $^{-1}$."
" Since the Galactic halo is ~10 Gyrs old and Lx declines with age through magnetic braking, these are almost certainly close binaries kept active through tidal interaction and tapping of orbital angular momentum to sustain strong dynamo activity."," Since the Galactic halo is $\sim$ 10 Gyrs old and $L_X$ declines with age through magnetic braking, these are almost certainly close binaries kept active through tidal interaction and tapping of orbital angular momentum to sustain strong dynamo activity."
" 7?) found that PopulationII binaries typically have lower X-ray luminosities than more metal-rich systems, but do exhibit a high-luminosity tail with log Lx erg s-!."," \citet{ottm97} found that Population binaries typically have lower X-ray luminosities than more metal-rich systems, but do exhibit a high-luminosity tail with log $L_X \sim 29 - 31$ erg $^{-1}$."
" We are unable to definitively identify halo members amongst our sample of moderately bright (log Lx~28—29 erg 1) sources, but our detection of two distant and highly-luminous sources with d=8 kpc and log Lx>31 erg s! suggests that a high-luminosity tail for the halo binary distribution does exist and at a higher X-ray luminosity than found by ?).."," We are unable to definitively identify halo members amongst our sample of moderately bright (log $L_X \sim 28 - 29$ erg $^{-1}$ ) sources, but our detection of two distant and highly-luminous sources with $d \gtrsim 8$ kpc and log $L_X > 31$ erg $^{-1}$ suggests that a high-luminosity tail for the halo binary distribution does exist and at a higher X-ray luminosity than found by \citet{ottm97}."
" If our sample is representative, a simple extrapolation suggests that the Galactic halo contains ~10° binaries with log Lx>91 erg s7!."," If our sample is representative, a simple extrapolation suggests that the Galactic halo contains $\sim 10^5$ binaries with log $L_X \geq 31$ erg $^{-1}$."
" Our most distant and highly luminous source, CID 1600, is a relatively faint detection with only ~5 net counts and a probability of being a background event of 0.026, though the COSMOS catalog lists it with ~12 net counts (?).."," Our most distant and highly luminous source, CID 1600, is a relatively faint detection with only $\sim$ 5 net counts and a probability of being a background event of 0.026, though the COSMOS catalog lists it with $\sim$ 12 net counts \citep{pucc09}. ."
 Optical and near-IR photometry suggests it is a Kl-type star (though no spectroscopy exists) that would put it at a distance of ~11.7 kpc and give it an X-ray luminosity of Lx5x10?! erg s-!., Optical and near-IR photometry suggests it is a K1-type star (though no spectroscopy exists) that would put it at a distance of $\sim$ 11.7 kpc and give it an X-ray luminosity of $L_X \sim 5 \times 10^{31}$ erg $^{-1}$.
" While the X-ray source is ~3” from the optical counterpart, the source is ~12’ off-axis, and its PSF is therefore similarly-sized."," While the X-ray source is $\sim$ $^{\prime\prime}$ from the optical counterpart, the source is $\sim$ $^\prime$ off-axis, and its PSF is therefore similarly-sized."
 There are no other suitable optical counterparts in either the deep Hubble Space Telescope observations of the field (thatextend or longer wavelength mid-IR observations., There are no other suitable optical counterparts in either the deep Hubble Space Telescope observations of the field \citep[that extend down to $I_{AB} or longer wavelength mid-IR observations.
" Considering this we are left with the choice that the X-ray source is either associated with the optical source, or is a background event."," Considering this we are left with the choice that the X-ray source is either associated with the optical source, or is a background event."
" We note that the chance of an X-ray source being within 3” of a source with r’<22 (as the optical source does) is 0.027, making the probability that a random background fluctuation would be found in such a location to be ~7x1074."," We note that the chance of an X-ray source being within $^{\prime\prime}$ of a source with $r' \lesssim 22$ (as the optical source does) is 0.027, making the probability that a random background fluctuation would be found in such a location to be $\sim 7 \times 10^{-4}$."
 Considering that theChandra COSMOS catalog contains ~1700 sources we might expect at least one spurious event such as this and this could be a possible candidate., Considering that the COSMOS catalog contains $\sim$ 1700 sources we might expect at least one spurious event such as this and this could be a possible candidate.
" If the source were real it is likely to be in an active binary that was observed to flare during the observations, an interpretation supported by the high median photon energy of the source (Ex3.8 keV) and the fact that all but one of the detected photons came in the second of two similar-length exposures."," If the source were real it is likely to be in an active binary that was observed to flare during the observations, an interpretation supported by the high median photon energy of the source $\bar{E_X} \sim 3.8$ keV) and the fact that all but one of the detected photons came in the second of two similar-length exposures."
 Deep spectroscopy will be necessary to confirm the stellar nature of the source and detect evidence of its binary nature., Deep spectroscopy will be necessary to confirm the stellar nature of the source and detect evidence of its binary nature.
" 'The second most distant and luminous source in our sample, CID 3205, is a more reliable detection with ~12 net counts and a close association with an optical source whose photometry suggests it is either an F8-type star at a distance of ~8 kpc or a much closer white dwarf (WD) or cataclysmic variable (sincethecolorsofthesesourcesoverlapintheSDSSsystem, ?).."," The second most distant and luminous source in our sample, CID 3205, is a more reliable detection with $\sim$ 12 net counts and a close association with an optical source whose photometry suggests it is either an F8-type star at a distance of $\sim$ 8 kpc or a much closer white dwarf (WD) or cataclysmic variable \citep[since the colors of these sources overlap in the SDSS system,][]{fan99}."
" If the source were an F8 star, then since the main sequence lifetime of such a star is ~5 Gyrs it could not be a member of the Galactic halo, but is more likely to be a member of the thin disk that has been forced onto a highly elliptical orbit through some sort of close encounter."," If the source were an F8 star, then since the main sequence lifetime of such a star is $\sim$ 5 Gyrs it could not be a member of the Galactic halo, but is more likely to be a member of the thin disk that has been forced onto a highly elliptical orbit through some sort of close encounter."
" While the X-ray luminosity of the source is high, it is not unfeasible if the source is relatively young, though it could also be evidence for the source being an active binary."," While the X-ray luminosity of the source is high, it is not unfeasible if the source is relatively young, though it could also be evidence for the source being an active binary."
" However, if the source were a WD it would likely be much closer and therefore have a lower X-ray luminosity."," However, if the source were a WD it would likely be much closer and therefore have a lower X-ray luminosity."
" The USNO-B catalog lists asmall proper motion for the source of 0.26"" /yr,"," The USNO-B catalog \citep{mone03} lists asmall proper motion for the source of $^{\prime\prime}$ /yr,"
" The USNO-B catalog lists asmall proper motion for the source of 0.26"" /yr,("," The USNO-B catalog \citep{mone03} lists asmall proper motion for the source of $^{\prime\prime}$ /yr,"
" The USNO-B catalog lists asmall proper motion for the source of 0.26"" /yr,(?"," The USNO-B catalog \citep{mone03} lists asmall proper motion for the source of $^{\prime\prime}$ /yr,"
" The USNO-B catalog lists asmall proper motion for the source of 0.26"" /yr,(?)"," The USNO-B catalog \citep{mone03} lists asmall proper motion for the source of $^{\prime\prime}$ /yr,"
where αρ1.,where $\alpha_6 \sim 1$.
 Disks do not break up directly into fragments of (hiis size., Disks do not break up directly into fragments of this size.
 Instead. the simulations show that the mode erows to nonlinear amplitude as a spiral wave.," Instead, the simulations show that the mode grows to nonlinear amplitude as a spiral wave."
 In the case of prompt fragmentation. [fragmentis appear first as a smaller-scale instability in Che dense post-shock region associated with the nonlinear wave.," In the case of prompt fragmentation, fragments appear first as a smaller-scale instability in the dense post-shock region associated with the nonlinear wave."
 If. we set ii=my in (10) and (13).M the thin-sheet criterion (1) becomes and criterion (3) For compression-induced stability. against Iragmentation becomes where The similar numerical quantities on the right hand sides of conditions (16) and (17) sugeest (hat. whenever the spiral shock compresses gas into a thin slab. it will be stable against fragmentation.," If we set $m = m_{\rm f}$ in (10) and (13), the thin-sheet criterion (1) becomes and criterion (3) for compression-induced stability against fragmentation becomes where The similar numerical quantities on the right hand sides of conditions (16) and (17) suggest that, whenever the spiral shock compresses gas into a thin slab, it will be stable against fragmentation."
 HLowever. this also means (hat the thinness and strong shock assumptions required for use of relations (3) and (7) break down together simultaneously near the CR.," However, this also means that the thinness and strong shock assumptions required for use of relations (3) and (7) break down together simultaneously near the CR."
 We cannot invert the inequalities to obtain conditions for instability because (here is al least one other stabilizing ellect. namely. shear.," We cannot invert the inequalities to obtain conditions for instability because there is at least one other stabilizing effect, namely, shear."
 So. strictly speaking. our analysis does not tell us whether or when prompt fragmentation does occur. but it implies that. if prompt fragmentation occurs. then it must happen in the vicinity of the CR for conditions under which the inequalities are reversed.," So, strictly speaking, our analysis does not tell us whether or when prompt fragmentation $does$ occur, but it implies that, $if$ prompt fragmentation occurs, then it must happen in the vicinity of the CR for conditions under which the inequalities are reversed."
 We can see from the Qi;-dependence of (17) that (his is morelikely to happen for low Q;., We can see from the $Q_{\rm i}$ -dependence of (17) that this is morelikely to happen for low $Q_{\rm i}$ .
In general. a GW background with a DC of unity or above is defined ascontinuous*.,"In general, a GW background with a $DC$ of unity or above is defined as."
. And non-conlinuous signals can be further categorized into popcorn noise (0.1<DC1) and shot noise (DOC< 0.1) tvpe.," And non-continuous signals can be further categorized into popcorn noise $0.1\leq DC <
1$ ) and shot noise $DC < 0.1$ ) type."
 As source rate evolution will increase out to large cosmological volumes. some studies investigate how too increases with z (Cowardllowelletal. 2011).," As source rate evolution will increase out to large cosmological volumes, some studies investigate how too increases with $z$ \citep{coward_regimbau_06,Eric2010}."
. In this study we are concerned with the value of DC as seen at the detector and thus equation (11)) is interpreted as a total value., In this study we are concerned with the value of $DC$ as seen at the detector and thus equation \ref{duty}) ) is interpreted as a total value.
 For a background signal produced by a source population with an average chirp mass (M)—8.T.M.. for rates rj and rs we find DC values of 0.01 and 0.2 for advanced detectors.," For a background signal produced by a source population with an average chirp mass $\langle M_{c} \rangle = 8.7 \hspace{0.5mm}
M_{\odot}$, for rates $r_{1}$ and $r_{2}$ we find $DC$ values of 0.01 and 0.2 for advanced detectors."
 Therefore (he signal will most likely be of (he shot noise category (or at most popcorn noise for the higher rate r3)., Therefore the signal will most likely be of the shot noise category (or at most popcorn noise for the higher rate $r_{2}$ ).
" For ET twpe detectors the signal will be continuous with DC values of 5.8 and 80.6 for rates r4 and ο respectively,", For ET type detectors the signal will be continuous with $DC$ values of 5.8 and 80.6 for rates $r_{1}$ and $r_{2}$ respectively.
" Although for advanced detectors. the SGWD caleulated in this paper is not continuous (Gaussian) even at the higheroO rate rs. we note thal Drasco&Flanagane»(2003). have found the eross correlation method nearly optimal [or a DC>10? over l-vearintegration""."," Although for advanced detectors, the SGWB calculated in this paper is not continuous (Gaussian) even at the higher rate $r_2$, we note that \citet{drasco} have found the cross correlation method nearly optimal for a $DC > 10^{-3}$ over 1-year."
 In the following sections we therefore consider the cross correlation statistic (ο assess the detectability of the estimated DDII background., In the following sections we therefore consider the cross correlation statistic to assess the detectability of the estimated BBH background.
 The optimum detection strategv for continuous GW background signals is cross-correlating the output of two neighbouring detectors (see.Allen&Romano1999;Maggiore2000).," The optimum detection strategy for continuous GW background signals is cross-correlating the output of two neighbouring detectors \citep[see,][]{SNR,Maggiore}."
. This requires that the detectors are separated by less (han one reduced wavelength. which is about 100 km for frequencies around 500 Iz where Oca(/) might peak.," This requires that the detectors are separated by less than one reduced wavelength, which is about 100 km for frequencies around 500 Hz where $\Omega _{\mathrm{GW}}(f)$ might peak."
 The detectors also need to be sufficiently. well separated that their noise sources are largely uncorrelated., The detectors also need to be sufficiently well separated that their noise sources are largely uncorrelated.
 We note that although (his may not be possible for ET. techniques are in development to remove environmental noise and instrumental correlations (Fotopoulos2008)..," We note that although this may not be possible for ET, techniques are in development to remove environmental noise and instrumental correlations \citep{Fotopoulos:2008yq}."
 Under these conditions. assuming Gaussian noise in each detector and optimal filtering. a filter function chosen to maximize the signal-to-noise ratio. SNR. [or two such detectors is," Under these conditions, assuming Gaussian noise in each detector and optimal filtering, a filter function chosen to maximize the signal-to-noise ratio, SNR for two such detectors is"
ooptical wavelengths.,optical wavelengths.
For cach QSO spectrum. we selected. 715 sections of Thr spectra at wavelengths. corresponcing to the observed. absorption lines in the QSO spectra.,"For each QSO spectrum, we selected $\sim$ sections of ThAr spectra at wavelengths corresponding to the observed absorption lines in the QSO spectra."
 Each Thr ~15 section contains several (typically up to 10) lines., Each ThAr $\sim$ section contains several (typically up to $\sim 10$ ) lines.
 ]irror arravs were generated assuming Poisson counting statistics., Error arrays were generated assuming Poisson counting statistics.
 We [fitted continua to these Thr. sections. dividing the raw spectra by the continua to obtain normalized spectra.," We fitted continua to these ThAr sections, dividing the raw spectra by the continua to obtain normalized spectra."
 The main steps in the procedure are then as follows., The main steps in the procedure are then as follows.
 We select several independent sets of Thr lines., We select several independent sets of ThAr lines.
" Each set of ""hr lines therefore corresponds to. and has been selected in an analogous way to. one QSO absorption system."," Each set of ThAr lines therefore corresponds to, and has been selected in an analogous way to, one QSO absorption system."
 In order to obtain the best estimate of any wavelength calibration errors. the average (Aaανν is determined [or each Thr set ((vpically 37 sets for cach QSO absorption system).," In order to obtain the best estimate of any wavelength calibration errors, the average $(\da)_{\rm ThAr}$ is determined for each ThAr set (typically 3–7 sets for each QSO absorption system)."
 The process of defining sets in this way also oovides a direct additional estimate of the errors on (NaaVrae, The process of defining sets in this way also provides a direct additional estimate of the errors on $(\da)_{\rm ThAr}$.
 To fit the Thr spectra. a mocified: version of was used (to fit Gaussian profiles. see MOIa). adopting Thr aboratory wavelengths from Palmer Eneleman Jr. (1983) or values of wy in equation 1..," To fit the ThAr spectra, a modified version of was used (to fit Gaussian profiles, see M01a), adopting ThAr laboratory wavelengths from Palmer Engleman Jr. (1983) for values of $\omega_0$ in equation \ref{eq:omega}."
" The q; and qe coellicients for he QSO absorption lines are applied to the corresponding ""Thr lines.", The $q_1$ and $q_2$ coefficients for the QSO absorption lines are applied to the corresponding ThAr lines.
 Phat is. we treat the set of Thor lines as if they were the QSO lines.," That is, we treat the set of ThAr lines as if they were the QSO lines."
 The free parameters involved. in cach it were line width (assumed the same for all Thr. lines in a given set. although the results were insensitive to this assumption). “redshift” (numerically close to zero) and. peak intensity.," The free parameters involved in each fit were line width (assumed the same for all ThAr lines in a given set, although the results were insensitive to this assumption), `redshift' (numerically close to zero) and peak intensity."
 We applied. the above method. to cach available PhaAr spectrum corresponding to each QSO absorption cloud considered and the results are illustrated in bie., We applied the above method to each available ThAr spectrum corresponding to each QSO absorption cloud considered and the results are illustrated in Fig.
 2., 2.
 Our first observation is that the scatter in these results is an order of magnitude less than that of the QSO absorption results themselves., Our first observation is that the scatter in these results is an order of magnitude less than that of the QSO absorption results themselves.
 This clearly demonstrates that wavelength calibration of the CCD has not mimicked a shift in à to any significant degree., This clearly demonstrates that wavelength mis-calibration of the CCD has not mimicked a shift in $\alpha$ to any significant degree.
 The results above also show that the data quality permit a precision of afa—10., The results above also show that the data quality permit a precision of $\da \sim 10^{-7}$.
 Secondlv. many points in Fig.," Secondly, many points in Fig."
 2 deviate significantly from zero., 2 deviate significantly from zero.
 There are three possible reasons for this: weak blended emission lines. Thr line mis-identifications. and errors in the Thr. laboratory wavelengths.," There are three possible reasons for this: weak blended emission lines, ThAr line mis-identifications, and errors in the ThAr laboratory wavelengths."
 Inspecting the Thr spectra does indeed reveal typically 720. easily identifiable. but very. weak. lines per 15 and we would expect this to add to the scatter in the final results for (Aa/opi ," Inspecting the ThAr spectra does indeed reveal typically $\sim$ 20 easily identifiable, but very weak, lines per $15$ and we would expect this to add to the scatter in the final results for $(\da)_{\rm ThAr}$."
"Furthermore. this elfect. would not wocuece correlated. deviations in (Aea), as observed."," Furthermore, this effect would not produce correlated deviations in $(\da)_{\rm ThAr}$, as observed."
 Consistent with this interpretation. we found that the value of X7 per degree of reedom for a fit to a given set of Thr ines was “>1.," Consistent with this interpretation, we found that the value of $\chi^2$ per degree of freedom for a fit to a given set of ThAr lines was $\gg 1$."
 Note that i£ Thr line mis-icentifications iwl occurred. one may expect systematically correlated deviations.," Note that if ThAr line mis-identifications had occurred, one may expect systematically correlated deviations."
 Finally. we do not expect any significant errors in the Thr laboratory wavelengths (as discussed in Section ?77)).," Finally, we do not expect any significant errors in the ThAr laboratory wavelengths (as discussed in Section \ref{sec:wavemiscal}) )."
 On inspection of some of the PhaAr spectra. we noted that some Thr lines towards the blue end of the spectrum. were slightly asymmetric.," On inspection of some of the ThAr spectra, we noted that some ThAr lines towards the blue end of the spectrum were slightly asymmetric."
 This may indicate a degradation in the polynomial fits near the edges of the fitting regions., This may indicate a degradation in the polynomial fits near the edges of the fitting regions.
" ""here may also be intrinsic ΜΕ in the IP (Section ?7)).", There may also be intrinsic asymmetries in the IP (Section \ref{sec:IP}) ).
 I£ such asvmnmetrey. variations exist then they should also apply to the QSO spectra., If such asymmetry variations exist then they should also apply to the QSO spectra.
 However. the results in Fig.," However, the results in Fig."
 2 show that such variations do not produce a systematic ellect in Aasa at any detectable level., 2 show that such variations do not produce a systematic effect in $\da$ at any detectable level.
 We therefore conclude this section bv. stating that the results in Fig., We therefore conclude this section by stating that the results in Fig.
 2 show unambiguously that wavelength calibration errors and variations in IP asvmmetries are not responsible for the observed. shifts in a in the QSO absorption line results., 2 show unambiguously that wavelength calibration errors and variations in IP asymmetries are not responsible for the observed shifts in $\alpha$ in the QSO absorption line results.
associated with ongoing accretion (e.g.?)..,associated with ongoing accretion \citep[e.g.][]{2006A&A...452..245N}.
" The H-band shows the [Fell] lines at 1.644, 1.669, 1.677, and 1.748 um (the latter might be blended with a hydrogen feature)."," The H-band shows the [FeII] lines at 1.644, 1.669, 1.677, and $\,\mu$ m (the latter might be blended with a hydrogen feature)."
" While the lines of molecular hydrogen can originate in the disk or in a jet, the [Fell] emission can only occur in the jet."," While the lines of molecular hydrogen can originate in the disk or in a jet, the [FeII] emission can only occur in the jet."
" Thus, our spectra prove ongoing accretion and outflow activity in IRAS04325C. From the way we have analysed the spectra it is certain that the emission originates in an area within the PSF of the point source, i.e. within a radius of 075 AAU) around the source."," Thus, our spectra prove ongoing accretion and outflow activity in IRAS04325C. From the way we have analysed the spectra it is certain that the emission originates in an area within the PSF of the point source, i.e. within a radius of $0\farcs5$ AU) around the source."
 The [Fell] lines provide a constraint on the electron density ne in the jet (e.g.??)..," The [FeII] lines provide a constraint on the electron density $n_e$ in the jet \citep[e.g.][]{2002A&A...393.1035N,2008A&A...487.1019G}."
" The flux ratio between the well-detected lines at 1.644 and um is 2.9, which gives ne> 10?ccm?."," The flux ratio between the well-detected lines at 1.644 and $\,\mu$ m is 2.9, which gives $n_e > 10^5$ $^{-3}$."
 Note that the [Fell] line ratio saturates for larger densities., Note that the [FeII] line ratio saturates for larger densities.
" Consistently, the lower limits to the flux ratios of 1.544 vs. um and 1.600 vs. 1.644um yield densities >0.5-10° and > 10*ccm7?, respectively."," Consistently, the lower limits to the flux ratios of 1.544 vs. $\,\mu$ m and 1.600 vs. $\mu$ m yield densities $>0.5 \cdot 10^5$ and $>10^4$ $^{-3}$, respectively."
" Such high values are usually found at the base of the jet, while the density typically drops further away from the source (e.g.??).."," Such high values are usually found at the base of the jet, while the density typically drops further away from the source \citep[e.g.][]{2006ApJ...641..357T,2008A&A...487.1019G}."
 Thus the analysis confirms that the [FeII] emission originates in a jet close to IRAS04325C. The Br 4 line is used to derive an estimate for the mass accretion rate following ?.., Thus the analysis confirms that the [FeII] emission originates in a jet close to IRAS04325C. The Br $\gamma$ line is used to derive an estimate for the mass accretion rate following \citet{2006A&A...452..245N}.
 The approximate equivalent width in this line is 7+2AA.., The approximate equivalent width in this line is $7 \pm 2$.
 Scaling with the K-band magnitude of mmag for a 3400KK star at 1MMyr from ? gives a line luminosity of log(Lpry/Lo)~—4.1., Scaling with the K-band magnitude of mag for a K star at Myr from \citet{1998A&A...337..403B} gives a line luminosity of $\log (L_{\mathrm{Br\gamma}} / L_{\odot}) \sim -4.1$.
 With the empirical relations by ? this translates into an accretion luminosity of log(Lacc/~—0.8 (seealso?).., With the empirical relations by \citet{2004AJ....128.1294C} this translates into an accretion luminosity of $\log (L_\mathrm{acc} / L_{\odot}) \sim -0.8$ \citep[see also][]{1998AJ....116.2965M}.
 The mass accretion rate is then MLco)=LaccR/GM., The mass accretion rate is then $\dot{M} = L_{\mathrm{acc}} R / G M$.
" Assuming a mass of MMo and a radius of RRo results in an accretion rate in the range of (1.0+0.3)-10? Mo yyr comparable to accretion rates derived for T Tauri stars but!, higher than those for young brown dwarfs (?).."," Assuming a mass of $_{\odot}$ and a radius of $_{\odot}$ results in an accretion rate in the range of $(1.0 \pm 0.3) \cdot 10^{-8}\,$ $_{\odot}$ $^{-1}$, comparable to accretion rates derived for T Tauri stars but higher than those for young brown dwarfs \citep{2005ApJ...626..498M}."
 We derived magnitudes/fluxes for the two components of IRAS04325 from the Gemini near- and mid-infrared images as well as from the submm continuum map., We derived magnitudes/fluxes for the two components of IRAS04325 from the Gemini near- and mid-infrared images as well as from the submm continuum map.
 For the JHK images we adopted an aperture of 122 (10 pixel) and a background annulus between 122 and 1:88 (10-15 pixel)., For the JHK images we adopted an aperture of 2 (10 pixel) and a background annulus between 2 and 8 (10-15 pixel).
" The field contains one 2MASS source outside IRAS04325 with J=17.98, H=16.41, and K=14.64 which was used to shift the instrumental magnitudes into the 2MASS system."," The field contains one 2MASS source outside IRAS04325 with $J=17.98$, $H=16.41$, and $K=14.64$ which was used to shift the instrumental magnitudes into the 2MASS system."
 The resulting values are listed in Table 1.., The resulting values are listed in Table \ref{phot}.
" In the J-band there is diffuse emission at the position of object C, but no clear point source; we consider the flux measurement to be an upper limit."," In the J-band there is diffuse emission at the position of object C, but no clear point source; we consider the flux measurement to be an upper limit."
 The previously published photometry of IRAS04325 has been obtained with larger apertures (??)..," The previously published photometry of IRAS04325 has been obtained with larger apertures \citep{1999AJ....118.1784H,2008AJ....135.2496C}."
" Hence, these literature fluxes are significantly larger than ours because they include extended emission."," Hence, these literature fluxes are significantly larger than ours because they include extended emission."
" For the L'-, M'-, and N’-band photometry we used apertures of 1700 or 122 and a sky annulus similar to the one for JHK."," For the L'-, M'-, and N'-band photometry we used apertures of 0 or 2 and a sky annulus similar to the one for JHK."
" In these bands the extended emission is less problematic, thus the choice of the sky annulus does not significantly affect the results."," In these bands the extended emission is less problematic, thus the choice of the sky annulus does not significantly affect the results."
" The L’/M’-band magnitudes were shifted to the standard system basedon the values for the calibration star HD36719, which has mmag throughout this wavelength regime (?).. "," The L'/M'-band magnitudes were shifted to the standard system basedon the values for the calibration star HD36719, which has mag throughout this wavelength regime \citep{2003MNRAS.345..144L}. ."
For the N’-band, For the N'-band
eqe.(8)).(9)). (10]): Second. an expression [or gradient of a? is derived [rom an integrated form of eq.(11)): where. ry is a certain reference point of integration Chat will be set later.,"\ref{eq:cntn}) \ref{eq:eqm}) ), \ref{eq:eqst}) ): Second, an expression for gradient of $a^2$ is derived from an integrated form of \ref{eq:egcns}) ): where, $r_0$ is a certain reference point of integration that will be set later."
 £4 is mathematically an integral constant. which must be conserved in the entire region of the calculation.," $E_{\rm tot}$ is mathematically an integral constant, which must be conserved in the entire region of the calculation."
 Onlv (ransonic solutions are allowed [or the flow speed. ο.," Only transonic solutions are allowed for the flow speed, $v$."
 The integration of three differential equations (20)).(21)). and (6)) is carried out simultaneously from (he sonic point. r=r.to both outward and inward directions by the fourth-order Runge-Kutta method. also deriving densitwv [rom eq.(8)) al each integration.," The integration of three differential equations \ref{eq:vlgr}) \ref{eq:svgr}) ),and \ref{eq:wvdrrd}) ) is carried out simultaneously from the sonic point, $r=r_{\rm s}$ to both outward and inward directions by the fourth-order Runge-Kutta method, also deriving density from \ref{eq:cntn}) ) at each integration."
 We have used variable sizes for (he erid of the integration. setting the smaller mesh size in the region where physical values change rapidly.," We have used variable sizes for the grid of the integration, setting the smaller mesh size in the region where physical values change rapidly."
 For instance. in the TR. we set a mesh as small as 107 cm (0.01. kim)for one grid.," For instance, in the TR, we set a mesh as small as $10^3$ cm (0.01 km)for one grid."
" To start the integration. we have to set nine variables. c.(4).a2.(“)..Vase(£s... f FA aad rn, (subseript s. denotes the sonic point)."," To start the integration, we have to set nine variables, $v_{\rm s}, (\frac{dv}{dr})_{\rm s}, 
a^2_{\rm s}, (\frac{da^2}{dr})_{\rm s}, \alpha_{\rm w,s}, 
(\frac{d \alpha_{\rm w}}{dr})_{\rm s},\rho_{\rm s}$ , $E_{\rm tot}$ and $r_{\rm s} $ (subscript 's' denotes the sonic point)."
" We can determine ry. (5t). and 1, for given a2 and (we. by a condition that both the nunerator aud the denominator of eq.(20)) are zero alr, (Parker1953)."," We can determine $v_{\rm s}$, $(\frac{dv}{dr})_{\rm s}$ and $r_{\rm s}$ for given $a^2_{\rm s}$ and $(\frac{da^2}{dr})_{\rm s}$ by a condition that both the numerator and the denominator of \ref{eq:vlgr}) ) are zero at $r_{\rm s}$ \citep{pkr58}."
. (os).| is also derived rom eq.(6)) on a given ay...," $(\frac{d \alpha_{\rm w}}{dr})_{\rm s}$ is also derived from \ref{eq:wvdrrd}) ) on a given $\alpha_{\rm w,s}$."
 Setting the reference point. rg—ry. For the integration of radiative cooling function. we obtain. £4 as where all (he variables are evaluated at the sonic point.," Setting the reference point, $r_0=r_{\rm s}$, for the integration of radiative cooling function, we obtain $E_{\rm tot}$ as where all the variables are evaluated at the sonic point."
" Now we have four variables. a2,(E).dp wa 0X pu remaining (o be regulated bythe four boundary conditions of eqs.(15)) (13))."," Now we have four variables, $a^2_{\rm s}, (\frac{da^2}{dr})_{\rm s}, 
\alpha_{\rm w,s}$ and $\rho_{\rm s}$ , remaining to be regulated bythe four boundary conditions of \ref{eq:bc1}) \ref{eq:bc4}) )."
 Concrete procedures for finding a unique solution on a given F4 are described below.," Concrete procedures for finding a unique solution on a given $F_{{\rm w},0 }$ are described below."
spectrum.,spectrum.
 Despite these problems. the model provides an adequate fit to the data. given that our goal is to determine a general picture of the cloud structure rather than obtain a precise evaluation of cloud sizes. locations. ete.," Despite these problems, the model provides an adequate fit to the data, given that our goal is to determine a general picture of the cloud structure rather than obtain a precise evaluation of cloud sizes, locations, etc."
 As a result we shall adopt the three-component model in further discussions., As a result we shall adopt the three-component model in further discussions.
 In this section we discuss implications of the models used to describe the 21 em absorption in 3C 286., In this section we discuss implications of the models used to describe the 21 cm absorption in 3C 286.
 In 5 5.3 we discussed the three-component model used to model the A7/7. AP/P. and Aw absorption spectra.," In $\S$ 5.3 we discussed the three-component model used to model the $\Delta I/I$, $\Delta P/P$, and $\Delta \chi$ absorption spectra."
 As illustrated in Fig., As illustrated in Fig.
 the three-component model is actually based on two clouds., the three-component model is actually based on two clouds.
 One cloud (cloud 2) completely covers the core source., One cloud (cloud 2) completely covers the core source.
 The other cloud (cloud 1) partially covers the jet source. and contains a velocity gradient such that the part of the cloud toward the polarized portion of the jet (component 3) causes the velocity centroid of the polarized spectra to be offset in velocity from the unpolarized AJ/7 spectrum.," The other cloud (cloud 1) partially covers the jet source, and contains a velocity gradient such that the part of the cloud toward the polarized portion of the jet (component 3) causes the velocity centroid of the polarized spectra to be offset in velocity from the unpolarized $\Delta I/I$ spectrum."
 We simplified the model by letting component 3 be a uniform sub-region of cloud 1 with a velocity centroid. velocity dispersion. and 21 em optical depth that differ from their values in the parent cloud.," We simplified the model by letting component 3 be a uniform sub-region of cloud 1 with a velocity centroid, velocity dispersion, and 21 cm optical depth that differ from their values in the parent cloud."
" Table 1 indicates that the polarization position angle for the part of the source absorbed by the foreground cloud i$ \,=43.6"".", Table 1 indicates that the polarization position angle for the part of the source absorbed by the foreground cloud is $\chi_{b}$ $^{\circ}$.
 Because the ratio ofthe Stokes parameters. U /Q=tan(2\ ). the predicted ratio U /Q-20.4.," Because the ratio ofthe Stokes parameters, $U/Q$ $\chi$ ), the predicted ratio $U/Q$ =20.4."
 By contrast. U /Q=0.046 toward the part of the jet that bypasses the absorbing gas where the model predicts \g=2.67°.," By contrast, $U/Q$ =0.046 toward the part of the jet that bypasses the absorbing gas where the model predicts $\chi_{a}$ $^{\circ}$."
 This is consistent with the independent Stokes U and Q absorption spectra. which show a clear detection of the absorption feature in Stokes U but not in Stokes Q.," This is consistent with the independent Stokes $U$ and $Q$ absorption spectra, which show a clear detection of the absorption feature in Stokes $U$ but not in Stokes $Q$."
 The absorption feature. which is detected at a signal-to-noise ratio of about 20:1 in Stokes U. is predicted to be below the |-σ noise level in Stokes Q.," The absorption feature, which is detected at a signal-to-noise ratio of about 20:1 in Stokes $U$, is predicted to be below the $\sigma$ noise level in Stokes $Q$."
" The model also predicts q,,,,2387 for the position angle integrated over the entire source. which is consistent with single-dish measurements of 335 at 1.5 GHz (Tabara Inoue 1980)."," The model also predicts $\chi_{cont}$ $^{\circ}$ for the position angle integrated over the entire source, which is consistent with single-dish measurements of $\pm$ $^{\circ}$ at 1.5 GHz (Tabara Inoue 1980)."
 Our model predicts that the intrinsic. polarization. position. angle of the jet decreases by Vo— \o= —40.9° in the SW direction (see Fig. 5)., Our model predicts that the intrinsic polarization position angle of the jet decreases by $\chi_{a}-\chi_{b}$ = $-$ $^{o}$ in the SW direction (see Fig. ).
 Polarization maps at 5 GHZ (Jiang 1996; Cotton 1997) do show shifts by about this amount in the sense of decreasing \ along the jet axis in the SW direction., Polarization maps at 5 GHZ (Jiang 1996; Cotton 1997) do show shifts by about this amount in the sense of decreasing $\chi$ along the jet axis in the SW direction.
" However. the value of \ at the extreme edge of the jet appears to be greater than our prediction of 4,22.67."," However, the value of $\chi$ at the extreme edge of the jet appears to be greater than our prediction of $\chi_{a}$ $^{\circ}$."
 These predictions could be checked with VLBI polarization measurements at lower standard frequencies such as 609 MHz., These predictions could be checked with VLBI polarization measurements at lower standard frequencies such as 609 MHz.
 Next we turn to the properties of the absorbing gas., Next we turn to the properties of the absorbing gas.
 The large projected areas subtended by the sources at the absorber (Ae 070< 35 per for the core and ÀAj42:280 «68 per for the jet) likely indicate that the gas causing 21 em absorption in 3C 286 is comprised of many interstellar clouds., The large projected areas subtended by the sources at the absorber $A_{\rm core}$$\approx$ ${\times}$ 35 $^{2}$ for the core and $A_{\rm jet}$$\approx$ ${\times}$ 60 $^{2}$ for the jet) likely indicate that the gas causing 21 cm absorption in 3C 286 is comprised of many interstellar clouds.
 In the Galaxy a beam with area directed perpendicular to the plane of the disk would subtendAj4 Nene 1000 CNM (ice. cold neutral-medium |Wolfire 1995]) clouds.," In the Galaxy a beam with area $A_{\rm jet}$ directed perpendicular to the plane of the disk would subtend $\cal N_{\rm CNM}$$\approx$ 1000 CNM (i.e., cold neutral-medium [Wolfire 1995]) clouds."
 This 1s because NcNMAICNMAj \2H where the space density of CNM clouds ncxaj © 107 pe? (McKee Ostriker 1977) and the disk half-thickness H ~ 100 pe., This is because $\cal N_{\rm CNM}$ $n_{CNM}$${\times}A_{\rm jet}$$\times$ $H$ where the space density of CNM clouds $n_{CNM}$ $\approx$ $\times$ $^{-4}$ $^{-3}$ (McKee Ostriker 1977) and the disk half-thickness $H$ $\approx$ 100 pc.
 This value of Δον is a lower limit. since Wenm=fcayA 2H/cos() for a disk with inclination angle i.," This value of $\cal N_{\rm CNM}$ is a lower limit, since $\cal N_{\rm CNM}$ $n_{CNM}$$\times$$A{\times}$ $H$ $i$ ) for a disk with inclination angle $i$."
 For the core source ον 260., For the core source $\cal N_{\rm CNM}$$\ge$ 60.
 As a result the Gaussian shape of the absorption profile is naturally explained as the optical-depth weighted sum of multiple Gaussians (1.e.. the central limit theorem).," As a result the Gaussian shape of the absorption profile is naturally explained as the optical-depth weighted sum of multiple Gaussians (i.e., the central limit theorem)."
 Note that the central velocities of the Gaussians are likely superposed on large-scale velocity gradients present. both in ‘cloud’ | toward the jet and ‘cloud’ 2 toward the core., Note that the central velocities of the Gaussians are likely superposed on large-scale velocity gradients present both in `cloud' 1 toward the jet and `cloud' 2 toward the core.
" While such a gradient in ‘cloud’ | is required to explain the velocity shifts between the stokes 7 and polarized spectra. as explained above. the necessity for a gradient in ""cloud? 2 stems from the approximate agreement between the optical redshift and redshift of ‘cloud’ 2 illustrated in Fig."," While such a gradient in `cloud' 1 is required to explain the velocity shifts between the stokes $I$ and polarized spectra, as explained above, the necessity for a gradient in `cloud' 2 stems from the approximate agreement between the optical redshift and redshift of `cloud' 2 illustrated in Fig."
 6., 6.
 This implies that the optical continuum source is. as expected. physically associated with the compact radio core.," This implies that the optical continuum source is, as expected, physically associated with the compact radio core."
" The small. but significant. difference between these redshifts further suggests that since the optical beam size is small compared to the dimensions of ""cloud? 2. the optical beam samples only a limited portion of the velocities spanned by the gradient in ""cloud"" 2."," The small, but significant, difference between these redshifts further suggests that since the optical beam size is small compared to the dimensions of `cloud' 2, the optical beam samples only a limited portion of the velocities spanned by the gradient in `cloud' 2."
" Therefore. although the UV resonance lines form in gas within ""cloud? 2. the velocity centroids formed by averaging over the optical and wider radio beams need not be equal."," Therefore, although the UV resonance lines form in gas within `cloud' 2, the velocity centroids formed by averaging over the optical and wider radio beams need not be equal."
 To determine whether the absorbing gas is CNM. WNM. or something else we need to determine its kinetic temperature. Τε.," To determine whether the absorbing gas is CNM, WNM, or something else we need to determine its kinetic temperature, $T_{k}$."
" Assuming the foreground gas covers the entire source with a uniform cloud characterized by σι = 3.75 km s! (Wolfe 2008). one finds an upper limit of 7""""*=1690 K. If we further assume that equals the UV-determined value of «107! em*(e.g. Wolfe Davis 1978: Kanekar 20034). then eq. ("," Assuming the foreground gas covers the entire source with a uniform cloud characterized by $\sigma_{v}$ = 3.75 km $^{-1}$ (Wolfe 2008), one finds an upper limit of $T_{k}^{\rm max}$ =1690 K. If we further assume that equals the UV-determined value of $^{21}$ $^{-2}$ (e.g. Wolfe Davis 1978; Kanekar 2003a), then eq. ("
"6) implies 7,21035 K. since the central 21] em optical depth ™=0.10 (Wolfe 2008).","6) implies $T_{s}$ =1035 K, since the central 21 cm optical depth $\tau_{0}$ =0.10 (Wolfe 2008)."
" Because 7;«T;Tj"" for diffuse warm gas (Liszt 2001). these results would indicate the gas is neither standard CNM where 7;2: 80 K nor standard WNM where 7; - 8000 K (Wolfire 1995). but rather resembles the thermally unstable phase found by Heiles Troland (2003) in the Galaxy ISM."," Because $T_{s} \ < T_{k} \ < T_{k}^{\rm max}$ for diffuse warm gas (Liszt 2001), these results would indicate the gas is neither standard CNM where $T_{k}{\approx}$ 80 K nor standard WNM where $T_{k}$ $\approx$ 8000 K (Wolfire 1995), but rather resembles the thermally unstable phase found by Heiles Troland (2003) in the Galaxy ISM."
 On the other hand the frequency-dependent change in polarization position angle across the absorption. feature (panel 3 in Fig. 5), On the other hand the frequency-dependent change in polarization position angle across the absorption feature (panel 3 in Fig. )
 provides strong evidence against the assumption of a single uniform cloud (see 5; 5.4)., provides strong evidence against the assumption of a single uniform cloud (see $\S$ 5.4).
" From the values of 7,; in Table | we find that 7;721450 K. 1115 K. and 1400 K for gas in velocity components |. 2. and 3."," From the values of $\sigma_{v,i}$ in Table 1 we find that $T_{k,i}^{\rm max}$ =1450 K, 1115 K, and 1400 K for gas in velocity components 1, 2, and 3."
" If we again assume that ;21.77 107! em in each case. the values of Τε) in Table | indicate that 7,2892 K 7,:22230 K. and 7, 3=394 K. However. since the values of /T;; inTable 1 are spatial averages over dimension much larger than the ~ 1 It."," If we again assume that $_{,i}$ $^{21}$ $^{-2}$ in each case, the values of $_{,i}$ $T_{s,i}$ in Table 1 indicate that $T_{s,1}$ =892 K $T_{s,2}$ =2230 K, and $T_{s,3}$ =394 K. However, since the values of $_{,i}$ $T_{s,i}$ inTable 1 are spatial averages over dimension much larger than the $\approx$ 1 lt."
 yr., yr.
 size of the UV continuum. it is unlikely that the UV determined value of applies to the gas in each component. nor that each component has the same value of My).," size of the UV continuum, it is unlikely that the UV determined value of applies to the gas in each component, nor that each component has the same value of ."
" It is equally plausible to assume that averaged across cloud | is a factor of 2 lower than the UV determined value. and C, equals 0.25 rather than 0.5."," It is equally plausible to assume that averaged across cloud 1 is a factor of 2 lower than the UV determined value, and $C_{1}$ equals 0.25 rather than 0.5."
" In that case 7,,2223 K. Similarly if 32 4. then T, 32197 K. Since these are physically plausible temperatures for CNM gas with the low metal abundances inferred for this absorber (Wolfe 2008). it is reasonable to assume they represent kinetic temperatures."," In that case $T_{s,1}$ =223 K. Similarly if $_{,3}$ $_{,1}$, then $T_{s,3}$ =197 K. Since these are physically plausible temperatures for CNM gas with the low metal abundances inferred for this absorber (Wolfe 2008), it is reasonable to assume they represent kinetic temperatures."
" Interestingly. the assumption o=l.77 107! env results in the unreasonable prediction that 7,2 is higher than 7/55: we take this as direct evidence that Ng» is lower than «107 em? and/or €» < I. thereby implying that cloud 2 could also be CNM."," Interestingly, the assumption $_{,2}$ $^{21}$ $^{-2}$ results in the unreasonable prediction that $T_{s,2}$ is higher than $T_{k,2}^{\rm max}$ : we take this as direct evidence that $_{,2}$ is lower than $^{21}$ $^{-2}$ and/or $C_{2}$ $<$ 1, thereby implying that cloud 2 could also be CNM."
 Of course. we cannot entirely rule out the possibility that the gas is in a thermally unstable phase.," Of course, we cannot entirely rule out the possibility that the gas is in a thermally unstable phase."
 But the above arguments in addition to the results from the VLBI line experiment. which indicates individual components toward the core and jet with kinetic temperatures 7; « 500 K (Wolfe 1976). make a compelling case that thegascausing 21] em absorption in 3C," But the above arguments in addition to the results from the VLBI line experiment, which indicates individual components toward the core and jet with kinetic temperatures $T_{k}$ $<$ 500 K (Wolfe 1976), make a compelling case that thegascausing 21 cm absorption in 3C"
presence of neutral hydrogen around the ionized regions of this high excitation PN.,presence of neutral hydrogen around the ionized regions of this high excitation PN.
 Qur spectrum also reveals an abnormally TOC feature redward of the 111] AGS27 line (Fig. 1)., Our spectrum also reveals an abnormally broad feature redward of the ] $\lambda6827$ line (Fig. \ref{raman}) ).
 A Svo-Ciussiau profile fit shows that the broad feature has a ceutral waveloneth of aud à ΕΤΠΑΛΙ of |. which ds iuuch broader han the A6827 line (FIVITALSS i)," A two-Gaussian profile fit shows that the broad feature has a central wavelength of and a $FWHM$ of $^{-1}$, which is much broader than the ] $\lambda6827$ line $FWHM\sim88$ $^{-1}$ )."
 The feature was also detected by Péquignot&Daluteau(1991) and was identified bv them as a line., The feature was also detected by \citet{pequignotba1994} and was identified by them as a line.
 Caven he narrow width of other lines. we regard this as a uus-icdentification.," Given the narrow width of other lines, we regard this as a mis-identification."
 Sclunid(1989) sugeested that Raman scattering of the ALO32.1005 resonance doublet by neutral hywdrosenu gives rise to two velocity-broacdened lines a GS30 aud that have con. Widely observed iu 1e spectra of sviubiotie stars.," \citet{schmid} suggested that Raman scattering of the $\lambda1032,1038$ resonance doublet by neutral hydrogen gives rise to two velocity-broadened lines at 6830 and that have been widely observed in the spectra of symbiotic stars."
 Therefore. based on its measured wavelength aud the aree FYVITAL. we ideutify the broad οσο. feature at as the Raman line atA.," Therefore, based on its measured wavelength and the large $FWHM$, we identify the broad emission feature at as the Raman line at."
" Our identification is strengthened by the detection of another cature, Imareinally above he detection limit. atÀ.. i. at the expected position of the other O Raman ie."," Our identification is strengthened by the detection of another feature, marginally above the detection limit, at, i.e. at the expected position of the other O Raman line."
 The measurements vield a ATUSS/AGSOSO0. imteusitv ratio of approximately L/7. sualler thui he ratio of 1/1 vpically found iu sviubiotic stars.," The measurements yield a $\lambda7088/\lambda6830$ intensity ratio of approximately 1/7, smaller than the ratio of 1/4 typically found in symbiotic stars."
" While our measured ATOSSA6830 intensitv ratio in lav suffer roni large uncertainties because of the weakness of the ATUSS eature, its small value seems to suggest that the enuvirous from which these Raman lines arise iu nav differ from those iu syiubiotic svstenis."," While our measured $\lambda7088/\lambda6830$ intensity ratio in may suffer from large uncertainties because of the weakness of the $\lambda$ 7088 feature, its small value seems to suggest that the environs from which these Raman lines arise in may differ from those in symbiotic systems."
 Iu. sviubiotic inanes. the ultraviolet doublet cussion from the vicinity of a verv hot white dwarf is Raman scattered in the atmosphere of a giant star companion.," In symbiotic binaries, the ultraviolet doublet emission from the vicinity of a very hot white dwarf is Raman scattered in the atmosphere of a giant star companion."
 The imcleus of is not known as a binary star., The nucleus of is not known as a binary star.
 By analoey with the Raman lines observed by Péquignotetal.(1997. 2003).. the Raman lines may be formed in the photodissociation region (PDR) that corresponds to the interface between the II! reeiou aud the large molecular euvelope of7027.," By analogy with the Raman lines observed by \citet{pequignotbm,pequignot2003}, the Raman lines may be formed in the photodissociation region (PDR) that corresponds to the interface between the $^+$ region and the large molecular envelope of."
.. The relatively small column density of this PDR. compare to the atmosphere of a giaut. may be compcusated by the large covering factor.," The relatively small column density of this PDR, compared to the atmosphere of a giant, may be compensated by the large covering factor."
 The theoretical ratio of the ÀA1032.1035 lues is 2 aud he ratio of he Raman cross sections is about 3.3.," The theoretical ratio of the $\lambda\lambda1032,1038$ lines is 2 and the ratio of the Raman cross sections is about 3.3."
 Ina sinall optical depth approxiuation. the iuteusity ratio ofthe Raman components should be of order 6.6. in agreement with the observed inteusities.," In a small optical depth approximation, the intensity ratio of the Raman components should be of order 6.6, in agreement with the observed intensities."
 The sinaller ratio observed m sviubioties may reflect a departure from the 2:1 ratio of the lines (a similar departure is observed for aud in these objects: see Sclunicd 1989)., The smaller ratio observed in symbiotics may reflect a departure from the 2:1 ratio of the lines (a similar departure is observed for and in these objects; see Schmid 1989).
 Iu7027.. the red Raman lines may allow one to indirectly determine the intensity of he UV lines. which prestunably arise from the high ionization region of the PN. and thus provide a new useful coustraimt ou the properties of the nucleus.," In, the red Raman lines may allow one to indirectly determine the intensity of the UV lines, which presumably arise from the high ionization region of the PN, and thus provide a new useful constraint on the properties of the nucleus."
 Raman features lave also been detected in the spectrum of the PN (Cirovesetal..20023., Raman features have also been detected in the spectrum of the PN \citep{groves}.
. Both aud have a very high excitation class and show strong molecular cussion., Both and have a very high excitation class and show strong molecular emission.
 They thus may share some como evolutionary properties., They thus may share some common evolutionary properties.
 It is generally believe that a significant amount of dust coexists with the iouized aud neutral eas in (sec.e.g...Woodwardetal.," It is generally believed that a significant amount of dust coexists with the ionized and neutral gas in \citep[see, 
e.g.,][]{woodward}."
1992).. Large variatious of extinction across the nebula. caused by local dust. have been detected by optical aud radio imagine (Waltonetal.. 1985).," Large variations of extinction across the nebula, caused by local dust, have been detected by optical and radio imaging \citep{walton}."
. It follows that emission lines arising frou ciffereut ionized zones mav suffer different amounts of reddeniug., It follows that emission lines arising from different ionized zones may suffer different amounts of reddening.
 A comprehensive treatment of the dust extinction for the Cluission lines observe from 77027 requires detailed photoionization modeling and au accurate troeatiuent of the dust component (its spatial and size distributions. chemical composition etc.}.," A comprehensive treatment of the dust extinction for the emission lines observed from 7027 requires detailed photoionization modeling and an accurate treatment of the dust component (its spatial and size distributions, chemical composition etc.),"
 which is bevoud the scope of the current work., which is beyond the scope of the current work.
 Efforts to correct for the effects of dust extinction on observed line fluxes assmuius a siuplified. geometry have been attempted previousA by Seaton(1979). aud Middleimass(1990)., Efforts to correct for the effects of dust extinction on observed line fluxes assuming a simplified geometry have been attempted previously by \citet{seaton} and \citet{middlemass}.
. Ou the other haucd. our previous study shows that it is stil (good approximation to use the standard Galactic extinction curve for the diffuse interstellar 1uiedium (ISAT) to dereddeu the integrated spectruu of (Zhangetal. 2003).," On the other hand, our previous study shows that it is still a good approximation to use the standard Galactic extinction curve for the diffuse interstellar medium (ISM) to deredden the integrated spectrum of \citep{zhang}."
. Accordingly. we have dereddenued all line fluxes by where f(A) is the standard Calactic extinction curve for a total-to-selective extinction ratio of R=3.2 (Towarth. 1983)..aud ο1) is the logarithimic extinction at 11.0.," Accordingly, we have dereddened all line fluxes by where $f(\lambda)$ is the standard Galactic extinction curve for a total-to-selective extinction ratio of $R=3.2$ \citep{howarth}, ,and $c({\rm
H}\beta)$ is the logarithmic extinction at $\beta$ ."
substantial amount of carbon burns stably before a superburst is triggered.,substantial amount of carbon burns stably before a superburst is triggered.
 Consequently. the one-zone approximation drastically underestimates (he superburst recurrence times.," Consequently, the one-zone approximation drastically underestimates the superburst recurrence times."
 It is our pleasure to thank Edward Brown. Jean int Zand. and Don Lamb for helpful discussions and (he referee for insightful comments aud suggestions.," It is our pleasure to thank Edward Brown, Jean in't Zand, and Don Lamb for helpful discussions and the referee for insightful comments and suggestions."
 This work was supported by NASA grant NNGOA4GL238CG., This work was supported by NASA grant NNG04GL38G.
Both C II and C III absorption lines were detected in the spectrum of the star.,Both C II and C III absorption lines were detected in the spectrum of the star.
 However. the number of these lines in 122023 is few.," However, the number of these lines in I22023 is few."
 Further. all the identified carbon lines are blends with the exception of the CIKKI6) line at.," Further, all the identified carbon lines are blends with the exception of the CII(16) line at."
 The carbon abundance was therefore estimated using spectrum synthesis alone., The carbon abundance was therefore estimated using spectrum synthesis alone.
 The region of the CII(16) lines (~ -SISIA ) was used for this purpose., The region of the CII(16) lines $\sim$ $-$ ) was used for this purpose.
 The largest number of absorption lines in the spectrum are those of O II., The largest number of absorption lines in the spectrum are those of O II.
" The O II lines with W, > may be sensitive to NLTE effects.", The O II lines with $_{\lambda}$ $\ge$ may be sensitive to NLTE effects.
 Such strong lines were neglected while estimating the oxygen abundance using WIDTHO9., Such strong lines were neglected while estimating the oxygen abundance using WIDTH9.
 For the derived atmospheric parameters and oxygen abundance. using the SYNSPEC code. we could not obtain a good fit to such lines.," For the derived atmospheric parameters and oxygen abundance, using the SYNSPEC code, we could not obtain a good fit to such lines."
 The (total) equivalent width of the O I triplet in the spectrum of 122023 isLOLA., The (total) equivalent width of the O I triplet in the spectrum of I22023 is.
 This is comparable to the equivalent width of the O I triplet in the spectrum of the BI.51Ia hot post- AGB star. LSII-34726 (Garcíaa- Lario et al.," This is comparable to the equivalent width of the O I triplet in the spectrum of the B1.5Ia hot $-$ AGB star, $+$ $^{\circ}$ 26 $-$ Lario et al."
 1997b; Arkhipova et al., 1997b; Arkhipova et al.
 2001)., 2001).
 The O I triplet at. 47773A is known to be sensitive to NLTE effects., The O I triplet at $\lambda$ is known to be sensitive to NLTE effects.
 Using log e(O) 2 8.90 derived from the O II lines in the spectrum of 122023 (Table 6a). we could not obtain a fit to the O I triplet.," Using log $\epsilon$ (O) = 8.90 derived from the O II lines in the spectrum of I22023 (Table 6a), we could not obtain a fit to the O I triplet."
 Several N II and two Ne I lines were identified in 192023., Several N II and two Ne I lines were identified in I22023.
 The abundances of these lines were estimated using WIDTH9., The abundances of these lines were estimated using WIDTH9.
" Once again. the mismatch between the strong (W, > ) observed and synthetic N II lines indicates the need for à NLTE analysis of the spectrum of this star."," Once again, the mismatch between the strong $_{\lambda}$ $\ge$ ) observed and synthetic N II lines indicates the need for a NLTE analysis of the spectrum of this star."
 Only one Mg II line could be identified in the spectrum of the star., Only one Mg II line could be identified in the spectrum of the star.
 This line is blended with ΑΙ III8)., This line is blended with Al III(8).
 Since the Al III abundance in 122023 is uncertain (see below). we did not attempt to estimate the magnesium abundance from the blended Mg I(4) line.," Since the Al III abundance in I22023 is uncertain (see below), we did not attempt to estimate the magnesium abundance from the blended Mg II(4) line."
 Only four Al HI lines could be identified., Only four Al III lines could be identified.
 Three of these lines are blends with other species., Three of these lines are blends with other species.
 We therefore. estimated the aluminium abundance from the single 5722.730 ΑΙ ΠΙΟ) line using spectrum synthesis analysis and derived log e(Alj=6.79.," We therefore, estimated the aluminium abundance from the single 5722.730 Al III(2) line using spectrum synthesis analysis and derived log $\epsilon$ (Al)=6.79."
" This abundance from a single line with W,.=78.6mA may be treated as an upper limit.", This abundance from a single line with $_{\lambda}$ may be treated as an upper limit.
611.,-CDM.
611.l,-CDM.
611.lo,-CDM.
611.loo,-CDM.
611.look,-CDM.
he skvspectra from the skvlenslets. taken simultaneously with our galaxy spectra.,"the skyspectra from the skylenslets, taken simultaneously with our galaxy spectra."
 Providing the sky spectra of our dank skvfields as templates gave worse results. even though hese fields covered a larger sky area than the skvlenslets ancl herefore had higher S/N.," Providing the sky spectra of our blank skyfields as templates gave worse results, even though these fields covered a larger sky area than the skylenslets and therefore had higher $S/N$."
 This indicates the importance of obtaining simultaneous skyspectra over high S/N. to avoid mismatch due to the variability of the night sky.," This indicates the importance of obtaining simultaneous skyspectra over high $S/N$, to avoid mismatch due to the variability of the night sky."
 Finally. we tested the influence of subtracting residual galaxy light that could be present in our sky. spectra. since the skvlenslets were pointing at 6 - 7 £4. in the galaxy.," Finally, we tested the influence of subtracting residual galaxy light that could be present in our sky spectra, since the skylenslets were pointing at 6 - 7 $R_e$ in the galaxy."
 The ealaxy light at these distances of the nucleus is however very laint. approximately 3-4 mag/arcsec?. fainter than in the regions where we measure our kinematics.," The galaxy light at these distances of the nucleus is however very faint, approximately 3-4 $^2$ fainter than in the regions where we measure our kinematics."
 We simulated the effect of subtracting such a weak Gaussian. absorption line from our observed line and found that the maximal error we can introduce in this way is S km/s in our measured velocity dispersions., We simulated the effect of subtracting such a weak Gaussian absorption line from our observed line and found that the maximal error we can introduce in this way is 8 km/s in our measured velocity dispersions.
 This is well within our error bars. and we conclude that this effect is negligible.," This is well within our error bars, and we conclude that this effect is negligible."
 lt turned out that one of our fields in NCC 3379. 3N. cid not have sullicient signal to measure the LOSVD.," It turned out that one of our fields in NGC 3379, 3N, did not have sufficient signal to measure the LOSVD."
 The results for the other fields can be found in Figure 3. ancl Table 3.., The results for the other fields can be found in Figure \ref{fig:ppxf_3379} and Table \ref{tab:kin3379}.
 The last four columns for ‘Table 3 show the results. [or the LOSVD if we restrict our fit to the first two moments. instead. of fitting up to fy.," The last four columns for Table \ref{tab:kin3379} show the results for the LOSVD if we restrict our fit to the first two moments, instead of fitting up to $h_4$."
 The results for both fits agree within the errors., The results for both fits agree within the errors.
 Correcting for. barvcentric motion. we find that the svstemic velocity oof the galaxy measured from the central field (presented by Iimsellem et al. 2004))," Correcting for barycentric motion, we find that the systemic velocity of the galaxy measured from the central field (presented by Emsellem et al. \nocite{2004MNRAS.352..721E}) )"
 is 930 + 2 km/s. This agrees with, is 930 $\pm$ 2 km/s. This agrees with
Supernova1993J.. in the galaxyΜδΙ.. has been one of the brightest supernovae ever in the radio band.,"Supernova, in the galaxy, has been one of the brightest supernovae ever in the radio band."
 The peak of emission at GGHz was ~100 mmly (e.g. Weiler et al. 2007)).," The peak of emission at GHz was $\sim100$ mJy (e.g. Weiler et al. \cite{Weiler2007}) ),"
 much larger than typical peak flux densities of radio supernovae., much larger than typical peak flux densities of radio supernovae.
 The large flux density of this supernova. together with its high declination. allowed for long observing campaigns with VLBI.," The large flux density of this supernova, together with its high declination, allowed for long observing campaigns with VLBI."
 Two research groups (one led by N. Bartel and the other one led by J.M. Marcaide) have monitored this supernova with the VLBI technique. from 1993 (Marcaide et al. 1994.," Two research groups (one led by N. Bartel and the other one led by J.M. Marcaide) have monitored this supernova with the VLBI technique, from 1993 (Marcaide et al. \cite{Marcaide1994},"
 Bartel et al. 199-44)), Bartel et al. \cite{Bartel1994}) )
 to 2005., to 2005.
 Different results on the structure and expansion of the radio-emitting region of 11993J have been reported by both groups. based on the subset of VLBI data taken by each group.," Different results on the structure and expansion of the radio-emitting region of 1993J have been reported by both groups, based on the subset of VLBI data taken by each group."
 Mareaide et al. (1997)), Marcaide et al. \cite{Marcaide1997}) )
" reported the first evidence of deceleration in the shell expansion (1.0... Rx7"", see Chevalier 1982a)). with an estimated expansion index of i =0.8640.02."," reported the first evidence of deceleration in the shell expansion (i.e., $R \propto t^m$, see Chevalier \cite{Chevalier1982a}) ), with an estimated expansion index of $m = $ $\pm$ 0.02."
 Bartel et al. (2002)), Bartel et al. \cite{Bartel2002}) )
 confirmed a deceleration. but claimed up to four changes in the value of ii corresponding to four different expansion periods and interpreted the changes in the expansion index as changes 1n the mass-loss wind of the progenitor star through the pre-supernova stage.," confirmed a deceleration, but claimed up to four changes in the value of $m$ corresponding to four different expansion periods and interpreted the changes in the expansion index as changes in the mass-loss wind of the progenitor star through the pre-supernova stage."
 However. from their set of VLBI observations. which range from day 182 to day 3867 after explosion. Mareaide et al. (2009))," However, from their set of VLBI observations, which range from day 182 to day 3867 after explosion, Marcaide et al. \cite{Marcaide2009}) )"
 found a wavelength-dependent expansion curve that can be modeled using only one expansion index Gu= 0.86) for their low-frequency data data GGHz) and two expansion indices (0.86 and 0.79. separated by a break time on day ~1500 after explosion) for the data at all higher frequencies.," found a wavelength-dependent expansion curve that can be modeled using only one expansion index $m = 0.86$ ) for their low-frequency data data GHz) and two expansion indices (0.86 and 0.79, separated by a break time on day $\sim$ 1500 after explosion) for the data at all higher frequencies."
 These authors interpret the frequency-dependent expansion curve as being caused by (the possible combination of) two effects: 1) a changing (and frequency-dependent) opacity to the radio emission by the supernova and 2) a radial drop in the amplified magnetic fields inside the radiating region combined with the finite sensitivity of the VLBI observations., These authors interpret the frequency-dependent expansion curve as being caused by (the possible combination of) two effects: 1) a changing (and frequency-dependent) opacity to the radio emission by the supernova and 2) a radial drop in the amplified magnetic fields inside the radiating region combined with the finite sensitivity of the VLBI observations.
 In this paper (Paper D. we report on a homogeneous analysis of the complete set of available VLBI observations of SN11993J (69 epochs). using different approaches to minimizing the effects of any possible bias in the data analysis.," In this paper (Paper I), we report on a homogeneous analysis of the complete set of available VLBI observations of 1993J (69 epochs), using different approaches to minimizing the effects of any possible bias in the data analysis."
 We studied the details of the expansion curve at several frequencies and the evolution. of the structure of the radio shell throughout the history of the 11993J radio emission. with unprecedented time resolution and coverage.," We studied the details of the expansion curve at several frequencies and the evolution of the structure of the radio shell throughout the history of the 1993J radio emission, with unprecedented time resolution and coverage."
 We confirm earlier findings reported in Marcaide et al. (2009)), We confirm earlier findings reported in Marcaide et al. \cite{Marcaide2009}) )
 and report a model of the expansion curve compatible with the shell sizes obtained using different approaches., and report a model of the expansion curve compatible with the shell sizes obtained using different approaches.
 We also present a study of the distribution and evolution of inhomogeneities inside the shell., We also present a study of the distribution and evolution of inhomogeneities inside the shell.
 In another publication (Martí-Vidal et al. 2010..," In another publication (Martí–Vidal et al. \cite{PaperII},"
 hereafter Paper ID). we present a new simulation code able to simultaneously model the expansion curve and the radio light curves of 11993J reported by Weiler et al. (2007)).," hereafter Paper II), we present a new simulation code able to simultaneously model the expansion curve and the radio light curves of 1993J reported by Weiler et al. \cite{Weiler2007}) )."
 We then present the extensions to the Chevalier model (Chevalier 1982a.. 1982b)) to satisfactorily fit all the radio data.," We then present the extensions to the Chevalier model (Chevalier \cite{Chevalier1982a}, \cite{Chevalier1982b}) ) to satisfactorily fit all the radio data."
 In Sect., In Sect.
 2. we describe the complete set of VLBI observations of 11993J. most of which we have re-analyzedinitio.," \ref{II} we describe the complete set of VLBI observations of 1993J, most of which we have re-analyzed."
 In Sect., In Sect.
 3.1 we report on the location and proper motion of the radio shell., \ref{III} we report on the location and proper motion of the radio shell.
 In Sect., In Sect.
 3.2. we report on the complete expansion curve. obtained with different approaches. and in Sect.," \ref{IV} we report on the complete expansion curve, obtained with different approaches, and in Sect."
 3.4 we discuss the departure in the evolution of the supernova structure from self-similarity., \ref{V} we discuss the departure in the evolution of the supernova structure from self-similarity.
 In Table | we show the complete set of available VLBI observations of 11993J. made from year 1993 through the end of year 2005.," In Table 1 we show the complete set of available VLBI observations of 1993J, made from year 1993 through the end of year 2005."
 There is à total of 69 observing epochs. many of them made at several frequencies.," There is a total of 69 observing epochs, many of them made at several frequencies."
 All these observations used global VLBI arrays., All these observations used global VLBI arrays.
 In nearly all of them (except for some epochs in 1993 and 1994). the whole VLBA (10 identical antennas of mm diameter spread over the USA) and the Phased-VLA (equivalent to a ~ mm antenna in New Mexico. USA) were used.," In nearly all of them (except for some epochs in 1993 and 1994), the whole VLBA (10 identical antennas of m diameter spread over the USA) and the Phased-VLA (equivalent to a $\sim$ m antenna in New Mexico, USA) were used."
 Other antennas observed less often (each antenna participated in around of the epochs): Green, Other antennas observed less often (each antenna participated in around of the epochs): Green
where and where we require that (he sign of the square root in equation (20)) is to be chosen such that its real part GV)>0.,where and where we require that the sign of the square root in equation \ref{k1_eq}) ) is to be chosen such that its real part $\Re(K_1) > 0$.
 In order to obtain equation (19)). we apply the boundary condition vy.—0 when —x.," In order to obtain equation \ref{v1zsol_eq}) ), we apply the boundary condition $v_{1z} \rightarrow 0$ when $|z| \rightarrow \infty$."
 In general. solutions that are either midplane-syiunetric or nudplane-antisvnunnelric are possible.," In general, solutions that are either midplane-symmetric or midplane-antisymmetric are possible."
 Some previous work exanple).. has explored only (he case in which (4. is an odd function of z. so (0)=0.," Some previous work \citep[for example]{ish02}, has explored only the case in which $v_{1z}$ is an odd function of $z$, so $v_{1z}(0)=0$."
 We shall locus much of our attention on (he case in which ey. is an even function. and explore the odd case only lor reference.," We shall focus much of our attention on the case in which $v_{1z}$ is an even function, and explore the odd case only for reference."
 With ey-(2)στὴν Ay= and ele=οale.," With $v_{1z}(z)=v_{1z}(-z)$, $A_1=A_3$ and $A_{2+}=A_{2-}=A_2$."
 To relate Aly and «ο. we use the fact that the vertical displacement of a lagrangian point. 02=104:/(—hyi). must be a continuous function of 2 across anv interface in the flow.," To relate $A_1$ and $A_2$, we use the fact that the vertical displacement of a lagrangian point, $\delta z = \mi v_{1z}/(\omega - k_y V)$, must be a continuous function of $z$ across any interface in the flow."
 Applving this condition at the laver discontinuitw. this implies By integrating equation (17)) across the z=J£; interface. and substituting tlie solutions (19)). we obtain the dispersion relation," Applying this condition at the layer discontinuity, this implies By integrating equation \ref{v1z_eq}) ) across the $z=H_d$ interface, and substituting the solutions \ref{v1zsol_eq}) ), we obtain the dispersion relation"
neutrinos [from (he disk are all electron (wpe since nucleon pair capture is a dominant neutrino process in the disk.,neutrinos from the disk are all electron type since nucleon pair capture is a dominant neutrino process in the disk.
 The mean free path. A. of neutrinos in the corona lor neutrino-electron scattering is then (Landan Lilshits 1976) where 2. p. f. ave the velocity. momentum. and the distribution function of electrons. respectively. and. ji([ds cosine of the angle between the direction of neutrino and electron.," The mean free path, $\lambda$, of neutrinos in the corona for neutrino-electron scattering is then (Landau Lifshits 1976) where $\beta$ , $p$, $f_\mathrm{e}$ are the velocity, momentum, and the distribution function of electrons, respectively, and $\mu$ is cosine of the angle between the direction of neutrino and electron."
 We assume (hat the corona consists of pure electron positron plasma in thermal equilibrium by electromagnetic process so Chat the distribution function of pair should be Fermi-Dirac (wpe. fo=fexp(\/p?+mz/T)1] I.," We assume that the corona consists of pure electron positron plasma in thermal equilibrium by electromagnetic process so that the distribution function of pair should be Fermi-Dirac type, $f_\mathrm{e} = [\exp({\sqrt{p^2 + m_\mathrm{e}^2}/T}) + 1] ^{-1}$ ."
 In the relativistie limit. where T.>my. the total coronal depth for neutrino-electron scattering is where {ες is the energy of incident neutrino in the laboratory frame.," In the relativistic limit, where $T_\mathrm{c} \gg m_\mathrm{e}$, the total coronal depth for neutrino-electron scattering is where $E_\nu$ is the energy of incident neutrino in the laboratory frame."
 Hence. higher energy neutrinos have more chances to collide with coronal electrons than lower ones.," Hence, higher energy neutrinos have more chances to collide with coronal electrons than lower ones."
 We calculate the neutrino spectra scattered by coronal electrons and positrons by using Monte Carlo method. following Pozdnvakov. Sobol. Sunvaev (1977) and Lin. Mineshige. Ohsuga (2003).," We calculate the neutrino spectra scattered by coronal electrons and positrons by using Monte Carlo method, following Pozdnyakov, Sobol, Sunyaev (1977) and Liu, Mineshige, Ohsuga (2003)."
 We first set the weight of neutrino wy=1 [or a given thermal neutrino with energyv. £5. which has the Fermi-Dirac distribution.," We first set the weight of neutrino $w_0=1$ for a given thermal neutrino with energy, $E_\nu$, which has the Fermi-Dirac distribution."
" We calculate the probability Lor neulrinos passing through the corona. /4,=exp(—7/cosa). where a is the angle between neutrino direction and the z-axis which is perpendicular to the coronal plain."," We calculate the probability for neutrinos passing through the corona, $P_0 = \exp (-\tau/\cos \alpha)$, where $\alpha$ is the angle between neutrino direction and the $z$ -axis which is perpendicular to the coronal plain."
 Then i411 is (he (ransmillecl portion of neutrinos and (the remaining wy=wo(l—πι) is the portion of neutrinos scattered al least once., Then $w_0 P_0$ is the transmitted portion of neutrinos and the remaining $w_1=w_0(1-P_0)$ is the portion of neutrinos scattered at least once.
" Let (6,=iw,4(1—P,4) be the portion of neutrinos experiencing the »-th scattering.", Let $w_{n} = w_{n-1} (1 - P_{n-1})$ be the portion of neutrinos experiencing the $n$ -th scattering.
 We continue the calculation until the weight. το. becomes sullicientlv small.," We continue the calculation until the weight, $w_n$, becomes sufficiently small."
 Repeating the same procedures for sufficiently large number of neutrinos. we can calculate emergent spectra by collecting neutrinos going upward through the coronal surface al 2=f. while downward neutrinos crossing (he lower boundary. are re-absorbed by the disk body. thereby. heating thedisk.," Repeating the same procedures for sufficiently large number of neutrinos, we can calculate emergent spectra by collecting neutrinos going upward through the coronal surface at $z=H$, while downward neutrinos crossing the lower boundary are re-absorbed by the disk body, thereby heating thedisk."
 We alsocalculate the neutrino pair emission by weak interaction in the pair plasnia., We alsocalculate the neutrino pair emission by weak interaction in the pair plasma.
"radiation background affects their arrangement; i.e., the dynamical properties of the galaxy.","radiation background affects their arrangement; i.e., the dynamical properties of the galaxy."
" The first column of Table 2. shows a significant increase in the stellar mass within a 5 kpc radius as one moves from the Old UV to the New UV to the New UV+X case for all three galaxies (and FG UV again lies between Old UV and New UV): the mass increase from Old UV to New UV is an average of ~10%,, and from New UV to New UV+X averages ~20%,, although this includes a slight decrease for Galaxy A. This increase is natural considering the enhanced gas flows we saw in Fig. 4::"," The first column of Table \ref{tab:stars} shows a significant increase in the stellar mass within a 5 kpc radius as one moves from the Old UV to the New UV to the New UV+X case for all three galaxies (and FG UV again lies between Old UV and New UV): the mass increase from Old UV to New UV is an average of $\sim$, and from New UV to New UV+X averages $\sim$, although this includes a slight decrease for Galaxy A. This increase is natural considering the enhanced gas flows we saw in Fig. \ref{fig:cgasacc-A100}:"
 once that gas reaches the center it must cool and form stars there., once that gas reaches the center it must cool and form stars there.
 Figure 7 shows the circular speed v2=GM(r)/r for the innermost 2 kpc of galaxy A at z—0., Figure \ref{fig:vcirc-A100-z0} shows the circular speed $v_c^2=GM(r)/r$ for the innermost 2 kpc of galaxy A at $z=0$.
 We see that the stars become more centrally concentrated as one moves up in background intensity or hardness (from Old UV to FG UV to New UV to New UV+X); this is a general feature for z«2.5., We see that the stars become more centrally concentrated as one moves up in background intensity or hardness (from Old UV to FG UV to New UV to New UV+X); this is a general feature for $z<2.5$.
 Residual gas also increases in the same way., Residual gas also increases in the same way.
" Figure 8 shows the number of discrete (i.e., non- stellar systems with masses >3x105 identified substructure)from the star particles with the HOP halo-M5finding algorithm for the galaxy A simulation (within 2 Mpc comoving) over time."," Figure \ref{fig:haloint-A100} shows the number of discrete (i.e., non-substructure) stellar systems with masses $\geq 3\times 10^8 M_\odot$ identified from the star particles with the HOP halo-finding algorithm for the galaxy A simulation (within 2 Mpc comoving) over time."
" Many past simulations have shown that photoionization can prevent the formation of smaller stellar systems, and indeed we find that the number of total stellar systems decreases from Old UV to New UV to New UV+X, although we find the effect of increased high-z UV on the number of independent halos doesn't persist past z—0.5, where merger events consume several small halos that have formed in the Old UV case."," Many past simulations have shown that photoionization can prevent the formation of smaller stellar systems, and indeed we find that the number of total stellar systems decreases from Old UV to New UV to New UV+X, although we find the effect of increased high-z UV on the number of independent halos doesn't persist past $z=0.5$, where merger events consume several small halos that have formed in the Old UV case."
" Interestingly, FG UV shows fewer small halos than New UV: we see the suppression of halo formation (1.5>z 0.75) combined with late mergers < "," Interestingly, FG UV shows fewer small halos than New UV: we see the suppression of halo formation $1.5>z>0.75$ ) combined with late mergers $z<0.5$ )."
In all cases star formation in smaller halos is preferentially(z0.5). suppressed., In all cases star formation in smaller halos is preferentially suppressed.
" A future paper (Hambrick Ostriker 2009, in prep.)"," A future paper (Hambrick Ostriker 2009, in prep.)"
 will deal in more detail with the effect of radiation backgrounds on the mass spectrum of galaxies., will deal in more detail with the effect of radiation backgrounds on the mass spectrum of galaxies.
" Figure 9 shows the half-mass radius for stars within 30kpc for the three radiation models with galaxy A. The major effect is secular growth from z=8 to the present, combined with the stochastic effects of major mergers."," Figure \ref{fig:hmr-A100} shows the half-mass radius for stars within 30kpc for the three radiation models with galaxy A. The major effect is secular growth from $z=8$ to the present, combined with the stochastic effects of major mergers."
" found in à detailed examination of galaxy A the Old UV background, and including SN (withthat minor mergers are primarily responsible for the feedback)factor of >3 increase in halfmass radius from z=3 to 0."," found in a detailed examination of galaxy A (with the Old UV background, and including SN feedback) that minor mergers are primarily responsible for the factor of $\gtrsim 3$ increase in half-mass radius from $z=3$ to 0."
" The radiation background does have an effect, however: New UV and New UV+X show substantially (> 25%) smaller radii compared to Old UV at all z< 2.5, modulo the intermittent merger effects; FG UV has negligible difference from Old UV, and in fact has a slightly larger radius at z=0."," The radiation background does have an effect, however: New UV and New UV+X show substantially $\gtrsim25\%$ ) smaller radii compared to Old UV at all $z<2.5$ , modulo the intermittent merger effects; FG UV has negligible difference from Old UV, and in fact has a slightly larger radius at $z=0$."
" For zZ3 the picture looks somewhat different: specifically, in our snapshot at z=2.9, the New UV model for galaxy A has a13%,, and FG UV a stellar half-mass radius than Old UV, (although New UV+X is smaller than Old UV by 13%))."," For $z\gtrsim 3$ the picture looks somewhat different: specifically, in our snapshot at $z=2.9$, the New UV model for galaxy A has a, and FG UV a stellar half-mass radius than Old UV, (although New UV+X is smaller than Old UV by )."
" Further, New UV and New UV+X both show a reduction in peak circular speed compared to Old UV at that redshift."," Further, New UV and New UV+X both show a reduction in peak circular speed compared to Old UV at that redshift."
" found, using a modified background similar to our “New UV with cutoff”, that simulated galaxies have half-light radii that are too small and peak circular speeds that are too large compared to observed galaxies at redshift 3."," found, using a modified background similar to our “New UV with cutoff”, that simulated galaxies have half-light radii that are too small and peak circular speeds that are too large compared to observed galaxies at redshift 3."
" However, these authors also used a more stringent star-formation criterion: when they repeated their simulations using the same density threshold as in this work, their results at z—3 agreed with obesrvations (Ryan Joung, priv."," However, these authors also used a more stringent star-formation criterion: when they repeated their simulations using the same density threshold as in this work, their results at $z=3$ agreed with obesrvations (Ryan Joung, priv."
 comm.)., comm.).
 We speculate that the increase in central concentration with increasing radiation is enhanced by reduced substructure in the galaxy (because small stellar halos aresuppressed) and correspondingly less dynamical friction., We speculate that the increase in central concentration with increasing radiation is enhanced by reduced substructure in the galaxy (because small stellar halos aresuppressed) and correspondingly less dynamical friction.
" To test this hypothesis, we have constructed a"," To test this hypothesis, we have constructed a"
Substituting tliese expressious iuto the equatious of motion. aud recognizing that zx. we obtain For small perturbations. we expect that z. ο. €. 7 aud A. will change much more slowly than A does.,"Substituting these expressions into the equations of motion, and recognizing that $f=\lambda-\varpi$ , we obtain For small perturbations, we expect that $\varpi$ , $\Omega$, $e$, $i$ and $\lambda_\Sun$ will change much more slowly than $\lambda$ does."
 Thus. to obtain the long-term secular evolution of the orbital elements. we μιαν average these expression over a single orbit.," Thus, to obtain the long-term secular evolution of the orbital elements, we may average these expression over a single orbit."
 However. in doiug this. we must take care to account for Saturu’s shadow. which blocks the light from theSuu during a fraction ol the particle's orbit e.," However, in doing this, we must take care to account for Saturn's shadow, which blocks the light from theSun during a fraction of the particle's orbit $\epsilon$ ."
 The appropriate orbit-averaged equatious of motiou are: where d(e)=1—e. f(e)=1—¢+sin(2z:e)/6x and gle)=siu(se)/5 (see Appeucix).," The appropriate orbit-averaged equations of motion are: where $d(\epsilon)=1-\epsilon$, $f(\epsilon)=1-\epsilon+\sin(2\pi\epsilon)/6\pi$ and $g(\epsilon)=\sin(\pi\epsilon)/\pi$ (see Appendix)."
 For au oblate planet like Saturn. these equatious of motion are incomplete because they do uot take into account the steady precession in tlie periceuter and node caused by the plauet's finite oblateness. which augments the motion of zaud Q.," For an oblate planet like Saturn, these equations of motion are incomplete because they do not take into account the steady precession in the pericenter and node caused by the planet's finite oblateness, which augments the motion of $\varpi$and $\Omega$ ."
 Thefull equatious of motion are therefore:, Thefull equations of motion are therefore:
Fig.,Fig.
 7 reproduces the key features displayed by the real data in. for example. Fig.," 7 reproduces the key features displayed by the real data in, for example, Fig."
 3., 3.
 Specifically. even for input? =—2. both the ERS and HUDF simulated samples yield galaxies with apparent values of 7 as blue as; &5 in the faintest luminosity/magnitude bin probed by each sample.," Specifically, even for input $\beta = -2$, both the ERS and HUDF simulated samples yield galaxies with apparent values of $\beta$ as blue as $\beta 
\simeq -5$ in the faintest luminosity/magnitude bin probed by each sample."
 In addition. several of these apparently ultra-blue sources are classified as ROBUST.," In addition, several of these apparently ultra-blue sources are classified as ROBUST."
 By contrast. while artificially red sources up to 70 are produced by the photometric uncertainties. ultra-red sources are much less prevalent. and red ROBUST sources are very rare (only one ROBUST source in this simulation is retrieved with 937—1).," By contrast, while artificially red sources up to $\beta \simeq 0$ are produced by the photometric uncertainties, ultra-red sources are much less prevalent, and red ROBUST sources are very rare (only one ROBUST source in this simulation is retrieved with $\beta > -1$ )."
 The effect of these distributions of retrieved 9«3 values on the average deduced value of (7) as a function of UV luminosity is shown in Fig., The effect of these distributions of retrieved $\beta$ values on the average deduced value of $\langle \beta \rangle$ as a function of UV luminosity is shown in Fig.
 8. again for both the 97= 2and 9;—2.5 input scenarios.," 8, again for both the $\beta = -2$ and $\beta = -2.5$ input scenarios."
 The upper panel of Fig., The upper panel of Fig.
 8 is remarkably similar to the 2στ7 points plotted in Fig., 8 is remarkably similar to the $z \simeq 7$ points plotted in Fig.
 | of Bouwens et al. (, 1 of Bouwens et al. (
010b). and to those given in Fig.,"2010b), and to those given in Fig."
 6 of Finkelstein et al. (, 6 of Finkelstein et al. (
2010).,2010).
 Here the analysis of our 1Ξ J2simulation has resulted in an entirely artificial. apparently monotonic luminosity dependence of (2. with 1:2? approaching 3in the faintest luminosity bin.," Here the analysis of our $\beta = -2$ simulation has resulted in an entirely artificial, apparently monotonic luminosity dependence of $\langle \beta \rangle$, with $\langle \beta \rangle$ approaching $-3$ in the faintest luminosity bin."
 Only in the brightest bin has the true input value of . been successfully retrieved., Only in the brightest bin has the true input value of $\beta$ been successfully retrieved.
" It is important to stress that the fact we recover (25=24 at Mie,c19.59 does not contradict the value of 63)~2.12 we measured from the =7 data in this bin. as given in Fig."," It is important to stress that the fact we recover $\langle \beta \rangle 
\simeq -2.4$ at $M_{UV} \simeq -19.5$ does not contradict the value of $\langle \beta \rangle \simeq -2.12$ we measured from the $z = 7$ data in this bin, as given in Fig."
 6 and Table |., 6 and Table 1.
 As already 2.discussed. to try to minimize bias. these measurements were limited to objects with at least one >5- σ΄ detection in the near-infrared. and even the A4c19.ut luminosity bin contains some less significant detections which can bias the result to the blue unless filtered out.," As already discussed, to try to minimize bias, these measurements were limited to objects with at least one $>8$ $\sigma$ detection in the near-infrared, and even the $M_{UV} \simeq -19.5$ luminosity bin contains some less significant detections which can bias the result to the blue unless filtered out."
 Thus. our simulation simply implies that. with the depth of WFC3/IR data analysed here. unless such quality control is applied. a true 9&—2 will result in an accurately measured 67)= Dat Ady= 205.8 somewhat biased measurement of 63)&— 2.1at Adyo19.5. and a severely biased measurement of 67)2.3 at My& 18.5.," Thus, our simulation simply implies that, with the depth of WFC3/IR data analysed here, unless such quality control is applied, a true $\beta \simeq -2$ will result in an accurately measured $\langle \beta \rangle = -2$ at $M_{UV} \simeq -20.5$, a somewhat biased measurement of $\langle \beta \rangle \simeq -2.4$ at $M_{UV} \simeq -19.5$, and a severely biased measurement of $\langle \beta \rangle \simeq -3$ at $M_{UV} \simeq -18.5$ ."
 Thus. our simulation suggests that the apparent luminosity dependence of with Αι reported by both Bouwens et al. (," Thus, our simulation suggests that the apparent luminosity dependence of $\beta$ with $M_{UV}$ reported by both Bouwens et al. ("
2010b) and Finkelstein et al. (,2010b) and Finkelstein et al. (
2010) (from the same depth of data) is at least inconsistent with a true value of 3—2. independen of luminosity.,"2010) (from the same depth of data) is at least with a true value of $\beta \simeq -2$, independent of luminosity."
 The lower panel in Fig., The lower panel in Fig.
 8 simply shows how even bluer values. with apparent (9)<3. inevitably result when the input value of ης 2.5.," 8 simply shows how even bluer values, with apparent $\langle \beta \rangle < -3$, inevitably result when the input value of $\beta$ is $-2.5$ ."
" However. this is clearly inconsistent with the data. as the input value of 7=2.5 is of course correctly recovered from the simulation in the brightest— luminosity bin. and this is inconsisten with the observed value of 7= Jat M4,=20.5."," However, this is clearly inconsistent with the data, as the input value of $\beta = -2.5$ is of course correctly recovered from the simulation in the brightest luminosity bin, and this is inconsistent with the observed value of $\beta = -2$ at $M_{UV} = -20.5$."
 Interestingly. the retrieved value of (7) in the faintes luminosity bin is not the full 0.5 lower in the lower panel of Fig.," Interestingly, the retrieved value of $\langle \beta \rangle$ in the faintest luminosity bin is not the full 0.5 lower in the lower panel of Fig."
 8 as compared to the upper panel., 8 as compared to the upper panel.
 This implies that one cannot easily correct for the bias in a unique way. and that a measured value of 723 in this luminosity bin could be consistent with a true 97= ors=2.5 within the errors.," This implies that one cannot easily correct for the bias in a unique way, and that a measured value of $\langle \beta \rangle \simeq -3$ in this luminosity bin could be consistent with a true $\beta = -2$ or $\beta = -2.5$ within the errors."
 This simply reinforces the need to improve the depth of the WFC3/IR data to enable higher signal:noise measurements of ;$ in this crucial faint luminosity bin , This simply reinforces the need to improve the depth of the WFC3/IR data to enable higher signal:noise measurements of $\beta$ in this crucial faint luminosity bin at $z \simeq 7$.
"To explore the origin of the 7.77 bias so clearly displayed by the analysis of our simulations. we take advantage of the fact that the “true” input UV luminosity of every simulated galaxy is known. and explore how derived «7 relates to the level of ""flux-boosting"" experienced by the simulated sources."," To explore the origin of the $\beta$ ” bias so clearly displayed by the analysis of our simulations, we take advantage of the fact that the “true” input UV luminosity of every simulated galaxy is known, and explore how derived $\beta$ relates to the level of “flux-boosting” experienced by the simulated sources."
 This is shown in Fig., This is shown in Fig.
 9. where the extracted 9«2 value for all of the reclaimed ERS and HUDF high-redshift ;;=—2 simulated galaxies is plotted against UV luminosity (= ραπ) flux boost. in magnitudes (here. a positive value of “Boost” means the reclaimed «ος magnitude isbrighter than the input value by the plotted magnitude difference).," 9, where the extracted $\beta$ value for all of the reclaimed ERS and HUDF high-redshift $\beta = -2$ simulated galaxies is plotted against UV luminosity $\equiv J$ -band) flux boost, in magnitudes (here, a positive value of “Boost” means the reclaimed $J_{125}$ magnitude is than the input value by the plotted magnitude difference)."
 Both the ERS and HUDF simulated galaxies behave in the same way and show that the extremely blue values of 7 almost all result fromsources which haveentered the sample because their “true” Jo; magnitudes have been boosted by a few tenths of a, Both the ERS and HUDF simulated galaxies behave in the same way and show that the extremely blue values of $\beta$ almost all result fromsources which haveentered the sample because their “true” $J_{125}$ magnitudes have been boosted by a few tenths of a
1 Introductio,the$\Lambda$ is significantly negatively polarized and $\bar{\Lambda}$ is not .
n, The
photons scattered. to energies 110 keV. for each 7.,"photons scattered to energies $1-10$ keV, for each $i$."
 Also tabulated is the average number of scatterings cach photon undergoes., Also tabulated is the average number of scatterings each photon undergoes.
 As the CMD seed photons have average energies lower than the clisk seed. photons. the average. number of scatterings needed to reach X-ray energies is expected to be higher.," As the CMB seed photons have average energies lower than the disk seed photons, the average number of scatterings needed to reach X-ray energies is expected to be higher."
 As ds evident in ‘Table 2.. the X-ray. polarization of the scattered. CALB photons follows the sume. trend with viewing inclination angle as for the case of disk Comptonization.," As is evident in Table \ref{CMBpolarization}, the X-ray polarization of the scattered CMB photons follows the same trend with viewing inclination angle as for the case of disk Comptonization."
 However. the predicted. P at cach 7 are slightly clifferent.," However, the predicted $P$ at each $i$ are slightly different."
 We attribute this to the dillerence in the angular distribution of photon injection and the polarization condition in the last scattering event., We attribute this to the difference in the angular distribution of photon injection and the polarization condition in the last scattering event.
 The disk photons are unpolarized and. emitted. near the base of the jet and renee the disk seed photons participate in the scattering oocess have a narrow range of initial pitch angles., The disk photons are unpolarized and emitted near the base of the jet and hence the disk seed photons participate in the scattering process have a narrow range of initial pitch angles.
 Given he relatively low scattering optical depth in the jet. οποίος that have a single scattering dominate ancl hence οποίος emerging at small 7 are mostly restricted to small-angle scatterings (analogous to total internal rellection).," Given the relatively low scattering optical depth in the jet, photons that have a single scattering dominate and hence photons emerging at small $i$ are mostly restricted to small-angle scatterings (analogous to total internal reflection)."
 The €MD photons are unpolarized. similar to the clisk οποίος.," The CMB photons are unpolarized, similar to the disk photons."
 However. the angular distribution of the CAIB seed. photons involved in the scattering djs. isotropically. around. the whole leneth of the jet. which is contrary o the more restrictive angular distribution of the clisk seed photons.," However, the angular distribution of the CMB seed photons involved in the scattering is isotropically around the whole length of the jet, which is contrary to the more restrictive angular distribution of the disk seed photons."
 For CAIB seattering. photons emerging at ow £4 consist of two main tvpes: single scattered photons with large deflection angles. and multiple scattered photons with small deflection angles in the last scattering.," For CMB scattering, photons emerging at low $i$ consist of two main types: single scattered photons with large deflection angles, and multiple scattered photons with small deflection angles in the last scattering."
 For the case of the single scattered: photons. the injected. photons are unpolarized. and the polarization is caused. by. large-angle scatterings.," For the case of the single scattered photons, the injected photons are unpolarized, and the polarization is caused by large-angle scatterings."
 In the case of multiple scattered photons. the pre-scattered. photons. which propagate approximately along the jet. are already polarized because of previous scatterings.," In the case of multiple scattered photons, the pre-scattered photons, which propagate approximately along the jet, are already polarized because of previous scatterings."
 Thus. the polarization is slightly higher for scattering of CMD. photons than disk photons.," Thus, the polarization is slightly higher for scattering of CMB photons than disk photons."
 Scattered CAIB photons emerging at large ; mostly. undergo. many scattering with a large dellection angle in the last scattering., Scattered CMB photons emerging at large $i$ mostly undergo many scattering with a large deflection angle in the last scattering.
 As the pre-seattered photons in the last scattering are polarized. it is not surprising that the resulting polarization is higher than those of the disk photon scattering. whereas the pre-seattered photons in the last scattering ds less polarized.," As the pre-scattered photons in the last scattering are polarized, it is not surprising that the resulting polarization is higher than those of the disk photon scattering, whereas the pre-scattered photons in the last scattering is less polarized."
 The differences in the N-ray. polarizations for the IC disk. and CAIB models are only at;=19 and <35% at i=SO , The differences in the X-ray polarizations for the EC disk and CMB models are only at $i = 10\degree$ and $<3.5$ at $i = 80\degree$ .
Such a small difference makes it clillicult to ciscriminate between these two mocels from rav polarization observations alone., Such a small difference makes it difficult to discriminate between these two models from X-ray polarization observations alone.
 We expect the SSC polarization signature to diller significantly. from the EC case. since svnchrotron photons emitted by jet electrons are intrinsically polarized.," We expect the SSC polarization signature to differ significantly from the EC case, since synchrotron photons emitted by jet electrons are intrinsically polarized."
 In the EC case. Compton scattering polarizes the seed. photons. whereas in the case of SSC. scattering may depolarize the photons.," In the EC case, Compton scattering polarizes the seed photons, whereas in the case of SSC, scattering may depolarize the photons."
 The polarization vector e of the incident. svnchrotron photons is chosen so that it is perpendicular to both the magnetic field of the jet and the photon propagation vector OQ in the PE., The polarization vector $\mathbf{e}$ of the incident synchrotron photons is chosen so that it is perpendicular to both the magnetic field of the jet and the photon propagation vector $\mathbf{\Omega}$ in the PF.
 We consider a simple model. where the magnetic field is directed along the jet axis and we assume that the jet electrons spiral around the z-axis in the PE.," We consider a simple model, where the magnetic field is directed along the jet axis and we assume that the jet electrons spiral around the $z$ -axis in the PF."
 The intrinsic polarization in this case will be perpendicular to the projection of the jet axis in the sky., The intrinsic polarization in this case will be perpendicular to the projection of the jet axis in the sky.
 Realistically. the magnetic field may have a more complicated: configuration which could produce dillerent. polarization signatures.," Realistically, the magnetic field may have a more complicated configuration which could produce different polarization signatures."
 Here we demonstrate that the polarization of the SSC) photons may be clistinguishable from that of the EC emission even though the actual polarization signatures may cilfer for more complicated geometries and magnetic fields., Here we demonstrate that the polarization of the SSC photons may be distinguishable from that of the EC emission even though the actual polarization signatures may differ for more complicated geometries and magnetic fields.
 We investigate the elleet of dillerent. emission sites for the primary svnchrotron. photons., We investigate the effect of different emission sites for the primary synchrotron photons.
 We first inject Che seed photons uniformly throughout the jet and then from a particular dimensionless height d. where ¢=1 and ¢=0 corresponds to the top and base of the jet. respectively.," We first inject the seed photons uniformly throughout the jet and then from a particular dimensionless height $\zeta$, where $\zeta = 1$ and $\zeta = 0$ corresponds to the top and base of the jet, respectively."
 Uniform emission of svnchrotron seed photons in the jet is consistent with some observations of radio jets (c.g.ΙΣΤ.lxovalevetal. 2007).," Uniform emission of synchrotron seed photons in the jet is consistent with some observations of radio jets \citep[e.g. M87,][]{Kovalev07}."
. It may be attributed to continuous acceleration of electrons throughout the jet as a result. of stochastic reconnection events in a random magnetic field component (c.g.Blandford&Eichler1987)., It may be attributed to continuous acceleration of electrons throughout the jet as a result of stochastic reconnection events in a random magnetic field component \citep[e.g.][]{Blandford87}.
. On the other hand. emission. of svachrotron photons at. localized. sites within the jet could occur as a result of discrete shock events.," On the other hand, emission of synchrotron photons at localized sites within the jet could occur as a result of discrete shock events."
 There is some evidence for this from observations of brigh spots in radio jets (e.g.Ixataoka, There is some evidence for this from observations of bright spots in radio jets \citep[e.g.][]{Kataoka05}.
&Sta, Fig.
warz2005).. Fig. 5 shows the simulated. spectra of the total (see svnchrotron plus SSC) photons emerging at three cdilferen observer viewing angles 7., \ref{SSCspectrum} shows the simulated spectra of the total (seed synchrotron plus SSC) photons emerging at three different observer viewing angles $i$.
 Phe seed synchrotron photons are emitted uniformly throughout the jet in this case., The seed synchrotron photons are emitted uniformly throughout the jet in this case.
 Boosting occurs along the jet axis. hence most. photons are observec at small inclinations (/=10 case in E ," Boosting occurs along the jet axis, hence most photons are observed at small inclinations $i = 10\degree$ case in Fig. \ref{SSCspectrum}) )."
Fie., Fig.
 6 shows the polarization degree £2? of photons with energies between 1.10 keV emerging from a jet at dilleren inclination angles / and for different svnchrotron injection sites ¢., \ref{SSCpol} shows the polarization degree $P$ of photons with energies between $1-10$ keV emerging from a jet at different inclination angles $i$ and for different synchrotron injection sites $\zeta$.
 Fig., Fig.
 7 shows the corresponding average number of scatterings each photon undergoes before escaping the jet a 1H, \ref{SSCScatnum} shows the corresponding average number of scatterings each photon undergoes before escaping the jet at $i$.
 There is a clear correlation between number of scatterings and. depolarization of the seed. photons., There is a clear correlation between number of scatterings and depolarization of the seed photons.
 As the see svnchrotron photons are emitted into a narrow emission cone about cach electron's instantaneous. propagation direction along the z-axis (in the PE). photons which are observed to emerge from the jet at large 7 must undergo a Larger number of scatterings (see Fig 7)).," As the seed synchrotron photons are emitted into a narrow emission cone about each electron's instantaneous propagation direction along the $z$ -axis (in the PF), photons which are observed to emerge from the jet at large $i$ must undergo a larger number of scatterings (see Fig \ref{SSCScatnum}) )."
 Photons emerging from. small { qnay also experience a large number of scatterings when emitted. [rom close to the jet base because these photons propagate along the entire jet axis and thus see a large optical depth., Photons emerging from small $i$ may also experience a large number of scatterings when emitted from close to the jet base because these photons propagate along the entire jet axis and thus see a large optical depth.
 lies., Figs.
 G6 and 7 show that the polarization of ος jet photons is sensitive to the emission site of the seed. photons., \ref{SSCpol} and \ref{SSCScatnum} show that the polarization of SSC jet photons is sensitive to the emission site of the seed photons.
 We investigate three dillerent cases:photons emitted from, We investigate three different cases:photons emitted from
accurately checked by means of the detailed follow-up of the density profile for CDM haloes in simulations.,accurately checked by means of the detailed follow-up of the density profile for CDM haloes in simulations.
 Zhaoet and Muiüoz-Cuartasetal.(2011) carried out such a follow-up and found that the r; and M; NFW shape of accreting haloes are not constant but vary slightly with time., \citet{ZJMB09} and \citet{Mea} carried out such a follow-up and found that the $\rs$ and $\Ms$ NFW shape of accreting haloes are not constant but vary slightly with time.
 This was interpreted as evidence that the inner structure of accreting haloes is changing., This was interpreted as evidence that the inner structure of accreting haloes is changing.
 But this is not the only possible interpretation., But this is not the only possible interpretation.
" As the density profile of haloesfitted by the NFWprofile, even if the inner density profile did not change, its fit by the NFW law out to progressively larger radii should result in slightly different best values of r; and Mg."," As the density profile of haloes by the NFW, even if the inner density profile did not change, its fit by the NFW law out to progressively larger radii should result in slightly different best values of $\rs$ and $\Ms$."
 'To confirm that our interpretation for this change is correct we have followed in the log-log plane the evolution of accreting halos forced to grow at the typical accretion rate given by the excursion set formalism according to Salvador-Soléetal.(2007) model., To confirm that our interpretation for this change is correct we have followed in the log-log plane the evolution of accreting halos forced to grow at the typical accretion rate given by the excursion set formalism according to \citet{smgh07} model.
" The result, in the same cosmology as used in Zhaoetal.(2009), is shown in Figure 1.."," The result, in the same cosmology as used in \citet{ZJMB09}, is shown in Figure \ref{f0}."
" Even if, by construction, haloes grow inside-out without changing their instantaneous inner structure, when their density profile is fitted by a NFW profile, they are found to move in the log-log plane along one fixed direction over a distance that increases with current halo mass."," Even if, by construction, haloes grow inside-out without changing their instantaneous inner structure, when their density profile is fitted by a NFW profile, they are found to move in the log-log plane along one fixed direction over a distance that increases with current halo mass."
" This behaviour isidentical to that found by Zhaoetal.:: accreting halos moved in the log-log plane along straight lines with exactly the same universal slope (Msοςτς5) as in our experiment and ended up at z=0 lying along another straight line, with a different slope (Msοςτῇ9) also identical to that found in our experiment (see their Fig."," This behaviour is to that found by \citeauthor{ZJMB09}: accreting halos moved in the log-log plane along straight lines with exactly the same universal slope $\Ms\propto \rs^{1.65}$ ) as in our experiment and ended up at $z=0$ lying along another straight line, with a different slope $\Ms\propto \rs^{2.48}$ ) also identical to that found in our experiment (see their Fig."
 22)., 22).
 The fact that our constrained model reproduces the observed trend convincinglydemonstrates haloes primarily grow inside-out., The fact that our constrained model reproduces the observed trend convincinglydemonstrates haloes primarily grow inside-out.
" As we will show, by assuming that haloes grow inside- via PA and that the spherical total energy in spheres of fixed mass is conserved in the absence of shell-crossing, we are able to infer the radius enclosing a given mass in a halo from the spherical energy distribution of its progenitor protohalo."," As we will show, by assuming that haloes grow inside-out via PA and that the spherical total energy in spheres of fixed mass is conserved in the absence of shell-crossing, we are able to infer the radius enclosing a given mass in a halo from the spherical energy distribution of its progenitor protohalo."
" To do this, we will consider the virial relation (86)) for a perfectly uniform sphere with mass M, which implies W(M)=—3GM?/(5R), with spherical total energy £(M) equal to that of the system at turnaround £i(M) and null spherical surface term S(M), R(M)-..(90) 'This equation-= is the often used estimate for the radius encompassing mass M in spherical systems with an unknown internal mass distribution (Bryan and Norman 1998)."," To do this, we will consider the virial relation \ref{vir0l}) ) for a perfectly uniform sphere with mass $M$, which implies ${\cal
 W}(M)=-3GM^2/(5R)$, with spherical total energy ${\cal E}(M)$ equal to that of the system at turnaround ${\cal E}\col (M)$ and null spherical surface term ${\cal S}(M)$, R(M)=. This equation is the often used estimate for the radius encompassing mass $M$ in spherical systems with an unknown internal mass distribution (Bryan and Norman 1998)."
" In principle, R(M) is not believed to give an exact measure of that mass in the real virialised object as this has non-uniform density profile, its spherical total energy is equal to E(M) instead of £i(M) and its surface term is not null."," In principle, $R(M)$ is not believed to give an exact measure of that mass in the real virialised object as this has non-uniform density profile, its spherical total energy is equal to ${\cal E}(M)$ instead of ${\cal E}\col(M)$ and its surface term is not null."
" Yet, as shown below, as a consequence ofPA, the inaccuracies above exactly cancel and both radii turn out to fully coincide."," Yet, as shown below, as a consequence of, the inaccuracies above exactly cancel and both radii turn out to fully coincide."
" 'To see this, we will deform the system since its shells reach turnaround so as to construct a virialised toy object satisfying the conditions leading to equation (90))."," To see this, we will deform the system since its shells reach turnaround so as to construct a virialised toy object satisfying the conditions leading to equation \ref{vir0}) )."
 This is only possible in PA where the equivalent radius of turnaround ellipsoids increases with increasing time and the central virialised object grows inside-out., This is only possible in PA where the equivalent radius of turnaround ellipsoids increases with increasing time and the central virialised object grows inside-out.
 Shells reaching turnaround can then bevirtually contracted one after the other without any crossing so as to match the mass profile M(r) (although not necessarily the ellipsoidal isodensity contours) of the real virialised object developed until some time t., Shells reaching turnaround can then be contracted one after the other without any crossing so as to match the mass profile $M(r)$ (although not necessarily the ellipsoidal isodensity contours) of the real virialised object developed until some time $t$.
" By “virtual” motion we mean a motion of shells that preserves their particle energy and angular momentum, but is disconnected from the real timing of the system."," By “virtual” motion we mean a motion of shells that preserves their particle energy and angular momentum, but is disconnected from the real timing of the system."
" Of course, such a motion will not recover at the same time the axial ratios of the isodensity contours of the virialised object, but we only need to recover the mass profile, which by conveniently contracting each new shell is guaranteed (there is one degree of freedom: the final equivalent radius of the contracted ellipsoidal shell and one quantity to match: the mass within the new infinitesimally larger radius)."," Of course, such a motion will not recover at the same time the axial ratios of the isodensity contours of the virialised object, but we only need to recover the mass profile, which by conveniently contracting each new shell is guaranteed (there is one degree of freedom: the final equivalent radius of the contracted ellipsoidal shell and one quantity to match: the mass within the new infinitesimally larger radius)."
" By construction, the new toy object so built has the same mass profile M(r) (although not the total energy profile E(r) due to the different ellipsoidal mass distribution) as the virialised object."," By construction, the new toy object so built has the same mass profile $M(r)$ (although not the total energy profile $E(r)$ due to the different ellipsoidal mass distribution) as the virialised object."
" But the spherical total energy in such a toy object is equal to that at turn-around, €i[M (r)], as there has been no"," But the spherical total energy in such a toy object is equal to that at turn-around, ${\cal E}\col[M(r)]$ , as there has been no"
jparr. therearelwofinalstaleswhichpermitsuccess fulidenti ficationofbolltaius : ep+Ep πε. Ey. where,"pair, there are two final states which permit successful identification of both taus: $e\mu + \displaystyle{\not} E_T$ and $\tau_h\ell+\displaystyle{\not} E_T$ , where."
 Finalslalesresulling from fidlyhadronicdecayshacealargebackgroundfromdi jelprocesseswith narrow jelsmi: a, Final states resulting from fully hadronic decays have a large background from dijet processes with narrow jets misidentified as taus.
nd Ey)) are also less useful due to the overwhelming SM background from Z/5*—(f processes., Final states involving two leptons of like flavor and ) are also less useful due to the overwhelming SM background from $Z/\gamma^\ast\rightarrow \ell^+\ell^-$ processes.
 A hadronically-decaying tau will decay into either a “one-prone” (approximately of the time) or “three-prone” (approximately of the time) final state., A hadronically-decaying tau will decay into either a “one-prong” (approximately of the time) or “three-prong” (approximately of the time) final state.
 These final states involve uarrow. well-collimated jets including one or three charged pious. respectively.," These final states involve narrow, well-collimated jets including one or three charged pions, respectively."
 The ideutificatiou ola jet as coming [rom a hadronically-decayiug).. csopposedtosomeQC Dprocess.isfarfroilrivial.," The identification of a jet as coming from a hadronically-decaying, as opposed to some QCD process, is far from trivial."
Oueoflhepriucipaldiscriniualionvariablesisjelradiusltfy (see [0]. for more details regarding 7 icentilication)., One of the principal discrimination variables is jet radius (see \cite{tauID} for more details regarding $\tau$ identification).
 At the Tevatron (Run II). )forappy GeV cut.," At the Tevatron (Run II), for a GeV cut."
 At the LHC. a 7 identification efficiency of around iserpected [0]..," At the LHC, a $\tau$ identification efficiency of around is expected \cite{tauID}."
 Processes involving direct decays of // to muon pairs cau also be of interest for Higgs discovery iu the L2HDMN., Processes involving direct decays of $h$ to muon pairs can also be of interest for Higgs discovery in the L2HDM.
. The disadvantage of such chaunels for Higgs searches. relative to those involving direct decays to taus. is the suppressed branching ratio.," The disadvantage of such channels for Higgs searches, relative to those involving direct decays to taus, is the suppressed branching ratio."
 Since Yukawa coupling universality cictates that )..botlintheSMandiutheL2HDAL.," Since Yukawa coupling universality dictates that, both in the SM and in the L2HDM,."
faw)).. However. this is compeusated for to a great exteut by the fact that the dinuou signal is exceptionally clean.," However, this is compensated for to a great extent by the fact that the dimuon signal is exceptionally clean."
 Indeed. the muon identification elliciency at the LHC is more that [5.6]..," Indeed, the muon identification efficiency at the LHC is more that \cite{ATLASTDR, CMSTDR}."
 Tn addition. the measurement of muon momenta allows [or a precise reconstruction of the Higes lnass within x2.5 GeV. This permits the implementation of an extremely efficient cut ou Mj. the iuvarlant mass of the muon pair. aid a substantial reductiou iu background levels for all caunels involving direct Higgs-bosonu decays to πιο pairs.," In addition, the measurement of muon momenta allows for a precise reconstruction of the Higgs mass within $\pm 2.5$ GeV. This permits the implementation of an extremely efficient cut on $M_{\mu\mu}$, the invariant mass of the muon pair, and a substantial reduction in background levels for all channels involving direct Higgs-boson decays to muon pairs."
 We now turn to address tle prospects for detecting a light SM-like CP-eveu Higgs bosou at the LHC ou a chanuel-by-chaunel basis., We now turn to address the prospects for detecting a light SM-like $CP$ -even Higgs boson at the LHC on a channel-by-channel basis.
 In the present work. as discussecl in Section 2.. we will assume veneration uulversality zinong the lepton Yukawa couplings.," In the present work, as discussed in Section \ref{sec:Param}, we will assume generation universality among the lepton Yukawa couplings."
 Therefore. we will ignore the //—ee channel and focus only on //—77 aud /i—pg.," Therefore, we will ignore the $h\rightarrow e e$ channel and focus only on $h\rightarrow \tau\tau$ and $h\rightarrow \mu\mu$."
 The channels of primary interest. then. are those in which the Higgs is produced by gluou fusion. weak-boson fusion. or ///j associated production aud dlecays to either poge or 7τ.," The channels of primary interest, then, are those in which the Higgs is produced by gluon fusion, weak-boson fusion, or $t\bar{t}h$ associated production and decays to either $\mu^+\mu^-$ or $\tau^+\tau^-$."
 Associated W— aud Z production processes generally have smaller rates. but may also potentially be of interest. and as such we briefly cliscuss them as well.," Associated $W^\pm$ and $Z$ production processes generally have smaller rates, but may also potentially be of interest, and as such we briefly discuss them as well."
Massive star-forming regions in the Galaxy are commonly seen to host OH maser emission. including masers from excited states.,"Massive star-forming regions in the Galaxy are commonly seen to host OH maser emission, including masers from excited states."
 However. maser emission in the highly-excited 1344] MHz transition is rare.," However, maser emission in the highly-excited 13441 MHz transition is rare."
 To date. only 11. sources are known to host 13441 MHz masers. most detected only recently (Turneretal.1970:Baudry&Desmurs2002:Caswell 2004).," To date, only 11 sources are known to host 13441 MHz masers, most detected only recently \citep{turner70,baudry02,caswell04}."
. In addition. several possible sources were detected by Balisteretal.(1976) but not subsequently redetected.," In addition, several possible sources were detected by \citet{balister76}
 but not subsequently redetected."
 Nearly all 13441 MHz maser sources display nearly equal fluxes in the left (LCP) and right circular. polarized (RCP) modes., Nearly all 13441 MHz maser sources display nearly equal fluxes in the left (LCP) and right circular polarized (RCP) modes.
 Conventional wisdom attributes this to the small Zeeman splitting coefficient (0.018 kmss mG)., Conventional wisdom attributes this to the small Zeeman splitting coefficient (0.018 $^{-1}$ $^{-1}$ ).
 Since a typical magnetic field of a few milligauss! does not split the LCP and RCP lines of a Zeeman pair by a full linewidth. velocity-coherent amplification of the LCP and RCP lines should be similar. giving a flux ratio of nearly unity.," Since a typical magnetic field of a few milligauss does not split the LCP and RCP lines of a Zeeman pair by a full linewidth, velocity-coherent amplification of the LCP and RCP lines should be similar, giving a flux ratio of nearly unity."
 Approximately equal LCP and RCP fluxes are observed in 10 of 11 13441 MHz OH maser sources. including for individual Zeeman pairs observed at very long baseline interferometric (VLBI) resolution in (Baudry&Diamond1998).," Approximately equal LCP and RCP fluxes are observed in 10 of 11 13441 MHz OH maser sources, including for individual Zeeman pairs observed at very long baseline interferometric (VLBI) resolution in \citep{baudry98}."
. The one exception is (hereafter GI11.90).. for which. the RCP/LCP flux ratio was 3.4 in the Caswell(2004) observations and larger in the Baudry&Desmurs(2002) observations. in. which the LCP feature was not detected.," The one exception is (hereafter G11.90), for which the RCP/LCP flux ratio was 3.4 in the \citet{caswell04} observations and larger in the \citet{baudry02} observations, in which the LCP feature was not detected."
 Caswell(2004) infers a possible magnetic field of —3.0 mG. which produces a negligible splitting compared to the line width (0.24 ss! RCP and 0.31 ss! LCP) and therefore would be expected to result in fairly equal RCP and LCP fluxes.," \citet{caswell04} infers a possible magnetic field of $-3.0$ mG, which produces a negligible splitting compared to the line width (0.24 $^{-1}$ RCP and 0.31 $^{-1}$ LCP) and therefore would be expected to result in fairly equal RCP and LCP fluxes."
 Observations of multiple OH maser transitions at VLBI resolution would also provide information on whether the several maser transitions are co-spatial. with implications for 13441 MHz maser modelling.," Observations of multiple OH maser transitions at VLBI resolution would also provide information on whether the several maser transitions are co-spatial, with implications for 13441 MHz maser modelling."
 Multitransition maser overlaps provide strong observational constraints to test maser models and to derive local physical conditions., Multitransition maser overlaps provide strong observational constraints to test maser models and to derive local physical conditions.
 The literature includes only one source. W3(OH). observed at 13441 MHz at VLBI resolution (Baudry&Diamond1998)..," The literature includes only one source, W3(OH), observed at 13441 MHz at VLBI resolution \citep{baudry98}."
 Further observations are needed in order to understand the conditions that produce these rare masers., Further observations are needed in order to understand the conditions that produce these rare masers.
 It is for these reasons that G11.90 was selected for study at higher spatial resolution., It is for these reasons that G11.90 was selected for study at higher spatial resolution.
 Results are reported in this Letter., Results are reported in this Letter.
 Four transitions of OH masers were observed with the Very Long Baseline Array (VLBA) in three different epochs (experiment code BF088)., Four transitions of OH masers were observed with the Very Long Baseline Array (VLBA) in three different epochs (experiment code BF088).
 The ground-state 1665.40184 and 1667.35903 MHz masers were observed simultaneously on 2006 Feb 22 and the 4765.562 MHz masers on 2006 Feb 27., The ground-state 1665.40184 and 1667.35903 MHz masers were observed simultaneously on 2006 Feb 22 and the 4765.562 MHz masers on 2006 Feb 27.
 Total observing time was 6.5 hr per run. with approximately 3 hr spent on G11.90.," Total observing time was 6.5 hr per run, with approximately 3 hr spent on G11.90."
 The 13441.4173 MHz masers were observed with both the VLBA and the Green Bank Telescope (GBT) on 2006 Aug 29 over a total of 4.5 hr. with approximately 2 hr spent on GI1.90.," The 13441.4173 MHz masers were observed with both the VLBA and the Green Bank Telescope (GBT) on 2006 Aug 29 over a total of 4.5 hr, with approximately 2 hr spent on G11.90."
 Most of the remaining time was used to observe the nearby calibrator J1825—1718 for phase calibration., Most of the remaining time was used to observe the nearby calibrator $-$ 1718 for phase calibration.
 The calibrator 3C286 was observed as a fringe-finder and polarization calibrator., The calibrator 3C286 was observed as a fringe-finder and polarization calibrator.
 The 1665. 1667. and 4765 MHz masers were observed in full polarization mode using 0.125 MHz bandwidths divided into 128 spectral channels for a spectral resolution of 0.18 kmss! at 1665/1667 MHz and 0.06 ss! at 4765 MHz.," The 1665, 1667, and 4765 MHz masers were observed in full polarization mode using 0.125 MHz bandwidths divided into 128 spectral channels for a spectral resolution of 0.18 $^{-1}$ at 1665/1667 MHz and 0.06 $^{-1}$ at 4765 MHz."
 The 13441 MHz masers were observed in dual circular polarization mode using a 1.0 MHz bandwidth divided into 512 spectral channels for a spectral resolution of 0.04 ss!., The 13441 MHz masers were observed in dual circular polarization mode using a 1.0 MHz bandwidth divided into 512 spectral channels for a spectral resolution of 0.04 $^{-1}$.
 Because GI1.90 is significantly scatter-broadened at long wavelengths. no signal was detectable on the longer baselines at 1.6 GHz.," Because G11.90 is significantly scatter-broadened at long wavelengths, no signal was detectable on the longer baselines at 1.6 GHz."
 The usable array at 1.6 GHz effectively consisted of the inner five antennas plus a small amount of data from North Liberty. IA.," The usable array at 1.6 GHz effectively consisted of the inner five antennas plus a small amount of data from North Liberty, IA."
 Blank sky noise in the LCP and RCP image cubes was 18 bbeam in a single channel with a synthesized beam size of 67«32 !mas., Blank sky noise in the LCP and RCP image cubes was 18 $^{-1}$ in a single channel with a synthesized beam size of $67 \times 32$ mas.
 At 4.7 GHz. sufficient signal was seen to determine adequate calibration on baselines to all antennas except Mauna Kea. HI and St. Croix. VI.," At 4.7 GHz, sufficient signal was seen to determine adequate calibration on baselines to all antennas except Mauna Kea, HI and St. Croix, VI."
 Blank sky noise in Stokes | was 8 bbeam! in a 93 mas synthesized beam.," Blank sky noise in Stokes I was 8 $^{-1}$ in a $9 \times
3$ mas synthesized beam."
 At 13441 GHz all 10 VLBA antennas and the GBT produced usable data. for a blank sky Stokes I noise level of 3 bbeam! in a 2.80.7 mas beam when averaged over 5 spectral channels. comparable to a maser line width.," At 13441 GHz all 10 VLBA antennas and the GBT produced usable data, for a blank sky Stokes I noise level of 3 $^{-1}$ in a $2.8 \times 0.7$ mas beam when averaged over 5 spectral channels, comparable to a maser line width."
 All data were phase-referenced to J1825—1718 (375 away from 01190) using a cycle time of 5 min at 1665/1667 MHz and 3 min at 4765 and 13441 MHz., All data were phase-referenced to $-$ 1718 $3\fdg5$ away from G11.90) using a cycle time of 5 min at 1665/1667 MHz and 3 min at 4765 and 13441 MHz.
 The phase-referenced 1665 and 1667 MHz image quality was poor. but the reference feature (1665 MHz LCP) was sufficiently bright to determine a position.," The phase-referenced 1665 and 1667 MHz image quality was poor, but the reference feature (1665 MHz LCP) was sufficiently bright to determine a position."
 The reference maser was then used to self- the 1665 and 1667 MHz LCP and RCP data for further imaging., The reference maser was then used to self-calibrate the 1665 and 1667 MHz LCP and RCP data for further imaging.
 Phase-referencing at 4765 MHz produced, Phase-referencing at 4765 MHz produced
with h(.r)—for. O0 and hGr)=1 for xr«0.,with $h(x)=1$for $x\geq0$ and $h(x)=-1$ for $x<0$ .
 In the large 3v limit g(£) is symmetric about the vertical axis at €=£i., In the large $\beta_V$ limit $g(\xi)$ is symmetric about the vertical axis at $\xi=\xi_1$.
 Asthe phase velocity approaches the luminal limit 71 the right-hand side of the curve steepens aud approached the vertical axis., Asthe phase velocity approaches the luminal limit $\beta_V\to1$ the right-hand side of the curve steepens and approached the vertical axis.
 In figure 3.. (27)) is shown as dotted lines. which gives a quite good approximation to the exact numerical result. (solid lines).," In figure \ref{fig:g}, \ref{eq:g2}) ) is shown as dotted lines, which gives a quite good approximation to the exact numerical result (solid lines)."
 The condition (26)) leads to upper and lower limits to Ho. given by poe Using (18)) the positron momentum 1 can be expressed in terms of Ην which leads to the same upper and lower limits for positrons as (28)).," The condition \ref{eq:Phi2}) ) leads to upper and lower limits to $u_-$, given by with Using \ref{eq:betapm}) ) the positron momentum $u_+$ can be expressed in terms of $u_-$, which leads to the same upper and lower limits for positrons as \ref{eq:umax}) )."
 For luminal waves with js—1. οπο has muac(1|ESín)/2 and manm ol.," For luminal waves with $\beta_V=1$, one has $u_{\rm max}\approx (1+\tilde{E}^2_0/n_-)/2$ and $u_{\rm min}\approx -1$ ."
 ln this case oscillations skew strongly in the direction of wave propagation., In this case oscillations skew strongly in the direction of wave propagation.
 For τι(ox. (28)) reducesto thins=(E|Hoo|U o)/2and my=(Ponyov0)3.," For $\beta_V\to\infty$, \ref{eq:umax}) ) reducesto $u_{\rm max}=
(\Gamma+u_{+0}+u_{-0})/2$ and $u_{\rm min}=-(\Gamma-u_{+0}-u_{-0})/2$."
" When ua=0. oscillations become svmmetric with ""ag= MiminE72."," When $u_{\pm0}=0$, oscillations become symmetric with $u_{\rm max}=-u_{\rm min}=\Gamma/2$."
" Since ""yas= Maine Oscillating electrons ancl positrons have a net drift velocity ~(ties|mind/28(PO€))fy."," Since $u_{\rm max}\neq u_{\rm min}$ , oscillating electrons and positrons have a net drift velocity $\sim (u_{\rm max}+u_{\rm min})/2\approx (\Gamma-\xi_1)/\beta_V$."
 An accurate evaluation of the drift velocity is given in Sec., An accurate evaluation of the drift velocity is given in Sec.
 3.3., 3.3.
 Numerical solutions to 4|=0 (Le. &(€)= 0) are shown as contours in figure 4..," Numerical solutions to $E_\parallel=0$ (i.e., $\Phi(\xi)=0$ ) are shown as contours in figure \ref{fig:bounce}."
 We express n. in terms of v. Juge and ges using (13)) and. (14)).," We express $n_+$ in terms of $\beta_V$, $j_{0\parallel}$, and $\eta_{GJ}$ using \ref{eq:j-eta1}) ) and \ref{eq:j-eta2}) )."
 In the subfigure on the left one assumes luminal waves with 3=I., In the subfigure on the left one assumes luminal waves with $\beta_V=1$.
 Each pair of lines defines an upper limit. (yas. and a lowerlimit. Hain. to a particles momentum such that one hase>0 [Or (ainυὉ ια ," Each pair of lines defines an upper limit, $u_{\rm max}$ , and a lowerlimit, $u_{\rm min}$ , to a particle's momentum such that one has$\Phi>0$ for $u_{\rm min}<u_-<u_{\rm max}$ ."
limits can be obtained for positrons., The similarupper and lower limits can be obtained for positrons.
 Particles oscillate with their momenta confined between these two limits., Particles oscillate with their momenta confined between these two limits.
 As, As
 As., As
"where d; is the luminosity distance and F, the 1 s peak energy flux.",where $d_l$ is the luminosity distance and $F_p$ the 1 s peak energy flux.
" We determine F, in the energy range between 30 keV and 10 MeV in the rest frame of the GRB (?)..", We determine $F_p$ in the energy range between 30 keV and 10 MeV in the rest frame of the GRB \citep{yonetoku04}.
 We use the GBM peak spectral catalogue (Goldstein et al., We use the GBM peak spectral catalogue (Goldstein et al.
 2011) to determine the best fit spectral parameters of the brightest 1 s bin., 2011) to determine the best fit spectral parameters of the brightest 1 s bin.
" We then determined the energy flux in the rest-frame of the GRB by integrating where 30/(1+4z)No is the normalization of the spectrum (in photons cm? s!), E the energy and ®(E) the form of the spectrum (either Band or COMP)."," We then determined the energy flux in the rest-frame of the GRB by integrating where $N_0$ is the normalization of the spectrum (in photons $^2$ $^{-1}$ ), $E$ the energy and $\Phi(E)$ the form of the spectrum (either Band or COMP)."
" In order to determine the error on the energy flux, we calculate 1000 flux values for each GRB by varying the input parameters, No,Ερ., α and B according to the lo error on each parameter."," In order to determine the error on the energy flux, we calculate 1000 flux values for each GRB by varying the input parameters, $N_0$, $\alpha$ and $\beta$ according to the $1 \sigma$ error on each parameter."
" We then use the median and the mean absolute deviation (MAD) around the median of the 1000 simulated bursts to determine the peak energy flux and the error on the latter, respectively."," We then use the median and the mean absolute deviation (MAD) around the median of the 1000 simulated bursts to determine the peak energy flux and the error on the latter, respectively."
 We present this relation for 30 GBM GRBs in Fig., We present this relation for 30 GBM GRBs in Fig.
 8 and find a best-fit power law index of 0.58+0.08., \ref{fig:yonetoku} and find a best-fit power law index of $0.58 \pm 0.08$.
 The Spearman’s rank correlation gives p=0.7 with a chance probability of 2.3x10., The Spearman's rank correlation gives $\rho=0.7$ with a chance probability of $2.3\times10^{-5}$.
 Our findings are in good agreement with ?? and ?..," Our findings are in good agreement with \citet{yonetoku04, ghina09} and \citet{ghirlanda10}."
 The rest-frame is plotted as a function of redshift in Fig. 9.., The rest-frame is plotted as a function of redshift in Fig. \ref{fig:t90vsz}.
" Contrary to ?,, who find a negative correlation between and z, we do not find any evidence in the GBM data for such a correlation."," Contrary to \citet{pel08}, who find a negative correlation between and $z$ , we do not find any evidence in the GBM data for such a correlation."
" Also,there is no correlation between and z when accounting for Εριοι."," Also,there is no correlation between and $z$ when accounting for ."
. A Spearman rank test for all bursts gives a correlation coefficient of p=0.04 with a chance probability of P=0.81., A Spearman rank test for all bursts gives a correlation coefficient of $\rho=0.04$ with a chance probability of $P=0.81$.
" However, thep-value in ? is &10 which makes it a weak case for such a correlation in the first place."," However, thep-value in \citeauthor{pel08} is $\approx 10^{-3}$ which makes it a weak case for such a correlation in the first place."
 Our results confirm the analyses with detected GRBs (?).. We do note a lack of short GRBs for z>2., Our results confirm the analyses with detected GRBs \citep{greiner11}.. We do note a lack of short GRBs for $z \ge 2$.
clistribution mmomoent) and the mean vvelocity. field. mmomoent) of NGC 1512 are shown in Fig.,distribution moment) and the mean velocity field moment) of NGC 1512 are shown in Fig.
 3., 3.
 The yvelocity. clispersion nimoment. not. shown) varies between iin the spiral/tidal arms of NGC 1512.," The velocity dispersion moment, not shown) varies between in the spiral/tidal arms of NGC 1512."
 As a first step. we compare the racially averaged &eas distribution with the critical density and the racially averaged LUV emission.," As a first step, we compare the radially averaged gas distribution with the critical density and the radially averaged $FUV$ emission."
" For a Uat rotation curve (ic. (Cr) = constant). which is a reasonable assumption for p15"" (see Fig."," For a flat rotation curve (i.e., $r$ ) = constant), which is a reasonable assumption for $r
\ga 15\arcsec$ (see Fig."
 6). and a velocity dispersion of ((as used in previous work). the above equation reduces to Moni = 0.6 A £r. where ris the radius in κρὸ," 6), and a velocity dispersion of (as used in previous work), the above equation reduces to $\Sigma_{\rm crit}$ = 0.6 $\alpha_{\rm Q}$ $r$, where $r$ is the radius in kpc."
 Using the derived iinclination (7 = 35°)) and position angle (224 = 265)). the de-projected rraciial surface density. May ofNGC 1512 is shown in Fig.," Using the derived inclination $i$ = ) and position angle $PA$ = ), the de-projected radial surface density, $\Sigma_{\rm HI}$ of NGC 1512 is shown in Fig."
 18., 18.
 The ουσα density was computed for ao = 0.7 and ==1., The critcal density was computed for $\alpha_{\rm Q}$ = 0.7 and =.
 We find that the radially averaged geas alone lies just below the computed critical density., We find that the radially averaged gas alone lies just below the computed critical density.
 At radi less than 10 kpe. an increasing amount of molecular gas is needed to reach critical gas density ancl fec the star formation in the nuclear region and the inner ring.," At radii less than 10 kpc, an increasing amount of molecular gas is needed to reach critical gas density and feed the star formation in the nuclear region and the inner ring."
 The racially integrated. £V. lux drops below the noise a r= 28 kpe. the likely SE threshold.," The radially integrated $FUV$ flux drops below the noise at $r$ = 28 kpc, the likely SF threshold."
 In the inner region of NGC 1512. the gas motions are stronely alfected by the stellar bar. which would allec the critical density estimate.," In the inner region of NGC 1512, the gas motions are strongly affected by the stellar bar, which would affect the critical density estimate."
" Using NGC 1512's angular velocity. O(r) = real) fr. we can also determine the locations of the inner and outer Lindblad resonances: O,=O(r)+e)/2. where O, is the bar pattern speed anc wr)=v3 μμ, M"," Using NGC 1512's angular velocity, $\Omega$ $r$ ) = $r$ $r$, we can also determine the locations of the inner and outer Lindblad resonances: $\Omega_{\rm p} = \Omega(r) \pm \kappa(r)/2$, where $\Omega_{\rm p}$ is the bar pattern speed and $\kappa(r) = \sqrt2$ $r$ ."
 rc4 kpe (bar radius) its pattern speed is ~50 1 σι sugeesting that the ILR(s) lie at rZ1.2 kpe. an 1w OLR at οGS kpc.," At $r \approx 4$ kpc (bar radius) its pattern speed is $\sim$ $^{-1}$ , suggesting that the ILR(s) lie at $r \la 1.2$ kpc, and the OLR at $r \sim 6.8$ kpc."
 Figure 19 shows the logarithmic ratio of the measures eoas clensity. Nyy. to the critical. density. Mor. Lor all analysed. CY -rich. clusters together anc within the previously defined: cistinet regions.," Figure 19 shows the logarithmic ratio of the measured gas density, $\Sigma_{\rm HI}$ to the critical density, $\Sigma_{\rm crit}$, for all analysed $UV$ -rich clusters together and within the previously defined distinct regions."
 Taking all clusters. a peak is found at <log(Xgir/YXau)>=0.06. indicating that local star formation is. on average. happening at the local critical density.," Taking all clusters, a peak is found at $<\log(\Sigma_{\rm HI}/\Sigma_{\rm crit})> = 0.06$, indicating that local star formation is, on average, happening at the local critical density."
 Phe measured dddensities are slightly higher than critical along Arm 1 (0.33) and within the NW debris (0.17). but. significantly. lower than critical in the inner star-forming ring (O44) where the molecular gas density must be high.," The measured densities are slightly higher than critical along Arm 1 (0.33) and within the NW debris (0.17), but significantly lower than critical in the inner star-forming ring (–0.44) where the molecular gas density must be high."
 Finally. Fig.," Finally, Fig."
 20 shows — on a logarithmic seale the deensity versus the eas density for CV -rich clusters. in the NGC 1512/1510. svstem (derived. here) and in. the nearby She galaxy ALDS51 (Ixennicutt et al., 20 shows – on a logarithmic scale – the density versus the gas density for $UV$ -rich clusters in the NGC 1512/1510 system (derived here) and in the nearby Sbc galaxy 51 (Kennicutt et al.
 2007: for r8 kpc)., 2007; for $r \la 8$ kpc).
 Because no molecular cata are currently available for NGC 1512. only the eas density is shown.," Because no molecular data are currently available for NGC 1512, only the gas density is shown."
 Regions within the inner ring of NGC 1512 (located at 7r = 907)) are. as stated before significantly ollset., Regions within the inner ring of NGC 1512 (located at $r$ = ) are – as stated before – significantly offset.
 Tripling the nunass of cach region achieves an approximate alignement., Tripling the mass of each region achieves an approximate alignement.
 For 551. molecular data are taken into account ancl are found to be essential for the observed correlation (Ixennicutt et al.," For 51, molecular data are taken into account and are found to be essential for the observed correlation (Kennicutt et al."
 2007)., 2007).
 Dong et al. (, Dong et al. (
2008) found that the CY-selected regions in two small fields within the large gaseous disk of ALSS3. ~20 kpc from its centre. follow a similar trend.,"2008) found that the $UV$ -selected regions in two small fields within the large gaseous disk of 83, $\sim$ 20 kpc from its centre, follow a similar trend."
 Alolecular gas was not taken into account (for comparison. see Martin Ixennicutt. 2001).," Molecular gas was not taken into account (for comparison, see Martin Kennicutt 2001)."
 lore. we estimate the metallicities of NGC 1512. and NGC 1510 using data available in the literature., Here we estimate the metallicities of NGC 1512 and NGC 1510 using data available in the literature.
 Calzetti et al. (, Calzetti et al. (
2007) ancl Aloustakas Kennicutt (2006). give an oxvgen abundance between 8.37 and S.SI. in units of 12|log(O/L1D. for the chemical abundance of NGC 1512.,"2007) and Moustakas Kennicutt (2006) give an oxygen abundance between 8.37 and 8.81, in units of 12+log(O/H), for the chemical abundance of NGC 1512."
 The first value was derived using optical spectroscopy ancl the Pilvugin Thuan (2005) calibration., The first value was derived using optical spectroscopy and the Pilyugin Thuan (2005) calibration.
 The second value was obtained comparing with the predictions given by the photoioniscd evolutionary svnthesis. models provided. by Ixobulnickv Ixewlev (2004)., The second value was obtained comparing with the predictions given by the photoionised evolutionary synthesis models provided by Kobulnicky Kewley (2004).
 Llowever. some recent analysis using direct estimates of the electron temperature (2;) of the ionised eas (e.g. Lóppez-Sánnchez 2006) suggest that these models overestimate the oxvgen abundance by ~0.2 dex.," However, some recent analysis using direct estimates of the electron temperature $T_e$ ) of the ionised gas (e.g., Lóppez-Sánnchez 2006) suggest that these models overestimate the oxygen abundance by $\sim$ 0.2 dex."
 We conclude that the metallicity of NGC 1512 is between S. ancl 8.6. slightly lower than for the Milkv Way. but within the range typical observed. for spiral galaxies (Henry Worthey 1999).," We conclude that the metallicity of NGC 1512 is between 8.4 and 8.6, slightly lower than for the Milky Way, but within the range typical observed for spiral galaxies (Henry Worthey 1999)."
 On the other hand. we have used the emission. line intensity data for NGC 1510. provided byStorchi-Dergmannet al. (," On the other hand, we have used the emission line intensity data for NGC 1510, provided byStorchi-Bergmannet al. ("
1995).to compute its chemical abundance.,"1995),to compute its chemical abundance."
 We used, We used
"positions similar to Sr, but with a discontinuity at the stellar equator even for the region of relative underabundance.","positions similar to Sr, but with a discontinuity at the stellar equator even for the region of relative underabundance."
" For all chemical elements presented in Fig. 7,,"," For all chemical elements presented in Fig. \ref{DImaps},"
 the abundance difference between zones of the highest and lowest concentrations does not exceed 1 dex., the abundance difference between zones of the highest and lowest concentrations does not exceed 1 dex.
" For Ba the range is 0.81 dex, for Y it is 0.71 dex."," For Ba the range is 0.81 dex, for Y it is 0.71 dex."
" For Ti and Sr the range is 0.36 dex and 0.59 dex, respectively."," For Ti and Sr the range is 0.36 dex and 0.59 dex, respectively."
 This is in a good agreement with the relative amplitude of variability in individual line profiles of these chemical elements (Fig. 6))., This is in a good agreement with the relative amplitude of variability in individual line profiles of these chemical elements (Fig. \ref{varprofs}) ).
" To exclude the possibility that spots are found by DI code at roughly one hemisphere of the star due to a spurious systematic effect, we applied the same DI analysis to ten spectral lines of Fe, which show no variability."," To exclude the possibility that spots are found by DI code at roughly one hemisphere of the star due to a spurious systematic effect, we applied the same DI analysis to ten spectral lines of Fe, which show no variability."
 The inferred abundance map of Fe is not similar to the maps of other chemical elements in its morphology., The inferred abundance map of Fe is not similar to the maps of other chemical elements in its morphology.
" In particular, it does not show an abundance difference between two hemispheres."," In particular, it does not show an abundance difference between two hemispheres."
" The difference between the regions of the highest and lowest abundances does not exceed «0.15 dex, which we take to be an upper limit of the possible inhomogeneity for chemical elements without obvious line profile variability."," The difference between the regions of the highest and lowest abundances does not exceed $\approx0.15$ dex, which we take to be an upper limit of the possible inhomogeneity for chemical elements without obvious line profile variability."
" Earlier studies of several HgMn stars found a variability in some spectral lines, attributing it to the presence of chemical spots."," Earlier studies of several HgMn stars found a variability in some spectral lines, attributing it to the presence of chemical spots."
" However, all previous attempts to find magnetic fields that can be responsible for these chemical inhomogeneities"," However, all previous attempts to find magnetic fields that can be responsible for these chemical inhomogeneities were unsuccessful \citep{Shorlin:2002, Wade:2006, Folsom:2010, Auriere:2010, Makaganiuk:2011}."
" To extend the group of spotted HgMn stars with known constraints on magnetic field, we performed time-series spectropolarimetric observations of 66 Eri."," To extend the group of spotted HgMn stars with known constraints on magnetic field, we performed time-series spectropolarimetric observations of 66 Eri."
 This is the first such analysis of this unique spectroscopic binary star., This is the first such analysis of this unique spectroscopic binary star.
 It allowed us to search for magnetic fields in both components and investigate their line profile variability., It allowed us to search for magnetic fields in both components and investigate their line profile variability.
" Taking the advantage of one of the best spectropolarimeters in the world, we recorded high S/N and high-resolution spectra of 66 Eri covering its full orbital period."," Taking the advantage of one of the best spectropolarimeters in the world, we recorded high $S/N$ and high-resolution spectra of 66 Eri covering its full orbital period."
" Despite the high polarimetric sensitivity of the instrument, we did not find any evidence of the longitudinal magnetic field in either component of 66 Eri."," Despite the high polarimetric sensitivity of the instrument, we did not find any evidence of the longitudinal magnetic field in either component of 66 Eri."
 Our measurements of uusing LSD profiles have error bars 10—24 G for both stars., Our measurements of using LSD profiles have error bars 10–24 G for both stars.
" This null result agrees with the outcome of our earlier work (?),, where the longitudinal magnetic field was measured in a sample of 47 HgMn stars."," This null result agrees with the outcome of our earlier work \citep{Makaganiuk:2011}, where the longitudinal magnetic field was measured in a sample of 47 HgMn stars."
 The analysis of these objects showed no evidence of weak magnetic field signatures in Stokes V spectra., The analysis of these objects showed no evidence of weak magnetic field signatures in Stokes $V$ spectra.
" Our magnetic field measurements provide an upper limit of 60—70 G for a dipolar field in either star and LSD V profiles provide no evidence for the presence of strong, complex magnetic field structures."," Our magnetic field measurements provide an upper limit of 60–70 G for a dipolar field in either star and LSD $V$ profiles provide no evidence for the presence of strong, complex magnetic field structures."
 This upper limit is significantly smaller than the minimum dipolar field diagnosed in Ap stars (?).., This upper limit is significantly smaller than the minimum dipolar field diagnosed in Ap stars \citep{Auriere:2007}.
" Thus, 66 Eri is clearly not an Ap star."," Thus, 66 Eri is clearly not an Ap star."
" As we deal with the binary system that shows two systems of spectral lines in its composite spectrum, it is necessary to disentangle the spectra to study each component"," As we deal with the binary system that shows two systems of spectral lines in its composite spectrum, it is necessary to disentangle the spectra to study each component"
the other parameters.,the other parameters.
 This clegeneracy was discovered by Smith.Mao&Paezviski(2009)., This degeneracy was discovered by \citet*{smp}.
 See Table 1., See Table 1.
 To understand (he nature of this degeneracy. I extend the approach of bv Taylor expancine u. the vector position of the lens relative to the source in the Einstein rng. where uy is (he vector impact parameter. w is (he vector inverse (tiniescale (Le. t!=hl with direction given bv the lens-source relative motion). and e and j are the apparent acceleration and jerk of the Sun relative to the Earth. both divided bv an AU.," To understand the nature of this degeneracy, I extend the approach of \citet{smp} by Taylor expanding $\bu$, the vector position of the lens relative to the source in the Einstein ring, where $\bu_0$ is the vector impact parameter, $\bomega$ is the vector inverse timescale (i.e., $\omega=t_\e^{-1}$ with direction given by the lens-source relative motion), and $\balpha$ and $\bj$ are the apparent acceleration and jerk of the Sun relative to the Earth, both divided by an AU."
 Note that w. a. andj ave all evaluated al /=0 and that all ave two-dimensional vectors.," Note that $\bomega$, $\balpha$, and $\bj$ are all evaluated at $t=0$ and that all are two-dimensional vectors."
 | impose uj:=0. which is equivalent to assuming (hat fy (the time of closest approach) can bedirectly. determined from the lighteurve anc so does not require an additional parameter.," I impose $\bu_0\cdot \bomega=0$, which is equivalent to assuming that $t_0$ (the time of closest approach) can bedirectly determined from the lightcurve and so does not require an additional parameter."
" Squaring equation (10)) vields. where. and where the subscripts ""| and 7.L7 indicate components parallel and perpendicular to the acceleration a."," Squaring equation \ref{eqn:vecu}) ) yields, where, where I have introduced the “jerk parallax”, and where the subscripts $\parallel$ ” and $\perp$ ” indicate components parallel and perpendicular to the acceleration $\balpha$."
 Note (hat I have made use of the fact that the Earth's orbit is basically circular to approximate the derivative of the jerk as —Q? a. where Q=2avr f.," Note that I have made use of the fact that the Earth's orbit is basically circular to approximate the derivative of the jerk as $-\Omega_\oplus^2\balpha$ , where $\Omega_\oplus=2\pi\,\rm yr^{-1}$ ."
Considerable theoretical ancl observational interest has in recent vears been focused. on asymptotic giant branch (AGB) stars but much remains mysterious about them.,Considerable theoretical and observational interest has in recent years been focused on asymptotic giant branch (AGB) stars but much remains mysterious about them.
 The AGB phase is enjoved by low and intermediate mass stars (approximate mass range O.S to SM. )., The AGB phase is enjoyed by low and intermediate mass stars (approximate mass range 0.8 to $_{\odot}$ ).
 Lis in this phase that a star experiences extensive. internal. nucleosvnthesis whose fruits are dredged to the stellar surface2005)., It is in this phase that a star experiences extensive internal nucleosynthesis whose fruits are dredged to the stellar surface.
 Furthermore. mass-Ioss ensures that the products of nucleosynthesis are dispersed into the circumstellar and subsequently the interstellar environment.," Furthermore, mass-loss ensures that the products of nucleosynthesis are dispersed into the circumstellar and subsequently the interstellar environment."
 Thus. AGB stars are likely major contributors of Li. €. N. E ancl s-process elements among others to Galactic chemical evolution2010).," Thus, AGB stars are likely major contributors of Li, C, N, F and $s$ -process elements among others to Galactic chemical evolution."
 Observational validation of theoretical investigations of how AGB stars achieve internal nucleosvnthesis. ddredge-up ancl mass-loss are generally. hampered by the fact that the more evolved ancl more interesting of these stars have low," Observational validation of theoretical investigations of how AGB stars achieve internal nucleosynthesis, dredge-up and mass-loss are generally hampered by the fact that the more evolved and more interesting of these stars have low"
So this appears to support the conclusions of ?:: the likelihood of finding a companion galaxy with £<17.5 within 50 kpe of a Sevfert galaxy is not statistically cillerent from that for an inactive galaxy.,So this appears to support the conclusions of \citet{1998ApJ...496...93D}: the likelihood of finding a companion galaxy with $R<-17.5$ within 50 kpc of a Seyfert galaxy is not statistically different from that for an inactive galaxy.
 ACN with sub-quasar luminosities have essentially identical environments — from 30 kpe up toO.5 Mpe to those of normal’. inactive galaxies.," AGN with sub-quasar luminosities have essentially identical environments – from 30 kpc up to 0.5 Mpc – to those of `normal', inactive galaxies."
 The typical environments. of our. CN. sample are. no dilferent to those of inactive galaxies in general., The typical environments of our AGN sample are no different to those of inactive galaxies in general.
 Aside from one example. all the AGN are found in sub-cluster richness regions. in contrast to the studies of high luminosity AGN. such as racio-Ioud QSOs.," Aside from one example, all the AGN are found in sub-cluster richness regions, in contrast to the studies of high luminosity AGN, such as radio-loud QSOs."
 Fherefore. any fuelling mechanism requiring the presence or proximity of a rich cluster is unlikely to be important for fuelling lower Iuminosity ACN.," Therefore, any fuelling mechanism requiring the presence or proximity of a rich cluster is unlikely to be important for fuelling lower luminosity AGN."
be estimated.,be estimated.
 A small amount of experimentation showed tha selecting all stars with more than 75 subtractions provided clean. well-defined CMDs.," A small amount of experimentation showed that selecting all stars with more than $75$ subtractions provided clean, well-defined CMDs."
 The second subtraction method was very similar. but involved using the outer (eld) CMD to subtract a matching CMD from the intermediate sample. leaving the cluster stars.," The second subtraction method was very similar, but involved using the outer (field) CMD to subtract a matching CMD from the intermediate sample, leaving the cluster stars."
 The required number of subtractions was again determined using the measured field star density in the outer region., The required number of subtractions was again determined using the measured field star density in the outer region.
 As before. detection completeness Was not accounted for. meaning the estimated number of field stars Cand hence subtractions) was over-estimated.," As before, detection completeness was not accounted for, meaning the estimated number of field stars (and hence subtractions) was over-estimated."
 The subtraction process was identical. except this time random stars from the outer region CMD were selected.," The subtraction process was identical, except this time random stars from the outer region CMD were selected."
 Once again. one hundred realizations of each clusters CMD were obtained and the stars appearing the most times in these CMDs selected as the most likely cluster members.," Once again, one hundred realizations of each cluster's CMD were obtained and the stars appearing the most times in these CMDs selected as the most likely cluster members."
 For this method. selecting all stars with more than 50 subtractions provided good CMDs.," For this method, selecting all stars with more than $50$ subtractions provided good CMDs."
 Because of the differential reddening present in the outer field regions of the NGC 1939 frame. this subtraction method was not as effective for this cluster.," Because of the differential reddening present in the outer field regions of the NGC 1939 frame, this subtraction method was not as effective for this cluster."
 Nonetheless. adequate results were obtained.," Nonetheless, adequate results were obtained."
 The final cleaned. tield-subtracted CMDs for NGC 1928. 1939. and Reticulum appear in Fig. 3..," The final cleaned, field-subtracted CMDs for NGC 1928, 1939, and Reticulum appear in Fig. \ref{f:cmdsub}."
 In this Figure. the CMDs for GC 1928 and 1939 represent the combined results of the two subtraction processes.," In this Figure, the CMDs for NGC 1928 and 1939 represent the combined results of the two subtraction processes."
" Stars within ry, have also been plotted for both clusters. since these are very predominantly cluster members."," Stars within $r_1$ have also been plotted for both clusters, since these are very predominantly cluster members."
 To the best of our knowledge these are the first published CMDs or NGC 1928 and 1939., To the best of our knowledge these are the first published CMDs for NGC 1928 and 1939.
 Tt can clearly be seen that these two clusters are very old. thus confirming the results obtainedby Dutra et al.," It can clearly be seen that these two clusters are very old, thus confirming the results obtainedby Dutra et al."
 (1999). using integrated spectroscopy., \shortcite{dutra} using integrated spectroscopy.
 Both clusters appear © possess well populated horizontal branches. consisting almost entirely of blue stars.," Both clusters appear to possess well populated horizontal branches, consisting almost entirely of blue stars."
 In this respect they strongly resemble the old LMC clusters Hodge 11 (Walker1993:Mighelletal.1996:Johnsonetal.1999). and NGC 2005 (Olsenetal.1998).," In this respect they strongly resemble the old LMC clusters Hodge 11 \cite{walker:h11,mighell,johnson:99} and NGC 2005 \cite{olsen}."
. In addition. NGC 1928 apparently. possesses an extended blue HB. falling to V.~23.," In addition, NGC 1928 apparently possesses an extended blue HB, falling to $V\sim 23$."
 The effects of differential reddening on the field subtraction for NGC 1939 can be seen in its lower main sequence. which exhibits a sharp cut-off to the red.," The effects of differential reddening on the field subtraction for NGC 1939 can be seen in its lower main sequence, which exhibits a sharp cut-off to the red."
 The remainder ofthe CMD has not been significantly affected by this problem., The remainder of the CMD has not been significantly affected by this problem.
 Reticulum is also seen to be an old cluster. with a horizontal branch primarily consisting of stars within and to the red of the instability strip.," Reticulum is also seen to be an old cluster, with a horizontal branch primarily consisting of stars within and to the red of the instability strip."
 A significant population of blue stragglers also appears to be present., A significant population of blue stragglers also appears to be present.
 This CMD confirms the earlier results of Walker €(19923:: Johnson et al. (20023:, This CMD confirms the earlier results of Walker \shortcite{walker:ret}; Johnson et al. \shortcite{johnson:02};
: Marconi et al. (2002):, Marconi et al. \shortcite{marconi};
 and Tonelli et al. (2003).., and Monelli et al. \shortcite{monelli}.
 Ideally. we would like to use the final CMDs to wovide photometric measurements of the cluster reddenings and metallicities. and to place some constraints on their ages.," Ideally, we would like to use the final CMDs to provide photometric measurements of the cluster reddenings and metallicities, and to place some constraints on their ages."
 However. such calculations generally require the photometry to be on a standard magnitude scale — in this case. V and Cousins-/.," However, such calculations generally require the photometry to be on a standard magnitude scale – in this case, $V$ and $I$."
 At the time of writing. no transformations rom the ACS/WFC STmag system to the Johnson-Cousins V/ system are yet available.," At the time of writing, no transformations from the ACS/WFC STmag system to the Johnson-Cousins $VI$ system are yet available."
 Nonetheless. we were able to determine an approximate transformation.," Nonetheless, we were able to determine an approximate transformation."
 One of the clusters from the oresent study — Reticulum — has previously been observed with HST'WFPC? through the F555W and FSIJW filters (ας part of program 5897)., One of the clusters from the present study – Reticulum – has previously been observed with /WFPC2 through the F555W and F814W filters (as part of program 5897).
 Calibrations for these filters to the Cousins V/ system exist and are well established (Holtzmanetal.1995:Dolphin2000b)..," Calibrations for these filters to the Johnson-Cousins $VI$ system exist and are well established \cite{holtzman,dolphin}."
 The relevant archived data-groups are labelled u2xj0605b (5... F555W frames — two with exposure durations of 260 s and three with exposure durations of 1000 ο) and uU2xJ0608b (6. FSIJW frames — two with exposure durations of 260 s and four with exposure durations of 1000 s)., The relevant archived data-groups are labelled u2xj0605b $5 \times$ F555W frames – two with exposure durations of $260$ s and three with exposure durations of $1000$ s) and u2xj0608b $6 \times$ F814W frames – two with exposure durations of $260$ s and four with exposure durations of $1000$ s).
 Unfortunately. these observations did not image the cluster core. but rather are centred approximately 1.5 to the north west.," Unfortunately, these observations did not image the cluster core, but rather are centred approximately $1.5\arcmin$ to the north west."
 Fig., Fig.
 4. shows the positions of the two observation sets overlaid on a DSS image of Reticulum., \ref{f:ret} shows the positions of the two observation sets overlaid on a DSS image of Reticulum.
 While there is a small overlap. the number of common stars was not large enough to derive a point-by-point photometric transformation.," While there is a small overlap, the number of common stars was not large enough to derive a point-by-point photometric transformation."
 Nonetheless. it was still possible to calculate a global transformation.," Nonetheless, it was still possible to calculate a global transformation."
 First. photometric measurements were performed on the archival ΜΕΡΟΣ images usingHSTPHOT (Dolphin2000:).," First, photometric measurements were performed on the archival WFPC2 images using \cite{hstphot}."
. The measurement procedure is described fully by Mackey (2004).. and is identical to the procedure used by Mackey Gilmore to measure RR Lyrae stars in the globular clusters of the Fornax dwarf galaxy.," The measurement procedure is described fully by Mackey \shortcite{thesis}, and is identical to the procedure used by Mackey Gilmore \shortcite{rrlyr} to measure RR Lyrae stars in the globular clusters of the Fornax dwarf galaxy."
 The resultant CMD may be seen in Fig. 5.., The resultant CMD may be seen in Fig. \ref{f:retmatch}.
 We next determined fidueials for both the CMD from the present study (in the ACS/WFC STmag system) and the CMD from the archival WFPC? data tin V. V.IY.," We next determined fiducials for both the CMD from the present study (in the ACS/WFC STmag system) and the CMD from the archival WFPC2 data (in $V$, $V-I$ )."
 The main sequences are well populated enough that the fiducials below the turn off could be calculated by forming magnitude bins and finding the mode in colour for each., The main sequences are well populated enough that the fiducials below the turn off could be calculated by forming magnitude bins and finding the mode in colour for each.
 This would typically also work for the RGB. but on neither CMD is this region particularly well populatec (Reticulum is a very sparse cluster).," This would typically also work for the RGB, but on neither CMD is this region particularly well populated (Reticulum is a very sparse cluster)."
 Thus. the RGB fiducials were determined by eye. as were the tiducials around the turn-off and sub-giant branch (SGB).," Thus, the RGB fiducials were determined by eye, as were the fiducials around the turn-off and sub-giant branch (SGB)."
 These latter could not be measured via a simple binning technique because they are neither horizonta nor vertical., These latter could not be measured via a simple binning technique because they are neither horizontal nor vertical.
 As can be seen in Figs., As can be seen in Figs.
 3. and 5. the dispersion in these regions of each CMD is small (particularly for the ACS measurements). so the by-eye procedure should not introduce any large errors.," \ref{f:cmdsub} and \ref{f:retmatch}, the dispersion in these regions of each CMD is small (particularly for the ACS measurements), so the by-eye procedure should not introduce any large errors."
 This algorithm is very similar to many used in previous studies (see e.g.. Johnson et al. (1999))).," This algorithm is very similar to many used in previous studies (see e.g., Johnson et al. \shortcite{johnson:99}) )."
 The horizontal branch (HB) levels were determined using the (very few) stars just to the blue of the instability strip., The horizontal branch (HB) levels were determined using the (very few) stars just to the blue of the instability strip.
 Due to the lack of multi-epoch observations (especially for the ACS data). the RR Lyrae regions possess a significant spread in magnitude due to the intrinsic stellar variability.," Due to the lack of multi-epoch observations (especially for the ACS data), the RR Lyrae regions possess a significant spread in magnitude due to the intrinsic stellar variability."
 On each CMD. the colour of the main sequence turn-off (MSTO) was determined by fitting a second-order polynomial to the the data in this region. and finding the bluest point of the fit.," On each CMD, the colour of the main sequence turn-off (MSTO) was determined by fitting a second-order polynomial to the the data in this region, and finding the bluest point of the fit."
 The magnitude of a main-sequence (MS) reference point σος (the point on the MS which is 0.05 mag redder than the MSTO) was then determined by interpolating along the tiducial line., The magnitude of a main-sequence (MS) reference point $V_{0.05}$ (the point on the MS which is $0.05$ mag redder than the MSTO) was then determined by interpolating along the fiducial line.
 These two values are traditionally used to register cluster CMDs for the purposes of differential age comparison (VandenBerg (see Section 4.3)): however our purpose here was to determine a linear shift between the ACS/WFC photometry. and the standard Johnson-Cousins scale.," These two values are traditionally used to register cluster CMDs for the purposes of differential age comparison \cite{vbs} (see Section \ref{ss:ages}) ); however our purpose here was to determine a linear shift between the ACS/WFC photometry, and the standard Johnson-Cousins scale."
 By moving the ACS/WFC fiducial so that the colour of its MSTO matched that for the WFPC? fiducial. and its Vou; matehed that for the ΜΕΡΟΣ fiducial. this transformation was calculated.," By moving the ACS/WFC fiducial so that the colour of its MSTO matched that for the WFPC2 fiducial, and its $V_{0.05}$ matched that for the WFPC2 fiducial, this transformation was calculated."
 We found that VHipnnnw=0.03 and (V. 1.31., We found that $V - m_{{\rm F555W}} = 0.03$ and $(V-I) - (m_{{\rm F555W}}-m_{{\rm F814W}}) = 1.31$ .
 The registered fiducials may be seen in the top panel of Fig. 5.. ," The registered fiducials may be seen in the top panel of Fig. \ref{f:retmatch}, ,"
while the shifted ACS tiducial is, while the shifted ACS fiducial is
 ming. a Ruugc-I&utta- method., 	 = 	4 (r) r^2 	 = 	G 	 = 	a ( - - 1 	 	 = 	- 			+ 			+ using a Runge-Kutta method.
 We- compare the profiles. obtained. in. this. way with. those from. the uiuuerical. simulations., We compare the profiles obtained in this way with those from the numerical simulations.
. The parameter Όρων Was already calculated by approximating. he VDPR. see Tab. 1..," The parameter $\Phi_{\rm{out}}$ was already calculated by approximating the VDPR, see Tab. \ref{table-halos}."
 We choose ry as second free parameter. since it causes only a stretch along the radial axis.," We choose $r_0$ as second free parameter, since it causes only a stretch along the radial axis."
" This xranmieter can casily. be adjusted. bv approximating. the aN. position ⋅⋅⋅iu the dispersion. profile. ,ez(r).", This parameter can easily be adjusted by approximating the maximum position in the dispersion profile $\sigma_r^2(r)$.
 First.. we asstuue that the velocity⋅ dispersion⋅ is ⋅⋅⋅isotropic.⋅ ο;= V.," First, we assume that the velocity dispersion is isotropic, $\beta=0$ ."
" The resulting. density. profiles approximate. roughly the results from the μπαΊος,", The resulting density profiles approximate roughly the results from the simulations.
 Clear deviatious occur in the iiuges of the halo., Clear deviations occur in the fringes of the halo.
 This are the regions where considerable anisotropy in the velocity dispersion is prescut., This are the regions where considerable anisotropy in the velocity dispersion is present.
" The anisotropy can be reasonably approximated by 3=yO/Oou, with 4j20.27."," The anisotropy can be reasonably approximated by $\beta = \beta_0 
\Phi / \Phi_{\rm{out}}$ with $\beta_0 \approx 0.27$."
 After integration. this clearly leads to a better concordance of the deusity profiles. see Fie. LL.," After integration, this clearly leads to a better concordance of the density profiles, see Fig. \ref{fig-Jeans}."
 For the fit of .J a dependence on the potential ® was assumed in order to avoid the inplemieutation of additional parameters., For the fit of $\beta$ a dependence on the potential $\Phi$ was assumed in order to avoid the implementation of additional parameters.
 It can be casily seen that ij roughly follows the shape of &. see Fig. 1..," It can be easily seen that $\beta$ roughly follows the shape of $\Phi$, see Fig. \ref{fig-profiles}."
 Tn adelition. the results are almost inscusitive with respect to variations of the depeudence Jon ας lone as J ds a monotonic function.," In addition, the results are almost insensitive with respect to variations of the dependence $\beta$ on $r$ as long as $\beta$ is a monotonic function."
 For the halo Cl6. the integration of the differcutial equations leads still only to a poor approxination of the nuuerical results.," For the halo Cl6, the integration of the differential equations leads still only to a poor approximation of the numerical results."
 This halo shows clear sigus for mudergoing still a lereine process., This halo shows clear signs for undergoing still a merging process.
 The dispersion profile shows a peal at zzL/Sia. Which is caused by a remnant of a merger between the cluster and a eroup about 1Cyr ago.," The dispersion profile shows a peak at $\approx 1/8 \, r_{\rm{vir}}$, which is caused by a remnant of a merger between the cluster and a group about $1\;\rm{Gyr}$ ago."
 Thus this halo is still in the process of relaxation aud therefore. can be only roughly approximated. by spherical.. relaxed svsten.," Thus this halo is still in the process of relaxation and therefore, can be only roughly approximated by spherical, relaxed system."
 the halos∙ CIÀThe and Gal areprofiles recovered of by the inteeration using the mean over all halos for the parameters 4 and ss, The profiles of the halos ClA and Gal are recovered by the integration using the mean over all halos for the parameters $a$ and $\kappa$.
 However. all profiles obtained from the simulations are slightly steeper in the ceuter than the inteerated: curves.," However, all profiles obtained from the simulations are slightly steeper in the center than the integrated curves."
 :? found: that the inuer: profiles⋅ of: disturbed. halos are steepened., \citet{dekel:02} found that the inner profiles of disturbed halos are steepened.
 Therefore.. à more cuspy central profile. compared to a perfectly relaxed. system. may be caused by: some recent niereer activity.," Therefore, a more cuspy central profile, compared to a perfectly relaxed system, may be caused by some recent merger activity."
 The: inteerated: profiles⋅ show a sharp break for. radiiον much larger than the virial radius., The integrated profiles show a sharp break for radii much larger than the virial radius.
 They do uot show the asvinptotic density profile expected from the discussion in Sec. 5.., They do not show the asymptotic density profile expected from the discussion in Sec. \ref{sec-asymp}.
" This is due to the fact that the outer potential o, is reached at a finite radius. at which the dispersion aud the deusitv vanishes."," This is due to the fact that the outer potential $\Phi_{\rm{out}}$ is reached at a finite radius, at which the dispersion and the density vanishes."
 However. since the radius where the dispersion vaulshes is nich larger than the virial radius of the considered halos. this outer slope cannot be determined by sinulatious designed to investigate structure formation.," However, since the radius where the dispersion vanishes is much larger than the virial radius of the considered halos, this outer slope cannot be determined by simulations designed to investigate structure formation."
 The unmerous umuerical studies of the structure of dark matter halos indicate that perfectly virialized svstenmis slow prestunably a universal density profile., The numerous numerical studies of the structure of dark matter halos indicate that perfectly virialized systems show presumably a universal density profile.
 The latter cau be approximated reasonably well bv the NEW-profile or a eoncralized. version. of it., The latter can be approximated reasonably well by the NFW-profile or a generalized version of it.
 ≯⋅⊲∙∙Similarly we have introduced. the relation. between velocity. dispersion.: aud poteutial. Eq. (1)) (, Similarly we have introduced the relation between velocity dispersion and potential Eq. \ref{eq-of-state}) ) (
VDPR):it characterizes the profile by very few parameters.,VDPR): it characterizes the profile by very few parameters.
 Iu coutrast to the approximations of the density profiles the VDPR is nuplicit iu the seuse that it as à function of the poteutial imstead of the radius., In contrast to the approximations of the density profiles the VDPR is implicit in the sense that it as a function of the potential instead of the radius.
 However. by introducing the .VDPR; we can put restrictions⋅⋅ even ou the inner⋅ asviuptotic⋅ slope of.. both. the velocity. dispersion. profile. aud the deusity. profile.," However, by introducing the VDPR we can put restrictions even on the inner asymptotic slope of, both, the velocity dispersion profile and the density profile."
 The data cau be approximated verv well by the newly introduced VDPR., The data can be approximated very well by the newly introduced VDPR.
" Almost all halos clearly show a power-law for. στ),xΦ near the ceuter. the best resolved lalos do so over one order of magnitude."," Almost all halos clearly show a power-law for $\sigma_r^2(\Phi) \propto \Phi^\kappa$ near the center, the best resolved halos do so over one order of magnitude."
 We, We
We consider the Supernovae Type Ia observation AAmanullah et al. (,We consider the Supernovae Type Ia observation Amanullah et al. (
2010) which is one of the direct probes for the late time acceleration.,2010) which is one of the direct probes for the late time acceleration.
" It measures the apparent brightness of the Supernovae as observed by us which is related to the luminosity distance dr(z) defined as With this, we construct the distance modulus 4 which is experimentally measured Where m and M are the apparent and absolute magnitudes of the Supernovae which are logarithmic measure of flux and luminosity respectively."," It measures the apparent brightness of the Supernovae as observed by us which is related to the luminosity distance $d_{L}(z)$ defined as With this, we construct the distance modulus $\mu$ which is experimentally measured Where m and M are the apparent and absolute magnitudes of the Supernovae which are logarithmic measure of flux and luminosity respectively."
 Another observational probe that has been widely used in recent times to constrain dark energy models is related to the data from the BAO distance measurements obtained at 2Ξ0.2 and z=0.35 from the joint analysis of the 2dFGRS and SDSS data EEisenstein et al., Another observational probe that has been widely used in recent times to constrain dark energy models is related to the data from the BAO distance measurements obtained at $z = 0.2$ and $z = 0.35$ from the joint analysis of the 2dFGRS and SDSS data Eisenstein et al.
 PPercival et al., Percival et al.
 PPercival et al. (, Percival et al. (
2010)).,2010)).
" In this case,one needs to calculate the parameter D, which is related to the angular diameter distance as follows For BAO measurements we calculate the ratio this ratio is a relatively model independent quantity and has a value 1.736+0.065."," In this case,one needs to calculate the parameter $D_{v}$ which is related to the angular diameter distance as follows For BAO measurements we calculate the ratio this ratio is a relatively model independent quantity and has a value $1.736 \pm 0.065$."
 The CMB is sensitive to the distance to the decoupling epoch via the locations of peaks and troughs of the acoustic oscillations., The CMB is sensitive to the distance to the decoupling epoch via the locations of peaks and troughs of the acoustic oscillations.
 We employ the “WMAP distance priors” given by the five-year WMAP observations., We employ the “WMAP distance priors” given by the five-year WMAP observations.
"This includes the“acoustic scale"" /A,the “ shift parameter "" R","This includes the“acoustic scale” $l_{A}$ ,the “ shift parameter ” R"
MIDI. and difference spectrum in a consistent manner.,"MIDI, and difference spectrum in a consistent manner."
 We approximated a continuum with a straight line intersecting the spectra at 8.3 and um. subtracted this from the spectra and then normalized the spectra with the flux level at m. The bottom panel of reffig:tenmicron shows the result.," We approximated a continuum with a straight line intersecting the spectra at 8.3 and $\mu$ m, subtracted this from the spectra and then normalized the spectra with the flux level at $\mu$ m. The bottom panel of \\ref{fig:tenmicron} shows the result."
 The shape of the silicate feature is very similar. which implies that the shoulc also be very similar.," The shape of the silicate feature is very similar, which implies that the should also be very similar."
 Before we describe a detailed modeling effort of the circumstellar material. we summarize the analyses of the various observations.," Before we describe a detailed modeling effort of the circumstellar material, we summarize the analyses of the various observations."
 This already gives an insight into the spatial distribution of the gas and the dust., This already gives an insight into the spatial distribution of the gas and the dust.
 The AMBER data probes the very inner parts of the disk., The AMBER data probes the very inner parts of the disk.
 The analysis shows that the K-band emission could be explained with an emitting ring at «0.4 AAU., The analysis shows that the K-band emission could be explained with an emitting ring at $\sim$ AU.
 A more physical model. which will be presented in refsec:SED has the inner rim of the disk at AAU.," A more physical model, which will be presented in \\ref{sec:SED} has the inner rim of the disk at AU."
" A part of the K-band emission was shown to This could be an indicator of ongoing accretion,", A part of the K-band emission was shown to This could be an indicator of ongoing accretion.
 The I1] emission line is formed by the photo-dissociation of OH molecules by UV photons(?)., The I] \\r{A} emission line is formed by the photo-dissociation of OH molecules by UV photons.
. The detection of the II] rA at large distances (from one to tens of AU) from the central star is thus an indication that the outer disk has an illuminated gas surface., The detection of the I] \\r{A} at large distances (from one to tens of AU) from the central star is thus an indication that the outer disk has an illuminated gas surface.
 The Spitzer data establish the presence of PAH emission and the VISIR spectrum pins it down to a circumstellar disk., The Spitzer data establish the presence of PAH emission and the VISIR spectrum pins it down to a circumstellar disk.
 The emission features of PAH molecules are caused by internal vibrational modes. which are mainly excited by UV photons.," The emission features of PAH molecules are caused by internal vibrational modes, which are mainly excited by UV photons."
 Since the PAH molecules are coupled to the gas. the PAH emission is another indicator of an illuminated gas surface.," Since the PAH molecules are coupled to the gas, the PAH emission is another indicator of an illuminated gas surface."
 The resolved VISIR spectrum sets the radial scale of this gas surface at ~ AAU., The resolved VISIR spectrum sets the radial scale of this gas surface at $\sim$ AU.
 The VISIR Q-band image displays a faint extended emission component ( of the total flux). that stretches out to large radit (r ~ chAU)).," The VISIR Q-band image displays a faint extended emission component $\sim$ of the total flux), that stretches out to large radii (r $\sim$ )."
 The MIDI correlated flux spectrum shows that the PAH features originate from radit much larger than ~2 AAU and that the silicate composition is quite similar in the inner and outer disk., The MIDI correlated flux spectrum shows that the PAH features originate from radii much larger than $\sim$ AU and that the silicate composition is quite similar in the inner and outer disk.
 In this section we present a for the disk around 995881 which quite accurately reproduces the observations deseribed above., In this section we present a for the disk around 95881 which quite accurately reproduces the observations described above.
 To obtain this model we use the Monte Carlo radiative transfer code MCMax by ?.., To obtain this model we use the Monte Carlo radiative transfer code MCMax by \cite{2009A&A...497..155M}.
 This code can compute a self-consistent disk structure and a full range of observables., This code can compute a self-consistent disk structure and a full range of observables.
 It has a build-in option that models the full PAH excitation. using the temperature distribution approximation (see e.g. ?:: ?)) including multi-photon events for the excitation.," It has a build-in option that models the full PAH excitation, using the temperature distribution approximation (see e.g. \citealt{1989ApJ...345..230G}; \citealt{1992A&A...266..501S}) ) including multi-photon events for the excitation."
 We use MCMax here to compute the temperature structure. the vertical density structure and the resulting SED. the Spitzer spectrum. the AMBER visibilities. the MIDI correlated flux. the VISIR images. and the FWHM as function of wavelength.," We use MCMax here to compute the temperature structure, the vertical density structure and the resulting SED, the Spitzer spectrum, the AMBER visibilities, the MIDI correlated flux, the VISIR images, and the FWHM as function of wavelength."
 The steps to come to the final model presented here were the following., The steps to come to the final model presented here were the following.
 First we fixed the composition. and the size and shape distribution of the silicate component of the dust to be equal to that obtained from the 10 micron silicate feature by ?.. ," First we fixed the composition, and the size and shape distribution of the silicate component of the dust to be equal to that obtained from the 10 micron silicate feature by \citet{2005A&A...437..189V} ."
For 995881 these are large um) pyroxene grains. large enstatite grains. large forsterite grains. and small jm) silica grains.," For 95881 these are large $\mu$ m) pyroxene grains, large enstatite grains, large forsterite grains, and small $\mu$ m) silica grains."
 In order to get the required continuum opacity needed we added amorphous carbon grains., In order to get the required continuum opacity needed we added amorphous carbon grains.
 To model the irregular shape of the carbon grains weDHS., To model the irregular shape of the carbon grains we.
. For the refractive index of carbon we adopted the data by ?.., For the refractive index of carbon we adopted the data by \citet{1993A&A...279..577P}.
 Note that continuum component is most likely not all in the form of camorphouse carbon., Note that continuum component is most likely not all in the form of amorphous carbon.
 Small grains of metallic iron and/or tron sulfide have extinction. properties similar to that of carbon., Small grains of metallic iron and/or iron sulfide have extinction properties similar to that of carbon.
 Also. large grains of various dust species could produce the observed continuum component.," Also, large grains of various dust species could produce the observed continuum component."
 The abundance of amorphous carbon is a fitting parameter., The abundance of amorphous carbon is a fitting parameter.
 The other free parameters all have to do with the geometry of the disk., The other free parameters all have to do with the geometry of the disk.
 As a first step we focused on the thermal dust grains. ignoring the PAH bands.," As a first step we focused on the thermal dust grains, ignoring the PAH bands."
 The density distribution of the dust disk was parameterized using a radial surface density (2)) for Rin«rRow., The density distribution of the dust disk was parameterized using a radial surface density \citealt{2008ApJ...678.1119H}) ) for $R_\mathrm{in} < r < R_\mathrm{out}$.
 Here Ro is the turnover point from where an exponential decay ofthe surface density sets in and p sets the powerlaw in the inner region., Here $R_0$ is the turnover point from where an exponential decay ofthe surface density sets in and $p$ sets the powerlaw in the inner region.
 We fixed this powerlaw to p= l.acommonly used value (see e.g. ?)).," We fixed this powerlaw to $p=1$ , a commonly used value (see e.g. \citealt{2006ApJ...640L..67D}) )."
purposes of this Letter.,purposes of this Letter.
 Using the fluxes presented in Table |. we can use several methods to constrain the redshift. reddening. infrared lummosity and dust temperature of SSG 1.," Using the fluxes presented in Table 1, we can use several methods to constrain the redshift, reddening, infrared luminosity and dust temperature of SSG 1."
 It can be seen in Figure 2 that the optical-NIR SED of SSG is characterized by a prominent bump associated with the continuum| emission of the stellar populations., It can be seen in Figure 2 that the optical-NIR SED of SSG 1 is characterized by a prominent bump associated with the continuum emission of the stellar populations.
 This bump peaks at jim in the rest-frame. providing a very good constraint on the redshift. and also indicates SSG | is more likely to be starburst rather than AGN-dominated (which. are usually associated with a featureless power-law spectrum: see. e.g.. Egami et al.," This bump peaks at $\mu$ m in the rest-frame, providing a very good constraint on the redshift, and also indicates SSG 1 is more likely to be starburst rather than AGN-dominated (which are usually associated with a featureless power-law spectrum; see, e.g., Egami et al."
 2004)., 2004).
 ImpZ (Babbedge et al., Z (Babbedge et al.
 2004) uses only the optical-NIR photometry when fitting templates with Bayesian statistics., 2004) uses only the optical-NIR photometry when fitting templates with Bayesian statistics.
 The best-fitting redshifts are 04. for SSG 1 and z=1.0+0.05 for SSG IE respectively (I0). and are listed in Table 3.," The best-fitting redshifts are $^{+0.1}_{-0.05}$ for SSG 1 and $\pm$ 0.05 for SSG 1E respectively $\sigma$ ), and are listed in Table 3."
 It is worth noting that if. as ImpZ suggests. both galaxies are at the same distance. and interacting. it would make the SSG pair a widely separated interacting ULIRG system (a local example of which is IRAS 09111-1007: Khan et al.," It is worth noting that if, as Z suggests, both galaxies are at the same distance, and interacting, it would make the SSG pair a widely separated interacting ULIRG system (a local example of which is IRAS 09111–1007; Khan et al."
 2005)., 2005).
 STARDUST? (Chantal et al..," STARDUST2 (Chanial et al.,"
 2005. in preparation) uses a mid-IR to radio spectral library constrained by the IR-radio correlation and a variety of local./RAS. and SCUBA color-color correlations in. addition to a_ stellar synthesis model for the FUV to NIR window (Devriendt et al.," 2005, in preparation) uses a mid-IR to radio spectral library constrained by the IR-radio correlation and a variety of local, and SCUBA color-color correlations in addition to a stellar synthesis model for the FUV to NIR window (Devriendt et al."
 1999)., 1999).
 This provides a simultaneous constraint on the thermal dust emission as well as the redshift., This provides a simultaneous constraint on the thermal dust emission as well as the redshift.
 The effective dust temperature is found following the methodology given in Chapman et al. (, The effective dust temperature is found following the methodology given in Chapman et al. (
2002; 2005).,2002; 2005).
" The model assumes Ho=70kms!Mpe!. Q,,=0.3. and O4=0.7. and that a single component is responsible for both the optical and IR emission,"," The model assumes $\rm H_0=70\,km\,s^{-1}\, Mpc^{-1}$ , $\rm \Omega_{m}=0.3$, and $\Omega_{\Lambda}=0.7$, and that a single component is responsible for both the optical and IR emission."
" Because of the degeneracy in fitting Ty, and the Total-IR Luminosity. Ly jm). it is the μπι flux in conjunction. with the radio upper limit that constrains the infrared luminosity."," Because of the degeneracy in fitting $\rm T_{dust}$ and the Total-IR Luminosity, $\rm L_{TIR}$ $\mu$ m), it is the $\mu$ m flux in conjunction with the radio upper limit that constrains the infrared luminosity."
" The radio upper limit (Table 1) ts used to find 47 as follows: C=2\7 with the radio v7 term being where f, is the modeled spectral energy distribution and 30 is the radio upper limit listed in Table I.", The radio upper limit (Table 1) is used to find $\chi^{2}$ as follows: $\chi^2=\sum \chi^2_i$ with the radio $\chi^2_i$ term being where $f_\nu$ is the modeled spectral energy distribution and $\sigma$ is the radio upper limit listed in Table 1.
 The best-fitting model for SSG | from STARDUST2 returns a redshift of 0.997201 (all errors are quoted to Io: note: STARDUST? uses the longer wavelengths. in addition to the optical-NIR photometry. to fit the extinction. given in Table 3).," The best-fitting model for SSG 1 from STARDUST2 returns a redshift of $^{+0.04}_{-0.03}$ (all errors are quoted to $\sigma$; note: STARDUST2 uses the longer wavelengths, in addition to the optical-NIR photometry, to fit the extinction given in Table 3)."
 This model also gives a dust temperature of 30.344.5 KK. and a log(Lyig) of LL... with of the total bolometric luminosity12.02455 radiated at wavelengths longward of j/m. The predicted jim flux is 2.7*2 mmly.," This model also gives a dust temperature of $\pm$ K, and a $\rm L_{TIR}$ ) of $^{+0.22}_{-0.24}$ $_{\odot}$, with of the total bolometric luminosity radiated at wavelengths longward of $\mu$ m. The predicted $\mu$ m flux is $^{+1.7}_{-0.7}$ mJy."
 Other redshift estimates were independently obtained following an approach that mostly relies on the fit of the jim feature (see. e.g.. Le Floc'h et al.," Other redshift estimates were independently obtained following an approach that mostly relies on the fit of the $\mu$ m feature (see, e.g., Le Floc'h et al."
 2004)., 2004).
 In agreement with the results derived from ImpZ and STARDUST?. these fits of the stellar bump led to a redshift z—1 using the photometric Arp220 SED. and z=1.25+0.25 using various templates from the library of Devriendt et al. (," In agreement with the results derived from Z and STARDUST2, these fits of the stellar bump led to a redshift $\sim$ 1 using the photometric Arp220 SED, and $\pm$ 0.25 using various templates from the library of Devriendt et al. ("
1999).,1999).
 We also use the radio-submillimeter correlation to estimate a minimum redshift for the source. as the radio flux is unlikely to be significantly AGN-enhanced.," We also use the radio-submillimeter correlation to estimate a minimum redshift for the source, as the radio flux is unlikely to be significantly AGN-enhanced."
 Formally. these correlations should be used as a statistical redshift indicator. so with that caveat we present the minimum redshift of SSG | based on the GGHz upper limit (Table 1) and an jim flux derived from the best-fitting STARDUST2 SED (Figure 2).," Formally, these correlations should be used as a statistical redshift indicator, so with that caveat we present the minimum redshift of SSG 1 based on the GHz upper limit (Table 1) and an $\mu$ m flux derived from the best-fitting STARDUST2 SED (Figure 2)."
 Using the relations of Dunne. Clements Eales (2000) and Carilli Yun (2000a: 2000b) we get zii of 1.2 and 1.5 respectively.," Using the relations of Dunne, Clements Eales (2000) and Carilli Yun (2000a; 2000b) we get $\rm z_{min}$ of 1.2 and 1.5 respectively."
 The photometric redshift is more secure than the radio-submillimeter correlation. since the latter is prone to systematics caused by uncertain dust temperatures (see Clements et al., The photometric redshift is more secure than the radio-submillimeter correlation since the latter is prone to systematics caused by uncertain dust temperatures (see Clements et al.
 2004)., 2004).
" The (R-Kyeu, color of 5.7 for SSG 1 classifies it as an Extremely Red Object (ERO). and both ImpZ and STARDUST? find this object to be highly reddened (Table 3)."," The $\rm K_{s})_{Vega}$ color of 5.7 for SSG 1 classifies it as an Extremely Red Object (ERO), and both Z and STARDUST2 find this object to be highly reddened (Table 3)."
 At least a third of the SMG population can be classified as EROs (Smail et al., At least a third of the SMG population can be classified as EROs (Smail et al.
 2002: Webb et al., 2002; Webb et al.
 2004: Frayer et al., 2004; Frayer et al.
 2004). which comprise two classes of galaxies: elliptical and dusty star-forming luminous infrared galaxies (but à significant submillimeter flux is usually indicative of the latter class).," 2004), which comprise two classes of galaxies: elliptical and dusty star-forming luminous infrared galaxies (but a significant submillimeter flux is usually indicative of the latter class)."
 EROs are thought to comprise a significant fraction of the cosmic star formation density at redshifts of one and higher (Cimatti et al., EROs are thought to comprise a significant fraction of the cosmic star formation density at redshifts of one and higher (Cimatti et al.
 2002). the epoch by which the majority of the universe's star formation has taken place (Dickinson et al.," 2002), the epoch by which the majority of the universe's star formation has taken place (Dickinson et al."
 2003: Rudnick et al., 2003; Rudnick et al.
 2003)., 2003).
where /2=B5; is the trace of the Ievnolds tensor. which is twice the turbulent kinetic energy. per unit mass. and C. C. C's and € are positive dimensionless coefficients oforder unity. of a universal nature. (,"where $R=R_{ii}$ is the trace of the Reynolds tensor, which is twice the turbulent kinetic energy per unit mass, and $C_1$ , $C_2$, $C_6$ and $C_7$ are positive dimensionless coefficients of order unity, of a universal nature. ("
Coellicients Cy. Cy and C5 are reserved for a magnetohyvdrodynamic extension of the model. see Ogilvie 2003) The justification for introducing non-linear terms of the above form. is similar to that used. in. the. model of maenetorotational turbulent stresses originally introduced by Oeilvie (2003).,"Coefficients $C_3$, $C_4$ and $C_5$ are reserved for a magnetohydrodynamic extension of the model, see Ogilvie 2003) The justification for introducing non-linear terms of the above form is similar to that used in the model of magnetorotational turbulent stresses originally introduced by Ogilvie (2003)."
" The term involving C', causes a dissipation of turbulent. kinetic energy. ancl allows for the free decay of hydrodynamic turbulence."," The term involving $C_1$ causes a dissipation of turbulent kinetic energy, and allows for the free decay of hydrodynamic turbulence."
 The term involving C» pedistributes energy among the components of #7). and corresponds to the tendeney of hydrodynamic turbulence to return to isotropy through the effect of the pressurestrain correlation.," The term involving $C_2$ redistributes energy among the components of $\bar R_{ij}$, and corresponds to the tendency of hydrodynamic turbulence to return to isotropy through the effect of the pressure–strain correlation."
 Both are constructed assuming that these effects occur on a timescale related to the eddy. turnover time. LAB. where L is delined as the typical scale of the largest turbulent οὐστον.," Both are constructed assuming that these effects occur on a timescale related to the eddy turnover time, $ L / \bar R^{1/2}$, where $L$ is defined as the typical scale of the largest turbulent eddies."
 Terms €; and C. related to the transport of heat. are advanced: by simple analogy.," Terms $C_6$ and $C_7$, related to the transport of heat, are advanced by simple analogy."
" The coefficients must satisfv certain conditions to ensure the realizability of the model. as ciscussed in Appendix A. The terms proportional to the microscopic. cilfusion cocllicients are. introduced. to allow a modelling. of the correlation. terms μηOy,u,ο te|ROGOnIO;O' and 25((G0;0')73 at moderate Revnolds number. ic. close to the onset of convection."," The coefficients must satisfy certain conditions to ensure the realizability of the model, as discussed in Appendix A. The terms proportional to the microscopic diffusion coefficients are introduced to allow a modelling of the correlation terms $2\nu\langle \partial_{k} u_i'\partial_{k }u_j'\rangle$, $(\nu+ \kappa) \langle \partial_j u_i' \partial_{j}\Theta'\rangle$ and $2\kappa\langle (\partial_{i}\Theta')^2 \rangle$ at moderate Reynolds number, i.e. close to the onset of convection."
 In such a situation a turbulent cascade does not form and the dissipative ternis are proportional to. rather than independent of. the diffusion coefficients.," In such a situation a turbulent cascade does not form and the dissipative terms are proportional to, rather than independent of, the diffusion coefficients."
 In a similar way. for turbulent shear flows. GO05 proposed. to mocel the momentum cillusion term as on dimensional grouncls.," In a similar way, for turbulent shear flows, GO05 proposed to model the momentum diffusion term as on dimensional grounds."
 Indeed. it is expected that near the onset of convection. most [uidi motions will be on the largest scales of the svstem (L).," Indeed, it is expected that near the onset of convection, most fluid motions will be on the largest scales of the system $L$)."
 By analogy. we mocel the other two terms here as Therefore the dissipative term in each of equations (25)) (27)) is modelled by a sum of two terms. one that is independent of the dilfusivity and. dominates at high ]tevnolds numbers. and another that is proportional to the cilfusivity and dominates at moderate Itevnolds numbers.," By analogy, we model the other two terms here as Therefore the dissipative term in each of equations \ref{dtrij}) \ref{dtq}) ) is modelled by a sum of two terms, one that is independent of the diffusivity and dominates at high Reynolds numbers, and another that is proportional to the diffusivity and dominates at moderate Reynolds numbers."
 This completes the justification for the form of the closure mocel proposed in equations (28)). (29)) and (30)).," This completes the justification for the form of the closure model proposed in equations \ref{eq:Rprop}) ), \ref{eq:Fprop}) ) and \ref{eq:Qprop}) )."
 We now apply the closure model to the problem of RayleighXnnard convection., We now apply the closure model to the problem of Rayleigh--B\'ennard convection.
 We consider a horizontally infinite. plane-parallel system. where the bottom plate is located at height >—0 and the top plate at height +=.," We consider a horizontally infinite, plane-parallel system, where the bottom plate is located at height $z=0$ and the top plate at height $z=h$."
 The relative temperature of the bottom plate is O=AT’ while that of the top plate is O=0., The relative temperature of the bottom plate is $\bar \Theta = \Delta T$ while that of the top plate is $\bar \Theta= 0$.
" In this setup. we look for statistically steadyand horizontally homogeneousIn] solutions assuminge that mean quantities anc correlations between [luctuating quantities vary only with ο,"," In this setup, we look for statistically steadyand horizontally homogeneous solutions assuming that mean quantities and correlations between fluctuating quantities vary only with $z$."
 We also assume that there are no mean Hows in the svstem., We also assume that there are no mean flows in the system.
 Equations (13))-(15)) and. (28))-(30)) reduce to a set of ordinary. differential equations. (ODIEs) which can be solved to obtain the temperature profile O(z) between the two plates. the profiles of the turbulent kinetic energy. A0(z)/2. and the temperature. variance. (Q(z) (for example).," Equations \ref{eq:meancont}) \ref{eq:meanenergy}) ) and \ref{eq:Rprop}) \ref{eq:Qprop}) ) reduce to a set of ordinary differential equations (ODEs) which can be solved to obtain the temperature profile $\bar \Theta(z)$ between the two plates, the profiles of the turbulent kinetic energy, $\bar R(z) / 2$, and the temperature variance, $\bar Q(z)$ (for example)."
 ὃν analogv with Prandtls mixine-leneth Formulation (Prandtl. 1932) we set L. the size of the largest eddies. to be equal to the distance to the nearest wall. Σπα (See GOOS for applications of the same principle to pipe lows and to Couette.Tavlor Lows).," By analogy with Prandtl's mixing-length formulation (Prandtl, 1932) we set $L$, the size of the largest eddies, to be equal to the distance to the nearest wall, $L(z)=\min(z,h-z)$ (see GO05 for applications of the same principle to pipe flows and to Couette–Taylor flows)."
" Lt can be shown with little ellort that Ry,=ft.—0. as well as f=£P,—0."," It can be shown with little effort that $\bar R_{xy} = \bar R_{xz} 
= \bar R_{yz} = 0$ , as well as $\bar F_x = \bar F_y = 0$."
 The remaining set of five seconc-order ODEs fully characterizes the svstem: where g=g-., The remaining set of five second-order ODEs fully characterizes the system: where $g = -g_z$.
 In the case of no-slip boundaries with fixed temperature on cach plate as listed above. I... E. and Q are zero on both boundaries.," In the case of no-slip boundaries with fixed temperature on each plate as listed above, $\bar R$, $\bar R_{zz}$, $\bar F_z$ and $\bar Q$ are zero on both boundaries."
 This system of ODEs with associated boundary conditions can be solved: with a two-point bouncdary-value solver., This system of ODEs with associated boundary conditions can be solved with a two-point boundary-value solver.
 JFvpical solutions are shown in Fie., Typical solutions are shown in Fig.
 1. for various Ravleigh numbers. defined here as We set the Prandtl number to 1 for the purposes of illustration.," \ref{fig:raplots} for various Rayleigh numbers, defined here as We set the Prandtl number to $1$ for the purposes of illustration."
 Note. the appearance of the characteristically Hat. temperature profile between the two plates as Ra»o and of the thin thermal boundary lavers., Note the appearance of the characteristically flat temperature profile between the two plates as Ra $\rightarrow \infty$ and of the thin thermal boundary layers.
 We now study in more detail thestructure of the solution., We now study in more detail thestructure of the solution.
 As in the case of shear [lows past a wall (see C:O05). we can derive a universal profile for convection away from a wall.," As in the case of shear flows past a wall (see GO05), we can derive a universal profile for convection away from a wall."
 Let us consider a semi-inlinite domain z2 0. in which case L— z. and let fy be the convective heat [ux through the," Let us consider a semi-infinite domain $z>0$ , in which case $L = z$ , and let $F_0$ be the convective heat flux through the"
The direct comparison of images taken in different enerev bands (e.g.. optical and can be achieved through accurate image superposition.,"The direct comparison of images taken in different energy bands (e.g., optical and X-rays) can be achieved through accurate image superposition."
 Following the approach used bv Caraveo οἱ ((2001b). we superimposed the Ἁνναν image onto the optical ones wilh respect to the absolute (0.9) reference Irae. relving on the astrometric solution of each image.," Following the approach used by Caraveo et (2001b), we superimposed the X-ray image onto the optical ones with respect to the absolute $\alpha$ $\delta$ ) reference frame, relying on the astrometric solution of each image."
 We used the combined ACIS image. where individual images were co-aligned with an accuracy of 077 (Pavlov οἱ 22003).," We used the combined ACIS image, where individual images were co-aligned with an accuracy of $0\farcs7$ (Pavlov et 2003)."
 The error of the absolute aspect is (wpically about 076., The error of the absolute aspect is typically about $0\farcs6$.
" Therefore. we adopt 1"" as a reasonable estimate for the uncertainty of the absolute X-ray astrometry."," Therefore, we adopt $1''$ as a reasonable estimate for the uncertainty of the absolute X-ray astrometry."
 For the WEPC? images. the default astrometric solution across the focal plane is derived from the coordinates of the (wo guide stars used to point the telescope. which are taken as a reference to compute the astrometric reference point aud the telescope roll angle.," For the WFPC2 images, the default astrometric solution across the focal plane is derived from the coordinates of the two guide stars used to point the telescope, which are taken as a reference to compute the astrometric reference point and the telescope roll angle."
 The accuracy of theHS astrometry is limited bv the intrinsic error on the absolute coordinates ol the GSCI.1 (Guide Star Catalog) stars (Lasker et 11990) which are used for the telescope pointing., The accuracy of the astrometry is limited by the intrinsic error on the absolute coordinates of the GSC1.1 (Guide Star Catalog) stars (Lasker et 1990) which are used for the telescope pointing.
 According to the current estimates. the mean uncertainty of the absolute positions quoted in the GSCI.1 is about 0788 per coordinate(Diretta et 22002).," According to the current estimates, the mean uncertainty of the absolute positions quoted in the GSC1.1 is about 8 per coordinate(Biretta et 2002)."
inner part of the disc.),inner part of the disc.)
 In this case. equations (10)) (13)) are. after the usual separation of variables. reduced to where the subseript I2 refers το eccentric mode quantities.," In this case, equations \ref{free1}) \ref{free2}) ) are, after the usual separation of variables, reduced to where the subscript E refers to eccentric mode quantities."
 This system admits a solution of the form where £(r) is the eccentricity of the disc at radius r. and satisfies Again. this equation is closely related. but not identical. to equation (21) of Goodchild&Ogilvie(2006).. which was derived from an analysis of global eccentricity in a two-dimensional disc.," This system admits a solution of the form where $E(r)$ is the eccentricity of the disc at radius $r$, and satisfies Again, this equation is closely related, but not identical, to equation (21) of \cite{goodchildogilvie2006}, which was derived from an analysis of global eccentricity in a two-dimensional disc."
 Racially propagating solutions are obtained because #7<O7 in a velativistic disc., Radially propagating solutions are obtained because $\kappa^2<\Omega^2$ in a relativistic disc.
 As for the warp tilt Wr). we solve this equation numerically. using a 4th order RungeIxutta method with boundary conditions Min)=Ly. and d£fdr(r)=0. where fy is an arbitrary value for the cecentricity at the inner boundary. Le. at the mareinally stable orbit.," As for the warp tilt $W(r)$, we solve this equation numerically, using a 4th order Runge–Kutta method with boundary conditions $E(r_\mathrm{in})=E_0$, and $\textrm{d}E/\textrm{d}r(r_\mathrm{in})=0$, where $E_0$ is an arbitrary value for the eccentricity at the inner boundary, i.e., at the marginally stable orbit."
 A tvpical solution for £(r) is shown in Fig. 6.., A typical solution for $E(r)$ is shown in Fig. \ref{ecce}.
 The eccentricity has an oscillatory behaviour. Wih the wavelength decreasing with radius. consistent with thie© local dispersion relation.," The eccentricity has an oscillatory behaviour, with the wavelength decreasing with radius, consistent with the local dispersion relation."
 Again. we defer to a second paer à more realistic treatment of the propagation of the eccentricity into the inner part of the disc.," Again, we defer to a second paper a more realistic treatment of the propagation of the eccentricity into the inner part of the disc."
 The excitation mechanism is similar to the one discussed in the previous section: a global deformation mode. which is now the eccentricity mode. characterized by (i.m.n)=(0. 1.0). interacts with a trapped r mode (uw.0.1) giving rise to an intermediate mode.," The excitation mechanism is similar to the one discussed in the previous section: a global deformation mode, which is now the eccentricity mode, characterized by $(\omega,m,n)=(0,1,0)$ , interacts with a trapped r mode $(\omega,0,1)$ giving rise to an intermediate mode."
 The latter then couples with the global mode to feedback on to the trapped r modo., The latter then couples with the global mode to feedback on to the trapped r mode.
 Coupling rules require the intermediate mode to have the same frequency as the trapped: wave and. as before. we have mj=1l. in the case where the trapped. r mode is axisvmimetric.," Coupling rules require the intermediate mode to have the same frequency as the trapped wave and, as before, we have $m_\textrm{I}=1$, in the case where the trapped r mode is axisymmetric."
 As for the vertical dependence. since the eccentric mode has »= 0. the intermediate mode can only have the same vertical mode number as the trapped. modo. be. mp=1.," As for the vertical dependence, since the eccentric mode has $n=0$ , the intermediate mode can only have the same vertical mode number as the trapped mode, i.e., $n_\textrm{I}=1$."
 The propagation region for this mode is the same as for the (v.1.2) intermediate mode. present in the interaction of the r mode with the warp.," The propagation region for this mode is the same as for the $(\omega,1,2)$ intermediate mode, present in the interaction of the r mode with the warp."
 As in that case. a relatively large «ampine term must be included. in the equations for the inermediate mode so that it dissipates on a resolved. scale beore reaching the corotation resonance.," As in that case, a relatively large damping term must be included in the equations for the intermediate mode so that it dissipates on a resolved scale before reaching the corotation resonance."
 In this case the energy. exchanges are similar to the ones discussed for the ineraction in a warped disc., In this case the energy exchanges are similar to the ones discussed for the interaction in a warped disc.
 lo [find the erowth rates that result [from this interaction. we solve equations (N34)). (A44)). using the same numerical mehod and boundary conditions as before.," To find the growth rates that result from this interaction, we solve equations \ref{inte1}) \ref{inte2}) ), using the same numerical method and boundary conditions as before."
 The variation of the growth rate with the innereccentricity. sound speed. spin of black hole and dissipation," The variation of the growth rate with the innereccentricity, sound speed, spin of black hole and dissipation"
ol DALQSOs.,of BALQSOs.
 Notably. in a eravitationally lensecl quasar sample. Chartas (2000) found a BALQSO fraction of35%... however. the sample size was small.," Notably, in a gravitationally lensed quasar sample, Chartas (2000) found a BALQSO fraction of, however, the sample size was small."
 Recently. Trump et al. (," Recently, Trump et al. ("
2006. TOG hereafter) identified more than 4.000 BALQSOs in the large SDSS DR3 quasar sample (Schneider et al.,"2006, T06 hereafter) identified more than 4,000 BALQSOs in the large SDSS DR3 quasar sample (Schneider et al."
 2005). and found a BALQSO fraction of26%.," 2005), and found a BALQSO fraction of."
. One important diffieultv in estimation of the true BALQSO Traction is the selection bias caused by the UV spectral difference between BALQSOs and non-BALQSOs. which includes both absorption troughs and continuum dillerences (Spravberry Foltz 1992: 03).," One important difficulty in estimation of the true BALQSO fraction is the selection bias caused by the UV spectral difference between BALQSOs and non-BALQSOs, which includes both absorption troughs and continuum differences (Sprayberry Foltz 1992; R03)."
 llowever. direct. moclelings of this selection bias lead to different. correction factors (e.g... Tlewelt Foltz 2003: R03). though Che issue is further complicated by different. sample selections.," However, direct modelings of this selection bias lead to different correction factors (e.g., Hewett Foltz 2003; R03), though the issue is further complicated by different sample selections."
 Allernatively. we can study the fraction of BALQSOs in a spectral regime that is nol significantlv affected by these biases to obtain the true fraction of BALQSOs.," Alternatively, we can study the fraction of BALQSOs in a spectral regime that is not significantly affected by these biases to obtain the true fraction of BALQSOs."
 In this paper. we examine the population of BALQSOs in the near infrared by analvzing the large sample of SDSS-BALQSOs (T06) combining 2421ASS (Skrutskie et al.," In this paper, we examine the population of BALQSOs in the near infrared by analyzing the large sample of SDSS-BALQSOs (T06) combining 2MASS (Skrutskie et al."
 2006) data., 2006) data.
 We note that the BALQSO fraction also depends on the definition of BALQSOs used by clillerent studies. and we adopt the definition used by TOG.," We note that the BALQSO fraction also depends on the definition of BALQSOs used by different studies, and we adopt the definition used by T06."
" We assume that Z/,=r0kms!Mpe |. Q4,= 0.3. and Q4=0.7 throughout the paper."," We assume that $H_0 = 70~\rm{km~s^{-1}~Mpc^{-1}}$ , $\Omega_{\rm m} = 0.3$ , and $\Omega_{\Lambda}= 0.7$ throughout the paper."
 We started [rom the SDSS DR3 quasar catalog (Schneider οἱ al., We started from the SDSS DR3 quasar catalog (Schneider et al.
 2005) and the catalog (TOG)., 2005) and the SDSS-BALQSO catalog (T06).
 In particular. we tareetecl the redshift range of 1.7xz<1383. where the majority of BALQSOs are identified with aabsorption in the observec-lrame optical band pass.," In particular, we targeted the redshift range of $1.7 \le z \le 4.38$, where the majority of BALQSOs are identified with absorption in the observed-frame optical band pass."
" We matched the quasars of both BALQSOs and non-BALQSOs (2"")) to entries [rom the [ull 24ÀSS release.", We matched the quasars of both BALQSOs and non-BALQSOs ) to entries from the full 2MASS release.
 We note that not all database entries from the full 2421ASS release in (he 2ATASS All-Skv. Point Source Catalog (PSC) satisly the conservative requirements to be part of the official PSC., We note that not all database entries from the full 2MASS release in the 2MASS All-Sky Point Source Catalog (PSC) satisfy the conservative requirements to be part of the official PSC.
 Therefore. the completeness of the database entries extend below the official PSC completenessIimits?.," Therefore, the completeness of the database entries extend below the official PSC completeness."
. The completeness levels of the database are J=16.1. 2—15.5. and A;=15.1 mag.," The completeness levels of the database are $J=16.1$, $H=15.5$, and $K_{s}=15.1$ mag."
 To ensure that matched database entries represent detections. we required an in-band detection ffle !," To ensure that matched database entries represent detections, we required an in-band detection flg !"
— 0). that there was no confusion. contamination [Πο == 0). or blending ffle < 1) in the source. and that no source was near an extended galaxy. (galecontam == 0 and kkev is null),"= 0), that there was no confusion, contamination flg == 0), or blending flg $\le$ 1) in the source, and that no source was near an extended galaxy contam == 0 and key is null)."
 Fig., Fig.
" 1. shows the A, mag versus / mag distributions lor BALQSOs and non-BALQSOs.", \ref{fig:ik} shows the $K_s$ mag versus $i$ mag distributions for BALQSOs and non-BALQSOs.
 The quasars are concentrated in a linear relation between (he 7 and Jv. magnitudeswithscatter., The quasars are concentrated in a linear relation between the $i$ and $K_s$ magnitudeswithscatter.
 The near infrared quasar sample is limited by the 2MASS flix, The near infrared quasar sample is limited by the 2MASS flux
"state spectra we fix Ai, at OA keV. Fits to 119 data sets (9054) result in the null hypothesisprobability. pau greater than5.","state spectra we fix $kT_{bb}$ at 0.4 keV. Fits to 119 data sets ) result in the null hypothesisprobability, $p_{\rm null}$, greater than."
.. The remaining 13 data sets with puuc5A show residuals around 67 keV which are reminiscent of an iron line emission. which is not computed together with the rellected component by the code.," The remaining 13 data sets with $p_{\rm null}>5$ show residuals around 6–7 keV which are reminiscent of an iron line emission, which is not computed together with the reflected component by the code."
 Thus in these 13 cases we include explicitly the relativistic line model of Laor (1991).LAOR.," Thus in these 13 cases we include explicitly the relativistic line model of Laor (1991),."
 We fix the emissivity of the line at 3. and the inner and outer radius at the values used to compute the rellected. component in.EQPAIR: 20 ancl 1000. gravitational racii. respectively.," We fix the emissivity of the line at 3, and the inner and outer radius at the values used to compute the reflected component in: 20 and 1000 gravitational radii, respectively."
 After addition of the line. the fits to 9 of 13 data sets resulted in pau5X.," After addition of the line, the fits to 9 of 13 data sets resulted in $p_{\rm null}>5$."
. We notice that it has been reported by eg. Sala οἱ al. (, We notice that it has been reported by e.g. Sala et al. (
2007) ancl Miller. et. al. (,2007) and Miller et al. (
2008) that NAIAI-Newton and Chandra spectra of GRO J1655-40. respectively. show features implying strong accretion disc winds.,"2008) that XMM-Newton and Chandra spectra of GRO J1655-40, respectively, show features implying strong accretion disc winds."
 A number of absorption lines in the 7Ὁ keV band have been observed., A number of absorption lines in the 7–9 keV band have been observed.
 Llowever. RAPE has a much poorer energy resolution than Chandra or NMM-Néeswton. and so we are not able to handle properly these absorption features.," However, RXTE has a much poorer energy resolution than Chandra or XMM-Newton, and so we are not able to handle properly these absorption features."
 Nevertheless. we conclude that we have found a mocel that fits well 128 of 132 considered data sets. which means that the moclel fits our set of data well from the statistical point of view (sec e.g. Sobolewska Papacakis 2009).," Nevertheless, we conclude that we have found a model that fits well 128 of 132 considered data sets, which means that the model fits our set of data well from the statistical point of view (see e.g. Sobolewska Papadakis 2009)."
 An example of soft/hard state models fitted to the GIRO data is shown in Figs., An example of soft/hard state models fitted to the GRO J1655-40 data is shown in Figs.
 2aa/c. Following Paper Lo we caleulate the disc-to-Comptonization index. agpy. for GRO J1655-40. based on the best fitting spectral models.," \ref{fig:sed}a a/c. Following Paper I, we calculate the disc-to-Comptonization index, $\alpha^{\prime}_{\rm GBH}$, for GRO J1655-40 based on the best fitting spectral models."
 The diagram of apy versus the (efids Uux in. keV. Pos teat 3 keV ds presented in Fig., The diagram of $\alpha^{\prime}_{\rm GBH}$ versus the $(ef_e)_3$ flux in keV $^{-2}$ $^{-1}$ at 3 keV is presented in Fig.
 Ibb. and the hard/soft spectral states that GRO 1655-40. entered during its evolution are indicated with different colors.," \ref{fig:loop}b b, and the hard/soft spectral states that GRO J1655-40 entered during its evolution are indicated with different colors."
 The values of acq for GRO J1655-40 range between | and 2. which is in general true also in the case of other (1111 (see Paper D. and in the case of AGN’s Clas.," The values of $\alpha^{\prime}_{\rm GBH}$ for GRO J1655-40 range between 1 and 2, which is in general true also in the case of other GBH (see Paper I), and in the case of AGN's $\alpha_{\rm ox}$."
 We simulate spectra of AGN by scaling the disc temperature and the bolometric luminosity of our collection of the fitting GRO J165540 models., We simulate spectra of AGN by scaling the disc temperature and the bolometric luminosity of our collection of the best-fitting GRO J1655–40 models.
 For à standard. Shakura-Sunvaeyv ecometrically thin optically thick acerction disc. the cise temperature scales. with the black hole mass as DunsoxALfo and the bolometric luminosity is proportional to the black hole mass. Li.xM.," For a standard Shakura-Sunyaev geometrically thin optically thick accretion disc, the disc temperature scales with the black hole mass as $T_{\rm disc} \propto M^{-1/4}$, and the bolometric luminosity is proportional to the black hole mass, $L_{\rm bol} \propto
M$."
 We assume that the accretion [low geometry. and hence the heating to cooling compactness ratio. μέν. varies in the same wav as the Function of the mass accretion rate both in Galactic and supermassive black holes.," We assume that the accretion flow geometry, and hence the heating to cooling compactness ratio, $\ell_h/\ell_s$, varies in the same way as the function of the mass accretion rate both in Galactic and supermassive black holes."
 We set the normalization of the iron line. when present in the baseline spectrum. to zero. since it does not contribute significantly to the total simulated AGN flux.," We set the normalization of the iron line, when present in the baseline spectrum, to zero, since it does not contribute significantly to the total simulated AGN flux."
 Our goal is to check if taking into account the mass distribution among the AGN samples can provide insights to the origins of the observed. dependence between the X-ray Ioudness. 654) and optical/UV Iuminosity at2500A.," Our goal is to check if taking into account the mass distribution among the AGN samples can provide insights to the origins of the observed dependence between the X-ray loudness, $\alpha_{\rm ox}$ ) and optical/UV luminosity at."
. For our study we select. the ACN samples with measured black hole mass., For our study we select the AGN samples with measured black hole mass.
 Black hole mass measurements in ACN are not trivial., Black hole mass measurements in AGN are not trivial.
 However. various methods (c.g. reverberation mapping and correlations with the widths of emission lines) resulted in black hole mass estimates ranging between 10? and LOM Solar masses.," However, various methods (e.g. reverberation mapping and correlations with the widths of emission lines) resulted in black hole mass estimates ranging between $^6$ and $^{10}$ Solar masses."
 Figure 3. shows mass distributions for samples of AGN compiled from the literature.," Figure \ref{fig:masses}
 shows mass distributions for samples of AGN compiled from the literature."
 In Fig., In Fig.
 3aa we show the mass distribution reported for S6 ‘Pype 1 Broad Line ACN by Alerloni ct al. (, \ref{fig:masses}a a we show the mass distribution reported for 86 Type 1 Broad Line AGN by Merloni et al. (
2010) as part of the COSMOS project.,2010) as part of the COSMOS project.
 The masses were consistently calculated based. on the Alell emission line. however the sample spans relatively narrow ranges of redshifts (1< 2.2) ancl luminosities (44.5<logLi. 46.5).," The masses were consistently calculated based on the MgII emission line, however the sample spans relatively narrow ranges of redshifts $1<z<2.2$ ) and luminosities $44.5<\log L_{\rm bol} < 46.5$ )."
 For comparison. in Fig.," For comparison, in Fig."
 3bb we show the results of Shen ct al. (, \ref{fig:masses}b b we show the results of Shen et al. (
2008). for Pype 1 Broad Line ACN in SDSS.,"2008), for Type 1 Broad Line AGN in SDSS."
 This sample contains 760.000 objects with redshifts and luminosities in the Q1<z4.5 and 44.5«logLia48 range. with masses caleulated based on Mell. CIV or Ll. lines.," This sample contains $\sim$ 60.000 objects with redshifts and luminosities in the $0.1<z<4.5$ and $44.5<\log L_{\rm bol} < 48$ range, with masses calculated based on MgII, CIV or $_{\beta}$ lines."
 This is the largest ACGN sample with black hole mass measurements., This is the largest AGN sample with black hole mass measurements.
 Llowever. the authors point out that the estimations based on the CIV line may be biased due to the line being disc wind dominated.," However, the authors point out that the estimations based on the CIV line may be biased due to the line being disc wind dominated."
 In addition to the Tvpe 1 Broad Line AGN. we consider a sample of LINER. galaxies.," In addition to the Type 1 Broad Line AGN, we consider a sample of LINER galaxies."
 Black hole mass measurements are challenging because of the host galaxy domination in he optical band., Black hole mass measurements are challenging because of the host galaxy domination in the optical band.
 We compile a LINERs sample galaxies rom Maoz (2007). Gu Cao (2009) ancl Eracleous et al. (," We compile a LINERs sample galaxies from Maoz (2007), Gu Cao (2009) and Eracleous et al. ("
2010) containing black hole mass estimates for 50 sources.,2010) containing black hole mass estimates for 50 sources.
 The estimated Ecdington ratios for LINERS (e.g. Gu Cao 2009. LExracleous et al.," The estimated Eddington ratios for LINERs (e.g. Gu Cao 2009, Eracleous et al."
 2010) indicate that they accrete at ow rates and thus may be the counterparts of the hare state GBIIs., 2010) indicate that they accrete at low rates and thus may be the counterparts of the hard state GBHs.
 We test this possibility in Sec., We test this possibility in Sec.
 4.2., 4.2.
 Finally. we consider Narrow Line Sevfert. 1. galaxies (NLSI).," Finally, we consider Narrow Line Seyfert 1 galaxies (NLS1)."
 We study a ΛΙ subsample from. Cirupe οἱ al. (, We study a NLS1 subsample from Grupe et al. (
2010) defined as ΗΛ)<2000 kms | (Osterbrock Pogee. 1985: Goodrich 1989).,"2010) defined as $_{\beta}) < 2000$ km $^{-1}$ (Osterbrock Pogge, 1985; Goodrich 1989)."
 These objects have reliable estimates of the accretion rates based on the simultaneous SEDs in optical and X-ray bands., These objects have reliable estimates of the accretion rates based on the simultaneous SEDs in optical and X-ray bands.
 NLS1 are believed to have higher Exlcington ratios (and lighter black holes) than their broad line counterparts., NLS1 are believed to have higher Eddington ratios (and lighter black holes) than their broad line counterparts.
 In Sec., In Sec.
 4.2 we study if they can be considered as counterparts of the soft state CLDIIS., 4.2 we study if they can be considered as counterparts of the soft state GBHs.
 Distribution of AGN masses in all the above samples can be satisfactorily described. by a lognormal functions with means and PWIIAIs listed in Table 1.., Distribution of AGN masses in all the above samples can be satisfactorily described by a lognormal functions with means and FWHMs listed in Table \ref{tab:tab1}. .
 Phus. in this paper we will assume that the masses of black holes. in AGN have a lognormal distribution.," Thus, in this paper we will assume that the masses of black holes in AGN have a lognormal distribution."
 Phe means of the above distributions are dillerent: log(Mig/M.)=6.8 and 8.9 for NLSIs and SDSS samples. respectivelv.," The means of the above distributions are different: $\log ({\rm M}_{\rm BH}/{\rm M}_{\odot}) = 6.8$ and 8.9 for NLS1s and SDSS samples, respectively."
 The ENIMS. vary, The FWHMs vary
arise from the accretion of smaller satellite halos at low redshift.,arise from the accretion of smaller satellite halos at low redshift.
 Since these mergers are nearby. they are easily detectable with LISA.," Since these mergers are nearby, they are easily detectable with LISA."
 Tlhüs is in stark coutrast to the predictions for assembling the massive eud of the SMDIT mass spectrum: ?/— showed that the detectable lnass ratios are equally distributed im the range of 1-LO., This is in stark contrast to the predictions for assembling the massive end of the SMBH mass spectrum; \citet{Sesana:07} showed that the detectable mass ratios are equally distributed in the range of 1-10.
 Iu fact. Figure to indicates that the most commonly observed black hole merecrs m our voluue have mass ratios of 1000 or more (depending on the black hole erowthn prescription).," In fact, Figure \ref{fig:histq} indicates that the most commonly observed black hole mergers in our volume have mass ratios of 1000 or more (depending on the black hole growth prescription)."
 Mergers of more equal mass dark matter halos (and subsequeutly. the coalescence of more equal mass black holes) occured in this ποιό at 2>8. and at this epoch. most of our black holes were too simall to be observed by LISA.," Mergers of more equal mass dark matter halos (and subsequently, the coalescence of more equal mass black holes) occurred in this volume at $z>8$, and at this epoch, most of our black holes were too small to be observed by LISA."
 Figures 1. and 23. bear this out by plotting mapping the number of resolvable sources as a function of mass ratio and redshift for oue black hole erowthn prescription over a LO vear observation., Figures \ref{fig:histmap} and \ref{fig:snrmaprez} bear this out by plotting mapping the number of resolvable sources as a function of mass ratio and redshift for one black hole growth prescription over a 10 year observation.
 The signals from an equal mass iuspiraling svstem aud hose with mass ratios of upwards of ten thousand can ο quite different., The signals from an equal mass inspiraling system and those with mass ratios of upwards of ten thousand can be quite different.
 From an analytical staudpoiut. the difference in signals cau be understood by investigating he post-Newtoniau expausion of the eravitational wave strain. which uses the orbital velocity as an expausion xuanieter.," From an analytical standpoint, the difference in signals can be understood by investigating the post-Newtonian expansion of the gravitational wave strain, which uses the orbital velocity as an expansion parameter."
 Each term iu the expansion results in juinonies of the orbital frequency ?.., Each term in the expansion results in harmonics of the orbital frequency \cite{Blanchet:1996}.
 Starting at the first full order. certain select frequencies are scaled by jj=mnainisngq Lima)? which suppresses those frequencies for svstenis with larec mass ratios.," Starting at the first full order, certain select frequencies are scaled by $\eta \equiv m_{1}m_{2} / (m_{1} +
m_{2})^{2}$ , which suppresses those frequencies for systems with large mass ratios."
 Figure 6 compares the oower spectral deusities for two SAIBIT svsteiis: one with wo 10*\E. SMDBIIS and the other with i4=10°M. and wie=105M, Figure \ref{fig:rubbo1} compares the power spectral densities for two SMBH systems: one with two $10^7~\textrm{M}_{\odot}$ SMBHs and the other with $m_{1}=10^{7}~\textrm{M}_{\odot}$ and $m_{2}=10^{4}~\textrm{M}_{\odot}$.
 These figures demoustrate how the equal mass svsteni sweeps through all frequencies while he large mass ratio system shows mich more structure. with many of the frequencies being suppressed.," These figures demonstrate how the equal mass system sweeps through all frequencies while the large mass ratio system shows much more structure, with many of the frequencies being suppressed."
 The LISA data streams prescut unique challenges or data nuning., The LISA data streams present unique challenges for data mining.
 Unlike electromagnetic observatories. LISA is simultaneously scustitive to eravitational wave sources located all throughout the ska.," Unlike electromagnetic observatories, LISA is simultaneously senstitive to gravitational wave sources located all throughout the sky."
 LISA will also © sensitive to a wide varietv of astroplivsical sources from supermassive black hole binaries. to millions of close white dwiuf binarics within our galaxy. to EMBIS.," LISA will also be sensitive to a wide variety of astrophysical sources – from supermassive black hole binaries, to millions of close white dwarf binaries within our galaxy, to EMRIs."
 The data analysis objective is to conchisively detect a signal aud to extract descriptive paraueter values., The data analysis objective is to conclusively detect a signal and to extract descriptive parameter values.
 Iu xeparatioun for the munieuse data analysis challenge. siauulated data has been produced aud distributed to he LISA community as part of the Mock LISA Data Challenge (MLDC) ?..," In preparation for the immense data analysis challenge, simulated data has been produced and distributed to the LISA community as part of the Mock LISA Data Challenge (MLDC) \cite{Vallisneri:2009}."
 So far the challenges have focused on order unity nass ratios for the simulated supermassive dack hole iuspirals., So far the challenges have focused on order unity mass ratios for the simulated supermassive black hole inspirals.
 While for most analysis techuiques an assumption about the mass ratio is not hardwired mto he routine. the routines have not been tested at the more extreme ratios of 107 as sueeested here.," While for most analysis techniques an assumption about the mass ratio is not hardwired into the routine, the routines have not been tested at the more extreme ratios of $10^{5}$ as suggested here."
 It is possible hat the surpressed hiruonies will cause false-positives or negatives in sole routines because the svstem will be wissecl or misideutified., It is possible that the surpressed harmonics will cause false-positives or negatives in some routines because the system will be missed or misidentified.
 Figure ϐ preseuts the total characteristic strain for 7 bright black hole mergers in our volume. asstuine a 10:1 black hole erowth recipe and a Bovlan-Nolchin dynamical friction treatment.," Figure \ref{fig:tracks} presents the total characteristic strain for 7 bright black hole mergers in our volume, assuming a 10:1 black hole growth recipe and a Boylan-Kolchin dynamical friction treatment."
 The differences iu the two classes of source is clear the brightest sources are high mass ratio mergers which coalesce iu the LISA baud at ow redshift. while the other bright class probes more equalinass ποσο of low mass black holes at high redshift.," The differences in the two classes of source is clear – the brightest sources are high mass ratio mergers which coalesce in the LISA band at low redshift, while the other bright class probes more equal-mass mergers of low mass black holes at high redshift."
 Note that the black holes in the bottoinost rack actually merge outside the LISA band aud will ve considered an iuspi;alb this was the ouly detectable inspiral in our volume., Note that the black holes in the bottom-most track actually merge outside the LISA band and will be considered an inspiral; this was the only detectable inspiral in our volume.
 Iu Table 1.. we determined the redshift at which hese 7 mergers would become undetectable im a 3 wear observation.," In Table \ref{tab:snrdist}, we determined the redshift at which these 7 mergers would become undetectable in a 3 year observation."
 Understaudably. the highly unequal mass nerecrs can only be detected out to a huninosity distance of roughly 3 Copc. which is simular to the distance probed wv extreme mass ratio inspirals (?)..," Understandably, the highly unequal mass mergers can only be detected out to a luminosity distance of roughly 3 Gpc, which is similar to the distance probed by extreme mass ratio inspirals \citep{Gair:04}."
 LISA observations of his class of high mass ratio merecr. then. may be useful o probe the fraction of black hole mass is accreted at ate times for this low mass SMDIT range.," LISA observations of this class of high mass ratio merger, then, may be useful to probe the fraction of black hole mass is accreted at late times for this low mass SMBH range."
 Figure 5. preseuts the total characteristic strain for 16 LISA detector from all the black hole merecrs in our 1000 ? Alpe? volue over a 3-vear observation span., Figure \ref{fig:hsum} presents the total characteristic strain for the LISA detector from all the black hole mergers in our 1000 $^{-3}$ $^3$ volume over a 3-year observation span.
 Each curve has a low frequency rise that drops steeply off before hitting a shallow plateau: the rise aud drop-off is the signature of the acciuulated. high mass ratio morects. While the shallow segment marks the su of the relatively equal mass mergers.," Each curve has a low frequency rise that drops steeply off before hitting a shallow plateau; the rise and drop-off is the signature of the accumulated high mass ratio mergers, while the shallow segment marks the sum of the relatively equal mass mergers."
 We caution that this nuplicitly assumes that all of these merecrs would take place during this 3-vear period. or alternatively that this volume is representative of the Universe.," We caution that this implicitly assumes that all of these mergers would take place during this 3-year period, or alternatively that this volume is representative of the Universe."
 Issues of cosmic variance aside. we can see that the black hole growth prescription strouelv infiuences the total strain iu the LISA baud: this is best seen when ανασα friction is described by ?. where the drop-off changes by a decade in frequency in response to a change in the typical mass of the local SAIBOUs.," Issues of cosmic variance aside, we can see that the black hole growth prescription strongly influences the total strain in the LISA band; this is best seen when dynamical friction is described by \citet{Boylan:08} where the drop-off changes by a decade in frequency in response to a change in the typical mass of the local SMBHs."
 Although this is plagued dy cosmic variance. it is instructive to estimate the uunuboer of merecrs observable by LISA in the Universe expected from asseubliug these lhehtest. SMDBIIs.," Although this is plagued by cosmic variance, it is instructive to estimate the number of mergers observable by LISA in the Universe expected from assembling these lightest SMBHs."
 Table 5 presents the extrapolated Universal LISA inereer rates., Table \ref{tab:first} presents the extrapolated Universal LISA merger rates.
 Omne surprising point is that the LISA imerecr rates depend so strongly ou the adopted form of the dynamical fiction force., One surprising point is that the LISA merger rates depend so strongly on the adopted form of the dynamical friction force.
 All previous LISA inerecr rate estimates have used a seni-analvtic technique that ciplovs Chandrasckhay dynamical friction to meree the dark matter halos and to usher the black holes to the iuner kiloparsec., All previous LISA merger rate estimates have used a semi-analytic technique that employs Chandrasekhar dynamical friction to merge the dark matter halos and to usher the black holes to the inner kiloparsec.
 However. both perturbation theory (7). and uunucrical shuulationus (777) have shown that Chaudrasekhar dynamical friction approximates the imnerger time to within a factor of two. at best.," However, both perturbation theory \citep{colpi:99} and numerical simulations \citep{weinberg:89,KHB:1999:sat,Boylan:08}
 have shown that Chandrasekhar dynamical friction approximates the merger time to within a factor of two, at best."
 When we enmplov a dvuamical friction foriualiu that is based ou fits to merecr timescales from nunierical simulations. we fud that the black hole merecrs are delaved to lower redshift.," When we employ a dynamical friction formalism that is based on fits to merger timescales from numerical simulations, we find that the black hole mergers are delayed to lower redshift."
 Naively. this would simply make every merger louder.," Naively, this would simply make every merger louder."
 ILowever. in our SMDII growth prescription. there is au additional effect: since the incoming SAIBIT drives eas inflow to the primary SMDII over a longer timespan. resulting SAIBID ultimately grows more massive than with Chandrasekhar dvuamical friction.," However, in our SMBH growth prescription, there is an additional effect: since the incoming SMBH drives gas inflow to the primary SMBH over a longer timespan, resulting SMBH ultimately grows more massive than with Chandrasekhar dynamical friction."
 This can place the SAIBUs in our volue just out of LISAS “sweet spot’ which will decrease the signal-to-noise ratio of the more massive black holes., This can place the SMBHs in our volume just out of LISA's 'sweet spot' which will decrease the signal-to-noise ratio of the more massive black holes.
 Paradoxically. then. the LISA event rate for more correct iierger times drops by as mich as an order of mmaenitude.," Paradoxically, then, the LISA event rate for more correct merger times drops by as much as an order of magnitude."
 For our most realistic mocel to, For our most realistic model to
" Finally, using Eq.(29)). we obtain the following relations. ors DC EEESNES OO PCI Tutroducing⋜↧↸⊳↑∪↥↴∖∐↕↑∐↸∖↴⋝⋜↧↕⋜⋯↸⊳↸∖⋜⋯∖↑↕∐∖∶↴∙↥⋜∥∐↸∖∐↑∪↕⋜⋯∶↴∙∏↕⋜∐↖↸∖↕∪↸⊳≓ the cimmensiouless Rossby muuber. Ro= the necessary conditions⋅⋅ for Diustabilitya: 050 Ro 5 Ro, ∶≺↕⊰∪↙↭⋅≺∢∔∶≩≻ Ó oso Ro = Ro.",", Finally, using \ref{vort}) ), we obtain the following relations, _r > 0: r _r - 2, _r < 0: r _r. Introducing the dimensionless Rossby number, $Ro = r \partial_r \Omega / 2 \Omega$ , andusing the property $C^{n} = C_r^{n} + C_\phi^{n}$ we obtain the necessary conditions for instability: _r > 0: Ro - = Ro_c ( Ro_c > 0, _r < 0: Ro - = Ro_c ( Ro_c < 0."
" The quantity Ro, can be seen as constant by asstuning that the ratio C/C"" is independent from he mean flow.", The quantity $Ro_c$ can be seen as constant by assuming that the ratio $C_r^{n}/C^{n}$ is independent from the mean flow.
 This ratio quautifv the redistribution of] the euergy extracted for] the mean flow between the velocity components. bythe uou-linearities.," This ratio quantify the redistribution of the energy extracted for the mean flow between the velocity components, bythe non-linearities."
"Note that he cocficients Ch. COC! aud C"" themsclf most likely depend ou the mean flow (see ?)).","Note that the coefficients $C_r^{n}$, $C_\phi^{n}$ $C_z^{n}$ and $C^{n}$ themself most likely depend on the mean flow (see \citet{speziale}) )."
 ? already noted that the Rossby uuuber should approach a constant value when a linearly. stable rotating shear flow becomes turbulent. by arguing that flaw.," \citet{Townsend} already noted that the Rossby number should approach a constant value when a linearly stable rotating shear flow becomes turbulent, by arguing that ""."
" Following. the same path as in. the inviscid∙∙∙ case. we now add the viscous coustraints to equations (25))-(28)): = DOR. CTLED vC!XU! FIL - d> παuulaC)CULE = ""E↽∕∖ MEL vO↽∕∖⇉≺⊔⋝ 9,"," Following the same path as in the inviscid case, we now add the viscous constraints to equations \ref{ur4}) \ref{k4}) ): = 2 ^l u^2 + - = - ^l u^2+ ^n } - = _z - _t k = - r _r u^2 + }.-"
magnitude intervals down to the completeness level of the data (that is lower than the one adopted).,magnitude intervals down to the completeness level of the data (that is lower than the one adopted).
 Having for field 2 a magnitude limit of 22.2 (see Table 1)) and field 6 à limit of 2].5. we also verified that for the latter our star counts are in correct proportion below the completeness level of505.," Having for field 2 a magnitude limit of 22.2 (see Table \ref{M55tab2}) ) and field 6 a limit of 21.5, we also verified that for the latter our star counts are in correct proportion below the completeness level of."
" In a second test. we generated two LFs dividing the whole cluster in two octants (dividing along the 15"" line that runs from the center of the cluster till the field 19 FFigure 1))."," In a second test, we generated two LFs dividing the whole cluster in two octants (dividing along the $45^\circ$ line that runs from the center of the cluster till the field 19 Figure \ref{M55frames}) )."
 For each of the two slices we generated three LFs in the same radial range as in Figure 4.., For each of the two slices we generated three LFs in the same radial range as in Figure \ref{M55fdl}.
 After comparing all of them we did not find any significant difference., After comparing all of them we did not find any significant difference.
 Therefore the differences among the three LFs in Figure + must be real.," Therefore the differences among the three LFs in Figure \ref{M55fdl}
 must be real."
 Another source of eror in the LF construction Is represented by the LF of the field stars., Another source of error in the LF construction is represented by the LF of the field stars.
 As will be shown in Section 5.. M55 has a halo of probably unbound cluster stars.," As will be shown in Section \ref{M55rprof}, M55 has a halo of probably unbound cluster stars."
 The field star LF constructed from the star counts just outside the cluster can be affected by some contamination of the cluster halo., The field star LF constructed from the star counts just outside the cluster can be affected by some contamination of the cluster halo.
 The consequence is that we might over-subtract stars when subtracting the field LF from the cluster LF. modifying in this way the slope of the mass function (the more affected magnitudes are the faintest ones).," The consequence is that we might over-subtract stars when subtracting the field LF from the cluster LF, modifying in this way the slope of the mass function (the more affected magnitudes are the faintest ones)."
 To test this possibility. we have extracted background LFs in two different anulit outside the cluster (in terms of +. 1.0—rx1.3 and r> 1.3).," To test this possibility, we have extracted background LFs in two different anulii outside the cluster (in terms of $r_t$, $1.0<r\le1.3$ and $r>1.3$ )."
 Comparing the two background/foregroounnd LFs we founc that the number of stars probably belonging to the cluster but outside the tidal radius must be less than ~25% of the adopted field stars in the worst case (the faintest bins)., Comparing the two nd LFs we found that the number of stars probably belonging to the cluster but outside the tidal radius must be less than $\sim25\%$ of the adopted field stars in the worst case (the faintest bins).
 The possible over-subtraction is not a problem for the inner and intermediate LFs. where the number of field stars (after rescaling for the covered area) is always less than ~3% of the stars counted in each magnitude bin.," The possible over-subtraction is not a problem for the inner and intermediate LFs, where the number of field stars (after rescaling for the covered area) is always less than $\sim3\%$ of the stars counted in each magnitude bin."
 For the outer LF. the total contribution of the measured field stars is larger. but it is still less than 25'4 of the cluster stars (the worst case applies to the faintest magnitude bin): this means that the possible M55 halo star over-subtraction in the field-corrected LF is always less than GC (25%« 25%). negligible for our purposes.," For the outer LF, the total contribution of the measured field stars is larger, but it is still less than $25\%$ of the cluster stars (the worst case applies to the faintest magnitude bin): this means that the possible M55 halo star over-subtraction in the field-corrected LF is always less than $6\%$ $25\% \times 25\%$ ), negligible for our purposes."
 In order to build a mass function for the stars of M55. we needed to adopt a distance modulus and an extinction coefficient.," In order to build a mass function for the stars of M55, we needed to adopt a distance modulus and an extinction coefficient."
 Shadeetal.(1988) give GiM)=11410. E(B|V)=0.11x0.02. while. more recently. Mandushevetal.(1996) give GyM)=13.90007. E(BVW)=0.11+ 0.02.," \cite{Shade88} give $(m-M)_V=14.10$, $(B-V)=0.14\pm0.02$, while, more recently, \cite{Mand96} give $(m-M)_V=13.90\pm0.07$, $(B-V)=0.14\pm0.02$ ."
 In the absence of an independent measure made by us. we adopted the values published by Mandushevetal.(1996) because they are based on the application. with updated data. of the subdwarfs fitting method.," In the absence of an independent measure made by us, we adopted the values published by \cite{Mand96} because they are based on the application, with updated data, of the subdwarfs fitting method."
 Using the LFs of the previous Section we build the corresponding mass functions using the mass-luminosity relation tabulated by VandenBerg&Bell(1985) for an isochrone of Z=3«10.! and an age of 16 Gyr Aleainoetal.(1992).," Using the LFs of the previous Section we build the corresponding mass functions using the mass-luminosity relation tabulated by \cite{VDB85}
 for an isochrone of $Z=3\times10^{-4}$ and an age of 16 Gyr \cite{Alcaino92}."
. The MFs for the three radial intervals are presented in Figure 4..., The MFs for the three radial intervals are presented in Figure \ref{M55mf}.
 The MFs are vertically shifted in order to make their comparison more clear., The MFs are vertically shifted in order to make their comparison more clear.
 The MFs are significantly different: the slopes of the MFs increase moving outwards as expected from the effects of the mass segregation and from the LFs of Figure 4.., The MFs are significantly different: the slopes of the MFs increase moving outwards as expected from the effects of the mass segregation and from the LFs of Figure \ref{M55fdl}.
" Figure 4 clearly shows that the MF starting from the center out to the outer envelope of the cluster is flat: the index .c of the power law. £=gin(1107, best fitting the data aree=2140.1. oS0,8d:0.3. and ec=0.73:0.1 going from the inner to the outer anulii: this means that the slope of the global MF (of all the stars in M55) should be extremely flat."," Figure \ref{M55mf} clearly shows that the MF starting from the center out to the outer envelope of the cluster is flat: the index $x$ of the power law, $\xi=\xi_0m^{-(1+x)}$, best fitting the data are: $x=-2.1\pm0.4$, $x=-0.8\pm0.3$, and $x=0.7\pm0.4$ going from the inner to the outer anulii; this means that the slope of the global MF (of all the stars in M55) should be extremely flat."
 Indeed. the slope of the global mass function obtained from the corresponding LF of all the stars of MSS is: c1.030.1 This result agrees with the results of Irwin&Trimble (1984). while the results," Indeed, the slope of the global mass function obtained from the corresponding LF of all the stars of M55 is: $x=-1.0\pm0.4$ This result agrees with the results of \cite{Irwin84}, , while the results"
development of nonlinear processes.,development of nonlinear processes.
 Thus. a correct understanding of the energy exchange channels between different modes in the inear regime is vital for a correct understanding of the nonlinear shenomena.," Thus, a correct understanding of the energy exchange channels between different modes in the linear regime is vital for a correct understanding of the nonlinear phenomena."
 Indications of the shear induced mode conversion can be ound in a number of previous studies., Indications of the shear induced mode conversion can be found in a number of previous studies.
 Barranco and Marcus 2005 report that vortices are able to excite inertial gravity waves during 3D spectral simulations., Barranco and Marcus 2005 report that vortices are able to excite inertial gravity waves during 3D spectral simulations.
 Brandenburg and Dintrans (2006) qaave studied the linear dynamics of perturbation SFH to analyze nonaxisymetric stability in the shearing sheet approximation., Brandenburg and Dintrans (2006) have studied the linear dynamics of perturbation SFH to analyze nonaxisymetric stability in the shearing sheet approximation.
 Temporal evolution of the perturbation gain factors reveal a wave nature after the radial wavenumber changes sign., Temporal evolution of the perturbation gain factors reveal a wave nature after the radial wavenumber changes sign.
 Compressible Waves are present. along with vortical perturbations. in the simulation by Johnson Gammie (2005b) but their origin is not particularly discussed.," Compressible waves are present, along with vortical perturbations, in the simulation by Johnson Gammie (2005b) but their origin is not particularly discussed."
 In parallel. there are a number of papers that focus on the investigation of the shear induced mode coupling phenomena.," In parallel, there are a number of papers that focus on the investigation of the shear induced mode coupling phenomena."
 The study of the linear coupling of modes in Keplerian flows has been conducted in the local shearing sheet approximation (Tevzadze et al., The study of the linear coupling of modes in Keplerian flows has been conducted in the local shearing sheet approximation (Tevzadze et al.
 2003.2008) as well as in 2D global numerical simulations (Bodo et al.," 2003,2008) as well as in 2D global numerical simulations (Bodo et al."
 2005. hereafter BOS).," 2005, hereafter B05)."
 Tevzadze et al. (, Tevzadze et al. (
2003) studied the linear dynamics of three-dimensional small scale perturbations (with characteristic scales much less then the dise thickness) in vertically (stably) stratified Keplerian dises.,2003) studied the linear dynamics of three-dimensional small scale perturbations (with characteristic scales much less then the disc thickness) in vertically (stably) stratified Keplerian discs.
 They show. that vortex and internal gravity wave modes are coupled efficiently.," They show, that vortex and internal gravity wave modes are coupled efficiently."
 BOS performed global numerical simulations of the linear dynamics of initially imposed two-dimensional pure vortex mode perturbations in compressible Keplerian dises with constant background pressure and density., B05 performed global numerical simulations of the linear dynamics of initially imposed two-dimensional pure vortex mode perturbations in compressible Keplerian discs with constant background pressure and density.
 The two modes possible in this system are effectively coupled: vortex mode perturbations are able to generate density-spiral waves., The two modes possible in this system are effectively coupled: vortex mode perturbations are able to generate density-spiral waves.
 The coupling is. however. strongly. asymmetric: the coupling is effective for wave generation by vortices. bu not viceversa.," The coupling is, however, strongly asymmetric: the coupling is effective for wave generation by vortices, but not viceversa."
 The resulting dynamical picture points out the importance of mode coupling and the necessity of considering compressibility effects for processes with characteristic scales of the order or larger than the dise thickness., The resulting dynamical picture points out the importance of mode coupling and the necessity of considering compressibility effects for processes with characteristic scales of the order or larger than the disc thickness.
 Bodo et al. (, Bodo et al. (
2007) extended this work to nonlinear amplitudes and found that mode coupling is an efficient channel for energy exchange and is no an artifact of the linear analysis.,2007) extended this work to nonlinear amplitudes and found that mode coupling is an efficient channel for energy exchange and is not an artifact of the linear analysis.
 BOS is particularly relevan to the present study. since it studies the dynamics of mode coupling in 2D unstratified flows and is a good starting point for a further extension to radially stratified flows.," B05 is particularly relevant to the present study, since it studies the dynamics of mode coupling in 2D unstratified flows and is a good starting point for a further extension to radially stratified flows."
 Later. Heinemann Papaloizou (20092) derived WKBJ solutions of the generated waves and performed numerical simulations of the wave excitation by turbulent fluctuations (Heinemann Papaloizou 2009b).," Later, Heinemann Papaloizou (2009a) derived WKBJ solutions of the generated waves and performed numerical simulations of the wave excitation by turbulent fluctuations (Heinemann Papaloizou 2009b)."
 In the present paper we study the linear dynamies of perturbations and analyze shear flow induced mode coupling in the local shearing sheet approximation., In the present paper we study the linear dynamics of perturbations and analyze shear flow induced mode coupling in the local shearing sheet approximation.
 We investigate the properties of mode coupling using qualitative analysis within the three-mode approximation., We investigate the properties of mode coupling using qualitative analysis within the three-mode approximation.
 Within this approximation we tentatively distinguish vorticity. entropy and pressure modes.," Within this approximation we tentatively distinguish vorticity, entropy and pressure modes."
 Quantitative results on mode conversion are derived numerically., Quantitative results on mode conversion are derived numerically.
 It seems that a weak radial stratifications. while being a weak factor for the dise stability. still provides an additional degree of freedom tan active entropy mode). opening new options for velocity shear induced mode conversion. that may be important for the system behavior.," It seems that a weak radial stratifications, while being a weak factor for the disc stability, still provides an additional degree of freedom (an active entropy mode), opening new options for velocity shear induced mode conversion, that may be important for the system behavior."
 One of the direct result of mode conversion is the possibility of linear generation of the vortex mode (i.e. potential vorticity) by compressible perturbations.," One of the direct result of mode conversion is the possibility of linear generation of the vortex mode (i.e., potential vorticity) by compressible perturbations."
 We want to stress the possibility of the coupling between high and low frequency perturbations. considering that high frequency oscillations have been often neglected during previous investigations in particular for protoplanetary dises.," We want to stress the possibility of the coupling between high and low frequency perturbations, considering that high frequency oscillations have been often neglected during previous investigations in particular for protoplanetary discs."
 Conventionally there are two distinct viewpoints commonly employed during the investigation of hydrodynamic astrophysical discs., Conventionally there are two distinct viewpoints commonly employed during the investigation of hydrodynamic astrophysical discs.
 In one case (self gravitating galactic dises) the emphasis is placed on the investigation of the dynamics of spiral-density waves and vortices. although normally present in numerical simulations. are thought to play a minor role in the overall dynamics.," In one case (self gravitating galactic discs) the emphasis is placed on the investigation of the dynamics of spiral-density waves and vortices, although normally present in numerical simulations, are thought to play a minor role in the overall dynamics."
 In the other case (non-self-gravitating hydrodynamic discs) the focus is on the potential vorticity perturbations and density-spiral waves are often thought to play a minor role., In the other case (non-self-gravitating hydrodynamic discs) the focus is on the potential vorticity perturbations and density-spiral waves are often thought to play a minor role.
 Here. discussing the possible (multi) mode couplings. we want to draw attention to the possible flaws of these simplified views (see e.g. Mamatsashvili Chagelishvili 20073.," Here, discussing the possible (multi) mode couplings, we want to draw attention to the possible flaws of these simplified views (see e.g. Mamatsashvili Chagelishvili 2007)."
 In many cases. mode coupling makes different perturbation to equally participate in the dynamical processes despite of a significant difference in their temporal scales.," In many cases, mode coupling makes different perturbation to equally participate in the dynamical processes despite of a significant difference in their temporal scales."
 In the next section we present mathematical formalism of our study., In the next section we present mathematical formalism of our study.
 We describe three mode formalism and give schematic picture of the linear mode coupling in the radially sheared and stratified. flow., We describe three mode formalism and give schematic picture of the linear mode coupling in the radially sheared and stratified flow.
 Numerical analysis of the mode coupling is presented in Sec., Numerical analysis of the mode coupling is presented in Sec.
 3., 3.
 We evaluate mode coupling efficiencies at different radial stratification scales of the equilibrium pressure and entropy., We evaluate mode coupling efficiencies at different radial stratification scales of the equilibrium pressure and entropy.
 The paper is summarized in Sec., The paper is summarized in Sec.
 4., 4.
" The governing ideal hydrodynamic equations of a two-dimensional. compressible disc flows in. polar coordinates are: where V, and Ἐν, are the flow radial and azimuthal velocities respectively. P(r)."," The governing ideal hydrodynamic equations of a two-dimensional, compressible disc flows in polar coordinates are: where $V_r$ and $V_\phi$ are the flow radial and azimuthal velocities respectively. $P(r,\phi)$,"
 X(r.0) and > are respectively the pressure. the surface density and the adiabatic index.," $\Sigma(r,\phi)$ and $\gamma~$ are respectively the pressure, the surface density and the adiabatic index."
" c, is the gravitational potential of the central mass. in the absence of self-gravitation (ey~ απ)."," $\psi_g$ is the gravitational potential of the central mass, in the absence of self-gravitation $~(\psi_g \sim -{1 / r})$ ."
" This potential determines the Keplerian angular velocity: We consider an axisymmetric (0/00=0). azimuthal (1).=0) and differentially rotating basic flow: V5,=(1η."," This potential determines the Keplerian angular velocity: We consider an axisymmetric $(\partial / \partial \phi \equiv 0),~$ azimuthal $(\bar {V}_{r} = 0)~$ and differentially rotating basic flow: $\bar {V}_{\phi}= \Omega(r)r$."
" In the 2D radially stratified equilibrium (see Klahr 2004). all variables are assumed to follow a simple power law behavior: where overbars denote equilibrium and XM, and £%) are the values of the equilibrium surface density and pressure at some fiducial radius r= ro."," In the 2D radially stratified equilibrium (see Klahr 2004), all variables are assumed to follow a simple power law behavior: where overbars denote equilibrium and $\Sigma_0$ and $P_0$ are the values of the equilibrium surface density and pressure at some fiducial radius $r = r_0$ ."
 The entropy can be calculatedas:, The entropy can be calculatedas:
2.,2.
 We [ind that the total intensity pc-scale images reveal à core-jet morphology. similar to that observed in LBLs.," We find that the total intensity pc-scale images reveal a core-jet morphology, similar to that observed in LBLs."
 3., 3.
 The VLBP observations of HDLs and LBLs reveal dillerences in their core fractional polarization. similar to (he trend observed on kpe-scales — the IIDLs have a lower core fractional polarization compared to the LDLs.," The VLBP observations of HBLs and LBLs reveal differences in their core fractional polarization, similar to the trend observed on kpc-scales – the HBLs have a lower core fractional polarization compared to the LBLs."
 High degrees of core polarization were observed for the two RGB sources 07494-540 (c1095) and 09254-504 (c1454). but both of these are. in fact. LDLEs.," High degrees of core polarization were observed for the two RGB sources 0749+540 $\simeq 10\%$ ) and 0925+504 $\simeq 14\%$ ), but both of these are, in fact, LBLs."
 4., 4.
 The jet fractional polarizations in WBLs do not diller svstematically [rom (hose in LDLs., The jet fractional polarizations in HBLs do not differ systematically from those in LBLs.
 In both cases. jet degrees of polarization of (ens of per cent can sometimes be observed. indicating the presence of well-ordered. D fields.," In both cases, jet degrees of polarization of tens of per cent can sometimes be observed, indicating the presence of well-ordered $B$ fields."
 5., 5.
 The relative VLBI core polarization angles do not show any systematic differences between IHIDLs aud LDLs., The relative VLBI core polarization angles do not show any systematic differences between HBLs and LBLs.
 6., 6.
 The most intriguing difference between the IIBLs and the LBLs is the orientation of the predominant οἱ [ields in their pe-seale jets the D fields in HBL jets tend to be aligned wilh (he local jet direction. while the 5 fields of LBLs tend to be perpendicular to the local jet direction.," The most intriguing difference between the HBLs and the LBLs is the orientation of the predominant $B$ fields in their pc-scale jets – the $B$ fields in HBL jets tend to be aligned with the local jet direction, while the $B$ fields of LBLs tend to be perpendicular to the local jet direction."
 This result needs (o re-examined wilh a larger sample of IIBLs., This result needs to re-examined with a larger sample of HBLs.
 7., 7.
 Some of the WBLs display evidence for jet magnetic field structures wilh transverse 1 lied in the inner region of the jet aud longitudinal 2 field at the edges. i.e. a spine-sheatli structure;," Some of the HBLs display evidence for jet magnetic field structures with transverse $B$ field in the inner region of the jet and longitudinal $B$ field at the edges, i.e, a `spine-sheath' structure."
 Although such D-field structures have been taken to indicate a fast spine+slow sheath velocity structure. (μον can also come about in a natural. simpler wav if the jet has a helical B field (without anv requirement for a (wo-Iuver velocity structure).," Although such $B$ -field structures have been taken to indicate a fast spine+slow sheath velocity structure, they can also come about in a natural, simpler way if the jet has a helical $B$ field (without any requirement for a two-layer velocity structure)."
 S., 8.
 We lind tentative evidence that the observed. component speeds in IIDLs are lower than the (vpical apparent speeds observed in LBLs (71— 5e) and FSRQs (25— 10e)., We find tentative evidence that the observed component speeds in HBLs are lower than the typical apparent speeds observed in LBLs $\simeq 1-5c$ ) and FSRQs $\simeq 5-10c$ ).
 This result suggests that either (he LBL jets have lower Lorentz [actors (han LBLs and FSIQs. or that their jets are oriented al larger angles to the line of sight.," This result suggests that either the HBL jets have lower Lorentz factors than LBLs and FSRQs, or that their jets are oriented at larger angles to the line of sight."
 Since the collected properties of HDBLs provide no clear evidence that their jets lie at systematically larger aneles to the line of sight than do LBL jets. this suggests that the outflow speeds are intrinsically lower in HDLs. than in LBLs.," Since the collected properties of HBLs provide no clear evidence that their jets lie at systematically larger angles to the line of sight than do LBL jets, this suggests that the outflow speeds are intrinsically lower in HBLs, than in LBLs."
 9., 9.
 One wav (o understand our collected results in the context of other previous results in the literature is if HDLs. LBLs and FSRQs form a sequence of increasing average jet Lorentz [actor — the same order as for the sequence of the svuchrotron peaks of these objects. which proceed from highest peak [recuency for HEDBEs to lowest peak frequency for FSRQs. although the physical connection between these (wo sequences is not clear.," One way to understand our collected results in the context of other previous results in the literature is if HBLs, LBLs and FSRQs form a sequence of increasing average jet Lorentz factor — the same order as for the sequence of the synchrotron peaks of these objects, which proceed from highest peak frequency for HBLs to lowest peak frequency for FSRQs, although the physical connection between these two sequences is not clear."
 201. (Pakull&Mirioui2002:Ikaaretetal.2001:Lane2007)..," $20 M_{\odot}$ \citep{Colbert99,Makishima00,Kaaret01} \citep{Pakull02,Roberts03,Kaaret04,Lang07}. \citep{Blair01}."
 Te κοςeresP (Schlegel1991).," $\alpha$ $2 \times 10^{38} \rm \, erg \, s^{-1}$ \citep{Schlegel94}."
 Colbert.(2003). cinission., \citet{Roberts03} emission.
 Ilowever. several of the lines µη. Πα. and [N uf) contain a narrow componecut. eenerally|O associated with plotoionization iustead of shocks.," However, several of the lines ], $\alpha$, and [N ]) contain a narrow component, generally associated with photoionization instead of shocks."
 Blair.Fesen.&Schlegel(2001) concluded that unusually strong photoionization is needed. in addition to shocks. o explain the spectrum. but did not determine the source of ioniziug photons.," \citet{Blair01} concluded that unusually strong photoionization is needed, in addition to shocks, to explain the spectrum, but did not determine the source of ionizing photons."
 They also reported detection of ine cuuission. but did not discuss the origin of the line.," They also reported detection of line emission, but did not discuss the origin of the line."
 Aboluasovctal.(2008). studied the optical ciission roni AIF 16 with particular attention to the A1686 ine., \citet{Abolmasov08} studied the optical emission from MF 16 with particular attention to the $\lambda 4686$ line.
 Their modeling shows that most of the power iu he optical line ciission comes from photoionization. although a component frou shocks is still importaut.," Their modeling shows that most of the power in the optical line emission comes from photoionization, although a component from shocks is still important."
 Thev find that an ultraviolet source more Ihuuiuous than he observed ULX is needed to power the emission aud hat the UV. source has a blackbody temperature of 107157919 N and a luminosity 1.2«10Meres|.," They find that an ultraviolet source more luminous than the observed ULX is needed to power the emission and that the UV source has a blackbody temperature of $10^{5.15 \pm 0.10}$ K and a luminosity $1.2 \times 10^{40} \rm \, erg \, s^{-1}$."
 The ΠΕ is a lower bound on the source hnunuimositv if he photoionized nebula is density bouuded., The luminosity is a lower bound on the source luminosity if the photoionized nebula is density bounded.
 The UV source proposed by Abohluasovetal.(2008) should be bright in the ΕΙΝ., The UV source proposed by \citet{Abolmasov08} should be bright in the FUV.
 We obtained ΕΝ observations using the solar blind chaunuel (SBC) of the Advanced Camera for Surveys (ACS) on the Iubble Space Telescope (IST) to search for the predicted FUV chussion., We obtained FUV observations using the solar blind channel (SBC) of the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST) to search for the predicted FUV emission.
 We describe the observatious and analysis iu &2., We describe the observations and analysis in 2.
 We discuss the implications in 82., We discuss the implications in 3.
 Observations were performed ou 2008 May 01 using the ACS/SBC on IIST ceutered on the position of the N-rav source in ME 16 (GO program 11200. PI: KIiuuct).," Observations were performed on 2008 May 01 using the ACS/SBC on HST centered on the position of the X-ray source in MF 16 (GO program 11200, PI: Kaaret)."
 The spectral clement used was the FLLOLP long-pass filter which covers the bbaud., The spectral element used was the F140LP long-pass filter which covers the band.
 This filter was chosen because ithas amich sanaller backeround sky rate.blocking out ecocoronal Lyman à emission. than the filters that exteud to shorter," This filter was chosen because ithas amuch smaller background sky rate,blocking out geocoronal Lyman $\alpha$ emission, than the filters that extend to shorter"
The equations (10)) and. (11)) relate the bispectrum to the distribution of the plase sum On+Op-gp.,"The equations \ref{eq10}) ) and \ref{eq11}) ) relate the bispectrum to the distribution of the phase sum $\theta_{\sbfm{k}} +
\theta_{\sbfm{k}'} - \theta_{\sbfm{k} + \sbfm{k}'}$."
 To see if this kind of phase information is practically robust. we numerically examine simple examples of nou-Gaussian fields.," To see if this kind of phase information is practically robust, we numerically examine simple examples of non-Gaussian fields."
 Iustead of examining cosmological sinuulatious. the followingOm simple example is enoughfae) to compare the numerical phase clistributious and theoretical predictions.," Instead of examining cosmological simulations, the following simple example is enough to compare the numerical phase distributions and theoretical predictions."
 Series of non-Ciaussian fields are simply generated by exponeutial mapping of a random Gaussian field: where © is a random Gaussiau field with zero mean. unit variauce. aud g is the uou-Ciaussiau parameter.," Series of non-Gaussian fields are simply generated by exponential mapping of a random Gaussian field: where $\phi$ is a random Gaussian field with zero mean, unit variance, and $g$ is the non-Gaussian parameter."
 We simply take a [lat power spectrum for the CaISslan field o., We simply take a flat power spectrum for the Gaussian field $\phi$.
 The field f has zero mean aud variance. (/7jbpd/=exp(g)—1. and is called the logiormal field (€‘oles&Jones1991).," The field $f$ has zero mean and variance $\langle f^2 \rangle = \exp(g^2) - 1$, and is called the lognormal field \citep{col91}."
. This field has quite similar statistical properties to gravitation:dly evolved uor-Caussian fields aud approximately follows the hierarchical model of ligher-order corelations., This field has quite similar statistical properties to gravitationally evolved non-Gaussian fields and approximately follows the hierarchical model of higher-order correlations.
 The parameter g controls the noun-Craussiauity. aud the raudom Caussian fiekl is recove'ed by taking he limit g+0.," The parameter $g$ controls the non-Gaussianity, and the random Gaussian field is recovered by taking the limit $g\rightarrow
0$."
 The random field f is generated on 61? erids iu a rectangular box with the periodic‘boundary coucition., The random field $f$ is generated on $64^3$ grids in a rectangular box with the periodic boundary condition.
 In Fig. l.," In Fig. \ref{fig1},"
" the distribution of the phase sum Op,+On,—Üp,Ro is slotted for a binned configuration of the wavevectors. Ay|=[0.1.0.5]. οι=[0.5.0.6]. 045=[50°ορ as an example. where 015 is the augle between Ay aud. Ae. aud the maguituces of the waveunber are in units of the Nyquist wavenumber."," the distribution of the phase sum $\theta_{\sbfm{k}_1} + \theta_{\sbfm{k}_2} - \theta_{\sbfm{k}_1 +
\sbfm{k}_2}$ is plotted for a binned configuration of the wavevectors, $|\bfm{k}_1| = [0.4,0.5]$, $|\bfm{k}_2| = [0.5,0.6]$, $\theta_{12} =
[50^\circ,60^\circ]$, as an example, where $\theta_{12}$ is the angle between $\bfm{k}_1$ and $\bfm{k}_2$, and the magnitudes of the wavenumber are in units of the Nyquist wavenumber."
 The phase stun is averaged over the wavevectors in a coufiguration bin., The phase sum is averaged over the wavevectors in a configuration bin.
 The points represent the distributions of the phase sum in each realization., The points represent the distributions of the phase sum in each realization.
 Poisson errorbars are sinaller (hau the size of the ;»oiuts., Poisson errorbars are smaller than the size of the points.
 The normalized bispectra p!(Rey.ee) are numerically evaluated from each realization. which are used to draw the theoretical curves in the first-order approximation of equation (11)).," The normalized bispectra $p^{(3)}(\bfm{k}_1,\bfm{k}_2)$ are numerically evaluated from each realization, which are used to draw the theoretical curves in the first-order approximation of equation \ref{eq11}) )."
 There i:J. 100 any fitting parameter at all., There is not any fitting parameter at all.
 The agreement is remarkable in weakly non-Craussian fields., The agreement is remarkable in weakly non-Gaussian fields.
" When he nou-Gaussianity becomes high. the cata points deviate from the fi‘st-order approximation. and the distribution of the phase sum is sharply peaked at Op,+0,—ko=Omod23."," When the non-Gaussianity becomes high, the data points deviate from the first-order approximation, and the distribution of the phase sum is sharply peaked at $\theta_{\sbfm{k}_1} + \theta_{\sbfm{k}_2} - \theta_{\sbfm{k}_1 +
\sbfm{k}_2} = 0\ {\rm mod}\ 2\pi$."
 Up ogc3.0. or ([2)?>~ 100. the distribution of the phase sum is accurately described by the irst-order approximation. and is determined only by the nornalized. bispectrum.," Up to $g \sim 3.0$, or $\langle
f^2\rangle^{1/2} \simeq 100$ , the distribution of the phase sum is accurately described by the first-order approximation, and is determined only by the normalized bispectrum."
 Even though he non-Ctaussiaunity g3.0 on scales of the Nyquist waveuunber is ονομά the perturbative 'eeime. the uormalized bispectrum on scales of the presently tested coufiguratiou is still within the perturbative regime. p!~0.25.," Even though the non-Gaussianity $g \sim 3.0$ on scales of the Nyquist wavenumber is beyond the perturbative regime, the normalized bispectrum on scales of the presently tested configuration is still within the perturbative regime, $p^{(3)} \sim 0.25$."
 This means that the phase sum is well approximated by first-order orumula of the present work even when the field is strougly nouliuear in dyuamies. as long as the yarameter P(k)/V on the relevant scales is small.," This means that the phase sum is well approximated by first-order formula of the present work even when the field is strongly nonlinear in dynamics, as long as the parameter $P(k)/V$ on the relevant scales is small."
 Increasing the power ou relevant scales and/or decreasing the volume drive the phase correlation large. due to the [act that the phase correlations are particularlydepeucdent ou significant features in tlie sample.," Increasing the power on relevant scales and/or decreasing the volume drive the phase correlation large, due to the fact that the phase correlations are particularlydependent on significant features in the sample."
[Inxes are probably only good to about due to unknown slit losses.,fluxes are probably only good to about due to unknown slit losses.
" The coacded Tenagra images were blinked against each other to detect changing point sources,", The coadded Tenagra images were blinked against each other to detect changing point sources.
 Nova candidates were required to be observed in at least (wo epochsand to be missing on an epoch of sullicient depth coverage to confirm i(s transient nature., Nova candidates were required to be observed in at least two epochs to be missing on an epoch of sufficient depth coverage to confirm its transient nature.
 We also used the raw images lor an epoch to confirm (he presence of the candidate in each individual frame., We also used the raw images for an epoch to confirm the presence of the candidate in each individual frame.
 As a further verification. we checked the images from the archives listed in Table 1.. and images from the Digitized SkvSurvev?.," As a further verification, we checked the images from the archives listed in Table \ref{tab_obs}, and images from the Digitized Sky."
. For the regions of our target galaxies near the nuclei. where the intensity gradient makes detection more difficult. we used Che spatial filtering techiique described in (2004).. allowing5 us to detect novae to within LO” of the nuclei of M32 and NGC 205 and to within 2 of the centers of NGC 147 and NGC 185.," For the regions of our target galaxies near the nuclei, where the intensity gradient makes detection more difficult, we used the spatial filtering technique described in \citet{nei04}, allowing us to detect novae to within 10"" of the nuclei of M32 and NGC 205 and to within 2"" of the centers of NGC 147 and NGC 185."
 This subtraction technique was performed on each coadded image after which thev were blinked against each other., This subtraction technique was performed on each coadded image after which they were blinked against each other.
 For each coaddec image of each galaxy we determined the frame limit by using artificial stars and the exact techniques outlined above for detecting the novae., For each coadded image of each galaxy we determined the frame limit by using artificial stars and the exact techniques outlined above for detecting the novae.
" Table 2. gives the positions and munber of detections for the novae discovered in (his survey,", Table \ref{tab_nov_pos} gives the positions and number of detections for the novae discovered in this survey.
 The nova in M32. was discovered. first and is shown in outburst in Figure 1.., The nova in M32 was discovered first and is shown in outburst in Figure \ref{m32n1}.
 The nova candidate in NGC 205 is shown in outburst in Figure 2.., The nova candidate in NGC 205 is shown in outburst in Figure \ref{n205n1}. .
svuchrotron sell-Conmpton models are viable without requiring external seed. photons to account for (he observed y-ray emission. (ii) internal absorption effects are small or negligible. and (ii) that the high-energv beams require so much power (hat power ellicient beam formation models are strongly preferred.,"synchrotron self-Compton models are viable without requiring external seed photons to account for the observed $\gamma$ -ray emission, (ii) internal absorption effects are small or negligible, and (iii) that the high-energy beams require so much power that power efficient beam formation models are strongly preferred."
 We concentrate here on electromagnetic models as they predict powerful large-scale magnetic and electric fields that might be able to accelerate particles to very. hiel energies and produce large-scale ordered motion., We concentrate here on electromagnetic models as they predict powerful large-scale magnetic and electric fields that might be able to accelerate particles to very high energies and produce large-scale ordered motion.
 Three geometries for accelerating particles in electromagnetic models have been discussed in the literature., Three geometries for accelerating particles in electromagnetic models have been discussed in the literature.
 The models assume that the black hole and accretion disk are embedded in a conducting magnetosphere., The models assume that the black hole and accretion disk are embedded in a conducting magnetosphere.
 Free charge carriers are created in (he magnetosphere through cascades involving curvature radiation or inverse Compton scattering and pair creation processes., Free charge carriers are created in the magnetosphere through cascades involving curvature radiation or inverse Compton scattering and pair creation processes.
 A quasi-stationary configuration is achieved if the charge carriers are. distributed such that E-B=0.," A quasi-stationary configuration is achieved if the charge carriers are distributed such that $\vec{E} \cdot \vec{B}\,=0$."
 If this condition is Fulfilled. charged particles move along magnetic field lines without eaining or loosing energy.," If this condition is fulfilled, charged particles move along magnetic field lines without gaining or loosing energy."
 The magnetosphere can act as a series of nearly parallel conductors with zero resistance along the magnetic field lines (hat can sustain a voltage drop far away [rom where il was eeneraled., The magnetosphere can act as a series of nearly parallel conductors with zero resistance along the magnetic field lines that can sustain a voltage drop far away from where it was generated.
 The voltage drop mav either be generated by the disk itself 1976).. or by the Ixerr black hole spinning in (he horizon-threading magnetic field supported by the disk (Blancllored&ZnajekLOTT:Thorneetal.1936).," The voltage drop may either be generated by the disk itself \citep{Lovelace1976,Blandford1976}, or by the Kerr black hole spinning in the horizon-threading magnetic field supported by the disk \citep{Blandford1977,Thorne1986}."
. IXatz (2006) argues that the electric field generated by a homopolar generator in a rotating magnetized [hid has curls in the observers frame and cannot be shorted out al all locations by a stationary charge distribution., Katz (2006) argues that the electric field generated by a homopolar generator in a rotating magnetized fluid has curls in the observers frame and cannot be shorted out at all locations by a stationary charge distribution.
 It thus seems likely (hat non-stationary vacuum gaps form. which are capable ol accelerating particles.," It thus seems likely that non-stationary vacuum gaps form, which are capable of accelerating particles."
 A possible geometry of the magnetosphere and the gap region is shown in Figure 4.., A possible geometry of the magnetosphere and the gap region is shown in Figure \ref{GEOM}.
 The formation of the gap close to the rotation axis might explain the initial collimation of the jet., The formation of the gap close to the rotation axis might explain the initial collimation of the jet.
 In electromagnetic models. the Povnting flux launches the jet and anchors the jet to the disk or black hole.," In electromagnetic models, the Poynting flux launches the jet and anchors the jet to the disk or black hole."
 For simplicity we focus here on disk models: some considerations are applicable to the Blancllord-Znajek model as well., For simplicity we focus here on disk models; some considerations are applicable to the Blandford-Znajek model as well.
 We assume that the coordinate svstem r. ©. 2 ls aligned with the jet and the black hole resides at its origin.," We assume that the coordinate system $r$, $\phi$, $z$ is aligned with the jet and the black hole resides at its origin."
 Directly above the disk. the Povuting fIux is where (he subscripts p aud ( denote the poloidal (i aud z) and toroidal (6) components.," Directly above the disk, the Poynting flux is where the subscripts p and t denote the poloidal $r$ and $z$ ) and toroidal $\phi$ ) components,"
1.0d0.1 xiehter than a flux density of µ.ν.,$1.0 \pm 0.1$ $^{-2}$ brighter than a flux density of $\mu$ Jy.
 The correspondings count predicted by the 35-. 40-. 45- and 5O-IX. models are 0.6. 0.7. 1.0 and > respectively.," The correspondings count predicted by the 35-, 40-, 45- and 50-K models are 0.6, 0.7, 1.0 and $^{-2}$ respectively."
 Thus all the models. discussed here are consistent with the ceep radio observations., Thus all the models discussed here are consistent with the deep radio observations.
 For comparison. a count of 2 is predicted. by the modified Gaussian mocel discussed. in. DBarger. et. ((1999b).. which is. mocified from the results of the simple luminosity evolution models presented by Blain et ((1999c).," For comparison, a count of $^{-2}$ is predicted by the modified Gaussian model discussed in Barger et (1999b), which is modified from the results of the simple luminosity evolution models presented by Blain et (1999c)."
 The redshift distributions of subniillimetre-selected sources at. or just below. the flux density limits of current surveys have been discussed by Blain et ((1999¢) in the context of models of a strongly evolving population of distant clusty galaxies. based on the low-redshiltALS galaxy. luminosity function.," The redshift distributions of submillimetre-selected sources at, or just below, the flux density limits of current surveys have been discussed by Blain et (1999c) in the context of models of a strongly evolving population of distant dusty galaxies, based on the low-redshift galaxy luminosity function."
 The first spectroscopic observations of a large fraction of the potential optical counterparts to SCUBA galaxies identified in deep multicolour optical images (Smail et 11998) have been made by Darger et ((1999b) (see 110)., The first spectroscopic observations of a large fraction of the potential optical counterparts to SCUBA galaxies identified in deep multicolour optical images (Smail et 1998) have been made by Barger et (1999b) (see 10).
 This redshift distribution is consistent with the optical identifications made by Lilly ct (1999) in à SCUBA survey of CanadaFrance Redshift Survey (CERS) fieldwa, This redshift distribution is consistent with the optical identifications made by Lilly et (1999) in a SCUBA survey of Canada–France Redshift Survey (CFRS) fields.
- The distribution shown in 110 is. however. subject to potential misidentifications of SCUBA galaxies.," The distribution shown in 10 is, however, subject to potential misidentifications of SCUBA galaxies."
 For example. recent deep near-infrared images show that two of the Smail et ((1998) SCUBA galaxies. which were originally identified. with low-redshift spiral galaxies. can more plausibly be associated with extremely red objects (ΙΟΣ) that were unidentified in optical images (Small et 11999).," For example, recent deep near-infrared images show that two of the Smail et (1998) SCUBA galaxies, which were originally identified with low-redshift spiral galaxies, can more plausibly be associated with extremely red objects (EROs) that were unidentified in optical images (Smail et 1999)."
 In two other cases. at 2=2.55 (Ivison et 11999) and z=2S1 (νίκος et 11905). the identifications have been confirmed. by detections of τουςιο CO emission (Eraver et 11998: 1999). and in another case spectroscopy ancl£50 observations (Soucail et 11999) strongly support the identification of a ring galaxy al 2=1.06.," In two other cases, at $z=2.55$ (Ivison et 1999) and $z=2.81$ (Ivison et 1998), the identifications have been confirmed by detections of redshifted CO emission (Frayer et 1998; 1999), and in another case spectroscopy and observations (Soucail et 1999) strongly support the identification of a ring galaxy at $z=1.06$."
 The preliminary redshift clistribution of SCUBA ealaxies. shown in 110. is broadly consistent with the predictions of the Gaussian model of Blain ct (199960).," The preliminary redshift distribution of SCUBA galaxies, shown in 10, is broadly consistent with the predictions of the Gaussian model of Blain et (1999c)."
 Α modified Caussian model as shown in 11. was described by Barger οἱ ((1999b): the values of the evolution parameters in the modified Gaussian mocel were explicitly fitted both to the background. radiation intensity. and count cata and. to the observed. median redshift.," A modified Gaussian model, as shown in 1, was described by Barger et (1999b); the values of the evolution parameters in the modified Gaussian model were explicitly fitted both to the background radiation intensity and count data and to the observed median redshift."
 In 1400. the observed τοσα distribution is compared with the redshift distributions predicted. in the Gaussian ane mocified Gaussian models. and. with the predictions of the hierarchical models ceveloped here (see 11).," In 10 the observed redshift distribution is compared with the redshift distributions predicted in the Gaussian and modified Gaussian models, and with the predictions of the hierarchical models developed here (see 1)."
 Median redshifts of about 2.2. : 27.3.2 and 3.5 are expected in the 35-. 40-. 45- ancl 50-I&. hierarchical. models respectively.," Median redshifts of about 2.2, 2.7, 3.2 and 3.5 are expected in the 35-, 40-, 45- and 50-K hierarchical models respectively."
 The redshift distributions predicted by the hierarchical models have median redshifts greater than that in the moclificd Gaussian model. but. less than those in either the other models. presented. by Blain et ((1999c) or the hierarchical model [5 [rom Guiderdoni οἱ ((1998). all of which provide a reasonable fit to both the background radiation intensity and source counts in the far-infrared/submillimetre. waveband.," The redshift distributions predicted by the hierarchical models have median redshifts greater than that in the modified Gaussian model, but less than those in either the other models presented by Blain et (1999c) or the hierarchical model E from Guiderdoni et (1998), all of which provide a reasonable fit to both the background radiation intensity and source counts in the far-infrared/submillimetre waveband."
 Based: on these results. the coolest 35-Ix model seems to be in best agreement with," Based on these results, the coolest 35-K model seems to be in best agreement with"
"The individual regions for spectral fitting were determined using the contour binning method of Sanders (2006), which groups neighboring pixels of similar surface brightness until a desired signal-to-noise threshold is met.","The individual regions for spectral fitting were determined using the contour binning method of Sanders (2006), which groups neighboring pixels of similar surface brightness until a desired signal-to-noise threshold is met."
" For data of the quality discussed here, regions outside the X-ray bright arms are small enough that the X-ray emission from each can be approximated usefully by a single temperature plasma model."," For data of the quality discussed here, regions outside the X-ray bright arms are small enough that the X-ray emission from each can be approximated usefully by a single temperature plasma model."
 Regions inside the X-ray bright arms show strong evidence for multi-phase gas (see Paper I)., Regions inside the X-ray bright arms show strong evidence for multi-phase gas (see Paper I).
" In this paper,"," In this paper,"
the cluster is not a flat surface deusitv of σε=5825715 ealaxies per steradiau.,the cluster is not a flat surface density of $\sigma_f=583775$ galaxies per steradian.
 This effect is further exasperated by systematic plateto uucertainties in the magnitude zero-point of the EDSCC plotoeraphic plates (Nichol Collius 1993)., This effect is further exasperated by systematic plate–to–plate uncertainties in the magnitude zero-point of the EDSGC photographic plates (Nichol Collins 1993).
 We present the effective area of the whole EDC'CIT in Table 3. based on our siuiulatiou results.," We present the effective area of the whole EDCCII in Table \ref{area}
 based on our simulation results."
 These data were obtained by zunmiues over the larger EDSC survey area. as defined in Section 2.. but weighted by our success rate in detecting artificial clusters as computed from the sinaller test area.," These data were obtained by summing over the larger EDSGC survey area, as defined in Section \ref{EDSGC}, but weighted by our success rate in detecting artificial clusters as computed from the smaller test area."
 The data given in Table 3. illustrate the power of this new EDCCII catalogue since we now know the selection function of an opticallyselected cluster catalogue to the same accuracy as a Xravselected cluster survey Nichol et al., The data given in Table \ref{area} illustrate the power of this new EDCCII catalogue since we now know the selection function of an optically–selected cluster catalogue to the same accuracy as a X–ray–selected cluster survey Nichol et al.
 1999)., 1999).
 This effective area will be used below when calculating the space deusity of clusters (sec Section 1) The final use of our simulations was to determine the likely spurious detection rate., This effective area will be used below when calculating the space density of clusters (see Section \ref{results}) ) The final use of our simulations was to determine the likely spurious detection rate.
 Again. we have tried to use the real galaxy data as much as possible iu this analysis so as to ninic the real uncertainties in the EDSCC catalogue.," Again, we have tried to use the real galaxy data as much as possible in this analysis so as to mimic the real uncertainties in the EDSGC catalogue."
 This is different frou previous attempts to estimate the spurious detection rate which have relied on simulated ealaxy catalogues (see Postman et al., This is different from previous attempts to estimate the spurious detection rate which have relied on simulated galaxy catalogues (see Postman et al.
 1996)., 1996).
" To achieve this eoal therefore. we perturbed cach galaxy in the EDSCC in a raudonmi distauce (between 20, aud 56, with a flat distribution) iu a random direction from its original position."," To achieve this goal therefore, we perturbed each galaxy in the EDSGC in a random distance (between $2 \theta_c$ and $5 \theta_c$ with a flat distribution) in a random direction from its original position."
 This procedure effectively sinoothes the ealaxv catalogue ou these particular scales removing all sanallscale structure in the catalogue while retainius the largescale features within the catalogue., This procedure effectively smoothes the galaxy catalogue on these particular scales removing all small–scale structure in the catalogue while retaining the large–scale features within the catalogue.
 We then applied our matched filter algorithm to these perturbed galaxy catalogues and calculated the number of clusters that would satisfv our sclection criteria., We then applied our matched filter algorithm to these perturbed galaxy catalogues and calculated the number of clusters that would satisfy our selection criteria.
 This was performed nany thousands of times to determuue the standard deviation iu our spurious detection rate., This was performed many thousands of times to determine the standard deviation in our spurious detection rate.
 Our spurious detection rates are shown in Fieure I., Our spurious detection rates are shown in Figure \ref{spuriousplot}.
" The ΠΠΡΟ of spurious detections was siguificaut oulv for high redslift δὲ=50 and R,,=100 clusters.", The number of spurious detections was significant only for high redshift $R_m=50$ and $R_m=100$ clusters.
" Based ou these slunulations therefore. we restrict ourselves to :x0.12 or HR,100 and :<0.15 for R,,>LOO to ensure hat the spurious detection rate remains msiguificaut."," Based on these simulations therefore, we restrict ourselves to $z\le0.12$ for $R_m\le100$ and $z\le0.15$ for $R_m>100$ to ensure that the spurious detection rate remains insignificant."
 For exanuple. for δὲz200 svstems. «1% of detected clusters are likely to be spurious below :=0.15.," For example, for $R_m\ge200$ systems, $<1\%$ of detected clusters are likely to be spurious below $z=0.15$."
 We also exclude all clusters at z«0.05 as our thresholds are not accurately calibrated below this redshift., We also exclude all clusters at $z<0.05$ as our thresholds are not accurately calibrated below this redshift.
 In Section 1. therefore. we restrict ourselves to δι2100 systems in the redshift range 0.05«τς0.15 and make uo correction for spurious detectious within this richness aud redshift rauge.," In Section \ref{results} therefore, we restrict ourselves to $R_m\ge100$ systems in the redshift range $0.05<z<0.15$ and make no correction for spurious detections within this richness and redshift range."
" We note here that the results of these simulations are in good agreement with enipirical determünatious of the conipleteuess of the EDSCC aud EDCC catalogues. based ou 777 galaxy redshift ποσολος, Nichol (1993) showed that the EDCC was complete out to 5=0.13 and ouly coutained <15% contannination from spurious clusters."," We note here that the results of these simulations are in good agreement with empirical determinations of the completeness of the EDSGC and EDCC catalogues based on 777 galaxy redshift measurements, Nichol (1993) showed that the EDCC was complete out to $z\simeq0.13$ and only contained $<15\%$ contamination from spurious clusters."
 At low redshift. Luusdeu et al. (," At low redshift, Lumsden et al. ("
1992) also had difficulty detecting z<0.03 clusters iu the original EDCCT catalogue because of the large aneular size subteuded by such clusters.,1992) also had difficulty detecting $z<0.03$ clusters in the original EDCCI catalogue because of the large angular size subtended by such clusters.
 Iu addition to randomizing the positions of the EDSCC. we also performed simulations which randomly slated the magnitudes of the galaxies throughout the EDSGC cata.," In addition to randomizing the positions of the EDSGC, we also performed simulations which randomly shuffled the magnitudes of the galaxies throughout the EDSGC data."
 This resulted iu galaxy catalogues with the same statistical properties sune angular clustering and nuubermagnitude relationship but removed all. correlations with maguitude., This resulted in galaxy catalogues with the same statistical properties – same angular clustering and number--magnitude relationship – but removed all correlations with magnitude.
 Again. these randomized catalogues were analyzed with our matched filter algorithm aud. resulted Ina very simular result as presented im Figure { the spurious detection rate was insignificant for lower redshift. ligher richness systems.," Again, these randomized catalogues were analyzed with our matched filter algorithm and resulted in a very similar result as presented in Figure \ref{spuriousplot} the spurious detection rate was insignificant for lower redshift, higher richness systems."
 We note however. that ou average we detected fewer rich svstems than with the real EDSCC data.," We note however, that on average we detected fewer rich systems than with the real EDSGC data."
 Iu other words. magnitude correlations ouly appear to aid iu the detection of rich clusters in the data.," In other words, magnitude correlations only appear to aid in the detection of rich clusters in the data."
 Once we lave the cluster detections. as a function of redshift and richness. we ust then remove duplicate cluster detections to produce a final catalogue of unique cluster candidates.," Once we have the cluster detections, as a function of redshift and richness, we must then remove duplicate cluster detections to produce a final catalogue of unique cluster candidates."
 This was achieved by grouping together all svsteius whose measured ceutroids were within 2«6 of cach other.," This was achieved by grouping together all systems whose measured centroids were within $2\,\times\,\theta_c$ of each other."
 We began this process by grouping togetler candidates as a function of their redshitt. followed by their richness.," We began this process by grouping together candidates as a function of their redshift, followed by their richness."
 If duplicates were found. we simply averaged their rehuess aud redshift estimates to obtain a final estimate of the candidate clusters’ richness aud redshift.," If duplicates were found, we simply averaged their richness and redshift estimates to obtain a final estimate of the candidate clusters' richness and redshift."
 Tn the section. we prescut the results of applying the match filter algorithm as outlined in Section 3. to the whole EDSCC (as defined in Section 2)).," In the section, we present the results of applying the match filter algorithm as outlined in Section \ref{method} to the whole EDSGC (as defined in Section \ref{EDSGC}) )."
 However. we must first establish a common framework within which to compare our results with previous studies.," However, we must first establish a common framework within which to compare our results with previous studies."
 To this cud. we will quote our results both as a function of A4. the cluster richness as defined by Postinau et al. (," To this end, we will quote our results both as a function of $\Lambda_{cl}$, the cluster richness as defined by Postman et al. ("
1996) for the PDCS catalogue. aud Πιν which is defined by us.,"1996) for the PDCS catalogue, and $R_m$, which is defined by us."
" As R,,,and Αι ave just different richness normalizations of the match filter model (see Section 3.1)}. it is easy to relate the two analytically."," As $R_m$and $\Lambda_{cl}$ are just different richness normalizations of the match filter model (see Section \ref{matched}) ), it is easy to relate the two analytically."
" This was achieved by intcerating a normalized Schtecter function of richness &,,, over the hnmuuünositv range discussed in Section 3.1. and dividing bv £*. the characteristic luminosity of the Schtecter function."," This was achieved by integrating a normalized Schtecter function of richness $R_m$ over the luminosity range discussed in Section \ref{matched} and dividing by $L^{\star}$ , the characteristic luminosity of the Schtecter function."
 This is equivalent to the original definition of A4 iu Posti et al. (, This is equivalent to the original definition of $\Lambda_{cl}$ in Postman et al. (
1996).,1996).
 Uxing a=1.1 (Postman et al., Using $\alpha=-1.1$ (Postman et al.
" 1996). we show that A4=20.3 is R,,=100. A;=10.6 is Πιν=200 and Aw =al2is R,=1400."," 1996), we show that $\Lambda_{cl}=20.3$ is $R_m=100$, $\Lambda_{cl}=40.6$ is $R_m=200$ and $\Lambda_{cl}=81.2$ is $R_m=400$."
 We quote here X4 iu the PDCS Vi passbaud as this is closest to the EDSQC b; passbaud., We quote here $\Lambda_{cl}$ in the PDCS $V_4$ passband as this is closest to the EDSGC $b_j$ passband.
 Tn addition to comparing with the PDCS. we wish to compare with the original Abell catalogue.," In addition to comparing with the PDCS, we wish to compare with the original Abell catalogue."
 We achieve this by combining the equality Αρο)δι.)2:0.7 (taken from Postman et al., We achieve this by combining the equality $N_R(1.0)/N_R(1.5)\simeq0.7$ (taken from Postman et al.
 1996) with Figure 17 of Postiman et al. (, 1996) with Figure 17 of Postman et al. (
1996) to obtain au empirical relationship of R44;~2«Avy for low redshift clusters (2< 0.2).,1996) to obtain an empirical relationship of $R_{abell} \sim 2\times \Lambda_{cl}$ for low redshift clusters $z\le0.2$ ).
 Therefore. the three Avr vichnesses οἴνοι above approximately correspond to Abell Richness Classes (RC) 0. 1 aud 2 respectively.," Therefore, the three $\Lambda_{cl}$ richnesses given above approximately correspond to Abell Richness Classes (RC) 0, 1 and 2 respectively."
 These conversions are in good agreement with those prescuted iu Lubin Postinan (1996 and references therein) aud used bv others (olden. private conuuunicatious: DPostiman. private communications).," These conversions are in good agreement with those presented in Lubin Postman (1996 and references therein) and used by others (Holden, private communications; Postman, private communications)."
" For the rest of this paper. we concentrate on the &,,2LOO svstenis (X220 or RC 0)"," For the rest of this paper, we concentrate on the $R_m\ge100$ systems $\Lambda_{cl}\ge20$ or ${\rm RC}\ge0$ )."
 Iu Table L. we preseut the cumulative surface deusities of EDCCII clusters usingthe methodology outlined im this paper.," In Table \ref{density}, , we present the cumulative surface densities of EDCCII clusters usingthe methodology outlined in this paper."
 To compute these surface deusities. as a function of rehuess. we have sunuued the number of unique clusters," To compute these surface densities, as a function of richness, we have summed the number of unique clusters"
"searched amongst “normal galaxies"" which can at least xwtiallv account for the cxtragalactic light emitted at he wavelengths under exam (sec Macau Pozzetti 2000: 'l'otani et al.",searched amongst “normal galaxies” which can at least partially account for the extragalactic light emitted at the wavelengths under exam (see Madau Pozzetti 2000; Totani et al.
 2001)., 2001).
 In. their work. Washlinsky ο al. (," In their work, Kashlinsky et al. ("
2002) and Oclenwalel ct al. (,2002) and Odenwald et al. (
"2003) remove from the maps all resolved sources up to KeLY"" which. they claim. corresponds to the removal of possibly all resolvable galaxies oredshifts z21 (with some allowance made for low-surlace xiehtness systems).","2003) remove from the maps all resolved sources up to $\simeq 19^m$ which, they claim, corresponds to the removal of possibly all resolvable galaxies to redshifts $z\simeq 1$ (with some allowance made for low-surface brightness systems)."
 However. we argue that the population of galaxies responsible — together with Popllls — for the observed NIRB Iluctuations must reside at redshifts οκ2ον," However, we argue that the population of galaxies responsible – together with PopIIIs – for the observed NIRB fluctuations must reside at redshifts $z\simgt 2$."
 The argument supporting this conclusion goes as follows., The argument supporting this conclusion goes as follows.
 The power spectrum of Ductuations as measured. by PALASS [features a remarkable power-law behaviour down to angular scales of 1 aresec., The power spectrum of fluctuations as measured by 2MASS features a remarkable power-law behaviour down to angular scales of 1 arcsec.
 Under the hypothesis for this signal to be generated. by. large-scale structure. one has that since the correlation function of dark matter haloes (see equation 5)) must drop to -1 on distances smaller than two virial radii due to spatial halo-halo exclusion (see e.g. Magliocchetti Porciani 2003) — the power-law trend exhibited by the data down to Ll aresee sets a strong limit on the minimum redshift at which these “normal galaxies? appear.," Under the hypothesis for this signal to be generated by large-scale structure, one has that – since the correlation function of dark matter haloes (see equation \ref{eq:xi}) ) must drop to -1 on distances smaller than two virial radii due to spatial halo-halo exclusion (see e.g. Magliocchetti Porciani 2003) – the power-law trend exhibited by the data down to 1 arcsec sets a strong limit on the minimum redshift at which these “normal galaxies” appear."
 More in detail. one has that for masses of the order of 107AZ; (value taken as a reasonable lower limit for the kind of objects under exam) and the cosmology adopted in this paper. the condition for a power-law behaviour to hok at the minimum scales probed by the data is only met for redshifts 2<2.," More in detail, one has that for masses of the order of $10^8 M_{\sun}$ (value taken as a reasonable lower limit for the kind of objects under exam) and the cosmology adopted in this paper, the condition for a power-law behaviour to hold at the minimum scales probed by the data is only met for redshifts $z\simgt 2$."
" What are then these zo2 “normal galaxies"". fain enough not to be resolved by a Ws19 survey. which can significantly contribute to the observed. (uctuations of he NI. especially in the LE and Ix bands?"," What are then these $z\simgt 2$ “normal galaxies”, faint enough not to be resolved by a $\simlt 19$ survey, which can significantly contribute to the observed fluctuations of the NIRB, especially in the H and K bands?"
 X natura candidate can be found. in the population of high-z. star-orming galaxies detected. in the Hubble. Deep Fields to fyp=29.," A natural candidate can be found in the population of $z$ , star-forming galaxies detected in the Hubble Deep Fields to $I_{\rm AB}$ =29."
" Magliocchetti Macldox. (1999). have measure he clustering signal for this class of galaxies ancl shown hat they are indeed strongly clustered. with an amplitude Ac105 Gvbere the angular correlation function was xwamoetrized as e(8)248. 07), eonstant within the errors at all recdshifts 2s4.8."," Magliocchetti Maddox (1999) have measured the clustering signal for this class of galaxies and shown that they are indeed strongly clustered, with an amplitude $A\simeq 7\times 10^{-3}$ (where the angular correlation function was parametrized as $w(\theta)=A\theta^{-0.8}$ ), constant within the errors at all redshifts $2\simlt z\simlt 4.8$."
 The contribution of this class of galaxies to intensity luctuations in the NIRB can then be written as Chat)& ipe(@). where £j is the residual background intensity after source removal.," The contribution of this class of galaxies to intensity fluctuations in the NIRB can then be written as $C_{\rm gal}(\theta)\simeq I_{\rm B}^2
w(\theta)$ , where $I_{\rm B}$ is the residual background intensity after source removal."
 Suitable values for fp at the wavelengths relevant to this work can be obtained by taking the estimates for the total extragalactic light. emitted. in the J. HE and lx bands from Madau Pozzetti (2000). ancl by then subtracting the contributions from resolved sources as given in Ixashlinsky et al. (," Suitable values for $I_{\rm B}$ at the wavelengths relevant to this work can be obtained by taking the estimates for the total extragalactic light emitted in the J, H and K bands from Madau Pozzetti (2000), and by then subtracting the contributions from resolved sources as given in Kashlinsky et al. ("
2002).,2002).
 Note that the resulting Zjs. can in principle only be considered. as upper limits. since the Ixashlinsky and collaborators figures were given by assuming source removal up to z221l. while we are considering xickeround Uuctuations mace by οZ2 galaxies.," Note that the resulting $I_{\rm B}$ 's, can in principle only be considered as upper limits, since the Kashlinsky and collaborators figures were given by assuming source removal up to $z\simeq 1$, while we are considering background fluctuations made by $z\simgt 2$ galaxies."
 The power spectrum of intensity Ductuations in the IRB obtained for these “normal” galaxies in the J. HE and Ix bands is shown in Figures 2.. 3. and 4. by the dotted ines.," The power spectrum of intensity fluctuations in the NIRB obtained for these “normal” galaxies in the J, H and K bands is shown in Figures \ref{fig:PkJ}, \ref{fig:PkH} and \ref{fig:PkK} by the dotted lines."
 The solid lines instead illustrate thefofal signal steaming from the sum of the contribution of this class of sources with the one deriving from PopllIs. “, The solid lines instead illustrate the signal steaming from the sum of the contribution of this class of sources with the one deriving from PopIIIs. “
Normal”. faint. üeh-redshift. galaxies are able to reproduce the observed power spectrum at all angular scales in the Ix. band. bot= in its amplitude and its slope.,"Normal”, faint, high-redshift galaxies are able to reproduce the observed power spectrum at all angular scales in the K band, both in its amplitude and its slope."
" Their contribution however comes less and less important as one moves towareWu. shorter wavelengths and ab A=l25ym not onlv can just account for ~10% of the observed. signal. but. als ws a cillerent (shallower) dependence on the angular sca han what the data ]t is then found that. since Popllls and “normal” ealaxies seem to be responsible for the observed Iuctuations with a dillerent relevance at the different wavelengths (Popllls mainly appearing in the J band. ""normal galaxies eiving the strongest contribution in the Ix band). their suni is able to reproduce the data within the errors ataff wavelengths. (with some allowance made in the II band)."," Their contribution however becomes less and less important as one moves towards shorter wavelengths and – at $\lambda=1.25 \mu$ m – not only can just account for $\sim 10$ of the observed signal, but also has a different (shallower) dependence on the angular scale than what the data It is then found that, since PopIIIs and “normal” galaxies seem to be responsible for the observed fluctuations with a different relevance at the different wavelengths (PopIIIs mainly appearing in the J band, “normal” galaxies giving the strongest contribution in the K band), their sum is able to reproduce the data within the errors at wavelengths (with some allowance made in the H band)."
 In more detail. bv considering both populations we can give account for the observed (roughly constant) amplitude of the observed. power spectrum at the diferent wavelengths. and can also naturally explain the steepening of its slope when moving to shorter wavelengths as due to the intervening contribution of Poplll sources.," In more detail, by considering both populations we can give account for the observed (roughly constant) amplitude of the observed power spectrum at the different wavelengths, and can also naturally explain the steepening of its slope when moving to shorter wavelengths as due to the intervening contribution of PopIII sources."
 Note that the match is particularly &ood. especially if one considers thatm," Note that the match is particularly good, especially if one considers that."
"odel Our predictions do not depend on the exact functional form for the chosen EME and star-formation cllicicney f,. as lone as their combination is able to reproduce the observed NIR.B (SE03)."," Our predictions do not depend on the exact functional form for the chosen IMF and star-formation efficiency $f_{\star}$, as long as their combination is able to reproduce the observed NIRB (SF03)."
 Obviously. one of the main uncertainties of this work is associated to the determination of the slope of the predicted power spectrum.," Obviously, one of the main uncertainties of this work is associated to the determination of the slope of the predicted power spectrum."
 This is due to the lack of solid. determinations of the two- correlation function. from both an observational and a theoretical point of view. at the high redshifts and small masses under exam.," This is due to the lack of solid determinations of the two-point correlation function, from both an observational and a theoretical point of view, at the high redshifts and small masses under exam."
 The fact that our predictions result within the errors in agreement. with the observed. trend however gives us some confidence that the assumptions made in the course of our analysis are (at least) plausible., The fact that our predictions result -- within the errors – in agreement with the observed trend however gives us some confidence that the assumptions made in the course of our analysis are (at least) plausible.
" As a last remark we note that. even if in the above analysis we have considered. Zi in the different bands. as an upper limit of the true background intensity. the [act that the data at A= 2.175m is perfectly reproduced: by the population of ""normal"" galaxies suggests this upper limit to be very close to the real value for otherwise the curves wouldhave fallen below the observations."," As a last remark we note that, even if in the above analysis we have considered $I_{\rm B}$ in the different bands as an upper limit of the true background intensity, the fact that the data at $\lambda=2.17\mu$ m is perfectly reproduced by the population of “normal” galaxies suggests this upper limit to be very close to the real value for otherwise the curves wouldhave fallen below the observations."
 Under this assumption. there is then no need for the emergence of a third. stil-unknown. population of objects. since. PopllHs.," Under this assumption, there is then no need for the emergence of a third, still-unknown, population of objects, since PopIIIs,"
,.
 Qur prediction is consistent with the BINA95% confideuce result., Our prediction is consistent with the BIMA$95\%$ confidence result.
 But the comparison betwee observations aud theoretical predictions weeds further investigation., But the comparison between observations and theoretical predictions needs further investigation.
 Ou the simulation part. as we discuss above. the detailed normalization depends ou oy. uou-eravitational heating aud resolutio effects.," On the simulation part, as we discuss above, the detailed normalization depends on $\sigma_8$, non-gravitational heating and resolution effects."
 Observatious will put strong constraius on these aspects., Observations will put strong constrains on these aspects.
 Ou tlie observational part. the BINLA result needs to be recousidered.," On the observational part, the BIMA result needs to be reconsidered."
 It covers the sky where no strong SZ temperature clistortio due to known galaxy clusters exists., It covers the sky where no strong SZ temperature distortion due to known galaxy clusters exists.
 Since clusters contribute a significant if not domiuaut fractio to the SZ power spectrum. the BINLA result may be smaller than the statistical mean value.," Since clusters contribute a significant if not dominant fraction to the SZ power spectrum, the BIMA result may be smaller than the statistical mean value."
 Furthermore. the error of the BINLA result is estimated under the Caussiau assumption. but as we will see below. the SZ ellect is hiehly non-Gaussian around the BIMA central multipole /=5530.," Furthermore, the error of the BIMA result is estimated under the Gaussian assumption, but as we will see below, the SZ effect is highly non-Gaussian around the BIMA central multipole $l=5530$."
 The stroug non-Ciaussiaulty increases the intrinsic error of the SZ power spectrum rueasurement., The strong non-Gaussianity increases the intrinsic error of the SZ power spectrum measurement.
 According to fig. 6..," According to fig. \ref{fig:ngs},"
 the error in the power spectrum measurement caused by the SZ ellect is abot 3 times larger thau the correspoudiug Gaussian case., the error in the power spectrum measurement caused by the SZ effect is about $3$ times larger than the corresponding Gaussian case.
 Nontheless. tliese results demoustrate the convergeuce between theory aud observation.," Nontheless, these results demonstrate the convergence between theory and observation."
 Iu the uear future. routine and accurate measurement of the SZ ellect will be possible in raucous fields.," In the near future, routine and accurate measurement of the SZ effect will be possible in random fields."
4. Π will put stronger requirement for our theoretical uicderstaudiug of the SZ ellect aud allows us to study the effects of uou-gravitational heating. radiative cooling aud the thermal history of the IGM.," It will put stronger requirement for our theoretical understanding of the SZ effect and allows us to study the effects of non-gravitational heating, radiative cooling and the thermal history of the IGM."
 Iu contrast to the primary CMB. the SZ ellect is noun-Ciaussian. arising Crom the non-linearity of the intervening gas.," In contrast to the primary CMB, the SZ effect is non-Gaussian, arising from the non-linearity of the intervening gas."
 This non-Caussianity affects the error analysis and may help to separate the SZ effect [rom the primary CMB iu observation., This non-Gaussianity affects the error analysis and may help to separate the SZ effect from the primary CMB in observation.
 To quantity these effects. we sinooth SZ maps uxing a top hat window of radius @ aud measure the kurtosis of the sinoothed y.," To quantify these effects, we smooth SZ maps using a top hat window of radius $\theta$ and measure the kurtosis of the smoothed $y$."
" The kurtosis Q,—=i ——PF--—3is generally a function of smoothing scale."," The kurtosis $\Theta_4\equiv \frac{\langle y-\bar{y} \rangle 
^4}{\sigma_y^4}-3$ is generally a function of smoothing scale."
" For a Gaussian signal. O,=0."," For a Gaussian signal, $\Theta_4=0$."
 Fig., Fig.
 6 shows that O;~200 at simall scales as 0.—0.," \ref{fig:ngs} shows that $\Theta_4\simeq 200$ at small scales as $\theta\rightarrow
0$."
 This result is consistent with the prediction ∐⋅∩⋯⋯⊔⋅∐↥≺↵↕⋅⋜⋃⋅∢∙∐∎∢∙⋜↕↥∐↕⋯⇂≺↵↥⋜↕↥↽≻↥↽≻⋅∩⋜↕∢∙∐∩⊓∐≺↵∺∠≺↵∐⇀≺↵∢∙↕↸⇝↕⋜⋃, This result is consistent with the prediction from our hierarchical model approach of the SZ effect \citep{Zhang01a}.
"∑≟⋅∖↽↥⋟≺↵∐⊇∩∪⋥⋝⋅⋅∖↾∖⊽≺↵↥↽∐⋅≺↵≺∐∢∙↕≺↲≺⇂⋅⋜↕↕⊳∖∐↕⋜↕∐ augular scales. O,~Siar:1) 200."," We predicted, at small angular scales, $\Theta_4\sim
S_4\sigma_g^2(z\sim 1) \sim 200$ ."
 Here. |=) ⋅ ⊤∿∐⊔⊳∖⋜↕↥∐≺↵⋅⋜⋃⋅∢∙↥∐∢∙⋜↕↥∐↕⋯⇂≺↵↥∢∙∩≺↲∐∐∙↓≺↵⋯⋅ ⋅⋅ ⋜↾∺∢∙≺⋈∙∢∙∐∐⋜⋃⋅↓⋅∩⋅∖↽∌⊲⋅↕≺↵⋯⋜↕∐⋯∪⋠⊔⋜↕∐≺⇂√⊼∙↙∕↸∶∿↽⊔∿⋅↴↥⊳∖↕∐≺↵∑≟⋜↕⊳∖≺⇂≺↵∐⊳∖∐⊽∖⇁≺∐⊳∖↥↽≻≺↲↥⋅⊳∖↥∩∐⋜↕↕∶∿↕⋅⇀−∖," Here, $S_4\equiv\frac{\langle\delta^4\rangle}{\langle \delta^2 
\rangle^3}\sim 40$ is a hierarchical model coefficient \citep{Scoccimarro99} and $\sigma_g(z\sim1)\sim 5$ is the gas density dispersion at $z\sim1$."
"↕↥⋜⋃⋅∑≟≺↵ augular scales 0~C,20. O;x-21 and reflects the strong non-Craussianity. at this. scale."," At large angular scales $\theta\sim 20^{'}$, $\Theta_4\gg 1$ and reflects the strong non-Gaussianity at this scale."
 It means that. at this augular scale. the dominant contribution to tlie y parameter is from highly non-linear regious auc there is still stroug correlation at augular scales down to £1000. as can be seen frou. the SZ power spectrum.," It means that, at this angular scale, the dominant contribution to the $y$ parameter is from highly non-linear regions and there is still strong correlation at angular scales down to $l\sim 1000$, as can be seen from the SZ power spectrum."
" The kurtosis. iu. multipole. space a,=¢/DIETP—.3."," The kurtosis in multipole space $a_4\equiv\langle
|a|^4 \rangle/\langle |a|^2 \rangle^2-3$."
" As usual. 4=«4, defined. by O(n)EN=M,«nM«pp(1)are multipole. inodes of. O and C;.=57lαἱ244qi1)LoHoqx (al."," As usual, $a\equiv
a_{lm}$ defined by $\Theta(\hat{n})\equiv \sum_{lm} a_{lm}Y_l^m(\hat{n})$are multipole modes of $\Theta$ and $C_l\equiv
\sum_{-l}^{l} |a_{lm}|^2/(2l+1)\equiv \langle |a|^2\rangle$ ."
" Here. Y,"" are the spherical harmonics."," Here, $Y_{l}^m$ are the spherical harmonics."
 C...) is averaged over all in=—L.... [andall maps.," $\langle \ldots \rangle$ is averaged over all $m=-l,\ldots,l$ andall maps."
" We calculate a, indirectly through the map-map varlauce 63;0) of Cy. which is related to ay; by"," We calculate $a_4$ indirectly through the map-map variance $\sigma_M(l)$ of $C_l$ , which is related to $a_4$ by"
Until differences between these three methods a'e beter understood. the issue of the degree of uiiformity of the ISM abuudauces in a dwarl gal:NY Cai1ot be properly acdressecl.,"Until differences between these three methods are better understood, the issue of the degree of uniformity of the ISM abundances in a dwarf galaxy cannot be properly addressed."
 The N/O ratio in ESO 173-C2[. which is a veΝ nelal poor galaxy. is fouιά {ο be in the normal ange for «warl irregular galaxies. out elevated when coupared to the very 1arrow range found by [T99 for low metallicity BCGs.," The N/O ratio in ESO 473-G24, which is a very metal poor galaxy, is found to be in the normal range for dwarf irregular galaxies, but elevated when compared to the very narrow range found by IT99 for low metallicity BCGs."
 TIis shows that rot all ealaxies with 12 + og (O/H) < 7.6 have identical eemiental abuudauce ratios. and this iiplies hat the Izotov Tjuan scenario for low uetallicity eagalaxies is not universal.," This shows that not all galaxies with 12 $+$ log (O/H) $\le$ 7.6 have identical elemental abundance ratios, and this implies that the Izotov Thuan scenario for low metallicity galaxies is not universal."
 Measiremeuts of the HIT region sin NGC 625 vield log (N/O) = —1.25., Measurements of the HII regions in NGC 625 yield log (N/O) $\approx$ $-$ 1.25.
 TIis relatively high N/O allo may ye juclicative of a lone qtuescent period prior to the recent active bust of star forαἱlon., This relatively high N/O ratio may be indicative of a long quiescent period prior to the recent active burst of star formation.
 Altogether. tle observations presented bere and iithe literature can be seen to suppo1 the scenario olN produce by intermediate mass stars.," Altogether, the observations presented here and in the literature can be seen to support the scenario of N produced by intermediate mass stars."
 The la‘oe scatter ΝΟ at a given Ο/Η is still cosistent with contamination w O followed yw later coutaumination by N. This is in agreemen with what is observed in the high redshift damped Lya systeiis (see Pettini et 22002 aud references therein) ancl resolves the conlict )ebween hiel and low redshilt observations ciscussed by Pivugiu (1999)., The large scatter in N/O at a given O/H is still consistent with contamination by O followed by later contamination by N. This is in agreement with what is observed in the high redshift damped $\alpha$ systems (see Pettini et 2002 and references therein) and resolves the conflict between high and low redshift observations discussed by Pilyugin (1999).
 The oxygen abulalces in the Sculptor Cirou» dis are in good agreement with tle relatiouship between metallicity a μαuuiiosity observed in the Local Group dls., The oxygen abundances in the Sculptor Group dIs are in good agreement with the relationship between metallicity and luminosity observed in the Local Group dIs.
 Takeu together tle observationS show a better relatioshi2 between metallicity ane luminosity than between metallicity aud galaxy central surface brightiess., Taken together the observations show a better relationship between metallicity and luminosity than between metallicity and galaxy central surface brightness.
 The gas contents of the Sculptor Group cls a'e large compared to the Local Group dls., The gas contents of the Sculptor Group dIs are large compared to the Local Group dIs.
 Λally of the Sculptor Croup dls lie close to the evolutionary path predicted by a simple closed box model with au ellective oxveen vield tliat is close to the theoretically favo'ed values., Many of the Sculptor Group dIs lie close to the evolutionary path predicted by a simple closed box model with an effective oxygen yield that is close to the theoretically favored values.
 It is possible that some of the Sculptor Ciroup dwarls are evolving nearly as closed systems., It is possible that some of the Sculptor Group dwarfs are evolving nearly as closed systems.
 Ou the other hand. if the abundances tu their exteuded HI disks are lower than in the HII regious. the appropriate value of the effective yield would be much smaller.," On the other hand, if the abundances in their extended HI disks are lower than in the HII regions, the appropriate value of the effective yield would be much smaller."
 The higher gas contents. lower average star formation rates. aid closer resemblance to closed box evolution for Sculptor Ciroip dis relative to Local Croup ds could all be indicative of evolution in a relatively lower ¢eusitv e1virouimnent.," The higher gas contents, lower average star formation rates, and closer resemblance to closed box evolution for Sculptor Group dIs relative to Local Group dIs could all be indicative of evolution in a relatively lower density environment."
 We are grateful for the help of οΓΙΟ stall members Ma1iel Hernandez aud Ricardo Venegas., We are grateful for the help of CTIO staff members Manuel Hernandez and Ricardo Venegas.
 We wish to thank BBothun. ΕΕ». Garιοί. HHeurys.. Whvennicutt. TTolstoy. aud L. vau Zee for many helpful couversations.," We wish to thank Bothun, Edmunds, Garnett, Henry, Kennicutt, Tolstoy, and L. van Zee for many helpful conversations."
 Joli1 Canuou. Heruy. Lee. auc Liese vau Zee proolreacl aud provided valuable comments ou his manusipt.," John Cannon, Herny Lee, and Liese van Zee proofread and provided valuable comments on this manuscript."
 We also tliauk the referee for a prompt aud careful readiug of the mauuscript and heir valuade quibbles., We also thank the referee for a prompt and careful reading of the manuscript and their valuable quibbles.
 This research has made use of the NASA/IPAC Extragalactic Database (NED) whic Lis operated by the Jet Propulsion Laboratory. California Institute of Technology. uier contrac with the National Aeronautics aud Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
 This research bas made use of NÀSA's Astrophysics Data System Abstract Service., This research has made use of NASA's Astrophysics Data System Abstract Service.
 EDS acknowledges partial support from a ASA LTSARP grant No., EDS acknowledges partial support from a NASA LTSARP grant No.
 NA...(-922] and the University of Minnesota., NAG5-9221 and the University of Minnesota.
 BWI is supported by the Gemini Observatory. which is operated by the Association of Universities for Research in Astronomy. Inc.. on bebalf of the international Gemini," BWM is supported by the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., on behalf of the international Gemini"
Also. we have considered only the role of the streaming instability in scattering cosmic rays: however. this is the dominant process for ~| cosmic rays (2)..,"Also, we have considered only the role of the streaming instability in scattering cosmic rays; however, this is the dominant process for $\sim 1$ cosmic rays \citep{Kulsrud2005}."
" It is important to ask if other forms of turbulence may help scatter cosmic rays: however. turbulent wave modes. especially on the very small scales of the cosmic-ray gyroradius (7,~10'? cem) are believed to lose energy very quickly in molecular clouds (e.g..22).."," It is important to ask if other forms of turbulence may help scatter cosmic rays; however, turbulent wave modes, especially on the very small scales of the cosmic-ray gyroradius $r_{\rm g} \sim 10^{12}$ cm) are believed to lose energy very quickly in molecular clouds \citep[e.g.,][]{MacLowEtAl1998, OstrikerEtAl2001}."
 In addition. those modes do not scatter cosmic rays efficiently (e.g..22).. so we do not expect the diffusivity to be affected by those processes.," In addition, those modes do not scatter cosmic rays efficiently \citep[e.g.,][]{Chandran2000b,YanLazarian2004}, so we do not expect the diffusivity to be affected by those processes."
 Finally. when considering the role of cosmic rays in galactic winds. we do not consider the drop in velocity of the cool cloud as cosmic rays enter the cloud.," Finally, when considering the role of cosmic rays in galactic winds, we do not consider the drop in velocity of the cool cloud as cosmic rays enter the cloud."
 Such a drop in velocity is likely to occur. as the hot-wind component tis usually inferred to move at thousands of kni/s (e.g..?).. and the clouds have velocities of =300—500 kkm/s (e.g..22).. so that the velocity difference could be of order a factor of four.," Such a drop in velocity is likely to occur, as the hot-wind component is usually inferred to move at thousands of km/s \citep[e.g.,][]{StricklandHeckman2009}, and the clouds have velocities of $\la 300-500$ km/s \citep[e.g.,][]{Martin2005,
 SteidelEtAl2010}, so that the velocity difference could be of order a factor of four."
 In our calculations. though. the speed already drops by an order of magnitude: a smaller drop in the bulk gas velocity would have the same. but smaller. effect on the transition of cosmic rays into the cool cloud.," In our calculations, though, the speed already drops by an order of magnitude; a smaller drop in the bulk gas velocity would have the same, but smaller, effect on the transition of cosmic rays into the cool cloud."
 We infer that such a change in the bulk velocity will therefore only further increase the cosmic-ray flux injected into the clouds. and will have a similar impact on cosmic-ray pressure and heating as increases in the magnetic field strength and cosmic-ray pressure have.," We infer that such a change in the bulk velocity will therefore only further increase the cosmic-ray flux injected into the clouds, and will have a similar impact on cosmic-ray pressure and heating as increases in the magnetic field strength and cosmic-ray pressure have."
 Therefore. the effects of increased velocity in the hot. ambient medium would not qualitatively change the results presented here.," Therefore, the effects of increased velocity in the hot, ambient medium would not qualitatively change the results presented here."
 How do our results compare to previous work?, How do our results compare to previous work?
 In the two sections below. we consider previous work on cosmic rays in cool clouds in the interstellar medium and in the galactic winds.," In the two sections below, we consider previous work on cosmic rays in cool clouds in the interstellar medium and in the galactic winds."
 The small drop in cosmic-ray pressure that we predict might seem similar to the “exclusion” of cosmic rays from clouds that was found by both ?. and ?.., The small drop in cosmic-ray pressure that we predict might seem similar to the “exclusion” of cosmic rays from clouds that was found by both \citet{SkillingStrong1976} and \citet{CesarskyVoelk1978}.
 However. our model differs from those previous works in significant ways. and therefore this new model brings new effects to light.," However, our model differs from those previous works in significant ways, and therefore this new model brings new effects to light."
 In those earlier papers. the exclusion of cosmic rays was driven by a drop in cosmic-ray density at E=300 due to ionization loses: as low-energy cosmic rays traversed the entire cloud. they were destroyed in collisions.," In those earlier papers, the exclusion of cosmic rays was driven by a drop in cosmic-ray density at $E \la 300$ due to ionization loses; as low-energy cosmic rays traversed the entire cloud, they were destroyed in collisions."
 It is also important that those works assumed a uniform cosmic-ray pressure outside of the cloud. and so assumed that the streaming instability only operated on the outskirts of the clouds.," It is also important that those works assumed a uniform cosmic-ray pressure outside of the cloud, and so assumed that the streaming instability only operated on the outskirts of the clouds."
 In comparison. our work assumes an initial cosmic-ray pressure gradient (see Section ??)). such that the streaming instability operates throughout the hot phase of the ISM.," In comparison, our work assumes an initial cosmic-ray pressure gradient (see Section \ref{boundaryConditions}) ), such that the streaming instability operates throughout the hot phase of the ISM."
 This assumption seems reasonable as we have found that even very small pressure gradients lead to the growth of the streaming instability. which seems plausible since the streaming instability can explain the observed near isotropy of cosmic rays at the Earth (see.e.g...2)..," This assumption seems reasonable as we have found that even very small cosmic-ray pressure gradients lead to the growth of the streaming instability, which seems plausible since the streaming instability can explain the observed near isotropy of cosmic rays at the Earth \citep[see, e.g., ][]{Kulsrud2005}."
 With this configuration. our models show a qualitatively new effect at higher energies: the transition of cosmic rays from the advective regime to the diffusively-dominated regime leads to a (small) cosmic-ray pressure drop in the outskirts of clouds.," With this configuration, our models show a qualitatively new effect at higher energies: the transition of cosmic rays from the advective regime to the diffusively-dominated regime leads to a (small) cosmic-ray pressure drop in the outskirts of clouds."
 This effect is also present at low energies (down to ΕΞΙ MeV)). and is more important in our models than ionization losses (at least for cosmic rays just entering the outskirts of the clouds).," This effect is also present at low energies (down to $E = 1$ ), and is more important in our models than ionization losses (at least for cosmic rays just entering the outskirts of the clouds)."
the pressure as function of the temperature is given bv the Stefan-Doltzmann form where the factor 3 accounts lor the three charge states of the pion.,the pressure as function of the temperature is given by the Stefan-Boltzmann form = 3 T^4 where the factor 3 accounts for the three charge states of the pion.
 The corresponding form lor an ideal quark-glion plasma with (wo flavours and three colours is In ((3)). the first temperature term in the curly brackets accounts for the (wo spin ancl eight colour degrees of freedom of the gluons. the second for the three colour. two flavour. two spin and (wo particle-antiparticle degrees of [reedom of the quarks. with 7/8 to obtain the correct statistics.," The corresponding form for an ideal quark-gluon plasma with two flavours and three colours is = 2 8 + (3 2 2 2) T^4 - B = 37 T^4 - B. In\ref{2.3b}) ), the first temperature term in the curly brackets accounts for the two spin and eight colour degrees of freedom of the gluons, the second for the three colour, two flavour, two spin and two particle-antiparticle degrees of freedom of the quarks, with 7/8 to obtain the correct statistics."
 The bag pressure D takes into account the dillerence between the physical vacuum ancl the ground state for quarks and eluons in a medium., The bag pressure $B$ takes into account the difference between the physical vacuum and the ground state for quarks and gluons in a medium.
 Since m thermodvuamics. à svstem chooses (he state of lowest [ree energv ancl hence hiehest pressure. we compare in το λα (thetemperature behaviour of ((3)) and (2)).," Since in thermodynamics, a system chooses the state of lowest free energy and hence highest pressure, we compare in \\ref{2_3}$ $~\!$a thetemperature behaviour of \ref{2.3a}) ) and \ref{2.3b}) )."
 Our simple model (hus leads to a (wo-phase picture of strongly interacting matter. wilh a haclronic phase up to," Our simple model thus leads to a two-phase picture of strongly interacting matter, with a hadronic phase up to T_c = ( 0.72"
manner.,manner.
 Iu addition. teniperature fluctuations iu the rauge 0.7SAToxl.l keV are observed out to ~30 kpe (Buote et al.," In addition, temperature fluctuations in the range $0.7 \lta kT \lta 1.4$ keV are observed out to $\sim 30$ kpc (Buote et al."
 2003)., 2003).
 The most important morphological attribute of these thermal and density ireguluities is that thev are svinunietrie about the ceuter of NGC 5011. with no evidence of jet-likeeps orientation.," The most important morphological attribute of these thermal and density irregularities is that they are symmetric about the center of NGC 5044, with no evidence of jet-like orientation."
 If ACN heating is the agenev that the eas from cooling. the heating in NGC 5041 aud other similar E salaxies must be more racially sviunietrie than the jet-heating considered iu the radio galaxy simulations discussed above.," If AGN heating is the agency that keeps the gas from cooling, the heating in NGC 5044 and other similar E galaxies must be more radially symmetric than the jet-heating considered in the radio galaxy simulations discussed above."
 Alternatively. )ecause of the short dynamical times in galaxy scale flows. he jet axis in would have to be randomly reoriented every ~109 vears to nuüntaiu the appearance of svuuetric ieating.," Alternatively, because of the short dynamical times in galaxy scale flows, the jet axis in would have to be randomly reoriented every $\sim 10^6$ years to maintain the appearance of symmetric heating."
 Finally. the outward mass fiux AJ ds likely to be uaxinized in spherically svuuuectric bubble outflows that Love away from the central ACUN inu OVCLY direction.," Finally, the outward mass flux ${\dot M}$ is likely to be maximized in spherically symmetric bubble outflows that move away from the central AGN in every direction."
 Ruszkowshki Degeliian (2002) aud Brigegen (20035) wave studied non-cooling. static models of cluster gas atmospheres.," Ruszkowski Begelman (2002) and Brügggen (2003b) have studied non-cooling, static models of cluster gas atmospheres."
 The gas is heated at large radii by thermal conduction and at siiall radii by a cloud of tiny expauding mibbles. a heating mechanisia proposed by Beeelman (20013.," The gas is heated at large radii by thermal conduction and at small radii by a cloud of tiny expanding bubbles, a heating mechanism proposed by Begelman (2001)."
 The heating is spherically saunietric. similar to he flows we cliscuss here.," The heating is spherically symmetric, similar to the flows we discuss here."
 These cluster-scale τοσο] are almost uniquely successful in prescrving the positive cluperature eradieut observed., These cluster-scale models are almost uniquely successful in preserving the positive temperature gradient observed.
 However. this combination of ceutral heating and conduction does not work for flows on galaxy/eroup scales: even with the maxinuun (Spitzer) conductivity it is impossible to eet both small AL and acceptable temperature gradients (Brighenti Mathews 2003).," However, this combination of central heating and conduction does not work for flows on galaxy/group scales; even with the maximum (Spitzer) conductivity it is impossible to get both small ${\dot M}$ and acceptable temperature gradients (Brighenti Mathews 2003)."
" In the heating scheme proposed by Ruszkowski Begchnan a large uuuber of very small ""offervescent? bubbles are produced near the ceutral AGN aud buoyautly flow into the hot gas throughout a large volume. acquiring a steady state radial distribution."," In the heating scheme proposed by Ruszkowski Begelman a large number of very small “effervescent” bubbles are produced near the central AGN and buoyantly flow into the hot gas throughout a large volume, acquiring a steady state radial distribution."
 Thebubbles do not mix thermally with the ambicut gas but expand adiabatically. heating the ambient How by Pd work as they expand.," The bubbles do not mix thermally with the ambient gas but expand adiabatically, heating the ambient flow by $PdV$ work as they expand."
" While thisbubble model has some features in common with that cousidercd here, we fiud that very simallbubbles are incousisteut with steady state circulating flows."," While this bubble model has some features in common with that considered here, we find that very small bubbles are inconsistent with steady state circulating flows."
 Tu the following discussion we investigate the possibility hat hot eas iu ealaxy/eroups similar to NGC 1172 or NGC. 50LL iutois hotheated near the center of the flow and convertedbubbles., In the following discussion we investigate the possibility that hot gas in galaxy/groups similar to NGC 4472 or NGC 5044 is heated near the center of the flow and converted into hot bubbles.
 The rising bubbles are required © CALEY lass outward at the same rate that nass arrives in he core by normal radiative lossesin the iterbublble flow., The rising bubbles are required to carry mass outward at the same rate that mass arrives in the core by normal radiative losses in the interbubble flow.
 The iuterbubble flow receives both inomenutum aud euergy roni the rising bubbles., The interbubble flow receives both momentum and energy from the rising bubbles.
 To optiuize the outward rate of ass transfer. the heating i asstuned tobe svuuuctric about the central AGN. not confined to a single jet.," To optimize the outward rate of mass transfer, the heating is assumed to be symmetric about the central AGN, not confined to a single jet."
 We explore here the conditions required to establish oppositely moving. two phase. steady state flows in galaxy/groups in which the gas does not cool to low temperatures.," We explore here the conditions required to establish oppositely moving, two phase, steady state flows in galaxy/groups in which the gas does not cool to low temperatures."
 A cool phase uear the virial temperature flows nsviud as in a normal cooling flow., A cool phase near the virial temperature flows inward as in a normal cooling flow.
" At some small radius nj wo assume this flow is heated and. traustormed iuto hotbubbles that buovantlv float upstream in the cooling flow. returning the nass fux received in the inconnue flow to some large circulation radius r,. where thebubble cutropy matches that of the local cooling flow eas."," At some small radius $r_h$ we assume this flow is heated and transformed into hot bubbles that buoyantly float upstream in the cooling flow, returning the mass flux received in the incoming flow to some large circulation radius $r_c$ where the bubble entropy matches that of the local cooling flow gas."
 Cas circulates in the eravitational field like a thermodvuamic eugime., Gas circulates in the gravitational field like a thermodynamic engine.
" The mechanisin that produces the bubbles is not specified and for simplicity the physical condition of the eas within rj, or bevoud r. are not cousidered.", The mechanism that produces the bubbles is not specified and for simplicity the physical condition of the gas within $r_h$ or beyond $r_c$ are not considered.
" Clearly. an idealized steady flow model like this caunot be strictly correct since the mass of circulating eas is expected to eraduallv increase due to stellar mass loss and inflow of gas from r>r,.."," Clearly, an idealized steady flow model like this cannot be strictly correct since the mass of circulating gas is expected to gradually increase due to stellar mass loss and inflow of gas from $r > r_c$."
 As the mass of circulating gas slowly iucreases. the creasing eas density should eventually lead to exteusive cooling.," As the mass of circulating gas slowly increases, the increasing gas density should eventually lead to extensive cooling."
 Nevertheless. a quasi-steady circulation model could be a reasonable represcutation for a limited time between cooling episodes.," Nevertheless, a quasi-steady circulation model could be a reasonable representation for a limited time between cooling episodes."
 Uulike the umnuerical simulations described earlier. uo new lass Is injected iuto the flow.," Unlike the numerical simulations described earlier, no new mass is injected into the flow."
 We also asstune that the hot bubbles retain their cohereney as they move through ~LO15 hubhble diauneters., We also assume that the hot bubbles retain their coherency as they move through $\sim 10 - 15$ bubble diameters.
 This degree of coherence is in fact observed iu the Perseus Cluster (Fabian et al., This degree of coherence is in fact observed in the Perseus Cluster (Fabian et al.
" 2000). asstuuing that the distaut ""ghost cavities” were formed closer to the central AGN."," 2000), assuming that the distant “ghost cavities” were formed closer to the central AGN."
 The hot bubbles simulated by Churazov et al. (, The hot bubbles simulated by Churazov et al. (
2001) float upward in the cooling flow atimosplicre and come to rest at some large radius where the ambicut chtropy is equal to that within the bubble.,2001) float upward in the cooling flow atmosphere and come to rest at some large radius where the ambient entropy is equal to that within the bubble.
 IToxcever. we iive found bubble evolution to be a dificult problem for ποασ] and we remain pical of our asstunption stainof bubble coherency.," However, we have found bubble evolution to be a difficult problem for numerical gasdynamics and we remain sceptical of our assumption of bubble coherency."
 The bubbles in Perseus nav be stabilized ly internal magnetic fields., The bubbles in Perseus may be stabilized by internal magnetic fields.
" It would jf relatively easy to modify our cireulation model to include bubble fragmentation according to somehoc svescription. but the effect of bubble fragmentation can casily be imagined by ου.” suuilar circulation flows uius different constant bubble masses,"," It would be relatively easy to modify our circulation model to include bubble fragmentation according to some prescription, but the effect of bubble fragmentation can easily be imagined by comparing similar circulation flows having different constant bubble masses."
 Iuour circulating flows the hot bubbles do not thermally uix (by thermal couduction for example) with the inconius ow before reaching the circulation radius ri where thebubble aud iuftiow entropies are approximately equal., In our circulating flows the hot bubbles do not thermally mix (by thermal conduction for example) with the incoming flow before reaching the circulation radius $r_c$ where the bubble and inflow entropies are approximately equal.
 Therimal müxiug is another effect that could be included in more complicated circulation flows. but our assuniption of no thermal wusing is a lauiting case of interest.," Thermal mixing is another effect that could be included in more complicated circulation flows, but our assumption of no thermal mixing is a limiting case of interest."
" Tf premature bubble-flow thermalization occurs. ιο inflowing gas will be heated aud the apparent enrperature profile is likely to 'econue ucgative, unlike the (siugle phase) temperature profiles observed."," If premature bubble-flow thermalization occurs, the inflowing gas will be heated and the apparent temperature profile is likely to become negative, unlike the (single phase) temperature profiles observed."
 Uniuteuded jernial qudxing Ίαν have occurred dn our ποΊσα] simulations of heated cooling flows because of the Buited munuerical resolution (Drighenti Mathews 2002: 2003)., Unintended thermal mixing may have occurred in our numerical simulations of heated cooling flows because of the limited numerical resolution (Brighenti Mathews 2002; 2003).
 When thebubbles aud iuterbubble Bow are viewed alous lio sue line ofsight. an apparent positive dT /drcan be uaintained if the filhug factor fj of the .- OVINE mbbles is relatively small and if thebubbles do not sienificantly heat the ambient eas.," When the bubbles and interbubble flow are viewed along the same line of sight, an apparent positive $dT/dr$ can be maintained if the filling factor $f_b$ of the rapidly moving bubbles is relatively small and if the bubbles do not significantly heat the ambient gas."
" We show that the COLSrait fp<1 also the heating thatbubbles can receive at the heating Tinitsradius rj and therefore also he radial extent à, of the circulation.", We show that the constraint $f_b < 1$ also limits the heating that bubbles can receive at the heating radius $r_h$ and therefore also the radial extent $r_c$ of the circulation.
" Finally, to keep our circulation models as simple as yossible. we lave resisted a temptation to include au additional pressure component within thebubbles due to (velativistic) cosmic ravs or magnetic fields."," Finally, to keep our circulation models as simple as possible, we have resisted a temptation to include an additional pressure component within the bubbles due to (relativistic) cosmic rays or magnetic fields."
" An additional ionthermal pressure could help increase the outer radius r,of the circulation. but may transport less uiass outward."," An additional nonthermal pressure could help increase the outer radius $r_c$ of the circulation, but may transport less mass outward."
 Nevertheless. may same constraints we derive here or purely thermal circulationsofthe also apply tobubbles with παbuoyaucy.," Nevertheless, many of the same constraints we derive here for purely thermal circulations also apply to bubbles with nonthermal buoyancy."
which fits one of the phases well can be used (o answer this question.,which fits one of the phases well can be used to answer this question.
 Hence we restrict ourselves to simple continuum models., Hence we restrict ourselves to simple continuum models.
" The high phase spectrum can be well modeled with a simple blackbody of temperature Ty,—0.1tra keV and a broken power-law with photon index Ty=3.60.1 below 1.9 keV and D»=1.9d above 1.9 keV (also see Table 1)).", The high phase spectrum can be well modeled with a simple blackbody of temperature $kT_{bb}= 0.1^{+0.03}_{-0.02}$ keV and a broken power-law with photon index $\Gamma_1= 3.6\pm0.1 $ below 1.9 keV and $\Gamma_2=1.9_{-0.9}^{+0.4}$ above 1.9 keV (also see Table \ref{tab:fits}) ).
 Hereafter this model will be referred to as Mhibbbkn., Hereafter this model will be referred to as $Mhi\_bb\_bkn$.
 While formally acceptable with == 125.6/141. the power-law component in this model dominates not only the high energies. but also Che low energies as shown in Fie. 3..," While formally acceptable with = 125.6/141, the power-law component in this model dominates not only the high energies, but also the low energies as shown in Fig. \ref{fig:hi}."
 Since only the peak of the blackbocdy contributes to the observed spectrum. replacing (the blackbody with accretion disc models does not change the fits significantly.," Since only the peak of the blackbody contributes to the observed spectrum, replacing the blackbody with accretion disc models does not change the fits significantly."
 The fact that the power-law dominated both low amd hish energies when using a single blackbody was also noted in a previous work bv Poundsοἱ (1995)., The fact that the power-law dominated both low and high energies when using a single blackbody was also noted in a previous work by \citet{pounds1995}.
. Following Poundsetal.(1995). we tried fitting two blackbocdies with different temperature and normalizations to fit the soft. (20.3 0.9 keV) aud verv soft. (0.3. keV) energies and a power-law for the high (220.9 keV) energies., Following \citet{pounds1995} we tried fitting two blackbodies with different temperature and normalizations to fit the soft $\sim$ 0.3–0.9 keV) and very soft $<$ 0.3 keV) energies and a power-law for the high $>$ 0.9 keV) energies.
 The best fit parameters for this model (hereafter Mhi.200 pow) are presented in Table 1.. and the model components ancl residuals are shown in Fig. 3..," The best fit parameters for this model (hereafter $Mhi\_2bb\_pow$ ) are presented in Table \ref{tab:fits}, and the model components and residuals are shown in Fig. \ref{fig:hi}. ."
 Note that the power-law model parameters D» for the model bbAhibin. aad D [or the model AMhi.256pow are largely constrained only by a lew of the hardest energve. bins. and consequently have lareee error bars.," Note that the power-law model parameters $\Gamma_2$ for the model $Mhi\_bb\_bkn$, and $\Gamma$ for the model $Mhi\_2bb\_pow$ are largely constrained only by a few of the hardest energy bins, and consequently have large error bars."
 The 0.255 keV flux in the highex phase spectrum is (1.2620.1)x10.' eres/eni? /s., The 0.25–5 keV flux in the high phase spectrum is $(1.26\pm0.1)\times10^{-11}$ $^2$ /s.
 First we test if the shape of the low phase spectrum is the same as that of the high phase., First we test if the shape of the low phase spectrum is the same as that of the high phase.
 For (his we multiply the best-fit high phase model by an overall normalization constant and search for minimum X? allowing only this constant to change., For this we multiply the best-fit high phase model by an overall normalization constant and search for minimum $\chi^2$ allowing only this constant to change.
 The best-fit normalization constant has a value in the range of 0.9040.02 for both models Mhibb.blen ancl Mhi25bpow. when applied to the low phase spectrum.," The best-fit normalization constant has a value in the range of $0.90\pm0.02$ for both models $Mhi\_bb\_bkn$ and $Mhi\_2bb\_pow$, when applied to the low phase spectrum."
 However the residuals in tliis case (see Fig. 4)), However the residuals in this case (see Fig. \ref{f:lo_renorm}) )
 show a svslemalic deviation in the ~0.851.1 keV range. stronglv suggestive of an absorption edee.," show a systematic deviation in the $\sim$ 0.85–1.1 keV range, strongly suggestive of an absorption edge."
 There may also be some systematic deviation al lower energies. especially an edge-like absorption feature between 0.30.4 keV. While not statistically significant in the current dataset. if this indeed is the case then the continuum level at these energies is higher than our estimate.," There may also be some systematic deviation at lower energies, especially an edge-like absorption feature between 0.3–0.4 keV. While not statistically significant in the current dataset, if this indeed is the case then the continuum level at these energies is higher than our estimate."
 Some of these spectral features were also noted by Middletonetal.(2009) in (he (me-averaged spectrum. but coadding spectra from both high and low phasesmost likely lowered the observability of the features (which are prominent only during the low phase).," Some of these spectral features were also noted by \citet{middletonetal2009} in the time-averaged spectrum, but coadding spectra from both high and low phasesmost likely lowered the observability of the features (which are prominent only during the low phase)."
which illustrates the validity of our general approach.,which illustrates the validity of our general approach.
" The Ha line independently measures the instantaneous SFR from young ionizing stars, whereas the SFH fits a model to the stellar continuum and is averaged over several Gyr."," The $\alpha$ line independently measures the instantaneous SFR from young ionizing stars, whereas the SFH fits a model to the stellar continuum and is averaged over several Gyr."
" It is interesting to note that given the cosmic spectrum to get a SFR from SFH fitting much lower than 0.03 rrequires high 8, high a AND low metallicity today."," It is interesting to note that given the cosmic spectrum to get a SFR from SFH fitting much lower than 0.03 requires high $\beta$ , high $\alpha$ AND a low metallicity today."
 We note that the Gallegoaetal.(1995) sample also gives a lower SFR density than the UV sample of etal. (1998).., We note that the \cite{Gallego95} sample also gives a lower SFR density than the UV sample of \cite{Treyer98}. .
" Both Treyeretal.(1998) and Tresse&Maddox(1998b) are measured at z~0.2, whereas the current sample is z~0.05 which is comparable to the redshifts of the Gallegoetal.(1995) sample."," Both \cite{Treyer98}
 and \cite{TresseMaddox98} are measured at $z\sim 0.2$, whereas the current sample is $z\sim 0.05$ which is comparable to the redshifts of the \cite{Gallego95} sample."
 It is only the latter which is discrepant., It is only the latter which is discrepant.
 The most likely explanation is that the Gallegoetal.(1995) survey covers only 1/5th our sky area and is only sensitive to star-forming galaxies as it is emission line selected.," The most likely explanation is that the \cite{Gallego95}
 survey covers only $1/5^{\rm th}$ our sky area and is only sensitive to star-forming galaxies as it is emission line selected."
 Our survey represents a ‘cosmic average’ and galaxies contribute to the final cosmic spectrum even if the Ho flux is not detectable in individual galaxies., Our survey represents a `cosmic average' and galaxies contribute to the final cosmic spectrum even if the $\alpha$ flux is not detectable in individual galaxies.
 We note that Ho is only sensitive to transient star-formation over ~20 Myr whereas a UV sample like those of Treyeretal.(1998) effectively average over longer times scales of up to ~ 1 Gyr (see for a discussion of this)., We note that $\alpha$ is only sensitive to transient star-formation over $\sim 20$ Myr whereas a UV sample like those of \cite{Treyer98} effectively average over longer times scales of up to $\sim$ 1 Gyr (see \cite{glaze99} for a discussion of this).
 A final comment on the other line ratios: the Ho/[OII] line ratio of 2.1 (region B) is entirely consistent with the median found by the sample of galaxies observed by Kennicutt(1992)., A final comment on the other line ratios: the $\alpha/$ [OII] line ratio of 2.1 (region B) is entirely consistent with the median found by the sample of galaxies observed by \cite{Kenn92b}.
" The other line ratios are also entirely consistent with those from star-forming galaxies (Veilleux&Osterbrock1987) including the weak [OI] line, indicating that the AGN contribution to the cosmic spectrum is indeed negligible."," The other line ratios are also entirely consistent with those from star-forming galaxies \citep{VeilleuxOsterbrock87}
 including the weak [OI] line, indicating that the AGN contribution to the cosmic spectrum is indeed negligible."
 It can at most be only a few percent according to the models of Kewleyetal.(2001)., It can at most be only a few percent according to the models of \cite{kewley}.
. Most of the optical light of the Universe does indeed come from stellar nucleosynthesis., Most of the optical light of the Universe does indeed come from stellar nucleosynthesis.
" Computing the light in the Universe today from the SDSSand2dFGRS surveys allows us to derive à cosmic optical spectrum, determine its robustness, and make more accurate determinations of allowable star-formation"," Computing the light in the Universe today from the SDSSand2dFGRS surveys allows us to derive a cosmic optical spectrum, determine its robustness, and make more accurate determinations of allowable star-formation"
burning of the mass loser. case B where the RLOF occurs during the hydrogen shell burning phase prior to central helium burning. and case C where the RLOF begins after helium has been depleted in the core.,"burning of the mass loser, case B where the RLOF occurs during the hydrogen shell burning phase prior to central helium burning, and case C where the RLOF begins after helium has been depleted in the core."
 Case B RLOF ts further divided into early case B or case Br where at the onset of RLOF the envelope of the mass loser is mostly radiative and late case B or case Be where the primary has a deep convective envelope at the beginning of the RLOF phase., Case B RLOF is further divided into early case B or case Br where at the onset of RLOF the envelope of the mass loser is mostly radiative and late case B or case Bc where the primary has a deep convective envelope at the beginning of the RLOF phase.
 A star that already went through a first phase of RLOF during hydrogen shell burning may fill its Roche lobe for a second time during helium shell burning and perform case BB RLOF (Delgado and Thomas. 1981).," A star that already went through a first phase of RLOF during hydrogen shell burning may fill its Roche lobe for a second time during helium shell burning and perform case BB RLOF (Delgado and Thomas, 1981)."
 Figure | summarizes the general evolutionary scenario of intermediate mass close binaries (consisting of two B-type stars)., Figure 1 summarizes the general evolutionary scenario of intermediate mass close binaries (consisting of two B-type stars).
 The original most massive star (the primary) reaches its critical Roche lobe., The original most massive star (the primary) reaches its critical Roche lobe.
 Depending on the period of the binary the RLOF is accompanied by mass transfer or the RLOF results in a common envelope phase., Depending on the period of the binary the RLOF is accompanied by mass transfer or the RLOF results in a common envelope phase.
 In most of the binaries this phase implies the loss of hydrogen rich mass that is not affected by hot bottom burning and therefore it can serve to dilute mass that was lost by single stars during the AGB phase., In most of the binaries this phase implies the loss of hydrogen rich mass that is not affected by hot bottom burning and therefore it can serve to dilute mass that was lost by single stars during the AGB phase.
 When the primary has become a WD. the further evolution of the binary is governed by the secondary star.," When the primary has become a WD, the further evolution of the binary is governed by the secondary star."
 When this secondary star reaches its critical Roche lobe. a common envelope phase sets in where most of the hydrogen rich mass of the secondary ts lost.," When this secondary star reaches its critical Roche lobe, a common envelope phase sets in where most of the hydrogen rich mass of the secondary is lost."
 Again in most of the binaries. this lost mass has not been affected by hot bottom burning and may also serve to dilute single star mass loss.," Again in most of the binaries, this lost mass has not been affected by hot bottom burning and may also serve to dilute single star mass loss."
 Below we discuss this evolution in more detail., Below we discuss this evolution in more detail.
 The Brussels binary code originates from. the Paezynsski (1965) code., The Brussels binary code originates from the Paczyńsski (1965) code.
 In the latter only the evolution of the donor (mass loser) was followed but it contained a detailed calculation of the mass transfer rate (imposing the condition that once the radius of the donor becomes larger than its Roche lobe. the mass loss rate is calculated so that the radius of the star equals the Roche radius).," In the latter only the evolution of the donor (mass loser) was followed but it contained a detailed calculation of the mass transfer rate (imposing the condition that once the radius of the donor becomes larger than its Roche lobe, the mass loss rate is calculated so that the radius of the star equals the Roche radius)."
 Of primary importance ts the fact that this code modeled the gravitational energy loss when mass leaves the star through the first Lagrangian point. energy loss that is responsible for the luminosity drop which ts typical for the evolution of the donor during its RLOF phase.," Of primary importance is the fact that this code modeled the gravitational energy loss when mass leaves the star through the first Lagrangian point, energy loss that is responsible for the luminosity drop which is typical for the evolution of the donor during its RLOF phase."
 At present our code is a twin code that follows the evolution of both components simultaneously (the code has been deseribed in detail in Vanbeveren et al..," At present our code is a twin code that follows the evolution of both components simultaneously (the code has been described in detail in Vanbeveren et al.,"
 1998a. b).," 1998a, b)."
 The opacities are taken from Iglesias et al.. (," The opacities are taken from Iglesias et al., ("
1992). the nuclear reaction rates from Fowler et al. (,"1992), the nuclear reaction rates from Fowler et al. ("
1975).,1975).
 Semi-convection ts treated according to the criterion of Schwarzschild and Harm (1958) and convective core overshooting ts included as described by Schaller et al. (, Semi-convection is treated according to the criterion of Schwarzschild and Harm (1958) and convective core overshooting is included as described by Schaller et al. (
1992).,1992).
 The twin code follows the evolution of the mass gainer and therefore an accretion model is essential., The twin code follows the evolution of the mass gainer and therefore an accretion model is essential.
 When the period of the binary is small enough so that the gas stream leaving the first Lagrangian point hits the mass gainer directly. the mass gain process is treated using the formalism of Neo et al. (," When the period of the binary is small enough so that the gas stream leaving the first Lagrangian point hits the mass gainer directly, the mass gain process is treated using the formalism of Neo et al. ("
1977) (we call this the standard aceretion model).,1977) (we call this the standard accretion model).
 During its RLOF. the mass loser may lose mass which has been nuclearly processed and which has a molecular weight that is larger than the molecular weight of the outer layers of the gainer.," During its RLOF, the mass loser may lose mass which has been nuclearly processed and which has a molecular weight that is larger than the molecular weight of the outer layers of the gainer."
 The aceretion of this mass initiates mixing. a process commonly known as thermohaline convection (Kippenhahn et al..," The accretion of this mass initiates mixing, a process commonly known as thermohaline convection (Kippenhahn et al.,"
 1980)., 1980).
 In our code we treat this process as an instantaneous one., In our code we treat this process as an instantaneous one.
 Note that the periods of most of the Algol binaries indicate that the mass transfer process happened and happens by direct hit and therefore it is conceivable that the standard mass gain process applies 1n most of these binaries., Note that the periods of most of the Algol binaries indicate that the mass transfer process happened and happens by direct hit and therefore it is conceivable that the standard mass gain process applies in most of these binaries.
 Mass transfer implies angular momentum transfer and the mass gainer spins up., Mass transfer implies angular momentum transfer and the mass gainer spins up.
 When mass transfer proceeds via a Keplerian disk it was shown by Packet (1981) that very soon after the onset of mass transfer. the mass gainer acquires the critical rotation velocity and therefore its evolution may be affected by rotational mixing.," When mass transfer proceeds via a Keplerian disk it was shown by Packet (1981) that very soon after the onset of mass transfer, the mass gainer acquires the critical rotation velocity and therefore its evolution may be affected by rotational mixing."
 To simulate this mixing process. Vanbeveren et al. (," To simulate this mixing process, Vanbeveren et al. ("
1994) and Vanbeveren and De Loore (1994) introduced the accretion induced full mixing model.,1994) and Vanbeveren and De Loore (1994) introduced the accretion induced full mixing model.
 In this way they were able to explain the helium and mass discrepancy in Vela Χ-]., In this way they were able to explain the helium and mass discrepancy in Vela X-1.
 The more sophisticated mass gainer models of Cantiello et al. (, The more sophisticated mass gainer models of Cantiello et al. (
2007) demonstrate that the simplified models are not too bad.,2007) demonstrate that the simplified models are not too bad.
 Remark however that rotational mixing ts still heavily debated and current evolutionary models of single stars including rotational mixing do not seem to explain the observations (Hunter et al..," Remark however that rotational mixing is still heavily debated and current evolutionary models of single stars including rotational mixing do not seem to explain the observations (Hunter et al.,"
 2008)., 2008).
 In our code we can switch off this process (e.g.. mass accretion Is treated in the standard way in all the case A/Br binaries).," In our code we can switch off this process (e.g., mass accretion is treated in the standard way in all the case A/Br binaries)."
 An interesting conclusion resulting from all. caseA/Br intermediate mass close binary evolutionary computations done in the past is that the overall evolution of the mass loser is largely independent from the details of the RLOF process. from the initial. period of the binary. from the mass of its companion and from the initial metallicity.," An interesting conclusion resulting from all caseA/Br intermediate mass close binary evolutionary computations done in the past is that the overall evolution of the mass loser is largely independent from the details of the RLOF process, from the initial period of the binary, from the mass of its companion and from the initial metallicity."
 The RLOF stops when both stars merge (see section 2.6) or in case A binaries when the mass loser reaches the main sequence hook (end of core hydrogen burning) or in case Br binaries when the loser has lost most of its hydrogen rich layers and helium starts burning in the core., The RLOF stops when both stars merge (see section 2.6) or in case A binaries when the mass loser reaches the main sequence hook (end of core hydrogen burning) or in case Br binaries when the loser has lost most of its hydrogen rich layers and helium starts burning in the core.
 The evolution of the mass gainer however. depends on the details of the RLOF. on the adopted aceretion model and on the amount of accreted mass.," The evolution of the mass gainer however, depends on the details of the RLOF, on the adopted accretion model and on the amount of accreted mass."
 Moreover. the simultaneous evolutionary computations of caseA/Br binaries reveal that in most of them both stars come into contact during ΚΙΟΕ.," Moreover, the simultaneous evolutionary computations of caseA/Br binaries reveal that in most of them both stars come into contact during RLOF."
 It is not unreasonable to assume that mass will leave the binary when a contact system is formed but since the physics of contact is poorly constrained it is uncertain how much mass will leave the binary., It is not unreasonable to assume that mass will leave the binary when a contact system is formed but since the physics of contact is poorly constrained it is uncertain how much mass will leave the binary.
 This uncertainty may imply an important uncertainty for all applications of intermediate mass close binary evolution., This uncertainty may imply an important uncertainty for all applications of intermediate mass close binary evolution.
 For this reason De Loore and Vanbeveren (1995) and Vanbeveren et al. (, For this reason De Loore and Vanbeveren (1995) and Vanbeveren et al. (
1998) calculated several thousands of evolutionary tracks of mass gainers for different initial chemical compositions. for different aceretion efficiency characterized by B=1/ΜΙ|. which ts the fraction,"1998) calculated several thousands of evolutionary tracks of mass gainers for different initial chemical compositions, for different accretion efficiency characterized by $\beta = \left|\dot{M_{2}} / \dot{M_{1}}\right|$, which is the fraction"
"σημα required for a detection as follows: Tere the factor 9&1.1 reflects losses due to hardware Iuuitations. yi,= Sis the threshold signal-to-noise ratio. Έως aud Tua ave the receiver and sky noise temperatures OX). G is the effective antenna gain (x JvLy, NV,=21s the nuuber of polarizations sumuued. Av=128«60 kIIz =7.68 MIIz is the total observing bandwidth. μι=1380 sis the inteeration tine. Wis the detected pulse width and P is the pulse period.","$S_{\rm min}$ required for a detection as follows: Here the factor $\beta \simeq 1.1$ reflects losses due to hardware limitations, $\sigma_{\rm min}=8$ is the threshold signal-to-noise ratio, $T_{\rm rec}$ and $T_{\rm sky}$ are the receiver and sky noise temperatures (K), $G$ is the effective antenna gain (K $^{-1}$ ), $N_p=2$ is the number of polarizations summed, $\Delta
\nu=128\times60$ kHz $=7.68$ MHz is the total observing bandwidth, $t_{\rm int}=1380$ s is the integration time, $W$ is the detected pulse width and $P$ is the pulse period."
 For our 130-MIITz search we can iusert the above values along with Ty.=100 K. Τις= OK. C=10 K J«! to find μμiso&1.507Py’? iid.," For our 430-MHz search we can insert the above values along with $T_{\rm rec}=100$ K, $T_{\rm
sky}=150$ K, $G=10$ K $^{-1}$ to find $S_{\rm min,430} \simeq 1.5
(W/P)^{1/2}$ mJy."
 For the ΠΟΝΠΙΣ search. Teo=35 K. Ta.2 TK. ο=8 WK | so that Syuuano—ΙΕ nds. ," For the 1410-MHz search, $T_{\rm rec}=35$ K, $T_{\rm sky}\simeq7$ K, $G=8$ K $^{-1}$ so that $S_{\rm min,1410}
\simeq 0.3 (W/P)^{1/2}$ mJy."
Iu both these cases. we have assunied that V.«P.," In both these cases, we have assumed that $W \ll P$."
 Note also that both the £30-MITz aud 1110-AIlIz observations were carried out at necessarily high zenith angeles (2 107) x) that the quoted antenna gains are lower than those applicable to observations closer to the zenith (see http:www.naic.edu/aomenunu.hltui for up-to-date telescope information)., Note also that both the 430-MHz and 1410-MHz observations were carried out at necessarily high zenith angles $>10^{\circ}$ ) so that the quoted antenna gains are lower than those applicable to observations closer to the zenith (see http://www.naic.edu/aomenu.htm for up-to-date telescope information).
 These search observations were carried out as part of a larger search for pulsars in supcruova roenimants at Arecibo., These search observations were carried out as part of a larger search for pulsars in supernova remnants at Arecibo.
 As part of this project. we undertook: a umber of calibration observations with the PSPAL for pulsars with well-known flux densitics and pulse widths.," As part of this project, we undertook a number of calibration observations with the PSPM for pulsars with well-known flux densities and pulse widths."
 The sigual-to-noise ratios predicted from equation L compare well with those obtained iu practice and eive us confidence im the above flux limits obtained from our blind search., The signal-to-noise ratios predicted from equation \ref{equ:smin} compare well with those obtained in practice and give us confidence in the above flux limits obtained from our blind search.
 Qur search was carried out before the aunounceioeut of the 5.16-s periodicity bv Turley ct al. (, Our search was carried out before the announcement of the 5.16-s periodicity by Hurley et al. (
1998).,1998).
 Following this discovery. aud Shitows (1999) detection. at 111 AMIIz. we de-dispersed aud folded both sets of search data at the predicted topoceutric period to look for evideuce of a pulsed. signal.," Following this discovery, and Shitov's (1999) \nocite{shi99} detection at 111 MHz, we de-dispersed and folded both sets of search data at the predicted topocentric period to look for evidence of a pulsed signal."
 We found uo convincing evidence for pulsar-like profiles above a signal-to-noise threshold of 6., We found no convincing evidence for pulsar-like profiles above a signal-to-noise threshold of 6.
" This cffectively reduces the above limits from the bliud periodicity searches by a factor of 6/8 to Sui)&LAVPy? uty and Sqg,4nocOUTPy? indy."," This effectively reduces the above limits from the blind periodicity searches by a factor of 6/8 to $S_{\rm min,430} \simeq 1.1
(W/P)^{1/2}$ mJy and $S_{\rm min,1410} \simeq 0.2 (W/P)^{1/2}$ mJy."
 Shitov. Pugachev Iutuzov (2000) report a pulse width of 100 as for the 5.16-s radio pulsations from SGR 1900|L1 observed at 111 MIIz.," Shitov, Pugachev Kutuzov (2000) \nocite{spk00} report a pulse width of 100 ms for the 5.16-s radio pulsations from SGR 1900+14 observed at 111 MHz."
 Based on the above sensitivitv estinatious. such a pulsed signal would be detectable down to a flux-deusity limit of 150 pJy at. 130 Mz and 30 pJy at 1110 MITzZ.," Based on the above sensitivity estimations, such a pulsed signal would be detectable down to a flux-density limit of 150 $\mu$ Jy at 430 MHz and 30 $\mu$ Jy at 1410 MHz."
 We note that these upper liuits. together with Shitov ct al.," We note that these upper limits, together with Shitov et al."
« flux measurenieut of SCR 1900]11 as à 503uJyv. radio pulsar at lll MIIz. constrain the power-law index of the radio spectrum to be steeper than 3.,"'s flux measurement of SGR 1900+14 as a 50-mJy radio pulsar at 111 MHz, constrain the power-law index of the radio spectrum to be steeper than –3."
 Clearhy. further radio observations at lower frequencies are required to confini or refute the AIIIz detection reported by Shitoy et.," Clearly, further radio observations at lower frequencies are required to confirm or refute the 111-MHz detection reported by Shitov et."
 al., al.
 Diving the analysis of our L1110-MIIz search observations towards SCR L900)1b we found a very pronusing τοις pulsar candidate., During the analysis of our 1410-MHz search observations towards SGR 1900+14 we found a very promising 113-ms pulsar candidate.
 Subsequent 1110-ΑΠΣ observations mace iu September 1998 confirmed the existence of the pulsar (J1907|0918). and identified its true period to be 226 iuis. (Nilouris ct al., Subsequent 1410-MHz observations made in September 1998 confirmed the existence of the pulsar (J1907+0918) and identified its true period to be 226 ms \nocite{xkl+98} (Xilouris et al.
 1998)., 1998).
 Iun order to accurately determine the spin aud astrometric properties of PSR J1907|0918. we have becu carving out reeular timine measureients using the PSPAL since October 1998.," In order to accurately determine the spin and astrometric properties of PSR J1907+0918, we have been carrying out regular timing measurements using the PSPM since October 1998."
 A preliminary ephemeris. based ou he discovery aud confirmation observations. was used to xediet the apparent pulse period initially.," A preliminary ephemeris, based on the discovery and confirmation observations, was used to predict the apparent pulse period initially."
 This period was sutficicutly accurate to be used by the PSPM in timing uode to fold the icoming data from each of the 128 yequency chanucls for typically 90 s before saving the xofiles to disk., This period was sufficiently accurate to be used by the PSPM in timing mode to fold the incoming data from each of the 128 frequency channels for typically 90 s before saving the profiles to disk.
 The individual frequency channels were ater de-dispersed to produce a time-tageed integrated oilse profile for cach 90-8 observation., The individual frequency channels were later de-dispersed to produce a time-tagged integrated pulse profile for each 90-s observation.
 For cach of these scans. the mean pulse time of arrival (TOA) at the observatory was then determined by cross correlating the xofile with a lieh signal-to-noise template (for further details. see Tavlor 1992).," For each of these scans, the mean pulse time of arrival (TOA) at the observatory was then determined by cross correlating the profile with a high signal-to-noise template (for further details, see Taylor \nocite{tay92} 1992)."
 The software package (which is freely. available at http://pulsar.princeton.edu/teuipo) aud the DE200 planetary ephenmieris were then used to correct these topoceutric ΤΟΑς for the Eartlis notion around the Sun and transform them to the frame of the solar svsteni barveenter before fitting thei to a simple isolated pulsar spin-down model — (see c.g. Manchester Taylor 1977)., The software package (which is freely available at http://pulsar.princeton.edu/tempo) and the DE200 planetary ephemeris were then used to correct these topocentric TOAs for the Earth's motion around the Sun and transform them to the frame of the solar system barycenter before fitting them to a simple isolated pulsar spin-down model \nocite{mt77} (see e.g. Manchester Taylor 1977).
 Multiple passes ofthe TOAs through were required in order to minimize the model minus observed timing residuals and resolve auy pulse nunmberimg ambiguities., Multiple passes of the TOAs through were required in order to minimize the model minus observed timing residuals and resolve any pulse numbering ambiguities.
 The resulting ephemeris based ou observations spaunuiug a 15anouth baseline between October 1998 aud Jauuuv 2000 vields a sub-aresecoud position and highly accurate spin-down parameters which are presented in Table 1.., The resulting ephemeris based on observations spanning a 15-month baseline between October 1998 and January 2000 yields a sub-arcsecond position and highly accurate spin-down parameters which are presented in Table \ref{tab:1907}.
 The post-it timing residuals for this epliemeris are free from systematic trends at the level of 108 js ruis., The post-fit timing residuals for this ephemeris are free from systematic trends at the level of 108 $\mu$ s rms.
 >rueclnai From the spin paraicters of PSR J1907|0918. we infer a dipole surface magnetic field. strength. of L7«107 C and a characteristic age of only 38 kyr.," 7truecm From the spin parameters of PSR J1907+0918, we infer a dipole surface magnetic field strength of $4.7\times10^{12}$ G and a characteristic age of only 38 kyr."
 This low characteristic age places J1907]0915 iu the vouugest of known Galactic radio pulsars., This low characteristic age places J1907+0918 in the youngest of known Galactic radio pulsars.
 Mauv pulsars with simular characteristic ages have relatively flat radio spectra (Johnston 1990: Lorimer et al., Many pulsars with similar characteristic ages have relatively flat radio spectra (Johnston 1990; \nocite{joh90} Lorimer et al.
 1995) and are strongly polarized sources at radio frequencies above 1 CdIz (vou IIoecusbroech 1999)., \nocite{lylg95} 1995) and are strongly polarized sources at radio frequencies above 1 GHz \nocite{hoe99} (von Hoensbroech 1999).
 Iun order to investigate the onüssion properties of PSR J1907]0918. in October aud November 1908 we carried out a series of quasi-sinultaucous nulti-frequeucyv observations at Arecibo usine the 130-MIIz line feed. along with the Ciegorian receivers centered at 11410. 2380 and 5000. MIIz.," In order to investigate the emission properties of PSR J1907+0918, in October and November 1998 we carried out a series of quasi-simultaneous multi-frequency observations at Arecibo using the 430-MHz line feed, along with the Gregorian receivers centered at 1410, 2380 and 5000 MHz."
 We used the Arecilho-Berkeley Pulsar Processor (ABPP) for these observations., We used the Arecibo-Berkeley Pulsar Processor (ABPP) for these observations.
 The ADPP is a 32-channel colhereut-cispersion-removal machine capable of lueh-precision timing aud polarimetry (for details sec Backer ct al., The ABPP is a 32-channel coherent-dispersion-removal machine capable of high-precision timing and polarimetry (for details see \nocite{bdz+97} Backer et al.
 1997 and 7 abpp) which allowed us to carry out high-quality observations spanning a bandwidth of up to 55 MIIz., 1997 and $^{\sim}$ abpp) which allowed us to carry out high-quality observations spanning a bandwidth of up to 35 MHz.
We now perform the same search for the observed QSO spectra.,We now perform the same search for the observed QSO spectra.
" The results are shown in Figure[3], which shows the distribution of the flux void population extracted from the observations (points with Poissonian error bars) and from our largest volume simulation."," The results are shown in Figure \ref{fig2}, which shows the distribution of the flux void population extracted from the observations (points with Poissonian error bars) and from our largest volume simulation."
" From the Figure it is clear that the fiducial (120,400) run is in excellent agreement with observations and that flux voids of sizes 35h~' MMpc are about 1000 times less common than voids of sizes larger than few h! MMpc."," From the Figure it is clear that the fiducial $(120,400)$ run is in excellent agreement with observations and that flux voids of sizes $> 35\,h^{-1}$ Mpc are about 1000 times less common than voids of sizes larger than few $h^{-1}$ Mpc."
" We also note a bump at around 20 Mpc/h in the voids fraction, but one would need a larger sample of QSOs in order to better test its statistical significance."," We also note a bump at around $20$ $h$ in the voids fraction, but one would need a larger sample of QSOs in order to better test its statistical significance."
 The continuous lines show how the void fraction changes when the effective optical depth is varied at a +30 level around the observed value., The continuous lines show how the void fraction changes when the effective optical depth is varied at a $\pm 3\sigma$ level around the observed value.
" In particular, the lower and higher tog values (reg=0.136 and Tes=0.19, respectively) have been chosen in such a way to conservatively embrace at a confidence level of 3c the values obtained in ?,, where the measured value for the effective optical depth was Tes=0.163+ 0.009."," In particular, the lower and higher $\tau_{\rm eff}$ values $\tau_{\rm eff}=0.136$ and $\tau_{\rm eff}=0.19$, respectively) have been chosen in such a way to conservatively embrace at a confidence level of $3\sigma$ the values obtained in \cite{viel04}, where the measured value for the effective optical depth was $\tau_{\rm
 eff}=0.163\pm0.009$ ."
" The recently determined power-law fit of ? το=(0.0023--0.0007)2)5:5959-?! by using QSO spectra in the range 1.7<z(14-4, is influenced by the scatter coming from QSO spectra over a large range of redshift, and would in principle allow for a larger range of effective optical depths."," The recently determined power-law fit of \cite{tkim} $\tau_{\rm eff}=(0.0023\pm 0.0007) (1+z)^{3.65\pm 0.21}$ by using QSO spectra in the range $1.7<z<4$, is influenced by the scatter coming from QSO spectra over a large range of redshift, and would in principle allow for a larger range of effective optical depths."
" However, by selecting only a subsample at a given redshift, smaller statistical errors on the effective optical depth can be obtained."," However, by selecting only a subsample at a given redshift, smaller statistical errors on the effective optical depth can be obtained."
 We stress that the definition of void used here is based on the transmitted flux and thus is very different from studies based on the distribution of voids as inferred from the galaxy distribution., We stress that the definition of void used here is based on the transmitted flux and thus is very different from studies based on the distribution of voids as inferred from the galaxy distribution.
 We now discuss the main cosmological and astrophysical parameters that can affect the population of voids in the flux distribution., We now discuss the main cosmological and astrophysical parameters that can affect the population of voids in the flux distribution.
" In Figure we plot the ratios of the flux void fractions in the fiducial[4], and in other simulations, where we changed το at z=2.2 within the observational bounds and the value of og."," In Figure \ref{fig3}, we plot the ratios of the flux void fractions in the fiducial and in other simulations, where we changed $\tau_{\rm
 eff}$ at $z=2.2$ within the observational bounds and the value of $\sigma_8$."
" As could be expected, the fraction of voids changes dramatically with a change in Teg: at z~2 the flux probability distribution function (pdf) is a steep function around the mean flux level."," As could be expected, the fraction of voids changes dramatically with a change in $\tau_{\rm eff}$: at $z\sim 2$ the flux probability distribution function (pdf) is a steep function around the mean flux level."
 Thus adopting a different void selection criterion has a large impact on the void distribution., Thus adopting a different void selection criterion has a large impact on the void distribution.
" In the following, we normalize the mock QSO sample to reproduce the same effective optical depth."," In the following, we normalize the mock QSO sample to reproduce the same effective optical depth."
 Changing og appears to have a smaller effect., Changing $\sigma_8$ appears to have a smaller effect.
" A simulation with og=1 results in a faster evolution of cosmic structures and in a 50% increase of the void population of size R=20—35h! MMpc, compared with the reference og=0.85 case."," A simulation with $\sigma_8=1$ results in a faster evolution of cosmic structures and in a $\%$ increase of the void population of size $R=20-35\,h^{-1}$ Mpc, compared with the reference $\sigma_8=0.85$ case."
" In the simulation with a lower value of σε, the void sizes are more evenly distributed."," In the simulation with a lower value of $\sigma_8$, the void sizes are more evenly distributed."
" Note, however, that the trend at the largest sizes MMpc/h) is somewhat different: here the low og simulation predicts more voids than the higher og, since the smaller amount of power in the former can occasionally produce large regions devoid of matter and absorption."," Note, however, that the trend at the largest sizes $/h$ ) is somewhat different: here the low $\sigma_8$ simulation predicts more voids than the higher $\sigma_8$ , since the smaller amount of power in the former can occasionally produce large regions devoid of matter and absorption."
" We also checked the impact of varying other parameters such as Qm, Ho and a change in the shape of the linear dark matter power spectrum in the initial conditions to account for a presence of warm dark matter (WDM) and adding some extra power (EP) at intermediate (from Mpc to tens of Mpc) scales, that results in an effect opposite to that of WDM."," We also checked the impact of varying other parameters such as $\Omega_m$, $H_0$ and a change in the shape of the linear dark matter power spectrum in the initial conditions to account for a presence of warm dark matter (WDM) and adding some extra power (EP) at intermediate (from Mpc to tens of Mpc) scales, that results in an effect opposite to that of WDM."
" The mass chosen for the WDM particle is 0.15 keV. A mass that is already ruled out from the flux power spectrum of the forest (see ?)) and that produces a suppression of power at scales around 0.4 Mpc/h, which are marginally sampled in the initial conditions of our (120,400) simulation."," The mass chosen for the WDM particle is 0.15 keV. A mass that is already ruled out from the flux power spectrum of the forest (see \cite{viel08}) ) and that produces a suppression of power at scales around 0.4 $/h$, which are marginally sampled in the initial conditions of our (120,400) simulation."
" However, for the qualitative purposes of this paper this simulation should show the impact of a WDM particle on the distribution of the largest voids in the flux."," However, for the qualitative purposes of this paper this simulation should show the impact of a WDM particle on the distribution of the largest voids in the flux."
" At the scale of 2.5 Mpc/h the power is 100 times higher (smaller) in the EP (WDM) simulation than the corresponding (120,400)."," At the scale of 2.5 $/h$ the power is 100 times higher (smaller) in the EP (WDM) simulation than the corresponding (120,400)."
" The results are plotted in Figure B], where all the simulations have the same value of og—0.85: it is clear that the WDM and EP runs have opposite effects."," The results are plotted in Figure \ref{fig4}, where all the simulations have the same value of $\sigma_8=0.85$: it is clear that the WDM and EP runs have opposite effects."
" WDM significantly enhance by a factor larger than 3 the fraction of largest voids, by erasing substructures below a given scale, while the EP simulation presents more collapsed structures that suppress the number of the largest voids."," WDM significantly enhance by a factor larger than 3 the fraction of largest voids, by erasing substructures below a given scale, while the EP simulation presents more collapsed structures that suppress the number of the largest voids."
 The effect of a larger value of Qm=0.4 is similar to that of the EP simulation and determines a reduction in the fraction voids whose sizes are larger than 30 Mpc/h., The effect of a larger value of $\Omega_m=0.4$ is similar to that of the EP simulation and determines a reduction in the fraction voids whose sizes are larger than 30 $/h$.
" Finally, we checked for the effect of a smaller value of Ho—45 km/s/Mpc a value proposed recently by ? to be the true value of the Hubble constant outside a local underdense region of very similar density to those explored here."," Finally, we checked for the effect of a smaller value of $H_0=45$ km/s/Mpc a value proposed recently by \cite{stephon} to be the true value of the Hubble constant outside a local underdense region of very similar density to those explored here."
 Having a smaller value for theHubble constant produces effects that are very similar to those of a smaller value of og=0.7(see Fig. [])):, Having a smaller value for theHubble constant produces effects that are very similar to those of a smaller value of $\sigma_8=0.7$(see Fig. \ref{fig3}) ):
 the evolution of the cosmic web is less pronounced and can occasionally produce, the evolution of the cosmic web is less pronounced and can occasionally produce
ittle dependence. on the amplitude of dark matter luctuations.,little dependence on the amplitude of dark matter fluctuations.
 and are therefore rather succesful in deblending absorption features into thermally ooadened Voigt profiles., and are therefore rather succesful in deblending absorption features into thermally broadened Voigt profiles.
 Individual Voigt) components often trace the substructure in broader density. peaks., Individual Voigt components often trace the substructure in broader density peaks.
 This is the reason why the Doppler parameter distribution shows ittle sensitivity to the overall width of the density. peaks. which themselves depend strongly on the amplitude of dark matter [uetuations.," This is the reason why the Doppler parameter distribution shows little sensitivity to the overall width of the density peaks, which themselves depend strongly on the amplitude of dark matter fluctuations."
 The dux two-point function. anc the flux power spectrum on small scales are also mainty determined: by thermal broadening and show likewise no strong dependence on the amplitude of dark matter fluctuations., The flux two-point function and the flux power spectrum on small scales are also mainly determined by thermal broadening and show likewise no strong dependence on the amplitude of dark matter fluctuations.
 This is at least partially due to the fact that we have chosen to scale our simulated spectra to match the mean observed Dux., This is at least partially due to the fact that we have chosen to scale our simulated spectra to match the mean observed flux.
 The one-point function of the Dux. however. depends on temperature and normalization of the DAL Iluctuation. amplitude in a similar wav.," The one-point function of the flux, however, depends on temperature and normalization of the DM fluctuation amplitude in a similar way."
 This degeneracy arises because the PDE of the density depends mainly on the fluctuation amplitude of the gas. whieh itself. depends on both the fluctuation amplitude of the DM density and on the scale on which the eas distribution is smoothed relative to the DAI clistribution.," This degeneracy arises because the PDF of the density depends mainly on the fluctuation amplitude of the gas, which itself depends on both the fluctuation amplitude of the DM density and on the scale on which the gas distribution is smoothed relative to the DM distribution."
 We conclude that the Doppler parameter clistribution. the [ux two-point correlation function and the flux power spectrum on small scales are sensitive probes of. the temperature of the photoionized LGAL. while the one-point function of the Hux can be used to probe the Iuctuation amplitude on small scales once the temperature is known.," We conclude that the Doppler parameter distribution, the flux two-point correlation function and the flux power spectrum on small scales are sensitive probes of the temperature of the photoionized IGM, while the one-point function of the flux can be used to probe the fluctuation amplitude on small scales once the temperature is known."
" When comparing our simulated. spectra with a high resolution spectrum of QSO 1422|231. obtained with the LURES spectrograph on the Ixeck telescope. we find that the Doppler parameter distribution. the two-point correlation ""unction and the flux. power spectrum. are best matched » models that have a temperature of the LGAL of Z 150001 at mean cdensitv at z= 3.25."," When comparing our simulated spectra with a high resolution spectrum of QSO 1422+231, obtained with the HIRES spectrograph on the Keck telescope, we find that the Doppler parameter distribution, the two-point correlation function and the flux power spectrum are best matched by models that have a temperature of the IGM of $\ga$ 15000K at mean density at $z=3.25$ ."
" ""Fhis temperature is  50 per cent higher than in our reference. model L1 hat used the heating rates calculated. lor the Laardt Alacdau spectrum. (", This temperature is $\sim$ 50 per cent higher than in our reference model L1 that used the heating rates calculated for the Haardt Madau spectrum. (
Note however. that in this paper we rave only varied the amplitude of the ppT. relation. and not its slope.,"Note however, that in this paper we have only varied the amplitude of the $\rho-T$ relation, and not its slope."
 Εις prevents us from making stronger statements about. the value of 75.), This prevents us from making stronger statements about the value of $T_0$ .)
 The flux. one-point ‘unction depends both on temperature and the level of small scale fluctuations., The flux one-point function depends both on temperature and the level of small scale fluctuations.
 A ACDAL model with this temperature matches the observed: one-point function reasonably well., A $\Lambda$ CDM model with this temperature matches the observed one-point function reasonably well.
 Once an accurate measurement of Zu is available. it should be possible to determine ex accurately from the forest.," Once an accurate measurement of $T_0$ is available, it should be possible to determine $\sigma_8$ accurately from the forest."
 We thank M. Rauch ancl W. Sargent for providing us with the LURES spectrum of the QSO 1422|231., We thank M. Rauch and W. Sargent for providing us with the HIRES spectrum of the QSO 1422+231.
 JS thanks the Isaac Newton Trust. St. John's College and PPARC for support.," JS thanks the Isaac Newton Trust, St. John's College and PPARC for support."
 We thank It. Carswell for helping us withVPELFL., We thank R. Carswell for helping us with.
. This work has been supported by the Formation ancl Evolution of network set. up by the European Commission under contract ERB of its TMIU programme., This work has been supported by the Formation and Evolution of network set up by the European Commission under contract ERB of its TMR programme.
 Research conducted. in cooperation with Silicon Ciraphies/Cray. Research utilisingthe Origin 2000 supercomputer at DAATPP. Cambridge.," Research conducted in cooperation with Silicon Graphics/Cray Research utilisingthe Origin 2000 supercomputer at DAMTP, Cambridge."
1—1045TAL...,$10^8 - 10^{11}\himsun$.
" Phe G6 run covers the even wider nass range Ada,=107ΤΝΤ. depending on the value of extinction.", The G6 run covers the even wider mass range $M_{\rm star}=10^8 - 10^{12}\himsun$ depending on the value of extinction.
 As the wind strength is increased from O3 to J3. and then to Q3-run. we observe that the distribution of galaxies becomes slightly sparser at the most. luminous (ic. massive) end of the distribution.," As the wind strength is increased from O3 to P3, and then to Q3-run, we observe that the distribution of galaxies becomes slightly sparser at the most luminous (i.e. massive) end of the distribution."
 We see the same cllect more clearly in the luminosity functions. discussed in Section 6..., We see the same effect more clearly in the luminosity functions discussed in Section \ref{section:lf}.
 Note that since the numerical resolution is identical for O3. P3. and. Q3. the total barvonic mass in the simulation box is the same for these runs. but the runs with stronger winds convert less of their barvonic mass Components Lato stars.," Note that since the numerical resolution is identical for O3, P3, and Q3, the total baryonic mass in the simulation box is the same for these runs, but the runs with stronger winds convert less of their baryonic mass components into stars."
 1n simulations with larger box size (D- and C-serics). he masses of the most luminous galaxies found are substantially larger compared with those in the Q-scrics.," In simulations with larger box size (D- and G-series), the masses of the most luminous galaxies found are substantially larger compared with those in the Q-series."
 This is primarily a result of a finite box size ellect: the oxiehtest galaxies have a very. low space-density. such that hey are simply not found in simulations with too small a volume.," This is primarily a result of a finite box size effect; the brightest galaxies have a very low space-density, such that they are simply not found in simulations with too small a volume."
 On the other hand. the mass resolution of the small ον size simulations is much better. so that they can resolve a much larger number of dim low-mass galaxies.," On the other hand, the mass resolution of the small box size simulations is much better, so that they can resolve a much larger number of dim low-mass galaxies."
 Note that he lower cut-olf of the distribution in each panel is given by he stellar particle mass of the corresponding simulation., Note that the lower cut-off of the distribution in each panel is given by the stellar particle mass of the corresponding simulation.
 Interestingly. we find that the amount of stellar mass contained in the LBs that satisfv. the colour selection criteria is in general far from being close to the total stellar mass.," Interestingly, we find that the amount of stellar mass contained in the LBGs that satisfy the colour selection criteria is in general far from being close to the total stellar mass."
" In the “D5-run. for example. the LBGs that pass the criteria contain only of all the stars for an extinction of A(RV)20,40."," In the `D5'-run, for example, the LBGs that pass the criteria contain only of all the stars for an extinction of $E(B-V)=0.0$."
 In the C5-run. the fraction is and for E(D1)=0.0 and 0.15. respectively.," In the `G5'-run, the fraction is and for $E(B-V)=0.0$ and 0.15, respectively."
 In the G6run. 1)20.0 and 0.15. respectively.," In the `G6'-run, the fraction is and for $E(B-V)=0.0$ and 0.15, respectively."
 This means that the current survevs of LBCs fail to account for more than half of the stellar mass in the Universe., This means that the current surveys of LBGs fail to account for more than half of the stellar mass in the Universe.
 Nagamineetal.(2004) reach a similar conclusion [rom cdillerent. theoretical arguments by comparing the results of two cillerent types of hydrodynamic simulations and the theoretical model of Llernquist&Springel(2003) with near-H observations of galaxies (c.g.Dickinsonctal.2003:Itucdnicketal. 2003).," \citet{Nachos} reach a similar conclusion from different theoretical arguments by comparing the results of two different types of hydrodynamic simulations and the theoretical model of \citet{Her03} with near-IR observations of galaxies \citep[e.g.][]{Dickinson, Rudnick03}."
. Such missed stellar masses might be hidden in the red. population of galaxies as suggested by Fransetal.(2003).. but in our simulations the missed stellar mass fractions are mainlv contained in the faint galaxies below the observational magnitude Limit.," Such missed stellar masses might be hidden in the red population of galaxies as suggested by \citet{Franx}, but in our simulations the missed stellar mass fractions are mainly contained in the faint galaxies below the observational magnitude limit."
 In Figure 5.. we show the rest-frame V-band. luminosity functions for all the simulated galaxies at z=3.," In Figure \ref{lf_V.eps}, we show the rest-frame V-band luminosity functions for all the simulated galaxies at $z=3$."
 We have also tried restricting the luminosity functions to only those ealaxies that satisfy the colour-selection criteria., We have also tried restricting the luminosity functions to only those galaxies that satisfy the colour-selection criteria.
 Most of the ealaxies that fall out of the colour-selection criteria are the less luminous ones as demonstrated in Figure 3.., Most of the galaxies that fall out of the colour-selection criteria are the less luminous ones as demonstrated in Figure \ref{Rmag_GRcol.eps}.
 “Phis only allects the faint-end of the luminosity function not. probed bv current LBC surveys., This only affects the faint-end of the luminosity function not probed by current LBG surveys.
 Therefore we show the laninosity funetions for all galaxies in our simulations in Figures 5.. 6.. 7 Lor clarity.," Therefore we show the luminosity functions for all galaxies in our simulations in Figures \ref{lf_V.eps}, , \ref{lf_wind.eps}, \ref{lf_Rmag_all.eps} for clarity."
 As before. we plot results for three dilferent values of extinction: {01)=00 (blue long-dashed). 0.15 (green solid). ancl 0.3. (red.dot-dashed).," As before, we plot results for three different values of extinction: $E(B-V)=0.0$ (blue long-dashed), 0.15 (green solid), and 0.3 (reddot-dashed)."
 Note that stronger extinction simply shifts the luminosity function to, Note that stronger extinction simply shifts the luminosity function to
The spectral type and. the high rotational velocity xesented by LER 7355 (esiné = 300 £15 kms +) hindered he use of silicon lines (e.g. Si11 412.5-413.1. nm and Si 455.3-457.5 nm) to determine Z;r.,"The spectral type and the high rotational velocity presented by HR 7355 $v \sin i $ = 300 $\pm$ 15 km $^{-1}$ ) hindered the use of silicon lines (e.g., Si 412.8-413.1 nm and Si 455.3-457.5 nm) to determine $T_{\rm{eff}}$."
 Alternatively. the intensity of the Balmer and He transitions were found o be good diagnostics (see for example. Dutton οἱ al.," Alternatively, the intensity of the Balmer and He transitions were found to be good diagnostics (see for example, Dufton et al."
 2006)., 2006).
 For the determination of logg. we used mainly Ils.," For the determination of $\log g$, we used mainly $\gamma$."
 Our best fit for Zi; was 17000 Ix. and for logg the rest value was 3.55.," Our best fit for $T_{\rm{eff}}$ was 17000 K, and for $\log g$ the best value was 3.75."
 After the comparison of several eid models to the observed. spectrum. we estimated the errors in d; and log g to be about +2000 IX and -Εὐ25 dex. respectively.," After the comparison of several grid models to the observed spectrum, we estimated the errors in $T_{\rm{eff}}$ and log $g$ to be about $\pm $ 2000 K and $\pm 0.25$ dex, respectively."
 These conservative errors were adopted mainly due to the uncertainties involved in the normalization of our echelle spectra. the observed. spectral variability. ancl the fast rotation presented by this star. which causes a large broadening and blends several features. complicating the analvsis.," These conservative errors were adopted mainly due to the uncertainties involved in the normalization of our echelle spectra, the observed spectral variability, and the fast rotation presented by this star, which causes a large broadening and blends several features, complicating the analysis."
 The physical parameters. derived. are. summarized. in ‘Table 2.., The physical parameters derived are summarized in Table \ref{params}.
 ουν good. agreement is obtained in both UV and optical regions., Very good agreement is obtained in both UV and optical regions.
 Since LUSTY directly provides only Tig and log g. we have computed. the Dluminosity. racdius. and the mass of LR 7355 in the following wav.," Since TLUSTY directly provides only $T_{\rm{eff}}$ and log $g$, we have computed the luminosity, radius, and the mass of HR 7355 in the following way."
 First. we computed the visual extinction ely from the apparent (V — 6.02) and ISM-corrected: visual magnitudes (Và= BSL: CGutiérrrez-Moreno Moreno. LOGS).," First, we computed the visual extinction $A_V$ from the apparent (V = 6.02) and ISM-corrected visual magnitudes $_0 = 5.81$ ; Gutiérrrez-Moreno Moreno 1968)."
 Thereafter. we inferred. the distance from the HIPPATRCOS parallax for his object (x=3.66+0.57: van Leeuwen 2007) and computed the absolute visual magnitude CVV).," Thereafter, we inferred the distance from the HIPPARCOS parallax for this object $\pi = 3.66 \pm
0.37$; van Leeuwen 2007) and computed the absolute visual magnitude $M_V$ )."
" By using the rolometric correction (BC) corresponding to our FLUSTY model (Lanz Llubeny 2007). we thus determined. the xometric magnitude (Also,=Ady| DC) and then the uminosity. taking the bolometric magnitude of the Sun to ο 4.75 (Allen 1976)."," By using the bolometric correction (BC) corresponding to our TLUSTY model (Lanz Hubeny 2007), we thus determined the bolometric magnitude $M_{\rm{BOL}} = M_V + BC$ ) and then the luminosity, taking the bolometric magnitude of the Sun to be 4.75 (Allen 1976)."
" Taking the distance and its associated uncertainty.. we find⋅ that log£L,/L.=3.06y⋅∩⇁∩∣∣Gay."," Taking the distance and its associated uncertainty, we find that $\log L_\star/L_
\odot = 3.06^{+0.09}_{-0.08}$."
 Prom thisH value and Zi. we computed the radius and from log qi (sce below) we derived the mass.," From this value and $T_{\rm{eff}}$, we computed the radius and from log $g_{\rm{true}}$ (see below) we derived the mass."
 These parameters are in good agreement with typical carly D. star. values (see e.g. ‘Trunelle et al., These parameters are in good agreement with typical early B star values (see e.g. Trundle et al.
 2007)., 2007).
 Note that io we use the webyd colors of LER 7355 to derive Yigg. we obtain a value of ~ 15000 Ix. in agreement with our mocel analysis.," Note that if we use the $uvby\beta$ colors of HR 7355 to derive $T_{\rm{eff}}$, we obtain a value of $\sim$ 18000 K, in agreement with our model analysis."
 Regarding log g. the inferred. value of 3.75 should be considered as ellective. Lc. he surface gravity reduced by rotation.," Regarding log $g$, the inferred value of 3.75 should be considered as effective, i.e., the surface gravity reduced by rotation."
 An estimate of the rue gravity can be obtained by the following expression (Repolust et al., An estimate of the true gravity can be obtained by the following expression (Repolust et al.
 2004): For UR 7355. we find loggan=3.95. which is more consistent with its spectral classification.," 2004): For HR 7355, we find $\log g_{\rm{true}} = 3.95$, which is more consistent with its spectral classification."
 The stellar mass ollows from this value (see Table 2))., The stellar mass follows from this value (see Table \ref{params}) ).
 As previously claimed in the literature. we found that his star must be enriched in helium.," As previously claimed in the literature, we found that this star must be enriched in helium."
 We could not fit any of the observed He profiles using the solar abundance number ratio Llef/ll = 0.1., We could not fit any of the observed He profiles using the solar abundance number ratio He/H = 0.1.
 ALL the computel lines were systematically weaker than the ones observed., All the computed lines were systematically weaker than the ones observed.
 \lodels with Le/Ll ratios from 0.2 to 1.0 presented good agreement with the observations (taking into account the dillerent spectra), Models with He/H ratios from 0.2 to 1.0 presented good agreement with the observations (taking into account the different spectra).
 A compatible He/L1L ratio (0.43) was derived by. Rivinius et al. (, A compatible He/H ratio (0.4) was derived by Rivinius et al. (
2008) from equivalent widths of He L lines in a spectrum acquired in 1999.,2008) from equivalent widths of He I lines in a spectrum acquired in 1999.
 LIowever. these same authors estimated a solar He/LE value [rom a second spectrum. obtained in 2004.," However, these same authors estimated a solar He/H value from a second spectrum, obtained in 2004."
 We used the multiline analysis method: of Least-Squares Deconvolution (LSD: Donati et al., We used the multiline analysis method of Least-Squares Deconvolution (LSD; Donati et al.
 1997) to produce mean Stokes Land V. profiles [rom our spectra., 1997) to produce mean Stokes I and V profiles from our spectra.
 Phese profiles are illustratecL in Fig. L., These profiles are illustrated in Fig. \ref{LSDprof}.
 The details of the LSD method. as applied to B stars are described by Silvester et. al. (, The details of the LSD method as applied to B stars are described by Silvester et al. (
2009).,2009).
 The LSD line mask was created. using a VALD line list for a star with Zij—18000 Ix. log g = 4.0. ancl solar abundances except for helium.," The LSD line mask was created using a VALD line list for a star with $T_{\rm{eff}}$ =18000 K, log $g$ = 4.0, and solar abundances except for helium."
 The helium abundance was set το 0.1 in logarithmic units. although as mentioned previously. the helium line strength varies stronglv with rotation phase.," The helium abundance was set to $-$ 0.1 in logarithmic units, although as mentioned previously, the helium line strength varies strongly with rotation phase."
 Most helium and metal lines in the range 405-800. nm were included in the line mask., Most helium and metal lines in the range 405-800 nm were included in the line mask.
 Lines located in regions with Balmer lines were removed as hvelrogen lines are not included in the mask., Lines located in regions with Balmer lines were removed as hydrogen lines are not included in the mask.
 Phe wavelength range was chosen in an attempt to exclude regions with strong blending with hydrogen lines in the blue-most. part of the spectrum ancl strong telluric contamination in the red. part of the spectrum., The wavelength range was chosen in an attempt to exclude regions with strong blending with hydrogen lines in the blue-most part of the spectrum and strong telluric contamination in the red part of the spectrum.
 Mean. profiles were binned to 2.6 kms in wavelength., Mean profiles were binned to 2.6 km $^{-1}$ in wavelength.
 The LSD velocity range was set from 600 kms  to 1600 km + and the line depth eutoll was of the continuum., The LSD velocity range was set from $-$ 600 km $^{-1}$ to +600 km $^{-1}$ and the line depth cutoff was of the continuum.
 The mean Stokes L. null. and mean Stokes V. profiles are shown in Fig. 1..," The mean Stokes I, null, and mean Stokes V profiles are shown in Fig. \ref{LSDprof}."
 The lack of signal in the null profiles indicate that there are no. important spurious contributions to the Stokes V profiles., The lack of signal in the null profiles indicate that there are no important spurious contributions to the Stokes V profiles.
 Each separate LSD Stokes V. profile produced a definite detection (2 99.9994)) according to the criteria described bv Donati et al. (, Each separate LSD Stokes V profile produced a definite detection $>$ ) according to the criteria described by Donati et al. (
1997).,1997).
 From the LSD mean Stokes Land V profiles. we have calculated the Longitudinal magnetic field from the first moment of Stokes V. as a function of velocity ο from: (Alathys 1989: Donati etal.," From the LSD mean Stokes I and V profiles, we have calculated the longitudinal magnetic field from the first moment of Stokes V as a function of velocity $v$ from: (Mathys 1989; Donati etal."
 1997: Wace et al., 1997; Wade et al.
 2000)., 2000).
 In Eq. (, In Eq. (
2). gis the mean Landé factor and A is the mean wavelength of all the included lines.,"2), $g$ is the mean Landé factor and $\lambda$ is the mean wavelength of all the included lines."
 Ehe integration range emploved in calculating the longitudinal magnetic field was from 350 kms το [350 km 1., The integration range employed in calculating the longitudinal magnetic field was from $-$ 350 km $^{-1}$ to +350 km $^{-1}$ .
 Phe derived values. along with —their," The derived values, along with their"
conrparisou will allow us to establish under what cireunstauces the umuerical auc analytic models agree and thus can be applied with reasonable coufideuce. aud o gain physical insight iuto the IIo formation process wv oNamining cases where they differ.,"comparison will allow us to establish under what circumstances the numerical and analytic models agree and thus can be applied with reasonable confidence, and to gain physical insight into the $_2$ formation process by examining cases where they differ."
 Au additional xactieal benefit is that. du cases where the analytic nodels are able to do a reasonable job reproducing he numerical ones. they may provide an approximation hat can be used ia simulations at ercatly reduced conmputationa cost.," An additional practical benefit is that, in cases where the analytic models are able to do a reasonable job reproducing the numerical ones, they may provide an approximation that can be used in simulations at greatly reduced computational cost."
 The structure of the remainder of this paper is as ollows., The structure of the remainder of this paper is as follows.
 Iu 2 we sununumizethe nuunuerica aud analytic uodels for the to Il transition., In \ref{sec:methods} we summarize the numerical and analytic models for the to $_2$ transition.
 In 3 WO COlLre the models in idealized simulations in which the radiation field aud netal content are held fixed. while il I we consider siunulatious im which the radiation fielk laud netal couteut are determined ina fully self-cousist¢ut manner.," In \ref{sec:fixedrad} we compare the models in idealized simulations in which the radiation field and metal content are held fixed, while in \ref{sec:scrad} we consider simulations in which the radiation field and metal content are determined in a fully self-consistent manner."
 Finally is 5 we discuss aud suuiunuuize otr conclusions., Finally in \ref{sec:summary} we discuss and summarize our conclusions.
 Tere we sunnunuarize how we estinate the Πο mass fraction iu. every computational cell of our test siuulatious., Here we summarize how we estimate the $_2$ mass fraction in every computational cell of our test simulations.
 Section 2.1 describes the time-dependent nuuerical chemistry method. auc Section 2.2. describes the analytic model aud how we apply it to the simulations.," Section \ref{sec:numerical} describes the time-dependent numerical chemistry method, and Section \ref{sec:analytic} describes the analytic model and how we apply it to the simulations."
 Simulations used in this paper are fully described in GINO., Simulations used in this paper are fully described in GK10.
" Tere we only remind a reader that a set of ""fixed. ISAT” siuulations was perfornie« with the Adaptive Refincen Tree (ART) code (??7).."," Here we only remind a reader that a set of “fixed ISM” simulations was performed with the Adaptive Refinement Tree (ART) code \citep{sims:k99,sims:kkh02,sims:rzk08}."
 Each siuulation was started frou a :=[1 outpu of a fully self-consistent cosmological simulation of a Milky Wav progenitor galaxy., Each simulation was started from a $z=4$ output of a fully self-consistent cosmological simulation of a Milky Way progenitor galaxy.
 However. in cach fixed ISM simulation the dust-to-gas ratio and the interstellar radiation field were fixed to pre-defined. coustaut in space and time values. effectively turning a cosmological simulation iuto a lighly realistic simulation of au isolated ealaxy with a live dark matter halo au natural accretion history.," However, in each fixed ISM simulation the dust-to-gas ratio and the interstellar radiation field were fixed to pre-defined, constant in space and time values, effectively turning a cosmological simulation into a highly realistic simulation of an isolated galaxy with a live dark matter halo and natural accretion history."
 Each fixed ISX simmlation was coutinued for 600 \Lvr. long enough to fully establish a new ISM structure and (nuimate the memory of the initial coiditious.," Each fixed ISM simulation was continued for 600 Myr, long enough to fully establish a new ISM structure and eliminate the memory of the initial conditions."
" The dark inatter niass resolution iu the sinnuation is 1.3«10°AL, and the spatial resolution ofthe highes resolved region of the simulation (that includes the whoe of the simulated ealaxy disk) is 65 pe.", The dark matter mass resolution in the simulation is $1.3\times10^6\Msun$ and the spatial resolution of the highest resolved region of the simulation (that includes the whole of the simulated galaxy disk) is 65 pc.
" That translaes into the mass resolution in the gas of between 10AL, and 10AL. . depending on the value of tlic! eas density."," That translates into the mass resolution in the gas of between $\sim 10^3\Msun$ and $\sim 10^6\Msun$ , depending on the value of the gas density."
 All fixed ISA sinmlatious iuclucο σας (Lyniadudes. star formation. and the time-dependent and spatiallv-idomogeneous 23D radiative transfer of ultraviolet and ionizing radiation from mdividual stellar particles that forie during the course of cach simulation.," All fixed ISM simulations include gas dynamics, star formation, and the time-dependent and spatially-inhomogeneous 3D radiative transfer of ultraviolet and ionizing radiation from individual stellar particles that formed during the course of each simulation."
 The simulations incorporate a nou-equilibriun chemical network of drogen and heliuu aud ποιοπι cooling aud heating rates. which make use of the local abundance of atomic. molecular. aud ionic species and UV intensity.," The simulations incorporate a non-equilibrium chemical network of hydrogen and helium and non-equilibrium cooling and heating rates, which make use of the local abundance of atomic, molecular, and ionic species and UV intensity."
 This network iucludes formation of molecular hydrogen both iu the primordial phase aud on dus erains., This network includes formation of molecular hydrogen both in the primordial phase and on dust grains.
 The abundances of the relevant atomic aud molecular species are therefore followed selfconsistently durug the course of the simulation., The abundances of the relevant atomic and molecular species are therefore followed self-consistently during the course of the simulation.
 The heating and cooling terms in the equation for the iuterua enerev ποιος all of the terms normally iucluded iu the sinulatious of first stars and iun the ΤΝΤ models. iucludiug cooling ou metals.," The heating and cooling terms in the equation for the internal energy include all of the terms normally included in the simulations of first stars and in the ISM models, including cooling on metals."
 A complete set of Chemica reactions and the values for the adopted reaction au cooling/heating rates are preseuted in the appendix of 10., A complete set of chemical reactions and the values for the adopted reaction and cooling/heating rates are presented in the appendix of GK10.
 Iu adelition to the simulations with the fixed ISA properties. we also use a fully self-consistent cosinologica sinulation of several galaxies iu a cosmological volume with 255+ Mpe on a side.," In addition to the simulations with the fixed ISM properties, we also use a fully self-consistent cosmological simulation of several galaxies in a cosmological volume with $25h^{-1}$ Mpc on a side."
 These sinmlations will be fully described clsewhere ZZemp et al. in preparation). but in physical modclueg they are identical to the simulation of ?:: the only difference from the carly siuulation is a larger box size aud sliehtlv different values of cosinological parameors. base ou the WALAPS cosinoloey (2)..," These simulations will be fully described elsewhere Zemp et al, in preparation), but in physical modeling they are identical to the simulation of \citet{ng:gk10a}; the only difference from the early simulation is a larger box size and slightly different values of cosmological parameters, based on the WMAP5 cosmology \citep{cosmo:dkns08}."
 The spatial and mass resolution of the cosinological simulation is mached exactly to the spatial and mass resolution of the simulation wih the fixed ISM., The spatial and mass resolution of the cosmological simulation is matched exactly to the spatial and mass resolution of the simulation with the fixed ISM.
 Due to computational expense. the fulv self-consisteut cosinological simulation was rot continued hevoud +=3.," Due to computational expense, the fully self-consistent cosmological simulation was not continued beyond $z=3$ ."
 For comparison with the analy iodel. we use several of the model galaxies from hat simulation that differ widely in their ISAL properties.," For comparison with the analytical model, we use several of the model galaxies from that simulation that differ widely in their ISM properties."
 T1 ΠΟΙΟΙ galaxies from a full cosmological simulation he properties of the ISAL (dust-to-eas ratio. gas metallicitv. the interstellar radiatou field. the sas density distribution. etc.)," In model galaxies from a full cosmological simulation the properties of the ISM (dust-to-gas ratio, gas metallicity, the interstellar radiaton field, the gas density distribution, etc.)"
 vary inside the galaxy iu a manner that is controlled. by their prior cosmüc history and the X»esen costological enviromuent., vary inside the galaxy in a manner that is controlled by their prior cosmic history and the present cosmological environment.
 Finally. we call the readers atteution o two details of the mücrophnvsical model of Πω formation in the simulations that will ρα iurportaut for our comparison.," Finally, we call the reader's attention to two details of the microphysical model of $_2$ formation in the simulations that will be important for our comparison."
 First. suce the Πο destruction rate depends on the amount of shiclding. one iust estimate a column deusitv or each cell.," First, since the $_2$ destruction rate depends on the amount of shielding, one must estimate a column density for each cell."
 In. the «ιαΊος one does this using ie Sobolev approximation: in each cel one coniputes i6 local density scale height 5h=δρ]. aud then akes the column deusitv to be X—ph.," In the simulations one does this using the Sobolev approximation: in each cell one computes the local density scale height $h = \rho/|\nabla \rho|$, and then takes the column density to be $\Sigma=\rho h$."
 Comparison of Us approxinuatiou wihi rav-tracing done im post-process shows hat it is generally quite accurae., Comparison of this approximation with ray-tracing done in post-process shows that it is generally quite accurate.
 Second. one uust increase the rate coefficient for Πω formation by a ‘hunping factor to account for density inhomoecueities jii are unresolved on the computational erid.," Second, one must increase the rate coefficient for $_2$ formation by a clumping factor to account for density inhomogeneities that are unresolved on the computational grid."
 Since ie II» formation rate varies as the square of density. οσο juhomogeneities increase the overall rate.," Since the $_2$ formation rate varies as the square of density, these inhomogeneities increase the overall rate."
 The sinulations use a clumping factor of 30. which is chosen o eive a good match between the simulations aud FUSE observatious of the Milky Wav iux Magellanic clouds.," The simulations use a clumping factor of 30, which is chosen to give a good match between the simulations and FUSE observations of the Milky Way and Magellanic clouds."
 The analytic iocdel to which we conrpare the saunulatious is based on the work. of 7.. *?.. aud ?.. hereafter known as the IKMT model.," The analytic model to which we compare the simulations is based on the work of \citet{krumholz08c}, \citet{krumholz09a}, and \citet{mckee10a}, hereafter known as the KMT model."
 WATT consider an idealized spherical cloud immersed in a uniform. isotropic Lyiman-Werner baud rachation fiek.," KMT consider an idealized spherical cloud immersed in a uniform, isotropic Lyman-Werner band radiation field."
 They then solve the coupled problems of radiative trausfer aud fformation-cissociation balance under the assuniptiou that the cloud is in steady state., They then solve the coupled problems of radiative transfer and formation-dissociation balance under the assumption that the cloud is in steady state.
 They do not consider eas-plase reactions (though see the Appendix of ?. for a discussion of how these nieht be included)., They do not consider gas-phase reactions (though see the Appendix of \citealt{mckee10a} for a discussion of how these might be included).
 The solution 15, The solution is
the modes.,the modes.
 We used sit-and-stare mode data observed with EIS., We used sit-and-stare mode data observed with EIS.
" The duration of the data was 2006 Dee 3 12:43:10-13:37:20 UT. and the field of view was 1"" in the .-direction and 256"" in the y-direction in heliocentric coordinates."," The duration of the data was 2006 Dec 3 12:43:10-13:37:20 UT, and the field of view was $1''$ in the $x$ -direction and $256''$ in the $y$ -direction in heliocentric coordinates."
 There were 288 data points in the time direction., There were 288 data points in the time direction.
" The exposure time was 10s, and the average time cadence was 11.35. with a range of 11.3—11.5s."," The exposure time was $10 \mspace{3mu} \mathrm{s}$, 	and the average time cadence was $11.3\mspace{3mu} \mathrm{s}$, with a range of $11.3-11.5\mspace{3mu}\mathrm{s}$."
 The standard Solar Software interactive data language (SSWIDL) procedure was used to calibrate the data., The standard Solar Software interactive data language (SSWIDL) procedure was used to calibrate the data.
" We used an Fe 195.12 emission line, whose line profile at cach pixel was fitted using a single Gaussian function."," We used an Fe $195.12\mspace{3mu}$ emission line, whose line profile at each pixel was fitted using a single Gaussian function."
 Then intensity (integrated line profile) and Doppler shift (difference between line centroid and 195.12 were calculated using the fitted parameters., Then intensity (integrated line profile) and Doppler shift (difference between line centroid and $195.12\mspace{3mu}$ were calculated using the fitted parameters.
" After removing the influence of the slit-tilt and the orbital variation using the procedure, we were able to investigate the average line centroid at cach exposure."," 	After removing the influence of the slit-tilt and the orbital variation using the procedure, 	we were able to investigate the average line centroid at each exposure."
" However, a 90—100min period component still remained, which was subtracted for accuracy."," 	However, a $90-100\mspace{3mu} \mathrm{min}$ period component still remained, 	which was subtracted for accuracy."
" Usually, the zero-point of Doppler velocity is removed through examination of a quict sun region, where Doppler velocity is much smaller and less dynamic."," 	Usually, the zero-point of Doppler velocity is removed through examination of a quiet sun region, 	where Doppler velocity is much smaller and less dynamic."
" However, with a short exposure time (10 9) in our data, the incident photons at the quict sun region were too small for analysis, so the zero-point of Doppler velocity was calculated by using all pixels along the slit."," 	However, with a short exposure time $10\mspace{3mu}\mathrm{s}$ ) in our data, 	the incident photons at the quiet sun region were too small for analysis, 	so the zero-point of Doppler velocity was calculated by using all pixels along the slit."
" Since no explosive events were observed in our data, this zero-point adjustment does not produce artificial oscillations."," 	Since no explosive events were observed in our data, this zero-point adjustment does not produce artificial oscillations."
" Note that there were no prominent wing components in the line profiles, and the weak line blend with Fe 195.18A (Youngetal.2009) can be ignored because the main focus of our research is the time variability of Doppler velocity (Wangetal.2009b)."," 	Note that there were no prominent wing components in the line profiles, and the weak line blend with Fe $195.18\mspace{3mu}$ \citep{you09} can be ignored because the main focus of our research is the time variability of Doppler velocity \citep{wan09b}."
. Panel (a) in Figure 1. shows a 171 image of NOAA active region 10926 taken by TRACE., Panel (a) in Figure \ref{ss trace} shows a $171\mspace{3mu}$ image of NOAA active region 10926 taken by TRACE.
 Panel (b) in Figure 1. shows an EIS Fe 284.16 raster scan image., Panel (b) in Figure \ref{ss trace} shows an EIS Fe $284.16\mspace{3mu}$ raster scan image.
 The white vertical line in each panel indicates the cotemporal position of the EIS slit in both panels., The white vertical line in each panel indicates the cotemporal position of the EIS slit in both panels.
 In the TRACE, In the TRACE
tthe star formation episode.,the star formation episode.
shown iu Fieure (1) where the planets orbital senui-niajor axis and the plauet-satellite distance are shown as functions of time.,shown in Figure (1) where the planet's orbital semi-major axis and the planet-satellite distance are shown as functions of time.
 The massless satellites are ejected from around the planet at 0.101. AU (Callisto). at 0.23 AU (Ganvinede). at O11 AU (Europa). aud at 0.017 AU (lo) iu agrecneut with the expression (3)).," The massless satellites are ejected from around the planet at $0.41$ AU (Callisto), at $0.23$ AU (Ganymede), at $0.14$ AU (Europa), and at $0.047$ AU (Io) in agreement with the expression \ref{acrit}) )."
" Two satellites remain on star-bouud orbits: Callisto with a,=OSAU.ο0.15. and Camuviede with a,=5OAU.ος=0.99 where ος is the satellites. orbita eccentricity."," Two satellites remain on star-bound orbits: Callisto with $a_s=0.8\,\mbox{AU}, \ e_s=0.45$, and Ganymede with $a_s=50\,\mbox{AU}, \ e_s=0.99$ where $e_s$ is the satellite's orbital eccentricity."
 Europa and Io. however. are ejected frou the planetary svstemi.," Europa and Io, however, are ejected from the planetary system."
 Whereas satellites are invariably ejected frou around. the planet. we find that their fina orbits in the planetary svstem are sensitive to the initia orbital paraiucters as notion 1s chaotic near the stability boundary.," Whereas satellites are invariably ejected from around the planet, we find that their final orbits in the planetary system are sensitive to the initial orbital parameters as motion is chaotic near the stability boundary."
 In agreemieut with the adiabatic character of planet nuüeration with respect to the satellites orbita motion. the satellite is oblivious to the planct’s shrinkine orbit uutil the stability boundary closes in on the satellite a few 107 years before ejection.," In agreement with the adiabatic character of planet migration with respect to the satellite's orbital motion, the satellite is oblivious to the planet's shrinking orbit until the stability boundary closes in on the satellite a few $10^3$ years before ejection."
 AIutual satellite interactions further disturb the svsten once instability. closes iu on the outermost satellite., Mutual satellite interactions further disturb the system once instability closes in on the outermost satellite.
 Figure (2) shows how Cauvinede's instability perturbs Europa and ejects it in its place., Figure (2) shows how Ganymede's instability perturbs Europa and ejects it in its place.
 Mutual iuteractions result in an carlicr perturbation (a few 105 years) before ejection than when iutual interactions were abseut., Mutual interactions result in an earlier perturbation (a few $10^4$ years) before ejection than when mutual interactions were absent.
 This perturbation forces lo orbit to assume a larecr eccentricity that results in a collisiou with the planet., This perturbation forces Io's orbit to assume a larger eccentricity that results in a collision with the planet.
" The remaining three satellites assuine star-bound orbits: (6;=L16ÀU.ος0:55) for Callisto. (a,=0.5AU.c,0.99) for Canvinede aud (a.=3.12AU.e,=0.91) for Europa."," The remaining three satellites assume star-bound orbits: $a_s=4.16\,\mbox{AU}, \ e_s=0.88)$ for Callisto, $a_s=49.5\,\mbox{AU}, \ e_s=0.99)$ for Ganymede and $a_s=3.12\,\mbox{AU}, \ e_s=0.91)$ for Europa."
 As with the massless satellite simulation. the instability outcome after ejection 1s sensitive to the initial conditions.," As with the massless satellite simulation, the instability outcome after ejection is sensitive to the initial conditions."
 This work illustrates bow the imigratiou-cjectiou instability. operatesfor exoplanet satellites., This work illustrates how the migration-ejection instability operatesfor exoplanet satellites.
 Iu doing so we asstuned that satellites formi at semi-nuajor axes in the range from 10° to 107Ry as observed in the solar system., In doing so we assumed that satellites form at semi-major axes in the range from $10^{-3}$ to $10^{-2}R_H$ as observed in the solar system.
 This assumption is reasonable as the three solar svstei planets with regular satellites (Table ID formed iun different conditious at different orbital radi from the sun., This assumption is reasonable as the three solar system planets with regular satellites (Table I) formed in different conditions at different orbital radii from the sun.
 Iu addition them large regular satellites are likely to have been subjected to inward (Type ID migration by the cireumplanctary disks torques hat resulted partly in the currently observed orbital resonances.," In addition, their large regular satellites are likely to have been subjected to inward (Type I) migration by the circumplanetary disk's torques that resulted partly in the currently observed orbital resonances."
 À number of factors may alter the ΓΣejection iustability characteristics such as the ejection iuescale and scui-major axis but they do not suppress it., A number of factors may alter the migration-ejection instability characteristics such as the ejection timescale and semi-major axis but they do not suppress it.
 For instance it is known that as the planet opens a eap in the disk. itserowth may be significantly reduced mut not completely halted (Isley et al 2001. Lubow ot al 2002).," For instance it is known that as the planet opens a gap in the disk, itsgrowth may be significantly reduced but not completely halted (Kley et al 2001, Lubow et al 2002)."
 The accretiou rate of order 102 to 10MAL Ll depeuding ou the cieunistelha disks physical oxoperties. mereases the size of the Till sphere aud may counter the effect of müeration.," The accretion rate of order $10^{-2}$ to $10^{-4}M_\oplus$ $^{-1}$, depending on the circumstellar disk's physical properties, increases the size of the Hill sphere and may counter the effect of migration."
 As the Till parameter depends on AL;55 dobling the planet's mass duriug uieration reduces the critical planet seniü-najor axis o» only.," As the Hill parameter depends on $M_p^{1/3}$, doubling the planet's mass during migration reduces the critical planet semi-major axis by only."
 The uneratiou-cjection instability may also be nütieated if the planctary svstenis snow line is closer to the star than the nominal value of 3 AU (Cirad and Lin 2007) allowing eiaut planet formation o occur closer to the star., The migration-ejection instability may also be mitigated if the planetary system's snow line is closer to the star than the nominal value of 3 AU (Garaud and Lin 2007) allowing giant planet formation to occur closer to the star.
 Type I uueration iu the cireunplauetaryv disk transports formüng the satellites owards the planet aud away from the stability boundary., Type I migration in the circumplanetary disk transports forming the satellites towards the planet and away from the stability boundary.
 Tn order for the satellites not to collide with the planet. he timing of the disks dispersal needs to be linked to he satellites’ arrival outside the Roche radius.," In order for the satellites not to collide with the planet, the timing of the disk's dispersal needs to be linked to the satellites' arrival outside the Roche radius."
 Lastly. eccentrieitv diupius bv the circtunplauctary disk of orniue satellite orbits (Ward 1988) would oppose the exciting effect of niutual satellite interactions and make he imigratiou-cjectiou instability effect simular to that on nassless satellites.," Lastly, eccentricity damping by the circumplanetary disk of forming satellite orbits (Ward 1988) would oppose the exciting effect of mutual satellite interactions and make the migration-ejection instability effect similar to that on massless satellites."
" As planetary satellites form outside he planet's Roche radius where differcutial tidal forces are not able to prevent plauctesimal accunulatiou. a conservative estimate of the limiting planet seni-nmiajor Axis. Cydia. at Which all moous are ejected is reached when the Roche radius. Rye=(241,2p.)? (for a rigid satellite). becomes comparable to the stability zone radius. Hs=fFRy. where p, is the satellites mean mass density."," As planetary satellites form outside the planet's Roche radius where differential tidal forces are not able to prevent planetesimal accumulation, a conservative estimate of the limiting planet semi-major axis, $a_{p,{\rm lim}}$, at which all moons are ejected is reached when the Roche radius, $R_R=({3M_p}/{2\rho_s})^{1/3}$ (for a rigid satellite), becomes comparable to the stability zone radius, $R_{SZ}= fR_H$, where $\rho_s$ is the satellite's mean mass density."
" It is given as: Notwithstanding the factors that iitieate the uueratiou-cjection instability. it appears unlikely that the innermost moon that formed on a prograde orbit around a close-in git exoplanct with e,<0.02 AU could survive the unieration-ejection instability."," It is given as: Notwithstanding the factors that mitigate the migration-ejection instability, it appears unlikely that the innermost moon that formed on a prograde orbit around a close-in giant exoplanet with $a_p\leq 0.02$ AU could survive the migration-ejection instability."
 If moons are observed by around close-in planets by precision photometry missions such as Kepler or CoRot. they arc likely to be have been captured planetary cuabrvos on retrograde orbits much like Triton around Neptune.," If moons are observed by around close-in planets by precision photometry missions such as or , they are likely to be have been captured planetary embryos on retrograde orbits much like Triton around Neptune."
"southern half of the stream peaks at a distance modulus of 14.9+0.2 mag, while the northern half of the stream peaks at 14.8+0.3 mag.","southern half of the stream peaks at a distance modulus of $14.9 \pm 0.2$ mag, while the northern half of the stream peaks at $14.8 \pm 0.3$ mag."
" This puts the southern end of the stream at a sun-centric distance of 9.7+0.9 kpc, while the northern end is at 9.4+1.4 kpc."," This puts the southern end of the stream at a sun-centric distance of $9.7 \pm 0.9$ kpc, while the northern end is at $9.4 \pm 1.4$ kpc."
 The portion of the stream visible in Figure 1 is evidently almost perpendicular to our line of sight., The portion of the stream visible in Figure 1 is evidently almost perpendicular to our line of sight.
" Integrating the background-subtracted, unfiltered counts of stars within 3o of the Z = 0.0003 isochrone along the length of the stream and over a width of 1.5° we find the total number of stars in the discernible stream to be 530+230."," Integrating the background-subtracted, unfiltered counts of stars within $3\sigma$ of the Z = 0.0003 isochrone along the length of the stream and over a width of $1.5\arcdeg$ we find the total number of stars in the discernible stream to be $530 \pm 230$."
" The large uncertainty simply reflects the Poisson statistics of the very high background, (=33,000 field stars in the same region of color and configuration space)."," The large uncertainty simply reflects the Poisson statistics of the very high background, $\approx 33,000$ field stars in the same region of color and configuration space)."
" Figure 4 shows a background-subtracted, longitudonal profile of the filtered star counts, normalized to yield an integrated total of 530 stars."," Figure 4 shows a background-subtracted, longitudonal profile of the filtered star counts, normalized to yield an integrated total of 530 stars."
" For stars with g«22 and a stream width of1.0°,, the average surface density is 30+13 stars deg?, with a peak of over 100 stars deg"""," For stars with $g < 22$ and a stream width of, the average surface density is $30 \pm 13$ stars $^{-2}$, with a peak of over 100 stars $^{-2}$."
" Like the Pal 5 and GD-1 streams, the EBS profile shows?. interesting peaks and troughs, fairly regularly spaced with a separation of 4.0?+0.2°."," Like the Pal 5 and GD-1 streams, the EBS profile shows interesting peaks and troughs, fairly regularly spaced with a separation of $4.0\arcdeg \pm
0.2\arcdeg$."
" While the uncertainties are large, the clumps and gaps that give rise to these features appear quite obvious in Figure 1."," While the uncertainties are large, the clumps and gaps that give rise to these features appear quite obvious in Figure 1."
" The regular spacing may suggest an origin tied to the orbit of the stream, perhaps a result of episodic stripping (e.g. major stripping pulses at the perigalacticon of a highly eccentric orbit (Grillmair1992;Johnston,Spergel,&Hernquist 1995)))."," The regular spacing may suggest an origin tied to the orbit of the stream, perhaps a result of episodic stripping (e.g. major stripping pulses at the perigalacticon of a highly eccentric orbit \citep{grillmair1992,
 johnston1995}) )."
" Alternatively, the undulations may be due to scattering by encounters with massive objects in the disk or halo (Murali&Du-binski1999;Yoon,Johns"," Alternatively, the undulations may be due to scattering by encounters with massive objects in the disk or halo \citep{mura99, yoon2010}."
"ton,&Hogg 2010).. Schlaufmanetal.(2009) recently detected a number of cold halo substructures in velocity space (4ECHOS"") using SEGUE data.", \citet{schlaufman2009} recently detected a number of cold halo substructures in velocity space (“ECHOS”) using SEGUE data.
" With a plate center at dec) = (132.6°, 6.1°), their B-7/PCL8/PCIL-20 (R.A.,detection overlaps the EBS (Figure 5) and the estimated distance of 10 kpc is almost identical to what we find for the stream."," With a plate center at (R.A., dec) = $132.6\arcdeg$, $6.1\arcdeg$ ), their B-7/PCI-8/PCII-20 detection overlaps the EBS (Figure 5) and the estimated distance of 10 kpc is almost identical to what we find for the stream."
 We have examined this region using isochrone filters with [Fe/H] ranging from -2.2 to 0.0 and find no evidence for other cold substructures at this distance., We have examined this region using isochrone filters with [Fe/H] ranging from -2.2 to 0.0 and find no evidence for other cold substructures at this distance.
 We conclude that is most likely sampling stars in theB-7/PCI-8/PCII-20 EBS stream., We conclude that B-7/PCI-8/PCII-20 is most likely sampling stars in the EBS stream.
" Schlaufmanetal.(2009) find à mean radial velocity for B-7/PCI-8/PCII-20 of 71 km s~!, and a dispersion of 13 km s~!."," \citet{schlaufman2009} find a mean radial velocity for B-7/PCI-8/PCII-20 of 71 km $^{-1}$, and a dispersion of 13 km $^{-1}$."
" This dispersion is significantly larger than that measured for known and presumed globular cluster streams (Odenkirchenetal.2009;WillettKoposov,Rix,&Hogg 2010)."," This dispersion is significantly larger than that measured for known and presumed globular cluster streams \citep{odenkirchen2009, willett2009, koposov2010}."
". On the other hand, it is quite similar to measurements of presumed dwarf galaxy streams (Grillmair,Carlin,&Majewski2008;Carlinetal.2010;Newberg 2010)."," On the other hand, it is quite similar to measurements of presumed dwarf galaxy streams \citep{grill2008, carlin2010, newberg2010}."
". Combining the width of the stream with its large apparent velocity dispersion, we might infer that the stream's progenitor was substantially more massive than Pal 5 or the globular clusters that produced GD-1, Acheron, Cocytos, or Lethe."," Combining the width of the stream with its large apparent velocity dispersion, we might infer that the stream's progenitor was substantially more massive than Pal 5 or the globular clusters that produced GD-1, Acheron, Cocytos, or Lethe."
" On the other hand, the stream is neither as broad nor as populous as streams associated with classical dwarf galaxies like Sagittarius or the progenitor of the Orphan Stream."," On the other hand, the stream is neither as broad nor as populous as streams associated with classical dwarf galaxies like Sagittarius or the progenitor of the Orphan Stream."
" The high velocity dispersion may be partly due to non-EBS stars in the sample, or it may be due to heating of the progenitor by disk or perigalactic shocking prior to the stripping of these stars."," The high velocity dispersion may be partly due to non-EBS stars in the sample, or it may be due to heating of the progenitor by disk or perigalactic shocking prior to the stripping of these stars."
 It may also be due to significant heating by encounters with either dark matter subhalos (Carlberg or massive structures (e.g. giant molecular clouds) 2009)in the disk., It may also be due to significant heating by encounters with either dark matter subhalos \citep{carlberg2009} or massive structures (e.g. giant molecular clouds) in the disk.
" Another possibility is that the EBS may be the remnant of an ultrafaint dwarf galaxy (Willmanetal.2005;Grillmair2006a;Zuckeretal. 2006a,b),, though as we discuss below, it is difficult to imagine how such an object could have retained a dark matter envelope for any length of time."," Another possibility is that the EBS may be the remnant of an ultrafaint dwarf galaxy \citep{willman2005, grill2006a,
 zucker2006a, zucker2006b}, though as we discuss below, it is difficult to imagine how such an object could have retained a dark matter envelope for any length of time."
" Figure 5 shows an expanded, lower-contrast view of the southern portion of the EBS."," Figure 5 shows an expanded, lower-contrast view of the southern portion of the EBS."
 With a contrast maximum at the same distance (9.7 kpc) as the EBS is an, With a contrast maximum at the same distance $\approx 9.7$ kpc) as the EBS is an
"The parameter n""(q) can be caleulated from the finctions that define the specific mocel. namely a(q). s(q). g(q). and f(q). through equations (5)) and (8)).","The parameter $n''(q)$ can be calculated from the functions that define the specific model, namely $a(q)$ , $s(q)$ , $g(q)$, and $f(q)$ , through equations \ref{equ_h}) ) and \ref{equ_n}) )."
 Additionally. however. the evaluation of equation (26)) requires the determination of «.," Additionally, however, the evaluation of equation \ref{equ_local4}) ) requires the determination of $\omega''$."
 This can be obtained as follows., This can be obtained as follows.
 A solution to the equation of motion must. of course. be such that the equation of motion F—0 holds feq. (6)):," A solution to the equation of motion must, of course, be such that the equation of motion $F=0$ holds [eq. \ref{equ_motion}) );"
 eq. (13))]., eq. \ref{equ_critical1}) )].
 Therefore. That is Thus. in general Bul. at the critical point. both the denominator ancl the numerator of equation 31)) are equaltozero θες. (19))," Therefore, That is Thus, in general But, at the critical point, both the denominator and the numerator of equation \ref{equ_local10}) ) are equaltozero [eqs. \ref{equ_critical2}) )"
and (20)). respectively].,"and \ref{equ_critical3}) ), respectively]."
 Therefore. Or," Therefore, or"
and probably a PMS star.,and probably a PMS star.
" We could not find a consistent Johnson UBV dataset for this cluster, and so used the Tycho-2 catalogue and its first supplement (?),, although it does not contain a magnitude for o Ori C. In this dataset a combined magnitude is given for c Ori A and B. As c Ori A is itself a binary, we removed the effect of σ Ori B on the combined magnitude by assuming the magnitude difference between the components is the mean of the values for the difference found from speckle (7) and adaptive optics (?) work."," We could not find a consistent Johnson $UBV$ dataset for this cluster, and so used the Tycho-2 catalogue and its first supplement \citep{2000A&A...355L..27H}, although it does not contain a magnitude for $\sigma$ Ori C. In this dataset a combined magnitude is given for $\sigma$ Ori A and B. As $\sigma$ Ori A is itself a binary, we removed the effect of $\sigma$ Ori B on the combined magnitude by assuming the magnitude difference between the components is the mean of the values for the difference found from speckle \citep{2001AJ....121.1583H} and adaptive optics \citep{2000AJ....119.2403T} work."
 c Ori B will have little effect on the combined colour., $\sigma$ Ori B will have little effect on the combined colour.
" The disadvantage of using the Tycho-2 data is that we cannot determine the extinction, as there are no U-band data."," The disadvantage of using the Tycho-2 data is that we cannot determine the extinction, as there are no $U$ -band data."
 We therefore simply adopted E(B—V)0.06 from ?.., We therefore simply adopted $E(B-V)=0.06$ from \cite{1994A&A...289..101B}.
" The resulting 7? contour does not close at low ages, and so we have only an upper limit on the age."," The resulting $\tau^2$ contour does not close at low ages, and so we have only an upper limit on the age."
" We therefore quote (in Table 1)) the upper limit below which 68 percent of the probability lies, but exclude this cluster from further analysis."," We therefore quote (in Table \ref{ages_table}) ) the upper limit below which 68 percent of the probability lies, but exclude this cluster from further analysis."
" We used the photo-electric data of ?,, as presented in his Table 1."," We used the photo-electric data of \cite{1956ApJS....2..365W}, , as presented in his Table 1."
" Fitting all stars blueward of B—V= for extinction in U—B/ B—V space gives Pr(7?)—0.37, implying0 uniform extinction over the field."," Fitting all stars blueward of $B-V=0$ for extinction in $U-B$ $B-V$ space gives $\Pr(\tau^2)$ =0.37, implying uniform extinction over the field."
 We then fitted in V/B—V and obtained a Pr(r*) of 0.06., We then fitted in $V$ $B-V$ and obtained a $\Pr(\tau^2)$ of 0.06.
" This is on the margins of acceptability, and there is a case that the two data points furthest redward from the sequence should be removed."," This is on the margins of acceptability, and there is a case that the two data points furthest redward from the sequence should be removed."
" However, in not doing so we simply enlarge our uncertainty estimate, and so are being conservative."," However, in not doing so we simply enlarge our uncertainty estimate, and so are being conservative."
 We used the data from ? taking only those stars within half a degree of A Ori., We used the data from \cite{1977MNRAS.181..657M} taking only those stars within half a degree of $\lambda$ Ori.
" We excluded objects with B—V>0.2, which results in a sample which is almost complete blueward of this colour, and has no stars redward of (B—V)o=—0.04."," We excluded objects with $B-V>0.2$, which results in a sample which is almost complete blueward of this colour, and has no stars redward of $(B-V)_0=-0.04$ ."
" After applying reddenings determined on a star-by- basis, we obtained a value of Pr(7?) of 0.52, provided we assumed the uncertainties were 0.01 mags in V and 0.008 mags in B—V (?,donotprovideerrorbars) and removed two objects (HD36881 and HD36913) which appear to be non-members based on their position in the Vo/(B —V)o diagram."," After applying reddenings determined on a star-by-star basis, we obtained a value of $\Pr(\tau^2)$ of 0.52, provided we assumed the uncertainties were 0.01 mags in $V$ and 0.008 mags in $B-V$ \citep[][do not provide error bars]{1977MNRAS.181..657M} and removed two objects (HD36881 and HD36913) which appear to be non-members based on their position in the $V_0$ $(B-V)_0$ diagram."
 The resulting fit is shown on the right-hand side of Figure 5.., The resulting fit is shown on the right-hand side of Figure \ref{lambda_ori}.
" We used the data and associated uncertainties for NGC2362 from ?,, which were taken as part of a programme to define what became the UBV system."," We used the data and associated uncertainties for NGC2362 from \cite{1953ApJ...117..313J}, which were taken as part of a programme to define what became the $UBV$ system."
 We used only those stars blueward of B—V0.04 and excluded stars noted as members by ?.., We used only those stars blueward of $B-V=0.04$ and excluded stars noted as non-members by \cite{1953ApJ...117..313J}.
 We also excluded the brightest star (r CMa) as it is clearly beyond the turnoff., We also excluded the brightest star $\tau$ CMa) as it is clearly beyond the turnoff.
" Finally we found that star 36 gave a high 7? in both U— B/B—V and V/B —V, and star 50 in V/ B—V, and so removed them from the fit as well."," Finally we found that star 36 gave a high $\tau^2$ in both $U-B$ $B-V$ and $V$ $B-V$, and star 50 in $V$ $B-V$, and so removed them from the fit as well."
" We then measured a global extinction from the U- B/B—V diagram, before fitting in V/B—V."," We then measured a global extinction from the $U-B$ $B-V$ diagram, before fitting in $V$ $B-V$."
 ? carried out a photometric survey of stars of spectral type AO and earlier identified from objective prism plates., \cite{1959ApJ...130...69B} carried out a photometric survey of stars of spectral type A0 and earlier identified from objective prism plates.
" We take the membership list from ?,, but exclude BHJ11 for which the measurement is a combined light measurement for arather wide Am—2.5 binary."," We take the membership list from \cite{2001PhDT........75P}, but exclude BHJ11 for which the measurement is a combined light measurement for a rather wide $\Delta m = 2.5$ binary."
 We de-reddened this sample on a star-by-star basis., We de-reddened this sample on a star-by-star basis.
" Using the uncertainties quoted in the paper, we obtain a just about acceptable value of Pr(7?)=0.05."," Using the uncertainties quoted in the paper, we obtain a just about acceptable value of $\Pr(\tau^2)$ =0.05."
" Our data and uncertainties are taken from ?,, who aimed to obtain photometry for as many stars as possible within the outline of the dark cloud, since this area will be the least contaminated by background stars."," Our data and uncertainties are taken from \cite{1969ApJ...155..447W}, who aimed to obtain photometry for as many stars as possible within the outline of the dark cloud, since this area will be the least contaminated by background stars."
" We removed stars redward of B— V=0.0, those marked as variables or visual doubles, and three stars which lie away from the sequence in the -- B/B—V diagram."," We removed stars redward of $B-V$ =0.0, those marked as variables or visual doubles, and three stars which lie away from the sequence in the $U-B$ $B-V$ diagram."
" We fitted the data for a single extinction in U— B/B—V space, and after removingtwo outliers in 7? obtained a good fit with Pr(7?)—0.46."," We fitted the data for a single extinction in $U-B$ $B-V$ space, and after removingtwo outliers in $\tau^2$ obtained a good fit with $\Pr(\tau^2)$ =0.46."
" At first this may sound counter-intuitive, since it implies uniform extinction, yet it is well known that the extinction of ONC members is highly variable."," At first this may sound counter-intuitive, since it implies uniform extinction, yet it is well known that the extinction of ONC members is highly variable."
 In fact it seems this only applies to stars in the central cluster., In fact it seems this only applies to stars in the central cluster.
 We then fitted in V/B—V to, We then fitted in $V$ $B-V$ to
peak to trough amplitudes of our models are within a factor of 2 of the Lanetou&Lauehlin(2008) and Tro&Douüus(2010) siunlatious.,peak to trough amplitudes of our models are within a factor of 2 of the \cite{Langton_2008} and \cite{Iro_2010} simulations.
" If we adopt the Trad=LAL2) hour of Lanetou&Laughliu(2008).. we obtain peak-to-trough amplitudes of ~3«10, still a actor of ~2 larecr than their models."," If we adopt the $\tau_{\rm rad}=4.5(2)$ hour of \cite{Langton_2008}, we obtain peak-to-trough amplitudes of $\sim3\times 10^{-3}$ , still a factor of $\sim 2$ larger than their models."
 That factor of 2 is siguificaut: it indicates that unlike the Laugton&Laughlin(200S) and Tro&Deming(2010) simulations in our nodel he planet cau completely radiate away it’s leat over a slele planetary orbit., That factor of 2 is significant: it indicates that —unlike the \cite{Langton_2008} and \cite{Iro_2010} simulations— in our model the planet can completely radiate away it's heat over a single planetary orbit.
 m eencral. nou-zero heat capacity reduces the Hani111 COtrast ratio and imuereases the mdininmun. nakine for a flatter light curve.," In general, non-zero heat capacity reduces the maximum contrast ratio and increases the minimum, making for a flatter light curve."
 Furthermore. the phase ap!itide generally diminishes for larger advective YCQUCLCICS Dt16 peak-to-trough amplitude is smaller for “rot=p5wus than for wu=2uwyas across the hoard.," Furthermore, the phase amplitude generally diminishes for larger advective frequencies: the peak-to-trough amplitude is smaller for $\omega_{\rm rot} = 1.5\omega_{\rm max}$ than for $\omega_{\rm rot} = 2\omega_{\rm max}$ across the board."
 For plajets with 907Tou270°. however. a seco order effect may also take place: the light curve maxi {κx moderate Taq as this allows the planet &) hold ou Ogtbsorbed power until its hot-spot has rotate into our viav (sce Figure 12)).," For planets with $90^{\circ}<\omega<270^{\circ}$, however, a second order effect may also take place: the light curve maximum for moderate $\tau_{\rm rad}$ as this allows the planet to hold onto absorbed power until its hot-spot has rotated into our view (see Figure \ref{sample_HD17156}) )."
" Short-perioc exoplauets Ol civcular orbits should have tidallv ocked corOs, because the tidal locking timescale Is considerabv shorter than the age of most planetirv svstenis (0.2o=Lubowetal.1997)."," Short-period exoplanets on circular orbits should have tidally locked cores, because the tidal locking timescale is considerably shorter than the age of most planetary systems \citep[e.g.,][]{Lubow_1997}."
. By the same token. auets on eccentric orbits should have cores im pseudo-svuchrouotS rotation (roughly speakiug. this means cine tidalv locked at periastron. when the tidal forces are stroug¢st).," By the same token, planets on eccentric orbits should have cores in pseudo-synchronous rotation (roughly speaking, this means being tidally locked at periastron, when the tidal forces are strongest)."
" Unlike detailed hydro. siniulations that use the plajets rotation rate as an iuput paracter. our uodel is agnostic about precise prescriptions for pseudo-svuchrouots rotation frequency, (e.g.Unt1981:Ivanov&Papaloizou 2007)."," Unlike detailed hydro simulations that use the planet's rotation rate as an input parameter, our model is agnostic about precise prescriptions for pseudo-synchronous rotation frequency, $\omega_{\rm ps}$ \citep[e.g.][]{Hut_1981, Ivanov_2007}."
". For example.wy. at a moderate eccentrieits~ of «0.5. the Unt(1981) formmlation gives uuu8m]0δώµμος while the Ivanov&Papaloizou(2007) fornulation gives ip,%Lelenas."," For example, at a moderate eccentricity of $e=0.3$, the \cite{Hut_1981} formulation gives $\omega_{\rm ps} \approx 0.8 \omega_{\rm max}$, while the \cite{Ivanov_2007} formulation gives $\omega_{\rm ps} \approx 1.4 \omega_{\rm max}$."
 These details about the planct’s interior are not imuportant for our purposes. however. because the planets in οιro study have thick eascous envelopes covering their cores aud lvdvodvuamic simulations Πιοισατο that wiud-s)eeds iu the vicinity of the photosphere are comparable| fo the rotational velocity.," These details about the planet's interior are not important for our purposes, however, because the planets in our study have thick gaseous envelopes covering their cores and hydrodynamic simulations indicate that wind-speeds in the vicinity of the photosphere are comparable to the rotational velocity."
 As a result. the parcels of eas are probably usot stationary with respect to the sustellar point. even for (prestmeadily idally locked) planets ou circular orldts.," As a result, the parcels of gas are probably not stationary with respect to the sub-stellar point, even for (presumably tidally locked) planets on circular orbits."
 As a first-order model we treat them as advecting longitudinally at a constaut angular velocity., As a first-order model we treat them as advecting longitudinally at a constant angular velocity.
 Tidal cireularization ls iudoubtedlv responsible or the5 eenerally low eccentrieityv of short-perioc planeE (e.g.Ford—&aso2008.andoreferences therein)., Tidal circularization is undoubtedly responsible for the generally low eccentricity of short-period planets \citep[e.g.][and references therein]{Ford_2008}.
 It is not clear. however. tljat tidal heating has an observable effect on licht curves.," It is not clear, however, that tidal heating has an observable effect on light curves."
" For «xaniple. Jacksonetal.(2008) calculate preseut-daao fida heatiis beween 2 and 5«Lah Ww for Cl L361) tw| orders of magnitude less power than the stellar Isoation. which ranges between 1 and 13«1019 NV, For short-period planets. the observational signature of internal heat (ether reninant gravitational energv frou fornation or ougoine tidal heating) will likely be an inflated radius rather than additional emereent flux."," For example, \cite{Jackson_2008} calculate present-day tidal heating between $2$ and $5\times 10^{17}$ W for Gl 436b, two orders of magnitude less power than the stellar insolation, which ranges between $7$ and $13\times 10^{19}$ W. For short-period planets, the observational signature of internal heat (either remnant gravitational energy from formation or ongoing tidal heating) will likely be an inflated radius rather than additional emergent flux."
 If hot Jupiters exist for which the tidal heating 1s : vsizeable fraction of the insolation. then our nodel will have 1uder-estimated the zeroetbl-order (DC) conrponeut of tlicav lighteurves.," If hot Jupiters exist for which the tidal heating is a sizeable fraction of the insolation, then our model will have under-estimated the zeroeth-order (DC) component of their lightcurves."
 By the same token. we also neglect the remnant heat of formation. which is ikelv less than 100 TS for a Covrs-oldl panet (Burrowsetal.2006).," By the same token, we also neglect the remnant heat of formation, which is likely less than 100 K for a Gyrs-old planet \citep[][]{Burrows_2006}."
. One uportaut consequencee of ueelecting auv interual heatiis and having a siuple ouc-laver ποσο. is that the wieht-side equilibri temperature is zero.," One important consequence of neglecting any internal heating and having a simple one-layer model, is that the night-side equilibrium temperature is zero."
 For the uost part this is uot important. since ouwcels pass throueh the ihuninated hemisphere often enough to never get very cool.," For the most part this is not important, since parcels pass through the illuminated hemisphere often enough to never get very cool."
 The exceytion is for our uodels of ΠΟ SOG06b.. where the pliuxt spends such a long time (>LOO days) far fro its ιο star.," The exception is for our models of HD 80606b, where the planet spends such a long time $>100$ days) far from its host star."
 It is vossible that the xx‘al planet never cools to the LOris seen iu our tov 106el., It is possible that the real planet never cools to the 400 K seen in our toy model.
 Our model neglects couduclon of heat from one cell to its neighbors., Our model neglects conduction of heat from one cell to its neighbors.
 As such. our model is onlyrelevant when advection aud radialon are the dominant euergv transport mechanisms.," As such, our model is onlyrelevant when advection and radiation are the dominant energy transport mechanisms."
 This is probably hne case for hot .hpitersο which are thoieht to have strong wills.," This is probably the case for hot Jupiters, which are thought to have strong winds."
 If hea conduction nmuporaut ou hot .hpiters. then thev nav ¢hibit sualler daxwelt teniperature οςuntrasts aud the relation between the thermal pise offset and the thermal phase amplitude will be bxken (see first Equation 12).," If heat conduction important on hot Jupiters, then they may exhibit smaller day--night temperature contrasts and the relation between the thermal phase offset and the thermal phase amplitude will be broken (see first Equation 12)."
 Our noctel assumes that all the incident ewerey from the hos star is either reflected or absorbed. and the euerey that is absorbe oulv gocs iuto Increasing the teryature of the pareels of gas.," Our model assumes that all the incident energy from the host star is either reflected or absorbed, and the energy that is absorbed only goes into increasing the temperature of the parcels of gas."
 In effect. this iuxnuts to neglecting fιο work term in the First Law of Thermodvuaiics.," In effect, this amounts to neglecting the work term in the First Law of Thermodynamics."
 In the steady-sate. the power absorbe by the planet must be halanced by the power entted.," In the steady-state, the power absorbed by the planet must be balanced by the power emitted."
 That is to sav. fιο ultimate effect of work is to Warla up parcels of Ooea8 (e.g.parcelsspeeduptoCoodman 2009).," That is to say, the ultimate effect of work is to warm up parcels of gas \citep[e.g., parcels speed up to super-sonic velocities, shock and deposit energy as heat;][]{Goodman_2009}."
. This nuplicitlv neelec‘ts four effects that may very well 1IC OCCTring: enerev nay instead co into latent heat to vaporize particulates: thermal tides (Arras&Socrates2010): increasing the atmospheric scale height (Caullot&5LOWLl2002): or spoeedius up winds., This implicitly neglects four effects that may very well be occurring: energy may instead go into latent heat to vaporize particulates; thermal tides \citep{Arras_2009}; increasing the atmospheric scale height \citep{Guillot_2002}; or speeding up winds.
 Any of these “work” ternis would tend to dampen the temperature fluctuations experienced by a parcel of Sgas: they act as energv⊓⋅ sinks near periastron and may :let as CLOTEV SOTPCOS Lear apoastron., Any of these “work” terms would tend to dampen the temperature fluctuations experienced by a parcel of gas: they act as energy sinks near periastron and may act as energy sources near apoastron.
 As such our mode will under-estimate Sas response time. over-estimate the temperature excursions. aud hence the peak-to-trough amplitiue of disc-iutegrated hehtcurves.," As such our model will under-estimate gas response time, over-estimate the temperature excursions, and hence the peak-to-trough amplitude of disc-integrated lightcurves."
 The plaiet mnost likely o show these effects is ΠΙΟ S0606b. sncο realistically uuch of the energy absorbed at periastrOu must eo inte) speeding up winds. inflating the atiosprere. etc.," The planet most likely to show these effects is HD 80606b, since realistically much of the energy absorbed at periastron must go into speeding up winds, inflating the atmosphere, etc."
 Indece. the peak-to-trough phase amplitudes wo5 compute for this planet are roughly a factor of 2 ereater than those from lyvdrodvuamic simulations (Laugto1&Laughin25).," Indeed, the peak-to-trough phase amplitudes we compute for this planet are roughly a factor of 2 greater than those from hydrodynamic simulations \citep{Langton_2008}."
 {)r toy model neglects advection towards the planet'* poles: effectively we have zonal but uo mucridiona winds., Our toy model neglects advection towards the planet's poles: effectively we have zonal but no meridional winds.
 At low (aubar) pressures. there almost certainly is iieridional flow away from the sub-stellar point. but the low presures lead that only a sinall fraction of absorbedpower is moved in this way.," At low (mbar) pressures, there almost certainly is meridional flow away from the sub-stellar point, but the low presures mean that only a small fraction of absorbedpower is moved in this way."
 But even when norklional flows are uot cficieut. latitudinal heat transport may still occur in the form of eddy fluxes (c.e..Choctal.20082: 2008).. which we have also inpicitly neglected.," But even when meridional flows are not efficient, latitudinal heat transport may still occur in the form of eddy fluxes \citep[e.g.,][]{Cho_2008, Cho_2008b}, , which we have also implicitly neglected."
 Efficicut equatorpole heat transport on a planet would have, Efficient equator–pole heat transport on a planet would have
mechanism is discussed in Sect.,mechanism is discussed in Sect.
 2? where we show that these high-energy tons can lead to power-law ton tails. however with a power-index of —4.," \ref{sect5} where we show that these high-energy ions can lead to power-law ion tails, however with a power-index of $-4$."
 The results are discussed and compared with observations in Sect., The results are discussed and compared with observations in Sect.
 2? Challenged by recent results concerning ton spectra at large distances (2222). we look into the problem of PUI phasesspace transport under these new given auspices.," \ref{sect6}
 Challenged by recent results concerning ion spectra at large distances \citep{mcdonald03,decker05,kiraly05,fisk06}, we look into the problem of PUI space transport under these new given auspices."
" First we discussthe argument given by ??. thàt pick-up tons under resonant interaction with ambient compressive turbulences enter a quasi-equilibrium state with saturated power-law distributions of a somehow sacrosanct spectral velocity index of y,=—5."," First we discussthe argument given by \citet{fisk06,fisk07} that pick-up ions under resonant interaction with ambient compressive turbulences enter a quasi-equilibrium state with saturated power-law distributions of a somehow sacrosanct spectral velocity index of $\gamma _{v}=-5$."
" The related. well known Kolmogorov formalism is based on à ""dimensional"" reasoning: The problem concerning the energy distribution in eddies of a typical wave number & is considered using two different dimensional quantities: namely the spectral energy density Ej and the wave number &."," The related, well known Kolmogorov formalism is based on a “dimensional” reasoning: The problem concerning the energy distribution in eddies of a typical wave number $k$ is considered using two different dimensional quantities: namely the spectral energy density $E_{k}$ and the wave number $k$."
 The spectral energy flux is then defined by and has to be constant for stationary cases., The spectral energy flux is then defined by and has to be constant for stationary cases.
 τε is the typical life time of eddies with scale ἐξ27/k and is given by with the typical vortex velocity v; given by Requiring a constant energy flux S. given by Eq.(1)). with τε given by Eq.(2)) then automatically leads to the Kolmogorov spectrum in the form E; ~£.," $\tau _{k}$ is the typical life time of eddies with scale $\lambda _{k}=2\pi /k$ and is given by with the typical vortex velocity $v_{k}$ given by Requiring a constant energy flux $S$ given by \ref{energy_flux}) ), with $\tau _{k}$ given by \ref{tauk}) ) then automatically leads to the Kolmogorov spectrum in the form $E_{k}$ $\sim k^{-5/3}$."
 In their thermodynamical approach to particle spectra in equilibrium with waves. ??.. however. only consider one single dimensional quantity. te. the kinetic energy 7. because in their case the energy distribution F has the dimension IT].," In their thermodynamical approach to particle spectra in equilibrium with waves, \citet{fisk06,fisk07}, however, only consider one single dimensional quantity, i.e. the kinetic energy $T$, because in their case the energy distribution $\digamma$ has the dimension }$."
 The flux combination is then expectedAf by them to be constant and. thus. the energy gain Is given by and hence the above requirements result in the requirement thatF~T7.," The flux combination is then expected by them to be constant and, thus, the energy gain is given by and hence the above requirements result in the requirement that $\digamma
\sim T^{-2}$."
 This shows that the analogy of the Fisk-Gloeckler approach with the Kolmogorov formalism is not complete. since in their theory the quantity Aris not specified.," This shows that the analogy of the Fisk-Gloeckler approach with the Kolmogorov formalism is not complete, since in their theory the quantity $\Delta t$ is not specified."
 Only if we play the same game with the Kolmogorov turbulence using instead of Eq.(2)) the assumption 7=Gar)|. where w is some independent external speed (especially independent on Κ). we will find from Eq.(1)) also the result analogous to the result obtained by Fisk-Gloeckler.," Only if we play the same game with the Kolmogorov turbulence using instead of \ref{tauk}) ) the assumption $\tau =(ku)^{-1}$, where $u$ is some independent external speed (especially independent on $k$ ), we will find from \ref{energy_flux}) ) also the result analogous to the result obtained by Fisk-Gloeckler."
 The problem may be briefly inspected here whether or not some spatial/temporal disturbances in the solar wind plasma canbe considered às waves., The problem may be briefly inspected here whether or not some spatial/temporal disturbances in the solar wind plasma canbe considered as waves.
 They should be considered às waves. If the convection time Tur/U is much greater than the passage period of these waves in the wind reference frame r.=2g/(kova) where r denotes the radial distance to the sun and U the solar wind bulk velocity.," They should be considered as waves, if the convection time $\tau
_{\mathrm{conv}}=r/U$ is much greater than the passage period of these waves in the wind reference frame $\tau _{\sim }=2\pi /(k_{0} v_{\mathrm A})$ where $r$ denotes the radial distance to the sun and $U$ the solar wind bulk velocity."
 The Alfvénn speed is designated as v4 with the corresponding wavevector Ko for the turbulence., The Alfvénn speed is designated as $v_{\mathrm A}$ with the corresponding wavevector $k_0$ for the turbulence.
" For aninequality of the kind τ.<fia. this then leads to the definition of the principal wave turbulence correlation scale given by Ly,«rva/U=0.1]AU."," For aninequality of the kind $\tau _{\sim }\ll t_{\mathrm{conv}}$, this then leads to the definition of the principal wave turbulence correlation scale given by $%
L_{\mathrm m}<rv_{\mathrm A}/U\approx 0.1\,\mathrm{AU."
 A more exact definition of the turbulence correlation scale Is given in the paper of ?.., A more exact definition of the turbulence correlation scale is given in the paper of \citet{chashei03}.
 This value is also i accordance with the value used by ? Του his estimation of the upper possible velocity border., This value is also in accordance with the value used by \citet{fahr07b} for his estimation of the upper possible velocity border.
 Some measurements indicate smaller values for the turbulence correlation length than calculated by us in this paper (seee.g.?).., Some measurements indicate smaller values for the turbulence correlation length than calculated by us in this paper \citep[see e.g.][]{matthaeus05}.
 However. the following conclusions in our paper drawn on the effects of turbulent heating of suprathermal protons become even less promising for an MHD-equilibrium to be established. if smaller correlation lengths prevail.," However, the following conclusions in our paper drawn on the effects of turbulent heating of suprathermal protons become even less promising for an MHD-equilibrium to be established, if smaller correlation lengths prevail."
 One can expect that at larger distances. ry>1AU. the value for {μι increases proportionally to r. since a similar dependence is also valid for the outer scale of turbulence [d (seeο).," One can expect that at larger distances, $r\geq 1\,\mathrm{AU}$, the value for $L_{\mathrm m}$ increases proportionally to $r$, since a similar dependence is also valid for the outer scale of turbulence $k_{0}^{-1}$ \citep[see][]{chashei03}."
" Here we estimate the maximum energy for protons resonating with these largest scales ζω at distances of about 100 AU and find with ving=O(IO0AU)-L4,(00ΑΙΥΕ)O(IAUD-Li(lΑΙyO nas). where vaa. ¥YO'max) and Q are the maximal ton speed. the associated Lorentz factor and the ion gyrofrequency. respectively."," Here we estimate the maximum energy for protons resonating with these largest scales $L_{\mathrm m}$ at distances of about 100 AU and find with $v_{\max }\simeq \Omega (100\,\mathrm{AU})\cdot
L_{\mathrm m}(100\,\mathrm{AU})/\gamma (v_{\max })\simeq \Omega (1\,\mathrm{AU})\cdot L_{\mathrm m}(1\,\mathrm{AU})/\gamma
(v_{\max })$ , where $v_{\max }$, $\gamma (v_{\max })$ and $\Omega $ are the maximal ion speed, the associated Lorentz factor and the ion gyrofrequency, respectively."
 Protons with these speeds have energies of about Εν«1GeV independent on distance.," Protons with these speeds have energies of about $E_{\max
}\approx 1\,\mathrm{GeV}$ independent on distance."
 Thus. ions with E.<Ei; can be expected to be scattered by waves. whereas tons with E>Ej44 can only be scattered by large scale velocity structures in the solar wind.," Thus, ions with $E\leq E_{\max }$ can be expected to be scattered by waves, whereas ions with $E\geq E_{\max }$ can only be scattered by large scale velocity structures in the solar wind."
 Since in any case. however. power-laws seem to be a fact well supported by observations within a certain range of solar distances and ion energies. we shall nevertheless take this finding as serious here as deserved and determine in the following the absolute spectral intensity of this PUL power-law distribution and its consequences.," Since in any case, however, power-laws seem to be a fact well supported by observations within a certain range of solar distances and ion energies, we shall nevertheless take this finding as serious here as deserved and determine in the following the absolute spectral intensity of this PUI power-law distribution and its consequences."
" If the PUI-distribution is given in the form with the velocity power index y,=—5. then the PUI density is given by where fnio=fpio() is a local normalization value. and vo and v4 are lower and upper velocity limits of the quasistationary PUL power law."," If the PUI-distribution is given in the form with the velocity power index $\gamma _{v}=-5$, then the PUI density is given by where $f_{\mathrm{pui},0}=f_{\mathrm{pui},0}(r)$ is a local normalization value, and $v_{0}$ and $v_{\infty }$ are lower and upper velocity limits of the quasistationary PUI power law."
 Using w= v/vo. one obtains for the power-law PUL pressure as the second moment of the distribution function the following result Evidently. the definition of Py(7) requires. the determination of all three local values. fpuioG). vo. and vod) Which we aim at below.," Using $\psi =v/v_{0}$ , one obtains for the power-law PUI pressure as the second moment of the distribution function the following result Evidently, the definition of $P_{\mathrm{pui}}(r)$ requires the determination of all three local values $f_{\mathrm{pui},0}(r)$ , $v_{0}(r)$ , and $v_{\infty }(r)$ which we aim at below."
light evinder and the propagation effects on (he scattered radiation would possibly release the requirement of a nearly orthogonal rotator ancl allow the scattered component to lag the MP by less than 1807 to account Lor the postcursors in the pulsar profiles.,light cylinder and the propagation effects on the scattered radiation would possibly release the requirement of a nearly orthogonal rotator and allow the scattered component to lag the MP by less than $180^\circ$ to account for the postcursors in the pulsar profiles.
 The induced Compton scattering of pulsar radio emission off the secondary plasma parücles inside (he open field line tube may play a significant. role., The induced Compton scattering of pulsar radio emission off the secondary plasma particles inside the open field line tube may play a significant role.
 The magnetic field of a pulsar affects the scattering process substantially., The magnetic field of a pulsar affects the scattering process substantially.
 At distances of order of the emission altitude. the magnetic field is tvpically strong enough for the scattering to occur in the longitudinal regime. ywfv«1.," At distances of order of the emission altitude, the magnetic field is typically strong enough for the scattering to occur in the longitudinal regime, $\gamma^2\nu^{\prime^4}/\nu_G^4\ll 1$."
 Then the photons of the pulsar radio beam are predominantly scaltered into the state with 84...cL/5 and νιev0752?>p. ie. the scattered radiation is almost aligned with the ambient magnetic field.," Then the photons of the pulsar radio beam are predominantly scattered into the state with $\theta_{1_{\rm max}}\sim 1/\gamma$ and $\nu_1\sim\nu\theta^2\gamma^2\gg \nu$, i.e. the scattered radiation is almost aligned with the ambient magnetic field."
 As the magnetic field strength decreases with distance from the neutron star. al higher altitudes the longitudinal scattering regime changes for the transverse one. which holds on condition that 1/5?<vhu.«1.," As the magnetic field strength decreases with distance from the neutron star, at higher altitudes the longitudinal scattering regime changes for the transverse one, which holds on condition that $1/\gamma^2\ll\nu^{\prime^4}/\nu_G^4\ll 1$."
" In this reeime. (he orientations of the scattered photons are mostly antiparallel to the magnetic field. (4,=z. and v,=v?/4«&v."," In this regime, the orientations of the scattered photons are mostly antiparallel to the magnetic field, $\theta_{1_{\rm max}}=\pi$, and $\nu_1=\nu\theta^2/4\ll\nu$."
 The radiation scattered in (he two regimes may [rom separate components in the pulse profile. which can be identified with the precursor and the IP.," The radiation scattered in the two regimes may from separate components in the pulse profile, which can be identified with the precursor and the IP."
 Apart from explaining the geometrical location of these components in the pulse profile. our model suggests the MP-IP ancl MP-precursor connections as well as can account [ον the peculiar properties of the emission components outside of the MP.," Apart from explaining the geometrical location of these components in the pulse profile, our model suggests the MP-IP and MP-precursor connections as well as can account for the peculiar properties of the emission components outside of the MP."
 In the geometrical aspect. our model is akin to the bidirectional model of (2005).," In the geometrical aspect, our model is akin to the bidirectional model of \citet{d05}."
. In that model. the precursor and IP originate at a certain location in the ouler magnetosphere and (he pulsar is approximately an orthogonal rotator.," In that model, the precursor and IP originate at a certain location in the outer magnetosphere and the pulsar is approximately an orthogonal rotator."
 Because of the ultrarelativistic outstream of the plasma particles along the field lines. anv conceivable emission mechanism would generate (he outward radiation directed along the magnetic field in (he corotating frame.," Because of the ultrarelativistic outstream of the plasma particles along the field lines, any conceivable emission mechanism would generate the outward radiation directed along the magnetic field in the corotating frame."
 The inward radiation is assumed to be antiparallel to the outwarel one., The inward radiation is assumed to be antiparallel to the outward one.
 Thus. the orientations of the emission components and their locations in the pulse prolile are roughly (he same as in our model.," Thus, the orientations of the emission components and their locations in the pulse profile are roughly the same as in our model."
 Note. however. that in our case (he precursor and the IP result from the scattering in the two distinct regimes and therefore originate at somewhat different altitudes.," Note, however, that in our case the precursor and the IP result from the scattering in the two distinct regimes and therefore originate at somewhat different altitudes."
 Besicles that. the scattering sites are restricted to the region of the radio bean passage in the rotating magnetosphere.," Besides that, the scattering sites are restricted to the region of the radio beam passage in the rotating magnetosphere."
 Correspondingly. (he positions of (he precursor and the IP in the pulse profile are tightly connected to that of the MP.," Correspondingly, the positions of the precursor and the IP in the pulse profile are tightly connected to that of the MP."
star was exposed lor a total of about 5 hours. vielding S/N per 0.2 Apixel of 100.300 in the co-adcded spectra. (,"star was exposed for a total of about 5 hours, yielding S/N per 0.2 pixel of 100–300 in the co-added spectra. ("
Ro = 18.000 = 1.9pix for the red fibers and 13.000 = 2.5pix for the blue fibers.),"R = 18,000 = 1.9pix for the red fibers and 13,000 = 2.5pix for the blue fibers.)"
 Data reduction emploved standard IRAF techniques dohyelra)., Data reduction employed standard IRAF techniques dohydra).
 Fits to the daytime skv spectra were used (o correct. for the throughput. variations between fibers and sky sublvaction used an average of ~20 skv fibers per exposure., Fits to the daytime sky spectra were used to correct for the throughput variations between fibers and sky subtraction used an average of $\sim20$ sky fibers per exposure.
 Effective temperatures were determined using WOCS BYRI CCD photometry (Delivannisetal.2004a2)., Effective temperatures were determined using WOCS BVRI CCD photometry \citep{d04}.
. An empirical relation was made between the /3—V color of the fiducial single-star sequence and the other colors (B—HR.DI.VR.IF.ER E). and an ellective D—V. was ealeulated for each.," An empirical relation was made between the $B-V$ color of the fiducial single-star sequence and the other colors $B-R,B-I,V-R,V-I,R-I$ ), and an effective $B-V$ was calculated for each."
 These six ellective B— Vs were then averaged. and this value was converted into T;jj using the (B.—V) [Fe/H] Tijg velation in DSJO2.," These six effective $B-V$ 's were then averaged, and this value was converted into $T_{eff}$ using the $(B-V)$ $[Fe/H]$ $T_{eff}$ relation in DSJ02."
 For an input [Fe/H] in that relation. we used ihe M35 cluster average of —0.14340.014 which agrees nicely with the value —0.2140.10 from Darradov.Navascués.Delivannis.& BDSOL)..," For an input [Fe/H] in that relation, we used the M35 cluster average of $-0.143\pm0.014$ which agrees nicely with the value $-0.21\pm0.10$ from \citet[][hereafter, BDS01]{b01}."
 Then. log ο is taken from Y? isochrones (Yietal.2001).. and £ is determined using Edvardssonetal.(1993)..," Then, log g is taken from $Y^2$ isochrones \citep{yy01}, and $\xi$ is determined using \citet{edvard}."
" The small contribution to the Li equivalent. widths. W(Li). trom the Fe I line at τοῦ, was removed using which was determined empirically from high resolution and very high S/N IIvacdes Li data of Thorburn.etal. (1993)..."," The small contribution to the Li equivalent widths, $W(Li)$, from the Fe I line at 6707.45 was removed using which was determined empirically from high resolution and very high S/N Hyades Li data of \citet{thdp}. ."
 Absolute Li abundances. A(Li).. were determined bv interpolating on the curves of growth of S04.," Absolute Li abundances, $A(Li)$, were determined by interpolating on the curves of growth of S04."
 The error in W(Li). ey. is derived using the relation from (1993).. and depends on the S/N and FWIIM of each Li line.," The error in W(Li), ${\sigma}_W$, is derived using the relation from \citet{dpd93}, and depends on the S/N and FWHM of each Li line."
 Upper limits are simply ὅσμ.., Upper limits are simply ${\sigma}_W$.
" The Li line is also sensitive to {νε but not to log g or £: in the gap. ATi,=100 implies (Li)=0.08decx."," The Li line is also sensitive to $T_{eff}$, but not to log g or $\xi$: in the gap, ${\Delta}{T_{eff}}=100K$ implies ${\Delta}A(Li)=0.08dex$."
 For the WOCS photometry of our stars in a given filter. (vpical σι from multiple measures are less than 0.01mag.," For the WOCS photometry of our stars in a given filter, typical ${\sigma}_{\mu}$ from multiple measures are less than 0.01mag."
" Tvpical internal errors in the 7;j. propagated from o, of the averaged B— Vs. ave 10.50IxX for known single stars."," Typical internal errors in the $T_{eff}$, propagated from ${\sigma}_{\mu}$ of the averaged $B-V$ 's, are 10–50K for known single stars."
" The reported σι in Table 1 has contributions from oy and 05;,,. added in quadrature."," The reported ${\sigma}_A$ in Table 1 has contributions from ${\sigma}_W$ and ${\sigma}_{T_{eff}}$, added in quadrature."
 Systematic errors due to E(D—V) and the Tipp scale itself may be at least as largeas LOOK.but since they would affect the whole ensemble in essentially the same (small) way. they do not affect our conclusions.," Systematic errors due to $E(B-V)$ and the $T_{eff}$ scale itself may be at least as largeas 100K,but since they would affect the whole ensemble in essentially the same (small) way, they do not affect our conclusions."
ines observed may have to do with their production at larger disc radii than usual.,lines observed may have to do with their production at larger disc radii than usual.
 Of course we cannot vet rule out the »ossibility that QU Car might be oriented nearly pole on., Of course we cannot yet rule out the possibility that QU Car might be oriented nearly pole on.
 We have to leave as an unsolved mystery the origin of he short-term variability within the spectral line profiles (section 4)) that has been revealed thanks to the HIST/S'TTIS TIME-EACG mode of data-taking., We have to leave as an unsolved mystery the origin of the short-term variability within the spectral line profiles (section \ref{s:timevar}) ) that has been revealed thanks to the HST/STIS TIME-TAG mode of data-taking.
 Since this variability is not present at all epochs and indeed. presents with cillerent o»operties on the two occasions we have noticed. it. it is unlikely to be connected. to anything as fundamental as the binary orbit.," Since this variability is not present at all epochs and indeed presents with different properties on the two occasions we have noticed it, it is unlikely to be connected to anything as fundamental as the binary orbit."
 Our third set of cata hints at a repetition timescale of around. an hour if. that is. the detected absorption profile disturbances do in fact. repeat.," Our third set of data hints at a repetition timescale of around an hour if, that is, the detected absorption profile disturbances do in fact repeat."
 To progress any further. another more intensive campaign of spectroscopic observation would be required.," To progress any further, another more intensive campaign of spectroscopic observation would be required."
 Lastlv. we have not vet been able to begin to find an explanation for the remarkable inverse P Cveni character of. most notably. the," Lastly, we have not yet been able to begin to find an explanation for the remarkable inverse P Cygni character of, most notably, the"
of nearby stars for low mass companions at large orbital separations.,of nearby stars for low mass companions at large orbital separations.
" Over the last decade, a handful of objects close to the planetary mass regime have been imaged with existing instruments, such as 2M 1207 b (?),, DH Tau B (?),, GQ Lup b (?),, AB Pic b (?),, CHXR Τὸ B (?),, and more recently Fomalhaut b (?),, 1RXS J1609 b (?),, 8 Pic b and the triple system around HR 8799 (?).."," Over the last decade, a handful of objects close to the planetary mass regime have been imaged with existing instruments, such as 2M 1207 b \citep{chauvin2005a}, DH Tau B \citep{itoh2005}, GQ Lup b \citep{neuhauser2005}, AB Pic b \citep{chauvin2005b}, , CHXR 73 B \citep{luhman2006}, and more recently Fomalhaut b \citep{kalas2008}, 1RXS J1609 b \citep{lafreniere2008}, $\beta$ Pic b \citep{lagrange2008} and the triple system around HR 8799 \citep{marois2008b}."
" However, the large uncertainty on the mass of these objects may place some of them in the sub-stellar rather than the planetary mass regime."," However, the large uncertainty on the mass of these objects may place some of them in the sub-stellar rather than the planetary mass regime."
 The next generation of planet finding instruments currently being built will combine: (i) extreme AO systems with a large number of actuators (??7) to reach very high corrections in the near-infrared (??);; and (ii) high-efficiency coronagraphs such as the apodized pupil Lyot coronagraph (? and references therein) or the achromatic 4-quadrant phase mask (??) to obtain optimal star extinction.," The next generation of planet finding instruments currently being built will combine: (i) extreme AO systems with a large number of actuators \citep{angel1994,stahl1995,langlois2001} to reach very high corrections in the near-infrared \citep{fusco2006,allercarpentier2008}; and (ii) high-efficiency coronagraphs such as the apodized pupil Lyot coronagraph \citealt{soummer2005} and references therein) or the achromatic 4-quadrant phase mask \citep{rouan2000,mawet2006} to obtain optimal star extinction."
 GPI (Gemini Planet Imager) for Gemini South (?) and SPHERE (Spectro-Polarimetric High-contrast Exoplanet REsearch) for the ESO-VLT (?) are the two leading instruments of that category., GPI (Gemini Planet Imager) for Gemini South \citep{macintosh2006} and SPHERE (Spectro-Polarimetric High-contrast Exoplanet REsearch) for the ESO-VLT \citep{beuzit2006} are the two leading instruments of that category.
" These will both start operation in 2011, along with HiCIAO (High-contrast Coronagraphic Imager for Adaptive Optics) for Subaru (?).."," These will both start operation in 2011, along with HiCIAO (High-contrast Coronagraphic Imager for Adaptive Optics) for Subaru \citep{hodapp2008}."
 They will aim at detecting exoplanets down to the Jupiter mass (M;up)) around nearby young stars by reaching contrast values of 15 to 17.5 mag (1078 to 1077) at angular separations of ~0.1”.., They will aim at detecting exoplanets down to the Jupiter mass ) around nearby young stars by reaching contrast values of 15 to 17.5 mag $10^{-6}$ to $10^{-7}$ ) at angular separations of $\sim$.
" Both GPI and SPHERE will incorporate diffraction limited integral field spectrographs (IFS) in the near infrared, allowing to obtain images simultaneously at several wavelengths."," Both GPI and SPHERE will incorporate diffraction limited integral field spectrographs (IFS) in the near infrared, allowing to obtain images simultaneously at several wavelengths."
" SPHERE will also incorporate a differential imager named IRDIS (InfraRed Dual Imaging Spectrograph, ?)) that will provide simultaneous images at two close wavelengths in either one of its 5 different filter pairs over the Y to Ks bands."," SPHERE will also incorporate a differential imager named IRDIS (InfraRed Dual Imaging Spectrograph, \citealt{dohlen2008a}) ) that will provide simultaneous images at two close wavelengths in either one of its 5 different filter pairs over the Y to Ks bands."
" These instruments will allow to use different observing strategies such as Spectral Differential Imaging (SDI, ?)) or Angular Differential Imaging (ADI, ?)) to obtain data which will be analyzed using advanced methods such as Spectral Deconvolution (?) or Localy Optimized Combination of Images (LOCI, ?)) for IRDIS data."," These instruments will allow to use different observing strategies such as Spectral Differential Imaging (SDI, \citealt{racine1999}) ) or Angular Differential Imaging (ADI, \citealt{marois2006a}) ) to obtain data which will be analyzed using advanced methods such as Spectral Deconvolution \citep{sparks2002} or Localy Optimized Combination of Images (LOCI, \citealt{lafreniere2007}) ) for IRDIS data."
 Specific methods of signal extraction have also been developed within the SPHERE consortium (??) to be used in the data reduction pipeline of IRDIS.," Specific methods of signal extraction have also been developed within the SPHERE consortium \citep{mugnier2008,smith2009} to be used in the data reduction pipeline of IRDIS."
 Data analysis methods are of extreme importance when it comes to the detection of faint objects in coronagraphic images dominated by speckles., Data analysis methods are of extreme importance when it comes to the detection of faint objects in coronagraphic images dominated by speckles.
" In particular, the precise estimation of the object flux after applying these methods is critical for calibration of planetary mass objects model atmospheres (???,, and in preparation; ??7)) based on effective temperature (Tess and surface gravity g)), and the corresponding evolutionary)) models (???7).."," In particular, the precise estimation of the object flux after applying these methods is critical for calibration of planetary mass objects model atmospheres \citealt{allard2001,allard2003,allard2007}, and in preparation; \citealt{burrows2006,ackerman2001,tsuji2005}) ) based on effective temperature ) and surface gravity ), and the corresponding evolutionary models \citep{burrows1997,chabrier2000,baraffe2003,fortney2005}."
" The only currently available method for mass estimation, when no dynamical mass estimations are known for the object, consists in incorporating photometric measurements and their error assessment into these physical models."," The only currently available method for mass estimation, when no dynamical mass estimations are known for the object, consists in incorporating photometric measurements and their error assessment into these physical models."
" Although spectroscopy is the method of choice for characterization, it may not be possible to obtain high-quality spectra of extremely faint sources (?7).."," Although spectroscopy is the method of choice for characterization, it may not be possible to obtain high-quality spectra of extremely faint sources \citep{vigan2008,janson2010}."
" In this paper we investigate the photometric limitations in high-contrast data obtained at different wavelengths with a dual-band imager like IRDIS, and we study how it translates in terms of characterization of planetary mass objects."," In this paper we investigate the photometric limitations in high-contrast data obtained at different wavelengths with a dual-band imager like IRDIS, and we study how it translates in terms of characterization of planetary mass objects."
 After reminding in Sect., After reminding in Sect.
" 2 the origin of the speckle noise which is the fundamental limitation in high-contrast imaging and the methods to overcome it, we present in Sect."," \ref{sec:limitations_high_contrast_imaging} the origin of the speckle noise which is the fundamental limitation in high-contrast imaging and the methods to overcome it, we present in Sect."
 3 the end-to-end simulations of IRDIS which were performed to obtain a realistic observing sequence., \ref{sec:end_to_end_simulations} the end-to-end simulations of IRDIS which were performed to obtain a realistic observing sequence.
 Section 4 briefly describes the detection limits obtained with ADI and SDI+ADI data analysis methods before studying the performances in terms of photometric accuracy., Section \ref{sec:photometric_accuracy} briefly describes the detection limits obtained with ADI and SDI+ADI data analysis methods before studying the performances in terms of photometric accuracy.
" Finally, in Sect."," Finally, in Sect."
 5 we advocate a filter pair procedure for IRDIS and present an analysis of the characterizations which will be possible from aperture photometry., \ref{sec:photometric_characterization} we advocate a filter pair procedure for IRDIS and present an analysis of the characterizations which will be possible from aperture photometry.
 Detecting very faint planetary objects requires to obtain diffraction limited images with high-order AO systems in order to overcome the large contrast ratio between the star and the planet with coronagraphy., Detecting very faint planetary objects requires to obtain diffraction limited images with high-order AO systems in order to overcome the large contrast ratio between the star and the planet with coronagraphy.
" In high-Strehl ratio coronagraphic images, the factor that limits the accessible dynamical range is the speckle noise (7) induced by atmospheric phase residuals and instrumental quasi-static aberrations not corrected by the AO system."," In high-Strehl ratio coronagraphic images, the factor that limits the accessible dynamical range is the speckle noise \citep{soummer2007} induced by atmospheric phase residuals and instrumental quasi-static aberrations not corrected by the AO system."
 The quasi-static speckles are caused by the instrumental aberrations that slowly change during a long exposure., The quasi-static speckles are caused by the instrumental aberrations that slowly change during a long exposure.
" Telescope orientation, temperature variations or rotating optical elements cause small mechanical variations in the optical elements which make the speckle pattern evolve."," Telescope orientation, temperature variations or rotating optical elements cause small mechanical variations in the optical elements which make the speckle pattern evolve."
" Long time exposures are typically decomposed in a series of short exposure images of a few seconds during which the atmospheric residuals are averaged out, forming a smooth halo over which the quasi-static speckles are superimposed, because their coherence time is much longer than the atmospheric residuals (???).."," Long time exposures are typically decomposed in a series of short exposure images of a few seconds during which the atmospheric residuals are averaged out, forming a smooth halo over which the quasi-static speckles are superimposed, because their coherence time is much longer than the atmospheric residuals \citep{langlois1998,macintosh2005,hinkley2007}."
" To optimize AO performances and speckle rejection, observations can be performed in pupil stabilized mode, leading to a very high stability of the star Point Spread Function (PSF) and a slow rotation of the field of view during the observations, at a rate which depends on the star position in the sky."," To optimize AO performances and speckle rejection, observations can be performed in pupil stabilized mode, leading to a very high stability of the star Point Spread Function (PSF) and a slow rotation of the field of view during the observations, at a rate which depends on the star position in the sky."
" The speckle noise can be reduced by subtracting a reference PSF from each science frame in order to remove the star halo and speckles, and possibly reveal a faint planetary object."," The speckle noise can be reduced by subtracting a reference PSF from each science frame in order to remove the star halo and speckles, and possibly reveal a faint planetary object."
 This reference PSF can be obtained by observing a reference star taken in the same observing conditions (parallactic angle and atmospheric conditions if possible) as the original target to reproduce a similar pattern of quasi-static speckles., This reference PSF can be obtained by observing a reference star taken in the same observing conditions (parallactic angle and atmospheric conditions if possible) as the original target to reproduce a similar pattern of quasi-static speckles.
" This is a very time consuming task because the time spent on the reference star is equal to the one spent on the target to precisely match both PSF, and the aberrations between the two stars cannot be exactly reproduced."," This is a very time consuming task because the time spent on the reference star is equal to the one spent on the target to precisely match both PSF, and the aberrations between the two stars cannot be exactly reproduced."
 The reference can also be built from the science frames by using either spectral or angular information., The reference can also be built from the science frames by using either spectral or angular information.
 The SDI method has first been proposed by ? for faint companion detection., The SDI method has first been proposed by \citet{racine1999} for faint companion detection.
" It has been extensively studied (?),, and subsequently tested on sky with TRIDENT on the CFHT (?) and with NACO on the VLT (?).."," It has been extensively studied \citep{marois2000}, and subsequently tested on sky with TRIDENT on the CFHT \citep{marois2005} and with NACO on the VLT \citep{lenzen2004}."
" The technique relies on the fact that planetary objects have large intrinsic molecular features in their spectrum, while the host star has a relatively flat spectrum."," The technique relies on the fact that planetary objects have large intrinsic molecular features in their spectrum, while the host star has a relatively flat spectrum."
"By taking simultaneously two images of a system at two close wavelengths located around one of these sharp features and subtracting them, the star contribution can be partially eliminated, and the planet","By taking simultaneously two images of a system at two close wavelengths located around one of these sharp features and subtracting them, the star contribution can be partially eliminated, and the planet"
"Lhe supernova SN 2002ap (aso00 = 017367 237.92: 040) = |15’ 45° 13""52) ) was discovered on 2002 January 29.4 UT ov X. Hirose at V14.5 mag. in the outer region of the nearby spira ALTA (Nakano et 22002).","The supernova SN 2002ap $\alpha_{2000}$ = $01^{h}36^{m}23^{s}.92 $ ; $\delta_{2000}$ = $+15^{\circ}$ $45^{\prime}$ $13^{\prime\prime}.3$ ) was discovered on 2002 January 29.4 UT by Y. Hirose at $V \sim 14.5$ mag, in the outer region of the nearby spiral M74 (Nakano et 2002)."
 Low resolution Esvectra of. SN. 2002ap obtained during 2002 January 30EJ31 ]xinugasao et al., Low resolution spectra of SN 2002ap obtained during 2002 January 30–31 (Kinugasa et al.
 2002a.b: Meikle et al.," 2002a,b; Meikle et al."
 2002: Filippenko ο.‘hornoc.x 202) were similar to those of type Ic supernovae., 2002; Filippenko Chornock 2002) were similar to those of type Ic supernovae.
 OWOVOL. the| spectral features were found to be extremely road reSeniling the tvpe Ic chvpernovae SN. 1997ef£ and N 1998' (Filippenko 1997: Nomoto ct al.," However, the spectral features were found to be extremely broad resembling the type Ic `hypernovae' SN 1997ef and SN 1998bw (Filippenko 1997; Nomoto et al."
 2002: Mazzali e al., 2002; Mazzali et al.
 20t, 2002).
 At a istanee of 7.3. Alpe. SN. 2002ap being the nearest IVpesrnova discovered. to date. was a good. target for a dealed monitoring. and has been subject to multi-wavelength observations.," At a distance of 7.3 Mpc, SN 2002ap being the nearest hypernova discovered to date, was a good target for a detailed monitoring, and has been subject to multi-wavelength observations."
 In addition to optical observations. it has been observed in the X-ray (Sutaria et 22002). racio (Berger. Kulkarni Chevalier 2002. Sutaria et 22002) and in the infrared (Alattila Aleikle 2002).," In addition to optical observations, it has been observed in the X-ray (Sutaria et 2002), radio (Berger, Kulkarni Chevalier 2002, Sutaria et 2002) and in the infrared (Mattila Meikle 2002)."
 Evidence for a high velocity asvnimetric explosion. has been indicated by spectropolarimetric observations (Ixawabata et 22002: Leonard et 22002: Wang et 22002)., Evidence for a high velocity asymmetric explosion has been indicated by spectropolarimetric observations (Kawabata et 2002; Leonard et 2002; Wang et 2002).
" The bro; spectral features in the spectrum and. the evolution of Si HE spectral line indicate a high expansion velocity (~ 3J0000 kn s 1j, which supports hypernova model for the SN 220Pap (Ixinugasa et al."," The broad spectral features in the spectrum and the evolution of Si II spectral line indicate a high expansion velocity $\sim$ 30000 km $^{-1}$ ), which supports hypernova model for the SN 2002ap (Kinugasa et al."
 2002b: Moiklo et al., 2002b; Meikle et al.
 2002: CGial-Yan et al., 2002; Gal-Yam et al.
 2002a.b: Filippenko Chornocx 2002).," 2002a,b; Filippenko Chornock 2002)."
 ‘The spectropolarimeric observations during the early. phase indicate a sinulewity with SN 19958bw (Leonard et , The spectropolarimetric observations during the early phase indicate a similarity with SN 1998bw (Leonard et 2002).
Aodeling of 1 Opical observations indicate SN * as an energetic event. with an explosion energy. of )e107 eres (Mazzali ct al.," Modeling of the optical observations indicate SN 2002ap as an energetic event, with an explosion energy of $\simeq 4-10\times 10^{51}$ ergs (Mazzali et al."
 2002)., 2002).
 However. the racio servalions SCCDL o indicate that SN 2002ap an ordinary pe Le supernova. without a jet (Berger ct al.," However, the radio observations seem to indicate that SN 2002ap an ordinary type Ic supernova, without a jet (Berger et al."
 2002)., 2002).
 A μον of the UBVRIHIN images of galaxy M4 obtained veral vears prior o the discovery of SN 2002ap (Smartt et ., A study of the $UBVRIH_{\alpha}K$ images of galaxy M74 obtained several years prior to the discovery of SN 2002ap (Smartt et al.
 2002) restted in a non-detection of the progenitor., 2002) resulted in a non-detection of the progenitor.
 The spec‘tral similarity to SN. 1998bw. the possible Linkbetween very| energetic supernova and eamma ray bursts," The spectral similarity to SN 1998bw, the possible linkbetween very energetic supernova and gamma ray bursts"
intervals. is illustrated in Figure 1..,"intervals, is illustrated in Figure \ref{f:spec}."
 The distribution of the dispersed photon events is consistent with a continuunr source: no narrow emission line features were discernible., The distribution of the dispersed photon events is consistent with a continuum source: no narrow emission line features were discernible.
 Our analvsis involved comparison of the observed LETGS spectrum of wwith the emergent spectra computed for a range of line blanketed non-LTE model atmospheres and for opticallv-thin collision-dominated coronal models., Our analysis involved comparison of the observed LETGS spectrum of with the emergent spectra computed for a range of line blanketed non-LTE model atmospheres and for optically-thin collision-dominated coronal models.
 Coronal radiative loss models were computed using PEINTofALE. emploving line and continuun emissivities computed using atomic data from the CHIANTI database version 4.02 (Youngetal.2003) and the ionization equilibrium of Mazzottaοἱal.(1993).," Coronal radiative loss models were computed using PINTofALE, employing line and continuum emissivities computed using atomic data from the CHIANTI database version 4.02 \citep{Young.etal:03} and the ionization equilibrium of \citet{Mazzotta.etal:98}."
 Model atmosphere computations were performed using the PRO? code (Werner2003) assuming hvedrostatic ancl radiative equilibrium., Model atmosphere computations were performed using the PRO2 code \citep{Werner.etal:03} assuming hydrostatic and radiative equilibrium.
 Model atoms and atomic data used for detailed non-LTE calculations have been described recently by Werneretal.(2004)., Model atoms and atomic data used for detailed non-LTE calculations have been described recently by \citet{Werner.etal:04}.
. Constraints for the input model parameters are provided by earlier analyses of UV. spectra obtained by Che HST Faint Object Spectrograph (FOS) and. Goddard. High Resolution Spectrograph (GIRS). and optical spectra obtained αἱ the 3.5 m telescope αἱ Calar Alto (Wernerοἱal.1994.1996).," Constraints for the input model parameters are provided by earlier analyses of UV spectra obtained by the HST Faint Object Spectrograph (FOS) and Goddard High Resolution Spectrograph (GHRS), and optical spectra obtained at the 3.5 m telescope at Calar Alto \citep{Werner.etal:94,Werner.etal:96}."
. Additionally. we have examined more recent Far Ultraviolet Spectroscopic Explorer (FUSE) spectra of (to provide further constraints on the abuuclauces of the elements O ancl Fe that exhibit lines in (he FUSE bandpass.," Additionally, we have examined more recent Far Ultraviolet Spectroscopic Explorer (FUSE) spectra of to provide further constraints on the abundances of the elements O and Fe that exhibit lines in the FUSE bandpass."
 From the lack of Fe lines an upper abundance limit was determined by Miksaοἱal.(2002).., From the lack of Fe lines an upper abundance limit was determined by \citet{Miksa.etal:02}.
 From O VI lines we have determined the O abundance with the models computed for the present paper., From O VI lines we have determined the O abundance with the models computed for the present paper.
 Our adopted nominal parameters for amare listed in Table 1., Our adopted nominal parameters for are listed in Table 1.
 For a single test caleulation including sulfur we assume a solar S abundance (by mass fraction)., For a single test calculation including sulfur we assume a solar S abundance (by mass fraction).
 Phe absolute flux calibration for our models is provided by IST UV observations to high. precision. aud for comparisons with spectra we normalise the models to the observed flux density of 40x10P erg ?s Fat 1200A.," The absolute flux calibration for our models is provided by HST UV observations to high precision, and for comparisons with spectra we normalise the models to the observed flux density of $4.0\times 10^{-12}$ erg $^{-2}$ $^{-1}$ $^{-1}$ at 1200."
 The reasoning of Flemingetal.(1993). leading to the conclusion that hhas a corona was based on the comparison of low resolution X-ray pulse height spectra with model predictions for atmospheres containing C. N and O. We argue in (his paper that diffieulties in matching the X-ray data with photospherie models was primarily due to the neglect of heavier elements that can contaminate the photosphere.," The reasoning of \citet{Fleming.etal:93} leading to the conclusion that has a corona was based on the comparison of low resolution X-ray pulse height spectra with model predictions for atmospheres containing C, N and O. We argue in this paper that difficulties in matching the X-ray data with photospheric models was primarily due to the neglect of heavier elements that can contaminate the photosphere."
 While settling under the strong eravily of the white cwarl tends to empty the atiosphere of heavier elements. the almospheres of hotter stars can be substantially enriched by radiative levitation Schuhetal. 2002).," While settling under the strong gravity of the white dwarf tends to empty the atmosphere of heavier elements, the atmospheres of hotter stars can be substantially enriched by radiative levitation \citep[e.g.][]{Chayer.etal:95,Schuh.etal:02}."
. Model fIuxes for atmospheres with an effective temperature 7;;;=120000 Ix. with and without Fe. are compared in Figure 2..," Model fluxes for atmospheres with an effective temperature $T_{eff}=120000$ K, with and without Fe, are compared in Figure \ref{f:fluxes}."
 Models include full non-LTE treatment of C. O. Ne. 5 and Fe. whose abundances were adopted from Table 1.," Models include full non-LTE treatment of C, O, Ne, S and Fe, whose abundances were adopted from Table 1."
 Also illustrated is the blackhbocly flux density distribution for the same temperature., Also illustrated is the blackbody flux density distribution for the same temperature.
 It is interesting (o note that the blackbodxy, It is interesting to note that the blackbody
ihe bipolar boundary. conditions.,the bipolar boundary conditions.
 Alternative bipolar boundary conditions would produce similar results., Alternative bipolar boundary conditions would produce similar results.
 The four plots in Figure 3. are labeled a. b. ὁ and d to match the four curves in Figure 2..," The four plots in Figure \ref{forcefree} are labeled a, b, c and d to match the four curves in Figure \ref{Pi1s}."
 Curve c has a shallower decrease and curve d a steeper decrease in axial [lux as one moves away [rom the magnetic axis than curve a. while curve b has the same decrease.," Curve c has a shallower decrease and curve d a steeper decrease in axial flux as one moves away from the magnetic axis than curve a, while curve b has the same decrease."
 Models a and b are very similar indeed., Models a and b are very similar indeed.
 There are differences but thev are too small to see in the plots., There are differences but they are too small to see in the plots.
 Model c is larger (han a and b while model d is smaller., Model c is larger than a and b while model d is smaller.
 These elfects are due to the different quantities of axial magnetic [Iux confined close to the magnetic axes of the moclels., These effects are due to the different quantities of axial magnetic flux confined close to the magnetic axes of the models.
 Figure 4 shows the angle between the magnetic vector and the magnetic axis in each model., Figure \ref{angles} shows the angle between the magnetic vector and the magnetic axis in each model.
 llere. (he difference between models a and b is much more obvious.," Here, the difference between models a and b is much more obvious."
 The magnetic augle of, The magnetic angle of
 , $\hbox{--}$ 
"the existence of the dusty inner ring of NGC 5194 (Pierce1986) as shown in Fig. 12,,","the existence of the dusty inner ring of NGC 5194 \citep{pie86} as shown in Fig. \ref{popmap},"
 which seems to be connected to the dusty populations in the spiral arms., which seems to be connected to the dusty populations in the spiral arms.
" The P5 fraction is very small (< 0.1%) in the NGC 5194 center, but it increases with the radius, particularly rapidly at R =150""(from 2% at R ~150"" to >20% at R > 200”), although the number fraction of pixels with V<21 mag arcsec? is very small at such a large radius, as shown in Fig. 11(("," The P5 fraction is very small $<0.1\%$ ) in the NGC 5194 center, but it increases with the radius, particularly rapidly at R $\gtrsim150''$(from $2\%$ at R $\sim150''$ to $>20\%$ at R $>200''$ ), although the number fraction of pixels with $V<21$ mag $^{-2}$ is very small at such a large radius, as shown in Fig. \ref{radfrac}( ("
b).,b).
" This result gives a hint about the origin of the P5 color, implying that the P5 may be related to background light."," This result gives a hint about the origin of the P5 color, implying that the P5 may be related to background light."
" For example, some background galaxies at high redshifts have colors like those of the P5 pixels (i.e,blueV—Ial. 2005)."," For example, some background galaxies at high redshifts have colors like those of the P5 pixels \citep[i.e., blue $V-I$ color but red $B-V$ color; e.g.][]{hil05}."
". The color of the P5 pixels, which is not easily explained using population synthesis models at z — 0, may be the result of the mixture of the faint extended light from NGC 5194 and the background light from high-redshift galaxies."," The color of the P5 pixels, which is not easily explained using population synthesis models at z = 0, may be the result of the mixture of the faint extended light from NGC 5194 and the background light from high-redshift galaxies."
 Fig., Fig.
" 13 displays the surface brightness distribution of each pixel population, showing that most of the P5 pixels are significantly biased to faint surface brightness (V>22 mag arcsec~?)."," \ref{pixelcount} displays the surface brightness distribution of each pixel population, showing that most of the P5 pixels are significantly biased to faint surface brightness $V\gtrsim22$ mag $^{-2}$ )."
 The results in Fig. 11((, The results in Fig. \ref{radfrac}( (
b) and Fig.,b) and Fig.
" 13 confirm that the P5 pixels are biased to the outskirt and faint pixels, which supports the possibility that the P5 pixels are contaminated by background light."," \ref{pixelcount} confirm that the P5 pixels are biased to the outskirt and faint pixels, which supports the possibility that the P5 pixels are contaminated by background light."
 The other (P1 — P4) pixels may also be partially contaminated by background light., The other (P1 – P4) pixels may also be partially contaminated by background light.
" It is very difficult to quantify by how much, but we discuss it qualitatively here."," It is very difficult to quantify by how much, but we discuss it qualitatively here."
" The color vector of such contaminations may be to the direction of bluer V—I and redder B—V colors (from lower-right to upper-left in Fig. 10)),"," The color vector of such contaminations may be to the direction of bluer $V-I$ and redder $B-V$ colors (from lower-right to upper-left in Fig. \ref{pccd}) ),"
" which is almost independent of the division between P1, P3 and P4 pixels."," which is almost independent of the division between P1, P3 and P4 pixels."
" On the other hand, the P2 pixels may appear to be P3 or P4 pixels by the contamination."," On the other hand, the P2 pixels may appear to be P3 or P4 pixels by the contamination."
" However, such an effect (1) will be very weak for bright pixels (e.g., V<21 mag arcsec2; the pixel surface brightness limit in refpopvartext)) and (2) will make the population variations be seen weaker rather than stronger."," However, such an effect (1) will be very weak for bright pixels (e.g., $V<21$ mag $^{-2}$; the pixel surface brightness limit in \\ref{popvartext}) ) and (2) will make the population variations be seen weaker rather than stronger."
" In short, the effect of the background light on our results may be very limited and our results may show the lower limits of the population variations."," In short, the effect of the background light on our results may be very limited and our results may show the lower limits of the population variations."
" Here, we focus on the population variations in two viewpoints: 1) the variation across the spiral arms and 2) the difference between the NGC 5195 direction and the opposite direction."," Here, we focus on the population variations in two viewpoints: 1) the variation across the spiral arms and 2) the difference between the NGC 5195 direction and the opposite direction."
" To investigate these variations, the spiral arms of NGC 5194 need to be defined reasonably."," To investigate these variations, the spiral arms of NGC 5194 need to be defined reasonably."
 We defined the spiral arms by fitting the spatial distribution of bright (V«20 mag arcsec3) pixels., We defined the spiral arms by fitting the spatial distribution of bright $V<20$ mag $^{-2}$ ) pixels.
" A spiral pattern is linearly distributed in the angle versus log radius plot, but the spiral arms of NGC 5194 are significantly distorted at the outskirt area due to the interaction with NGC 5195."," A spiral pattern is linearly distributed in the angle versus log radius plot, but the spiral arms of NGC 5194 are significantly distorted at the outskirt area due to the interaction with NGC 5195."
" Thus, we fit the spiral arms only in the radius range of 40 — 120 arcseconds (1.6 — 4.8 kpc) as shown in Fig. 14,,"," Thus, we fit the spiral arms only in the radius range of 40 – 120 arcseconds (1.6 – 4.8 kpc) as shown in Fig. \ref{spiral}, ,"
 which returns the spiral arm coordinate equations of: and where O is the clockwise position angle from the NGC 5195 direction in the unit of degree, which returns the spiral arm coordinate equations of: and where $\Theta$ is the clockwise position angle from the NGC 5195 direction in the unit of degree
"redshift distortious is merely to nmultiplv the amplitude of the linear power spectrum by a factor (1|f)?: that is. we can absorb the factor (1|f)? iuto a, ly aud D. or more simply iuto Pe(Ak). as was obvious from Eq.(83)).","redshift distortions is merely to multiply the amplitude of the linear power spectrum by a factor $(1+f)^2$; that is, we can absorb the factor $(1+f)^2$ into $\sigma_v^2$, $I_0$ and $I_2$ , or more simply into $P_L(k)$, as was obvious from \ref{sd-perp}) )."
 Therefore. we directly obtain from Eq.(18)) the expression Pi) E yr ug. since we only need to multiply each term that involves Pith) bv a factor (1|f).," Therefore, we directly obtain from \ref{Pkjn}) ) the expression (k) = (kq), since we only need to multiply each term that involves $P_L(k)$ by a factor $(1+f)^2$."
" In a similar fashion. the standard perturbative expausion (21)) becomes thy DUUM) pe) — PLO), whereas the vrenormalized” perturbative expansion (32)) becomes DU. PU eyτα Ue)."," In a similar fashion, the standard perturbative expansion \ref{Pstd}) ) becomes (k) = (k) (k) = (k), whereas the “renormalized” perturbative expansion \ref{Psigv}) ) becomes (k) = (k), (k) = (k)."
" Thus. cach order η of the perturbative expansions gets nultiplied by a factor (1|f£)?"" as we go from real space ο redshift space. for modes that are aligned with the line of sight."," Thus, each order $n$ of the perturbative expansions gets multiplied by a factor $(1+f)^{2n}$ as we go from real space to redshift space, for modes that are aligned with the line of sight."
 Iu particular. at lowest order we have Pith)=CLLPPPEA) |... so that we recover the well-known boost actor of 7.. associated with the iufallof galaxies within eravitational potential wells.," In particular, at lowest order we have $P^s_{\parallel}(k) = (1+f)^2 P_L(k)+...$ , so that we recover the well-known boost factor of \citet{Kaiser1987}, associated with the infallof galaxies within gravitational potential wells."
 On the other haud. the factor (11f£)? also sharpens the Caussiuu damping prefactor in he expansion (97)).," On the other hand, the factor $(1+f)^2$ also sharpens the Gaussian damping prefactor in the expansion \ref{sPsigv}) )."
 We now turn to the nounperturbative correction associated with the vsticky model” introduced in Sect. 2.1., We now turn to the nonperturbative correction associated with the “sticky model” introduced in Sect. \ref{nonperturbative}.
 Again we focus on the redshift-space power spectrum for wavemuubers that are parallel to the line of sight., Again we focus on the redshift-space power spectrum $P^s_{\parallel}(k)$ for wavenumbers that are parallel to the line of sight.
 Pith)From DEaq.(81)) we now need to compute the average {έik-Asre, From \ref{Pkq-s}) ) we now need to compute the average $\lag e^{\ii\vk\cdot\Delta\vs}\rag_{\vq}$.
" Before shell crossing. or more precisely. as long as Aurpy>0 for à Lagrangian-space separation vector along the ei-axis. the “sticky model” follows the Zeldovich dvuauiucs. k=k eq ve A(99) > q: =k | (11f£) vPsi,. whereas once Αρα<0 we affect the same position. Avy=0. aud the same peculiar velocity. ο=0. alone e£-axis. for both particles as in (16)). whence < gt Asy =0. = (LIE) | kV, | SV, ("," Before shell crossing, or more precisely, as long as $\Delta x_{L1}>0$ for a Lagrangian-space separation vector along the $\ve_1$ -axis, the “sticky model” follows the Zeldovich dynamics, k _z, q _1 : > -q: = + (1+f) _L, whereas once $ \Delta x_{L1}<0$ we affect the same position, $\Delta x_1=0$, and the same peculiar velocity, $\Delta v_1=0$, along the $\ve_1$ -axis, for both particles as in \ref{Deltax-def}) ), whence < -q : s_1 =0 , = (1+f) [ k_2 + k_3 ]. ("
The line-ofsight axis e. aud the axis ei of the Withherangian separation vector q are not related.),The line-of-sight axis $\ve_z$ and the axis $\ve_1$ of the Lagrangian separation vector $\vq$ are not related.)
" At the shell-crossing time we have a discontinuitv in the redshitt-space separation As. since As, jumps from fq to zero."," At the shell-crossing time we have a discontinuity in the redshift-space separation $\Delta\vs$, since $\Delta s_1$ jumps from $-f q$ to zero."
 Tudeed. just before contact. the two particles move closer. with the finite relative velocity Aey=(aD/D)q. while after collision their relative velocity is set to zero.," Indeed, just before contact, the two particles move closer, with the finite relative velocity $\Delta v_1=-(a\dot{D}/D)q$, while after collision their relative velocity is set to zero."
" This should actually be understood in a loose sense. since within the ""stickv model” we consider neither trausverse directions nor thecloud-du-cloud problem."," This should actually be understood in a loose sense, since within the “sticky model” we consider neither transverse directions nor thecloud-in-cloud problem."
 Then. as in Sect.," Then, as in Sect."
 2.1 the redshift-space power spectra of the “sticky model reads as LAO) 1 SI). and from 108) j-(113)) we obtain for waveuuuboers along the line of sight (UY f= Ps," \ref{nonperturbative} the redshift-space power spectrum of the “sticky model” reads as) = ) + ) , and from \ref{ks-sticky1})\ref{ks-sticky2}) ) we obtain for wavenumbers along the line of sight (k) = _0^1"
The Spitzer mid- and far-infrared photometric observations were performed during Cycle-2 as part of the General Observer program with identification number 30607 (PI Georgakakis).,The Spitzer mid- and far-infrared photometric observations were performed during Cycle-2 as part of the General Observer program with identification number 30607 (PI Georgakakis).
 The mid- and far-IR flux densities of the target sources are presented in Table 2.., The mid- and far-IR flux densities of the target sources are presented in Table \ref{tab_obs2}.
 The 3.6BAUM dati were taken with the TRAC (Infrared Array Camera) between July 9 and December 28 2006., The $\rm 3.6-8.0\mu m$ data were taken with the IRAC (Infrared Array Camera) between July 9 and December 28 2006.
 Each observation consists of 5 separate integrations of 2ss euch organised in a random Gaussian dither pattern., Each observation consists of 5 separate integrations of s each organised in a random Gaussian dither pattern.
 The median separation between successive offsets is aaremin., The median separation between successive offsets is arcmin.
 The observations were run through the SI+4.0 version of the Spitzer Science Center pipeline deseribed in IRAC data handbook., The observations were run through the S14.4.0 version of the Spitzer Science Center pipeline described in IRAC data handbook.
 The basic calibration data (BCD) products are processed for dark current. bias offset. linearisation. flat-tielding. cosmic ray detection and the flux calibration.," The basic calibration data (BCD) products are processed for dark current, bias offset, linearisation, flat-fielding, cosmic ray detection and the flux calibration."
 The individual frames of a given source were then registered and co-added into the final mosaiced image., The individual frames of a given source were then registered and co-added into the final mosaiced image.
 The uncertainty in the calibration factor applied to the final mosaics is about 10 per cent., The uncertainty in the calibration factor applied to the final mosaics is about 10 per cent.
 Following the TRAC data handbook flux densities are estimated by using apertures of 6Oaaresee radius and then applying aperture correction factors to account for the extended Spitzer point spread function (PSF) in different wavebands., Following the IRAC data handbook flux densities are estimated by using apertures of arcsec radius and then applying aperture correction factors to account for the extended Spitzer point spread function (PSF) in different wavebands.
 These are estimated |.06 for TRAC 3.6. 4.5sam bands and 1.10 for the 5.8 and S.07n bands.," These are estimated 1.06 for IRAC 3.6, $\rm 4.5\,\mu m$ bands and 1.10 for the 5.8 and $8.0\,\mu m$ bands."
" Colour corrections are expected to be small as the mid-IR spectra of the sources in the sample approximate a power-law with οντεconstant,", Colour corrections are expected to be small as the mid-IR spectra of the sources in the sample approximate a power-law with $\nu f\nu \rm \approx constant$.
 We therefore choose not to apply colour corrections., We therefore choose not to apply colour corrections.
 The 24. 70. and 160; data were taken with MIPS (Multiband Imaging Photometer for Spitzer) between July 14 2006 and September 15 2007.," The 24, 70, and $\rm 160 \mu m$ data were taken with MIPS (Multiband Imaging Photometer for Spitzer) between July 14 2006 and September 15 2007."
 The observations used the MIPS photometry mode with an exposure time of 35s for individual frames., The observations used the MIPS photometry mode with an exposure time of s for individual frames.
 The 24 and TOs: data were obtained in one cycle. while the 160//m observations were carried out in 3 cycles.," The 24 and $\rm 70\mu
m$ data were obtained in one cycle, while the $\rm 160 \mu m$ observations were carried out in 3 cycles."
 The total source exposure times were 48. 38 and ss at 24. 70. and 16054 respectively.," The total on-source exposure times were 48, 38 and s at 24, 70, and $\rm 160 \mu m$ respectively."
 The MIPS images were created from raw data frames using he MIPS Data Analysis Tools (MIPSDAT:?) version 3.10 along with additional processing steps., The MIPS images were created from raw data frames using the MIPS Data Analysis Tools \citep[MIPS DAT;][]{Gordon2005} version 3.10 along with additional processing steps.
 The processing steps for the 70 and 1607/0 data are similar. but the steps for the j/m data are significantly different from these other two bands.," The processing steps for the 70 and $\rm 160 \mu m$ data are similar, but the steps for the $\rm \mu m$ data are significantly different from these other two bands."
" The srr data processing is described first followed by descriptions of the 70 and 160,/m data processing.", The $\rm \mu m$ data processing is described first followed by descriptions of the 70 and $\rm 160 \mu m$ data processing.
 The individual jm frames were first processed through a droop correction (removing an exeess signal in each pixel that is woportional to the signal in the entire array) and were corrected for non-linearity in the ramps., The individual $\mu$ m frames were first processed through a droop correction (removing an excess signal in each pixel that is proportional to the signal in the entire array) and were corrected for non-linearity in the ramps.
 The dark current was then subtracted., The dark current was then subtracted.
 ext. scan-mirror-position dependent flats were applied to the data. atent images were removed. and scan-mirror-position independent flats were applied to the data.," Next, scan-mirror-position dependent flats were applied to the data, latent images were removed, and scan-mirror-position independent flats were applied to the data."
 Following this. the zodiacal light was measured by fitting planes to the background regions in each rame. and these planes were then subtracted from the data.," Following this, the zodiacal light was measured by fitting planes to the background regions in each frame, and these planes were then subtracted from the data."
 Next. a robust statistical analysis was applied to cospatial pixels from different frames in which statistical outliers (e.g. pixels affected by cosmic rays) were masked out.," Next, a robust statistical analysis was applied to cospatial pixels from different frames in which statistical outliers (e.g. pixels affected by cosmic rays) were masked out."
" After this. tinal mosaics were made with pixel sizes of 1.5"". residual backgrounds in the images were subtracted. and the data were calibrated into astronomical units."," After this, final mosaics were made with pixel sizes of $1.5^{\prime\prime}$, residual backgrounds in the images were subtracted, and the data were calibrated into astronomical units."
 The calibration factor for the jan data is given by ? as (4.5440.18).lo7 My + (MIPS instrumental A [O., The calibration factor for the $\rm \mu m$ data is given by \citet{Engelbracht2007} as $(4.54\pm0.18) \times 10^{-2}$ MJy $^{-1}$ [MIPS instrumental $^{-1}$.
Saaresec radius aperture is used for the photometry., A arcsec radius aperture is used for the photometry.
 The correction factor that accounts for PSF losses is estimated to be 1.22., The correction factor that accounts for PSF losses is estimated to be 1.22.
 The mid-IR spectra of the sample sources is close to a, The mid-IR spectra of the sample sources is close to a
where fy is a constaut of iuteeration.,where $t_0$ is a constant of integration.
 Considering a(t) iu Eq.C11)). the following coustraiut equations are fonud It is explicitly scen here that if fj=0. then the first integral £3 is equa to zero.," Considering $a(t)$ in \ref{frstI-33}) ), the following constraint equations are found It is explicitly seen here that if $t_{0}=0$, then the first integral $I_3$ is equal to zero."
 Also. using Eqs.(7)). (8)) and (9)). we haveadditional coustraint relatious as Thus. the constraint Eqs. (13))," Also, using \ref{E-L}) ), \ref{de-1}) ) and \ref{de-2}) ), we haveadditional constraint relations as Thus, the constraint Eqs. \ref{ceq-1}) )"
 and (15)) give the following simple relation The deceleration parameter. which is au important quantity. iu. cosmologv.is. defineda by 4—/—a0fàmm». where the positive sign of q indicates the standard decelerating models whereas the negative sign corresponds to models aud q=ϐ corresponds to expansion with constant velocity.," and \ref{ceq-3}) ) give the following simple relation The deceleration parameter, which is an important quantity in cosmology, is defined by $q =-a \ddot{a}/\dot{a}^2$, where the positive sign of $q$ indicates the standard decelerating models whereas the negative sign corresponds to accelerating models and $q=0$ corresponds to expansion with constant velocity."
 It takes the following form iu this inodel The effective equatiou of state parameter defiued by (opp=1ME-3L (Capozziclloetal.2003:Alleinaudietal.200142:Allexuanidi20015). cau be obtained as where 1 is Hubble pariueter.," It takes the following form in this model The effective equation of state parameter defined by $w_{eff}=-1-\frac{2\dot{H}}{3H^2}=\frac{2q-1}{3}$ \cite{capo03,all04a,all04b} can be obtained as where $H$ is Hubble parameter."
 Astrophivsical data indicate that a lies in a very narrow strip close to a= 1., Astrophysical data indicate that $w$ lies in a very narrow strip close to $w = -1$ .
 The case w=l corresponds to the cosinological constant., The case $w = -1$ corresponds to the cosmological constant.
 For uw< lthe phantom phase is observed. and for lo<ow<—1/3 the phase is described bx quintessence.," For $w<-1$ the phantom phase is observed, and for $-1 < w < -1/3$ the phase is described by quintessence."
 Thus. iu the iterval 160 we have quintessence phase.," Thus, in the interval $-1 < \ell < 0$ we have quintessence phase."
 ox<fd. then the phautom phase occurs. where the universe is both expaucdiue aud accelerating.," If $-\infty < \ell < -1$, then the phantom phase occurs, where the universe is both expanding and accelerating."
 If Ry=Ll. ie o=1 from the relation o=frRy. then the action is reduced to the ΡΕΠΕ action with cosmological coustaut ( FUR)=RA).," If $R_{0} = 1$, i.e. $\phi =1$ from the relation $\phi =
f_{\mathcal{R}} = R_0$, then the action is reduced to the Einstein-Hilbert action with cosmological constant ( $f(\mathcal{R})=\mathcal{R}-\lambda$ )."
 We note here that Palati aud metric foriualisni are comede., We note here that Palatini and metric formalism are coincide.
 Iu this case. the Nocther first integrals (33)) and (31)) cal be written as The modified Friedmann equation (7)) for this case reduces to the form From Eq. (503) ," In this case, the Noether first integrals \ref{frstI-2}) ) and \ref{frstI-3}) ) can be written as The modified Friedmann equation \ref{E-L}) ) for this case reduces to the form From Eq. \ref{frstII-0}) )"
the scale factor is solved as Tuserting the scale factor (53)) iuto Eq. (51)), the scale factor is solved as Inserting the scale factor \ref{sol-2}) )into Eq. \ref{frstI-3-ii}) )
 aud Eq. (52)), and Eq. \ref{mfeq}) )
" oue ects A=0. 6=1/2 aud as a coustraint equations. &p,,g=I5ο18 and ty=[ο1ο."," one gets $\lambda =0$, $\ell=1/2$ and as a constraint equations $\kappa \rho_{m0}=
I_{2}^2/48$ and $t_{1}=I_{3}/I_{2}$."
 Theretore. for accelerating’ οase the scale factor has the form which is obtained from (53)) when (=1δν," Therefore, for this case the scale factor has the form which is obtained from \ref{sol-2}) ) when $\ell=1/2$."
 Taking (=Ty lin Eq.(50)). we have a(t}=exptlafm) which eives pog=0 and A=E3.," Taking $\ell=-1/4$ in \ref{frstII-0}) ), we have $a(t)=\exp(\frac{I_{2}
t}{3})$ which gives $\rho_{m0}=0$ and $\lambda = 2 I_2^2 /3$."
 Usiug these results for (= lfliuthe Eq. (51)), Using these results for $\ell = -1/4$ in the Eq. \ref{frstI-3-ii}) )
 one can find Z;= 0., one can find $I_{3}=0$ .
" The Palatiui approach consider the metric goa, aud the connection TY as independeut field variables. but the spacetime ietric gy, is onlv indepeucdent variable in metric formalisin."," The Palatini approach consider the metric $g_{ab}$ and the connection $\Gamma^{a}_{bc}$ as independent field variables, but the spacetime metric $g_{ab}$ is only independent variable in metric formalism."
 The Palatini formatisim can be seen as contaimiug two indepeudent metrics Gab aid hunfriap rather that a metric aud independent connection.," The Palatini formalism can be seen as containing two independent metrics $g_{ab}$ and $h_{ab} = f_{\mathcal{R}}
g_{ab}$ rather that a metric and independent connection."
" In Palatini FUR) eravity the second metric μι determine the geodesic structure with the connection L7. which is the Levi-Civita connection of new metric 5,54.", In Palatini $f(\mathcal{R})$ gravity the second metric $h_{ab}$ determine the geodesic structure with the connection $\Gamma^{a}_{bc}$ which is the Levi-Civita connection of new metric $h_{ab}$ .
" Ia BD theory of gravity the second metric fy, is related tothe non-nüuimual coupling of the DD scalar. ic. 0= fg."," In BD theory of gravity the second metric $h_{ab}$ is related tothe non-minimal coupling of the BD scalar, i.e. $\phi = f_{\mathcal{R}}$ ."
 In Palatini approach. F(R) eravity eiven by the action (1)) is equivalent to a special," In Palatini approach, $f(\mathcal{R})$ gravity given by the action \ref{action}) ) is equivalent to a special"
The errors of the coefficients. are given in. brackets.,The errors of the coefficients are given in brackets.
 ‘These relations are obtained on the basis of the observations [rom 2008 July 6., These relations are obtained on the basis of the observations from 2008 July 6.
 They are valid. over the range 10.20.«R<10.35 mag., They are valid over the range $10.20 \le R \le 10.35$ mag.
 The Spearmans (rho) rank correlation gives p—0.85 (significanceS10. ) for αι]. p=0.66 (1.10) for E«q.2. p20.57 (2.10 1) for Ex3.," The Spearman's (rho) rank correlation gives $\rho=0.85$ (significance $ 8.10^{-6}$ ) for Eq.1, $\rho=0.66$ $1.10^{-7}$ ) for Eq.2, $\rho=0.57$ $2.10^{-10}$ ) for Eq.3."
 The significance in Exq.1-]5e.3 is always <<0.001 indicating that all these correlations are highly significant and changes in the colours of RS Oph are correlated with brightness variations., The significance in Eq.1-Eq.3 is always $<<0.001$ indicating that all these correlations are highly significant and changes in the colours of RS Oph are correlated with brightness variations.
 Auch (1992) proposed that the light curve of CVs can be separated into two parts constant. light ane variable (Llickering) source., Bruch (1992) proposed that the light curve of CVs can be separated into two parts – constant light and variable (flickering) source.
 We assume —that all the variability is clue to flickering., We assume that all the variability is due to flickering.
 In these suppositions the Hickering light source is considered. modulated., In these suppositions the flickering light source is considered modulated.
 Following these suggestions. we calculate the Dux of the Hickering light source as £y;Pas PF. Where νο is the average Dux during the run and P; is the minimum flux during the run (corrected for the tvpical error of the observations).," Following these suggestions, we calculate the flux of the flickering light source as $F_{fl}=F_{av}-F_{min}$ , where $F_{av}$ is the average flux during the run and $F_{min}$ is the minimum flux during the run (corrected for the typical error of the observations)."
 £j has been calculated for cach band. using the values given in Table 1: and Bessel (1979) calibration for the Uuxes of a zero magnitude star.," $F_{fl}$ has been calculated for each band, using the values given in Table 1 and Bessel (1979) calibration for the fluxes of a zero magnitude star."
 Aclopting these results. we find that the Iickering light source contributes about of the light in the 7? band. in V. in D. and in €.," Adopting these results, we find that the flickering light source contributes about of the light in the $R$ band, in $V$, in $B$, and in $U$."
 Following Snijders (1987) we assume interstellar extinction fpy=0.73 and an extinction law as given in Cardelli et al.(1989)., Following Snijders (1987) we assume interstellar extinction $E_{B-V} =0.73$ and an extinction law as given in Cardelli et al.(1989).
 Using the colour relations (I5q.1.150.3. derived in Sect.3.1) in the interval #=10.28.10.34 mae (in other words average 4?=10.28 and we calculate average values for other bands. minimal brightness in /7=10.34. and we derive minimal fluxes in other bands). we obtain colours of the flickering light source in RS Oph as," Using the colour relations (Eq.1–Eq.3, derived in Sect.3.1) in the interval $R=10.28-10.34$ mag (in other words average $R=10.28$ and we calculate average values for other bands, minimal brightness in $R=10.34$, and we derive minimal fluxes in other bands), we obtain colours of the flickering light source in RS Oph as"
XX-3 was first observed in 1967 (2?) and later identified as a pulsating high-mass X-ray binary system based on the observations (22)..,"X-3 was first observed in 1967 \citep{1967PhRvL..19..681C} and later identified as a pulsating high-mass X-ray binary system based on the observations \citep{1971ApJ...167L..67G, 1972ApJ...172L..79S}."
 The pulsation period of the source was found to be ~4.8ss (2).. the orbital period (determined from regularX-ray eclipses) 1s ~2.08 ddays (?)..," The pulsation period of the source was found to be $\sim$ s \citep{1971ApJ...167L..67G}, the orbital period (determined from regularX-ray eclipses) is $\sim$ days \citep{1972ApJ...172L..79S}."
 Measuring the eclipse times. ? derived a mass of Mc; for the compact object and of Me for the companion star. well in accordance with ? who determined the masses from the optical spectroscopic observations of the companion.," Measuring the eclipse times, \citet{1974ApJ...192L.139A} derived a mass of $M_{\sun}$ for the compact object and of $M_{\sun}$ for the companion star, well in accordance with \citet{1979ApJ...229.1079H} who determined the masses from the optical spectroscopic observations of the companion."
 The projected orbital radius was determmed from the analysis of pulse arrival times to be ~39.75+0.04 Ilight-seconds (?).., The projected orbital radius was determined from the analysis of pulse arrival times to be $\sim$ $\pm$ light-seconds \citep{1972ApJ...172L..79S}.
 The system ts located at a distance of about κκρο (2).. with a lower limit of kkpe (?)..," The system is located at a distance of about kpc \citep{1974ApJ...192L.135K}, with a lower limit of kpc \citep{1974ApJ...192L.135K}."
 XX-3 shows a non-periodic alternation of high and low states with a characteristic time of reccurence of ddays. derived from analysis of data (2)..," X-3 shows a non-periodic alternation of high and low states with a characteristic time of reccurence of days, derived from analysis of data \citep{1983ApJ...273..709P}."
 This long-term variability is sometimes attributed to ἃ precessing accretion disk (?).., This long-term variability is sometimes attributed to a precessing accretion disk \citep{1983ApJ...273..709P}.
 Our analysis was triggered by the work of ?.., Our analysis was triggered by the work of \citet{2005A&A...442L..15P}.
 Based on their analysis of RXTE/ASM data of XX-3. the authors reported the presence of two distinct spectral states/modes in the high state of the source distinguished by the hardness ratio.," Based on their analysis of RXTE/ASM data of X-3, the authors reported the presence of two distinct spectral states/modes in the high state of the source distinguished by the hardness ratio."
 In the high state the peak flux is up to a factor of 40 larger than during the low state (?).., In the high state the peak flux is up to a factor of 40 larger than during the low state \citep{2005A&A...442L..15P}.
 It was found that when the source makes a transition from a low to a high state. 1t adopts one of these two spectral modes and during the entire high-intensity phase remains in that mode.," It was found that when the source makes a transition from a low to a high state, it adopts one of these two spectral modes and during the entire high-intensity phase remains in that mode."
 ? showed that during all high states between December 2000 and April 2004 the source was in the hard spectral mode as indicated by large hardness ratio. while in all high states prior ad subsequent to this period it was in the soft spectral mode.," \citet{2005A&A...442L..15P}
 showed that during all high states between December 2000 and April 2004 the source was in the hard spectral mode as indicated by large hardness ratio, while in all high states prior and subsequent to this period it was in the soft spectral mode."
 To exclude systematic effects the authors analyzed ASM data on three other sources: XX-1. XX-] and XX-1.," To exclude systematic effects the authors analyzed ASM data on three other sources: X-1, X-1 and X-1."
 But none of the sources showed the behavior similar to that reported for XX-3., But none of the sources showed the behavior similar to that reported for X-3.
 ? interpreted their results as a manifestatio of two accretion modes which are at work in XX-3 at different times., \citet{2005A&A...442L..15P} interpreted their results as a manifestation of two accretion modes which are at work in X-3 at different times.
 We repeated the analysis of ? using the same data set and following the analysis procedures described 1n their paper.," We repeated the analysis of \citet{2005A&A...442L..15P}
 using the same data set and following the analysis procedures described in their paper."
 We then extended our analysis to the entire ASM data on the source available by now and included the data obtained with the MAXI. INTEGRAL/JEM-X and INTEGRAL/ISGRI instruments.," We then extended our analysis to the entire ASM data on the source available by now and included the data obtained with the MAXI, INTEGRAL/JEM-X and INTEGRAL/ISGRI instruments."
 We used data from the All Sky Monitor (ASM) onboard the Rossi X-ray Timing Explorer (RXTE). from the Monitor of All Sky X-ray Image (MAXI) and from the JEM-Xand ISGRI instruments of the INTEGRAL (INTErnational Gamma-Ray Astrophysics Laboratory) The ASM instrument is an X-ray monitor covering the 1.5-I2kkeV energy range (?)..," We used data from the All Sky Monitor (ASM) onboard the Rossi X-ray Timing Explorer (RXTE), from the Monitor of All Sky X-ray Image (MAXI) and from the JEM-Xand ISGRI instruments of the INTEGRAL (INTErnational Gamma-Ray Astrophysics Laboratory) The ASM instrument is an X-ray monitor covering the keV energy range \citep{1996ApJ...469L..33L}."
 It. consists of three Scanning Shadow Cameras (SSCs) each with a position-sensitive proportional counter (?).., It consists of three Scanning Shadow Cameras (SSCs) each with a position-sensitive proportional counter \citep{1996ApJ...469L..33L}.
 The counts detected in three energy bands. 1.5-3. 3-5 and kkeV. are accumulated in dwells of about 90ss. In our analysis we used the data available from the ASM team on the web (iuttpste.mit.edulasmlc/).," The counts detected in three energy bands, 1.5–3, 3–5 and keV, are accumulated in dwells of about s. In our analysis we used the data available from the ASM team on the web )."
 We checked that these data differ only in format from the data available at the NASA's High Energy Astrophysics Science Archive Center (HEASARC) which were used by ?.., We checked that these data differ only in format from the data available at the NASA's High Energy Astrophysics Science Archive Center (HEASARC) which were used by \citet{2005A&A...442L..15P}.
 MAXI is an X-ray monitor onboard of the International Space Station (ISS) (?).., MAXI is an X-ray monitor onboard of the International Space Station (ISS) \citep{2009PASJ...61..999M}.
 It scans the sky each orbit of minutes observing a particular source for about sseconds (?).. depending on the source position.," It scans the sky each orbit of minutes observing a particular source for about seconds \citep{2011arXiv1102.0891S}, depending on the source position."
 MAXI has two cameras. the Gas Slit Camera (GSC) and the Solid-state Slit Camera (SSC).," MAXI has two cameras, the Gas Slit Camera (GSC) and the Solid-state Slit Camera (SSC)."
 The main instrument. GSC. consists of proportional counters with slit and slat collimators (2) with a total effective area of env (?)..," The main instrument, GSC, consists of proportional counters with slit and slat collimators \citep{2009PASJ...61..999M} with a total effective area of $\mathrm{cm}^2$ \citep{2009PASJ...61..999M}."
 Every orbit. about of the sky is scanned in the kkeV energy band (?)..," Every orbit, about of the sky is scanned in the keV energy band \citep{2011arXiv1102.0891S}."
 MAXI lightcurves with a time resolution of one orbit are available on the web riken., MAXI lightcurves with a time resolution of one orbit are available on the web ).
jpp. The lighteurves are produced for three energy bands. kkeV. kkeV. and The INTEGRAL satellite observes objects simultaneously in gamma rays. X-rays and visible light.," The lightcurves are produced for three energy bands, keV, keV, and The INTEGRAL satellite observes objects simultaneously in gamma rays, X-rays and visible light."
 JEM-X 1s one of its instruments and consists of two telescopes with coded aperture masks (?).. , JEM-X is one of its instruments and consists of two telescopes with coded aperture masks \citep{2003A&A...411L.231L}. .
JEM-X obtains X-ray spectra and imaging in the 3-35kkeV energy band (?).. , JEM-X obtains X-ray spectra and imaging in the keV energy band \citep{2003A&A...411L.231L}. .
IBIS is the gamma-ray imager of INTEGRAL and consists of two detector layers.," IBIS is the gamma-ray imager of INTEGRAL and consists of two detector layers,"
well as LSB galaxy rotation curves from the literature.,well as LSB galaxy rotation curves from the literature.
 Their method is able to describe the variety present in data and simulations and from the distributions of respective fitting parameters HO4 conclude that the observed curves consistent with the CDM paradigm: a conclusion that contracicts a laree body οἱ observational work., Their method is able to describe the variety present in data and simulations and from the distributions of respective fitting parameters H04 conclude that the observed curves consistent with the CDM paradigm; a conclusion that contradicts a large body of observational work.
 For their analysis. 104 make use of the rotation curves presented in MeGaughetal. (2001).. deBlok&Dosma(2002) ancl Swatersetal.(2003).. ancl derive cireular. velocity curves lor haloes from their simulations.," For their analysis, H04 make use of the rotation curves presented in \citet{mcgaugh2001}, , \citet{dBB02} and \citet{swaters2003}, and derive circular velocity curves for haloes from their simulations."
 To quantify both sets of observed and simulated rotation curves they use a (ree-parameter fitting formula (see citealteourteau)): llere Vy is the asymptotic velocity of the flat part of (he rotation curve. 7; is a scale radius (the transition radius between (he rising and flat part of the rotation curve) and the parameter 5 describes the abruptuess of the turn-over between the rising and flat parts of the rotation eurve.," To quantify both sets of observed and simulated rotation curves they use a three-parameter fitting formula (see \\citealt{courteau}) ): Here $V_0$ is the asymptotic velocity of the flat part of the rotation curve, $r_t$ is a scale radius (the transition radius between the rising and flat part of the rotation curve) and the parameter $\gamma$ describes the abruptness of the turn-over between the rising and flat parts of the rotation curve."
 A higher value of > results in an abrupt transition. a low value in a more eradual turn-over.," A higher value of $\gamma$ results in an abrupt transition, a low value in a more gradual turn-over."
 Note that this [uncetion was designed to fit observed. rotation curves: it specifically assumes a solic-body rise in the inner parts and a flat rotation curve in (he outer parts., Note that this function was designed to fit observed rotation curves: it specifically assumes a solid-body rise in the inner parts and a flat rotation curve in the outer parts.
 1104 fit the (vo sets of rotation curves with this Iunction. constraining the fit parameters slightly by insisting that with Vy the maximum observed velocity in (he rotation curve.," H04 fit the two sets of rotation curves with this function, constraining the fit parameters slightly by insisting that with $V_{\rm max}$ the maximum observed velocity in the rotation curve."
 Fits with 5> [rejected by condition (1)] correspond with very abrupt transitions between the rising and flat parts of the curve., Fits with $\gamma > 5$ [rejected by condition (1)] correspond with very abrupt transitions between the rising and flat parts of the curve.
 Condition (3) prevents some of the solid-body rotation curves where the asymptotic velocily is not well-constrained. from producing fits with extremelv. laree V...," Condition (3) prevents some of the solid-body rotation curves where the asymptotic velocity is not well-constrained, from producing fits with extremely large $V_0$ ."
 Conditions (1)(3) thus help removing unrealistic values for the fit parameters., Conditions (1)–(3) thus help removing unrealistic values for the fit parameters.
 Figure 1. (after LIO4. their Fie.," Figure \ref{fig:hayashidata} (after H04, their Fig."
 9) shows the results presented in HO04., 9) shows the results presented in H04.
 The LSB rotation curves show a broad distribution in 5 with two pronouncedpeaks: onearound 5~1 and a," The LSB rotation curves show a broad distribution in $\gamma$ with two pronouncedpeaks: onearound $\gamma
\sim 1$ and a"
Ihuninosity is a fairly robust measure of the stellar mass (i.e.. short-term variations in the SFI are averaged out).,"luminosity is a fairly robust measure of the stellar mass (i.e., short-term variations in the SFH are averaged out)."
 ILowever. for smaller galaxies or galaxies with unusual SEIIs. the 4.5 Ihuminosity may be less directly connected to the stellar mass.," However, for smaller galaxies or galaxies with unusual SFHs, the 4.5 luminosity may be less directly connected to the stellar mass."
 vanLoonetal.(2005) provided a method of estimating the masses of stellar clusters based on the number of stars brighter than the TRGB in the L/ band (3.76 jan) and the age ol the cluster., \citet{van05} provided a method of estimating the masses of stellar clusters based on the number of stars brighter than the TRGB in the $^\prime$ band (3.76 ) and the age of the cluster.
 Performing this caleulation for WLM is somewhat difficult because. unlike stellar clusters. its stellar population is not single aged.," Performing this calculation for WLM is somewhat difficult because, unlike stellar clusters, its stellar population is not single aged."
 ILowever. by adopting an age of 2 Gvr for the stellar population (i.e..theageofrecentstarformingeventthat.formedsignificantfractionofthestars:Dolphin 2000).. we arrive at atotal stellar mass of 1.1x 10*AL..," However, by adopting an age of 2 Gyr for the stellar population \citep[i.e., the
age of the recent star forming event that formed a significant
fraction of the stars;][]{dol00}, we arrive at atotal stellar mass of $\times$ $^7$."
 Because of the range in age of the stellar population and our incomplete sky coverage (his value is probably uncertain by at least505: however. it is reasonably consistent with the Leeetal.(2006) value of 1.8x10* M...," Because of the range in age of the stellar population and our incomplete sky coverage this value is probably uncertain by at least; however, it is reasonably consistent with the \citet{lee06} value of $\times$ $^7$."
 We have presented Spizer/IRAC imaging al 3.6 and 4.5 jan. V and Ioptical data from the LGGS (Masseyetal.2006) were combined with these observations to help determine the stellar twpes of objects detected in the IR.," We have presented /IRAC imaging at 3.6 and 4.5 V and I optical data from the LGGS \citep{mas06}
 were combined with these observations to help determine the stellar types of objects detected in the IR."
 We find that above the TRGB at 3.6sam. of the objects were undetected in the optical aud were undetected or misidentilied. even though the Masseyetal.(2006) cata ave complete 3 magnitudes below the TRGB.," We find that above the TRGB at 3.6, of the objects were undetected in the optical and were undetected or misidentified, even though the \citet{mas06}
 data are complete 3 magnitudes below the TRGB."
" It is likely that many of these objects were not detected in the optical due to dust production in the winds of the AGB stars. which efficiently absorb visible photons while IR radiation is relatively unallected,"," It is likely that many of these objects were not detected in the optical due to dust production in the winds of the AGB stars, which efficiently absorb visible photons while IR radiation is relatively unaffected."
 Comparing our photometry with the narrow-band optical study of WLM Demers 2004).. we detect nearly all of the objects they. classify as carbon stars with a few exceptions. mostly due to crowding effects.," Comparing our photometry with the narrow-band optical study of WLM \citep{bat04}, we detect nearly all of the objects they classify as carbon stars with a few exceptions, mostly due to crowding effects."
 However. of the 691 stars we detect ancl classify as AGBs only were detected by the narrow-band study.," However, of the 691 stars we detect and classify as AGBs only were detected by the narrow-band study."
 It is unclear what effect. this mav have on (he assumed. C/A ratio. since there is evidence (hat both mass-losing carbon stars and M-tvpe AGDs both escaped detection in the Battinelli&Demers(2004). ειπαν.," It is unclear what effect this may have on the assumed C/M ratio, since there is evidence that both mass-losing carbon stars and M-type AGBs both escaped detection in the \citet{bat04} study."
 Additionally. we detect a significant population of objects with positions on the IR. CMD consistent with mass-Iosing AGDs.," Additionally, we detect a significant population of objects with positions on the IR CMD consistent with mass-losing AGBs."
 For these objects we derive a conservative total MLB [ου the galaxy of 0.7-1.6x 10.7 Loassuming the entire population is composed of carbon stars.," For these objects we derive a conservative total MLR for the galaxy of $\times$ $^{-3}$ $^{-1}$ , assuming the entire population is composed of carbon stars."
If we include all of our,If we include all of our
where pis the density. Ap is the Boltzmann constant. peas the mean molecular weight aud ος is the isothermal sound-speed.,"where $\rho$ is the density, $k_\mathrm{B}$ is the Boltzmann constant, $\mu$ is the mean molecular weight and $c_\mathrm{s}$ is the isothermal sound-speed."
 Now the evolution is characterized by an isothermal shock followed by a weaker. D-tvpe ionization front.," Now the evolution is characterized by an isothermal shock followed by a weaker, 'D-type' ionization front."
 Thefrout velocity is now ep«dion:, Thefront velocity is now $v_\mathrm{D}<a_\mathrm{ion}$.
 For a full analysis see e.g. ?.., For a full analysis see e.g. \citet{1991pagd.book.....S}.
 As the hot gas expands. its density is reduced.," As the hot gas expands, its density is reduced."
 At the same time. the cold smroundius eas is compressed.," At the same time, the cold surrounding gas is compressed."
 Under the assumption that the homogeneously ionized region cousunies all UV- of the source it follows from Eq., Under the assumption that the homogeneously ionized region consumes all UV-photons of the source it follows from Eq.
 2. that the density of the hot eas fora constant flux aud a coustaut teiiperature Dor at any given finie is To calculate the front position «(t) we follow the approach of 7.., \ref{x_Stroem} that the density of the hot gas for a constant flux and a constant temperature $T_\mathrm{hot}$ at any given time is To calculate the front position $x(t)$ we follow the approach of \cite{2003adu..book.....D}.
 Under the assuuption of a thin shock which is traveling at a speed ος the rai pressure in the hot. ionized eas has to be equal to the ram pressure iu the cold gas where The pressure on the ionized side of the shock is mainly even by the thermal pressure of the hot gas where we have already introduced a coustaut fitting factor f to account for the approximations made (e.g. the one leading to Eq. 13).," Under the assumption of a thin shock which is traveling at a speed $v_\mathrm{s}$ the ram pressure in the hot, ionized gas has to be equal to the ram pressure in the cold gas where The pressure on the ionized side of the shock is mainly given by the thermal pressure of the hot gas where we have already introduced a constant fitting factor $f$ to account for the approximations made (e.g. the one leading to Eq. \ref{rho_Stroem}) )."
 Combining Eq., Combining Eq.
 6 and 7 and using Eqs., \ref{P_cold} and \ref{P_ion} and using Eqs.
 2 and L vields With the initial condition οσο)=or. we can solve by inteeration and obtain Using Eq., \ref{x_Stroem} and \ref{rho_Stroem} yields With the initial condition $x(t_0)=x_\mathrm{s}$ we can solve by integration and obtain Using Eq.
 1 it follows that for a plane-parallel iufall of a constant flux outo a homogeueous medium., \ref{rho_Stroem} it follows that for a plane-parallel infall of a constant flux onto a homogeneous medium.
" Iu order to investigate the effect of different initial conditions aud levels of UV-radiatiou on the formation of pilus we couduct a pavaimeter study,", In order to investigate the effect of different initial conditions and levels of UV-radiation on the formation of pillars we conduct a parameter study.
 All simulations were performed with WINE (609a). an implementation of ionizing radiatiou iu the trec-SPIT code VINE (??).. ," All simulations were performed with iVINE (G09a), an implementation of ionizing radiation in the tree-SPH code VINE \citep{2009ApJS..184..298W,2009ApJS..184..326N}. ."
As the mean free path of the atoms auc electrous iu the cold aud the hot phase is of order AzLO?cm and AmLottem. respectively (sce c.g. 24). which is much smaller than the leught-scales involved iu our simulations. the eas can be treated by the means of fiuid dvnamucs.," As the mean free path of the atoms and electrons in the cold and the hot phase is of order $\lambda\approx10^{12}\cm$ and $\lambda\approx10^{14}\cm$, respectively (see e.g. \citet{1991pagd.book.....S}) ), which is much smaller than the lenght-scales involved in our simulations, the gas can be treated by the means of fluid dynamics."
 The evolution of a turbulent molecular cloud under the influence of ionization spans several orders of magnitude in deusitv., The evolution of a turbulent molecular cloud under the influence of ionization spans several orders of magnitude in density.
 We therefore chose to solve the hydrodynamic equations with the method SPILL. a Laerangian approach with adaptive resolution.," We therefore chose to solve the hydrodynamic equations with the method SPH, a Lagrangian approach with adaptive resolution."
 To prevent artificial fragiieutation (2) the Jeans leusth has always to be resolved with at least 50 particles., To prevent artificial fragmentation \citep{1997MNRAS.288.1060B} the Jeans length has always to be resolved with at least 50 particles.
 This leads to a resolution limit which we eusure to be small enough by using a sufficient amount of total particles (see 3.1))., This leads to a resolution limit which we ensure to be small enough by using a sufficient amount of total particles (see \ref{RES_SF}) ).
 Iun iVINE. the ionizing radiation is assunied to inpiuge plane-parallel outo the simulated volume from the neeative x-direction.," In iVINE, the ionizing radiation is assumed to impinge plane-parallel onto the simulated volume from the negative x-direction."
 From the surface of iufall the radiation is propagated along the x-direction bv a rav-shooting aleoritlin., From the surface of infall the radiation is propagated along the x-direction by a ray-shooting algorithm.
 Alone these rays the ionization degree a is calculated for cach particle i. According to the ionization degree the pressure 2 of the particle is caleulated by ai linear interpolation between the temperature των ofthe hot. ionized aud the teniperature Toga of the cold. wn-ionized gas.," Along these rays the ionization degree $\eta_\mathrm{i}$ is calculated for each particle i. According to the ionization degree the pressure $P_\mathrm{i}$ of the particle is calculated by a linear interpolation between the temperature $T_\mathrm{hot}$ of the hot, ionized and the temperature $T_\mathrm{cold}$ of the cold, un-ionized gas."
 Hore. we assume both eas components to be isothermal. since for the deusitv range considered in our simulations heatius aud cooling should balance cach other to approximate isothermality (seee.g.7).," Here, we assume both gas components to be isothermal, since for the density range considered in our simulations heating and cooling should balance each other to approximate isothermality \citep[see
e.g.][]{1998ApJ...504..835S}."
 Following Eq., Following Eq.
" 3. the new pressure iu our sinulation is given as where pj is the SPIT-deusity of the particle / aud 65,4,=0.5 and ον=1.0 are the mean molecular weights of the ionized and the undonized gas in the case of pure hydrogen. respectively."," \ref{P_iso} the new pressure in our simulation is given as where $\rho_\mathrm{i}$ is the SPH-density of the particle $i$ and $\mu_\mathrm{ion}=0.5$ and $\mu_\mathrm{nion}=1.0$ are the mean molecular weights of the ionized and the un-ionized gas in the case of pure hydrogen, respectively."
 As a first test we verify Eq., As a first test we verify Eq.
 9 aud fit a value for f bw ionizing a slab ofatomic lyvdrogen with a constant homogeneous deusity of pectSOQCla3 and temperature of ωμά=LOTS., \ref{x_front} and fit a value for $f$ by ionizing a slab ofatomic hydrogen with a constant homogeneous density of $\rho_\mathrm{cold}=300\dens$ and temperature of $T_\mathrm{cold}=10\K$.
" We perform2, three different runs. correspouding to a low flux 1.6G«105em27s 1) an intermediate flux σονdOscm278 1) aud a high fux (FL.=1.5«&LOM~cm28 1)"," We perform three different runs, corresponding to a low flux $F_\mathrm{Ly}=1.66\times 10^9\Lyunit$ ), an intermediate flux $F_\mathrm{Ly}=5\times 10^9\Lyunit$ ) and a high flux $F_\mathrm{Ly}=1.5\times 10^{10}\Lyunit$ )."
 This corresponds to the ionization penetrating imuueciately iuto the first 0.55%. 1.67% and 5% of the reeion. respectively (see Eq. 2)).," This corresponds to the ionization penetrating immediately into the first $0.55\%$ , $1.67\%$ and $5\%$ of the region, respectively (see Eq. \ref{x_Stroem}) )."
 At the same time this is equal to placing the simulation volume further away or closer to the source. e.g. the O-star.," At the same time this is equal to placing the simulation volume further away or closer to the source, e.g. the O-star."
 The simulations are conducted with the same accuracy aud setup as iu the parameter study eiven below (see §2.3))., The simulations are conducted with the same accuracy and setup as in the parameter study given below (see \ref{ICs}) ).
 Fie., Fig.
 shows the resulting evolution of the frout., \ref{frontspeed} shows the resulting evolution of the front.
 As oue can 1.clearly see the approximations leacing to Eq., As one can clearly see the approximations leading to Eq.
 2 (f=1. the dotted lines in Fie. 1))," \ref{x_Stroem}
 $f=1$, the dotted lines in Fig. \ref{frontspeed}) )"
 do not produce satistactory results., do not produce satisfactory results.
" Dustead. assuming (ie. f= qe the dashed dues in Fig.t)) perfectlymatches the simulationsV during the entire simulated time of fai,— HOOMawr."," Instead, assuming (i.e. $f=\sqrt{\frac{5}{4}}$ , the dashed lines in \ref{frontspeed}) ) perfectlymatches the simulations during the entire simulated time of $t_\mathrm{sim}=500\kyr$ ."
 Thus. we keep this assuuption for this work.," Thus, we keep this assumption for this work."
 To produce different turbulent initial conditions we use the same approach asin (1000., To produce different turbulent initial conditions we use the same approach asin G09b.
" We set up 2« particles to resemble a homoecucous medium with Peold=30005,01 in a voluune of (1pcJ).", We set up $2\times 10^6$ particles to resemble a homogeneous medium with $\rho_\mathrm{cold}=300\dens$ in a volume of $(4\pc)^3$ .
 Then. we," Then, we"
2000).,.
. The most obvious one might be that the correlation is related. to. the source size., The most obvious one might be that the correlation is related to the source size.
" One might expect that the characteristic time scales or the “excess variance"" should. scale or relate with the central black hole mass (Edelson&Nandra1999).", One might expect that the characteristic time scales or the “excess variance” should scale or relate with the central black hole mass \cite{en}.
. However. the distribution. of the other parameters. (probably the inclination. the accretion rate and the absorption properties) in the sample can also make the correlation.," However, the distribution of the other parameters (probably the inclination, the accretion rate and the absorption properties) in the sample can also make the correlation."
 In fact. Bao Abramowicz (1906). suggested that the X-ray. variability is produced by bright-spots on a rotating accretion disk. and that the correlation is primarily a projection ellect.," In fact, Bao Abramowicz \shortcite{ba} suggested that the X-ray variability is produced by bright-spots on a rotating accretion disk, and that the correlation is primarily a projection effect."
 In their moclel. X-ray variability was enhanced due to the increased relativistic elfects with the change of inclination from a face-on to an edege-on disk.," In their model, X-ray variability was enhanced due to the increased relativistic effects with the change of inclination from a face-on to an edge-on disk."
 However. this model. contliets with the unified model of AGNs in which tvpical Sevfert 1 &ealaxies are observed. close to face-on.," However, this model conflicts with the unified model of AGNs in which typical Seyfert 1 galaxies are observed close to face-on."
 Recently. Abrassart Czerny (2000). proposed. an alternative explanation for the variability of ACGINs based on a cloud mocel of accretion onto a black hole.," Recently, Abrassart Czerny \shortcite{ac} proposed an alternative explanation for the variability of AGNs based on a cloud model of accretion onto a black hole."
" Phe small random rearrangement of the cloud distribution can happen on time-scales on the order ol 107 10""s and may lead to relatively high. amplituce variability in X-ray based on the large mean covering factor.", The small random rearrangement of the cloud distribution can happen on time-scales on the order of $10^2-10^6$ s and may lead to relatively high amplitude variability in X-ray based on the large mean covering factor.
 However. the number of the clouds ancl the mean covering factor mav be determined. by the accretion rate and. the central black hole mass. and thus the X-ray variability may be related to some fundamental parameters of the central engine in this regime.," However, the number of the clouds and the mean covering factor may be determined by the accretion rate and the central black hole mass, and thus the X-ray variability may be related to some fundamental parameters of the central engine in this regime."
 Due to the endeavors of the past decade. the pursuing of measuring the central black hole mass in the galaxies nuclei has resulted in fruitful achievements.," Due to the endeavors of the past decade, the pursuing of measuring the central black hole mass in the galaxies nuclei has resulted in fruitful achievements."
 First. the central massive dark possibly the supermassive black have been measured in the centre of some nearby galaxies based on high spatial resolution observations of stellar dynamics (Ixormendy&Richstone1995:Magorrianetal.," First, the central massive dark $-$ possibly the supermassive black $-$ have been measured in the centre of some nearby galaxies based on high spatial resolution observations of stellar dynamics \cite{kr,magorr}."
 1998). X tight correlation has been found between the central black hole mass and the stellar velocity. dispersion which links the inner nuclei with the large scale galactic environment (Lerrarese&Alerritt20007Gebhardtοἱal. 2000a).," A tight correlation has been found between the central black hole mass and the stellar velocity dispersion which links the inner nuclei with the large scale galactic environment \cite{fm,g2000a}."
. Secondly. reverberation mapping cata can give a reliable estimation of the central black hole mass in ACiNs (Peterson Wandel 1998: Wandel. Peterson Malkan 1999. jereafter. WPM: EIxaspi et al.," Secondly, reverberation mapping data can give a reliable estimation of the central black hole mass in AGNs (Peterson Wandel 1998; Wandel, Peterson Malkan 1999, hereafter WPM; Kaspi et al."
 2000. hereafter. ΝΑΙ. which cannot be determined by the stellar dynamics due to he bright nuclei.," 2000, hereafter KSNMGJ), which cannot be determined by the stellar dynamics due to the bright nuclei."
 Vhe interesting thing is that the relation oetween the mass of black hole and the bulge gravitationa jxotential also exists in ACNs (Gebhardt et al., The interesting thing is that the relation between the mass of black hole and the bulge gravitational potential also exists in AGNs (Gebhardt et al.
 2000b: elson 2000)., 2000b; Nelson 2000).
 This consisteney may. further support tha he reveberation mapping method is reliable to determine he black hole mass as the stellar. dvnamics., This consistency may further support that the reveberation mapping method is reliable to determine the black hole mass as the stellar dynamics.
" ALL these measurements have somewhat distangled the fundamenta xuvameters of the central engine. and allow us to study the relations between the ""excess variance” and the funcdamenta xuvamoeters (e.g. the black hole mass and the accretion rate). and thus shed new light on our understanding of the centra engine."," All these measurements have somewhat distangled the fundamental parameters of the central engine, and allow us to study the relations between the “excess variance” and the fundamental parameters (e.g. the black hole mass and the accretion rate), and thus shed new light on our understanding of the central engine."
" In this paper. we investigate the links between the Cexcess variance"" and the fundamental parameters of the central engine. ic. the central black hole mass."," In this paper, we investigate the links between the “excess variance” and the fundamental parameters of the central engine, i.e. the central black hole mass."
 Phe sample anc data reduction. are presented. in. 82: the statistical analvsis and results are presented in E3 ancl discussed. in 84: and finally. the conclusions are summarized in 85.," The sample and data reduction are presented in 2; the statistical analysis and results are presented in 3 and discussed in 4; and finally, the conclusions are summarized in 5."
 νο masses of the central black holes for a bunch of Sevlerts and QSOs have been measured using reverberation mapping data (WPM. Ho 1998. and ΙΝΔΙΟ). based on the assumption that the linc-emission material is gravitationally bound and hence has a near-Keplerian velocity. dispersion.," The masses of the central black holes for a bunch of Seyferts and QSOs have been measured using reverberation mapping data (WPM, Ho 1998, and KSNMJG), based on the assumption that the line-emission material is gravitationally bound and hence has a near-Keplerian velocity dispersion."
 This assumption is supported by the recent investigation on the Sevfert 1 galaxy NGC5548 which demonstrates that the kinematics of the broad line region is Ixeplerian (Peterson&Wandel1999)., This assumption is supported by the recent investigation on the Seyfert 1 galaxy NGC5548 which demonstrates that the kinematics of the broad line region is Keplerian \cite{pw}.
 There are two kinds of methods to estimate a typical velocity and hence the central black hole mass: one is to measure the Full Width at Half Maximum (EWHAL of each line in all spectra and caleulate the mean FWIIM. the other uses the the rms spectrum to compute the ΛΗΛ of the lines (WAIP and KSNALIC): the Virial niass is then determined to be AdjGnean) orfand Adj (ems) using EWAGuean) or/and FWAGms) together with the time lags of the lines corresponding to the variance of the ionizing continuum (see WPM and ΙΝΔΙΟ for details).," There are two kinds of methods to estimate a typical velocity and hence the central black hole mass: one is to measure the Full Width at Half Maximum (FWHM) of each line in all spectra and calculate the mean FWHM, the other uses the the rms spectrum to compute the FWHM of the lines (WMP and KSNMJG); the Virial mass is then determined to be $M_{\rm bh}$ (mean) or/and $M_{\rm bh}$ (rms) using FWHM(mean) or/and FWHM(rms) together with the time lags of the lines corresponding to the variance of the ionizing continuum (see WPM and KSNMJG for details)."
 We searched in the public ASCA archive up to Oct. 1999 for the sample objects of ΡΑΔΙΟ ancl Llo (1998) (the objects in \WPAL sample are included in the KWSNALIC sample). of which the masses of the central black holes 1tve been measured using the reverberation data. and found wt 24 objects have been. observed by ASCA.," We searched in the public ASCA archive up to Oct. 1999 for the sample objects of KSNMJG and Ho \shortcite{ho} (the objects in WPM sample are included in the KSNMJG sample), of which the masses of the central black holes have been measured using the reverberation data, and found that 24 objects have been observed by ASCA."
 Except two objects. PCLT|442 and POG1700|518. which cannot mect 1ο criteria to do the timingoo analysis (see the followinee ala processing procedures). all the others are listed. in able 1 with the measured. Virial reverberation mass (both Afw(mean) and. AMuii(rms): IAKSNMJCG and Ho 1998).," Except two objects, PG1411+442 and PG1700+518, which cannot meet the criteria to do the timing analysis (see the following data processing procedures), all the others are listed in table \ref{tab1} with the measured Virial reverberation mass (both $M_{\rm bh}$ (mean) and $M_{\rm bh}$ (rms); KSNMJG and Ho 1998)."
 For consistent with the former investigations. we acloptect the same method used by Nandra et al.," For consistent with the former investigations, we adopted the same method used by Nandra et al."
 (1997). anc ‘Turner et al., \shortcite{nandra} and Turner et al.
" (1999). to do the timing analvsis and extrac the “excess variance"" as follows."," \shortcite{turner}
 to do the timing analysis and extract the “excess variance” as follows."
 We used a 256-second time bin for the light curve., We used a 256-second time bin for the light curve.
 We combined the SISO ancl SISI data to increase the signal-to-noise ratio and required: al time bins to be at least exposed., We combined the SIS0 and SIS1 data to increase the signal-to-noise ratio and required all time bins to be at least exposed.
 We required that the time series in our analysis have at least 20 counts per bin and at least 20 bins in the final light curve., We required that the time series in our analysis have at least 20 counts per bin and at least 20 bins in the final light curve.
 Phe backgrouim was not subtracted from the light curves since it has been demonstrated by Turner et al. (, The background was not subtracted from the light curves since it has been demonstrated by Turner et al. (
"1999) and Leiehly (1999) that in practice the background is only a small fraction of the photons and it only has a small clfects on the value of ""excess variance"" for light curves with SISO|SISI counts rate greater than 0.5 counts per second.",1999) and Leighly (1999) that in practice the background is only a small fraction of the photons and it only has a small effects on the value of “excess variance” for light curves with SIS0+SIS1 counts rate greater than 0.5 counts per second.
 Even for the objects (DOGOSA44|349.. D€:0953|415. and NOGC4579). which have SISO|SISI counts rate less than 0.5 counts per second in our sample. we found that the dillerence is not large.," Even for the objects (PG0844+349, PG0953+415 and NGC4579) which have SIS0+SIS1 counts rate less than 0.5 counts per second in our sample, we found that the difference is not large."
 Finally. we," Finally, we"
"strong CO emission (Haoetal.""n2005:Schweitzer2006:Far-rahetal.2007a.b:Netzer 2007).","strong CO emission \citep{Hao05, Schweitzer06, Farrah07a, Farrah07b, Netzer07}."
". At high redshifts (z  6). observations at sub-mm and mm wavelengths (rest-frame far-infrared (FIR) to sub-mm. depending on wavelength and redshift range) of optically luminous quasars suggest that up to of quasars also fall within the ULIRG range. vei touM IR quasars Caoetal.ae8) (Carilliet»""Ho:OmontCominStudied2001.Mtul2003:Coxetal.2005:Beelen2200 "," At high redshifts (z $\sim$ 2-6), observations at sub-mm and mm wavelengths (rest-frame far-infrared (FIR) to sub-mm, depending on wavelength and redshift range) of optically luminous quasars suggest that up to of quasars also fall within the ULIRG range, similar to the IR quasars studied by \citet{Cao08} \citep{Carilli01, Omont01, Omont03, Cox05, Beelen06, Hao08}. ."
i et(2008) compare dynamical. gas and SMBH masses of ten z zz2 sub-mm detected quasars and z 2:2. sub-mm galaxies (SMGs). which are high redshift counterparts to ULIRGs. though ess extreme with more evenly distributed star formation instead of a localised starburst (Menéndez-Delmestreetal.2009).," \citet{Coppin08} compare dynamical, gas and SMBH masses of ten z $\approx 2$ sub-mm detected quasars and z $\approx 2$ sub-mm galaxies (SMGs), which are high redshift counterparts to ULIRGs, though less extreme with more evenly distributed star formation instead of a localised starburst \citep{Menendez-Delmestre09}."
. They tind that the fainter half of their quasar sample could likely be ‘transition objects’ between SMGs and luminous quasars. based on heir SMG-like surface densities and their proximity to the local relation. since luminous quasars tend to lie above this relation TgH/Mypnand typical SMGs tend to lie below.," They find that the fainter half of their quasar sample could likely be `transition objects' between SMGs and luminous quasars, based on their SMG-like surface densities and their proximity to the local $_{\mathrm{BH}}$ $_{\mathrm{sph}}$ relation, since luminous quasars tend to lie above this relation and typical SMGs tend to lie below."
 The highest redshift sub-mm-selected source currently known has also been observed © posses similar stellar and gas masses to this z zz2. sample of ransition objects. as well as a small AGN contribution (Coppinetal. 2009).," The highest redshift sub-mm-selected source currently known has also been observed to posses similar stellar and gas masses to this z $\approx 2$ sample of transition objects, as well as a small AGN contribution \citep{Coppin09}."
. These quasars eventually may produce enough energy Tom accretion to effectively blow out the surrounding gas. halting Palar formation and clearing the view to an optically bright quasar before hi their own growκ CURithe&Loeb2003:DiMatteoetal.2005:accretingitingHopkins 2008).," These quasars eventually may produce enough energy from accretion to effectively blow out the surrounding gas, halting star formation and clearing the view to an optically bright quasar before halting their own growth \citep{Wyithe03, DiMatteo05, Hopkins08}."
. MMThis regulation of star formation by the SMBH could produce the observed black 1ole-galaxy mass correlation. and different stages of the process should present different SFRs. with initially high SFRs declining s the AGN becomes more dominant. culminating in à phase of Palrong SMBH accretion little or no ongoing star formation (Kauffmann&Haehnelt2000:MGranatoetal.2004).," This regulation of star formation by the accreting SMBH could naturally produce the observed black hole-galaxy mass correlation, and different stages of the process should present different SFRs, with initially high SFRs declining as the AGN becomes more dominant, culminating in a phase of strong SMBH accretion and little or no ongoing star formation \citep{Kauffmann00, Granato04}."
. FIR luminosities can used to determine the star formation rates in quasar host galaxies., FIR luminosities can be used to determine the star formation rates in quasar host galaxies.
 The FIR emission is due to dust heated either by star formation or quasar emission (see Haasetal.(2003) for discussion)., The FIR emission is due to dust heated either by star formation or quasar emission (see \citet{Haas03} for discussion).
 Lutzetal.(2007.2008) find evidence that FIR emission in quasar host galaxies at all redshifts is caused by dust heated by star formation and not the quasar.," \citet{ Lutz07, Lutz08} find evidence that FIR emission in quasar host galaxies at all redshifts is caused by dust heated by star formation and not the quasar."
 The strength of PAH emission. which is found almost exclusively in star forming regions. tightly corresponds to the strength of FIR emission in n same hosts (Daleetal.2001:Calzetti2007:Lutz200hi," The strength of PAH emission, which is found almost exclusively in star forming regions, tightly corresponds to the strength of FIR emission in the same hosts \citep{Dale01, Calzetti07, Lutz08}."
 Beelenetal.(2006) measure FIR emission from six high quasars and derive FIR to radio spectral indexes consistent with local star forming galaxies without AGN., \citet{Beelen06} measure FIR emission from six high redshift quasars and derive FIR to radio spectral indexes consistent with local star forming galaxies without AGN.
 The FIR emission is seemingly uncontaminated by hotter dust potentially heated by the AGN., The FIR emission is seemingly uncontaminated by hotter dust potentially heated by the AGN.
 These results along with others (see for example. Efstathiou&Rowan-Robinson(1995):SerjeantHatziminaoglou (2009))) urther substantiate the claim that the FIR luminosity in quasars is dominated by star formation and not by the AGN.," These results along with others (see for example, \citet{Efstathiou95, Serjeant09}) ) further substantiate the claim that the FIR luminosity in quasars is dominated by star formation and not by the AGN."
" FIR luminosity traces ongoing star formation, but past star ormation also can be observed indirectly by measuring chemica abundances."," FIR luminosity traces ongoing star formation, but past star formation also can be observed indirectly by measuring chemical abundances."
 Several studies have found that high-redshift quasars ypically have metallicities greater than or equal to solar metallicity in the broad emission line region (BLR). which requires significan orevious star formation in the host (Hamann&Ferland1999:Di-etrichetal.2003:Warner2004:Nagao 2006)..," Several studies have found that high-redshift quasars typically have metallicities greater than or equal to solar metallicity in the broad emission line region (BLR), which requires significant previous star formation in the host \citep{Hamann99, Dietrich03, Warner04, Nagao06a}."
 These BLR studies rely on emission line ratios such as and that are sensitive to the selective enrichmen of nitrogen as a secondary element and/or to the decreasing temperatures and increasing metal-line saturations that occur in metal-rich gas., These BLR studies rely on emission line ratios such as and that are sensitive to the selective enrichment of nitrogen as a secondary element and/or to the decreasing temperatures and increasing metal-line saturations that occur in metal-rich gas.
 The main result for typically solar or higher metallicities near quasars has been corroborated by independent studies of the narrow emission lines (Grovesetsuoetal.2006) and narrow absorption lines (Hamann&Fer- in quasar spectra.," The main result for typically solar or higher metallicities near quasars has been corroborated by independent studies of the narrow emission lines \citep{Groves06, Nagao06b} and narrow absorption lines \citep{Hamann99, Dodorico04, Gabel06, Simon10} in quasar spectra."
 The metal-rich BLR result is true even for the highest redshifts studied. e.g. Pentericcietal.(2002):Jiangetal.(2007):Juarez(2009)... with redshifts out fo z = 6.4.," The metal-rich BLR result is true even for the highest redshifts studied, e.g. \citet{Pentericci02, Jiang07, Juarez09}, with redshifts out to z = 6.4."
 There is no known change in metallicity with redshift (see Matsuokaetal.(2009) and references above)., There is no known change in metallicity with redshift (see \citet{Matsuoka09} and references above).
 Simple chemical evolution models for quasars and elliptieal galaxies find that galactic centres tend to be more metal-rich than their halos. and a centralised starburst can enrich the galactic centre to supersolar abundances in a short time (109vr) MM&TerlevichMGranatoetal.2004:Hamann2007:me 2009).," Simple chemical evolution models for quasars and elliptical galaxies find that galactic centres tend to be more metal-rich than their halos, and a centralised starburst can enrich the galactic centre to supersolar abundances in a short time $\le 10^8~yr$ ) \citep{Friaca98, Hamann99, Granato01, Hamann02, Granato04, Hamann07, Juarez09}."
. These models. combined with the consistently super-solar gas abundances observed in quasar environments. imply that quasars tend to emerge after or near the end of (potentially) short centralised star formation epochs.," These models, combined with the consistently super-solar gas abundances observed in quasar environments, imply that quasars tend to emerge after or near the end of (potentially) short centralised star formation epochs."
 If the quasar phase emerges when star formation is on the decline. less advanced environments might have higher SFRs and lower abundances (Georgakakisetal.2009).," If the quasar phase emerges when star formation is on the decline, less advanced environments might have higher SFRs and lower abundances \citep{Georgakakis09}."
 The absorption lines in quasar spectra provide additional information about outflows that might be related to the blowout of gas from the host galaxies. and perhaps. about the gaseous remnants of recent galaxy mergers.," The absorption lines in quasar spectra provide additional information about outflows that might be related to the blowout of gas from the host galaxies, and perhaps, about the gaseous remnants of recent galaxy mergers."
 These absorption lines may be more common in quasars with higher SFRs if mergers and interactions trigger the starMM as in M manGeorgakakisetal.20€ Q0)..," These absorption lines may be more common in quasars with higher SFRs if mergers and interactions trigger the star formation as in ULIRGs \citep{ Weymann91, Becker00, Richards01a, Rupke05, Georgakakis09}."
" Quiar ""associated! sorption lines (AALS). near the emission redshifts with velocity widths less than 500 and broad absorption lines (BALs). with velocity widths greater than a few thousands... are examples of these potential merger. interaction and outflow signatures."," Quasar `associated' absorption lines (AALs), near the emission redshifts with velocity widths less than 500 and broad absorption lines (BALs), with velocity widths greater than a few thousand, are examples of these potential merger, interaction and outflow signatures."
 In systems with recent mergers. remnants from the M may manifest as a greater incidence of low-velocity 2000kms !)) AALs.," In systems with recent mergers, remnants from the interaction may manifest as a greater incidence of low-velocity $< 2000$ ) AALs."
 About of bright quasars contain AALS w‘th rest equivalent widths (REWs) of NLA.. and around of quasars have BALs (GangulyetNestoral.2001:Vestergaard2003:MisawaNAHidalgoetal. 2010)..," About of bright quasars contain AALs with rest equivalent widths (REWs) of $>$ 0.3, and around of quasars have BALs \citep{Ganguly01, Vestergaard03, Misawa03, Trump06, Nestor08, Wild08, Gibson08, Paola08}."
. A higher incidence of AALS or BALs may occur in the quasars with higher SFRs if these quasars have more recently experienced an interaction and/or there is a progression in quasar outflow characteristics with time., A higher incidence of AALs or BALs may occur in the quasars with higher SFRs if these quasars have more recently experienced an interaction and/or there is a progression in quasar outflow characteristics with time.
 In this paper. we present an exploratory observational study examining whether SFR in the host galaxies correlates with metallicity in the near-quasar environment.," In this paper, we present an exploratory observational study examining whether SFR in the host galaxies correlates with metallicity in the near-quasar environment."
 We measure metallicity in the quasar BLR from the rest frame ultraviolet (UV) spectrum and estimate galactic SFRs from the FIR luminosities., We measure metallicity in the quasar BLR from the rest frame ultraviolet (UV) spectrum and estimate galactic SFRs from the FIR luminosities.
 The data and analysis are deseribed in 2 and 3.., The data and analysis are described in \ref{sec:data} and \ref{sec:anal}. .
 Theresults and discussion are presentedin 4+.. with a brief summary in 5..," Theresults and discussion are presented in \ref{sec:disc}, with a brief summary in \ref{sec:sum}."
" We adopt cosmological parameters Hy=70kms /. Qj,=0.3. and Qa=0.7 throughout this work."," We adopt cosmological parameters $_0~=~70$ $^{-1}$, $\Omega_m~=~0.3$, and $\Omega_{\Lambda}~=~0.7$ throughout this work."
" We selectquasars for this study from the observed with MAMBO at IRAM at 1.2 mm by Carillietal.MIN""001). and and SCUBA at JCMT at tum by Isaaketal. (2002).. McMahonetal.(1999) and Priddeyetal.(2003) and all compiled by Haoetal. (2008).."," We selectquasars for this study from the sample observed with MAMBO at IRAM at 1.2 mm by \citet{Carilli01} and \citet{Omont01, Omont03} and SCUBA at JCMT at $\mu m$ by \citet{Isaak02}, , \citet{McMahon99} and \citet{Priddey03} and all compiled by \citet{Hao08}. ."
 The quasars all were selected to be optically bright with absolute B-band Magnitude Mz26.1 for the Carilli et al., The quasars all were selected to be optically bright with absolute B-band Magnitude $_B < -26.1$ for the Carilli et al.
 objects. Mp27.0 for the Omont etal.," objects, $_B < -27.0$ for the Omont etal."
Iu this work. our interesting is mainly focused on microwave ZP structures.,"In this work, our interesting is mainly focused on microwave ZP structures."
 There are 3 segments of mucrowave ZP structures recistered ou the microwave spectrograms the first oue occurred at frequency of about 6.70 GIIz CZP1). the secoud oue is at about 2.68 GIIz CZP2). and the third one is at about 1.08 CHIz (ZP3).," There are 3 segments of microwave ZP structures registered on the microwave spectrogram: the first one occurred at frequency of about 6.70 GHz (ZP1), the second one is at about 2.68 GHz (ZP2), and the third one is at about 1.08 GHz (ZP3)."
 With a simple glance. we find that the higher frequency of ZP structure appears at the earlier flare phase. while the lower frequeney of ZP structures take place in the later fare phase.," With a simple glance, we find that the higher frequency of ZP structure appears at the earlier flare phase, while the lower frequency of ZP structures take place in the later flare phase."
 The following is the details. (, The following is the details. (
1) Microwave ZP Structure at 6.10 ~ 7.00 Cz Iu the frequency of 5.20 ~ 7.60 GIIz. the microwave cussion always behaves as coutiunous type IV burst. aud the fine structure is very rare during solar flares.,"1) Microwave ZP Structure at 6.40 $\sim$ 7.00 GHz In the frequency of 5.20 $\sim$ 7.60 GHz, the microwave emission always behaves as continuous type IV burst, and the fine structure is very rare during solar flares."
 Towever. the most interesting plenomenon isan ZP structure occurred at frequency of around 6.70 GIIz iu the N2.2 flare.," However, the most interesting phenomenon is an ZP structure occurred at frequency of around 6.70 GHz in the X2.2 flare."
 The left panel of Fig.2 is the spectrogram of the ZP structure observed at SDRS/IInairou., The left panel of Fig.2 is the spectrogram of the ZP structure observed at SBRS/Huairou.
 It shows that the frequency range of the structure is occured. from 6.10 to 7.00 GITz with a central frequency of 6.70 GIIz., It shows that the frequency range of the structure is occurred from 6.40 to 7.00 GHz with a central frequency of 6.70 GHz.
 The time interval is from 01:19:50.2 to 01:19:51.5 UT. just at about { 1nimutes after the onset of the flare (01:16 UT). aud about 6 iuinutes before the GOES flare peak (01:56 UT).," The time interval is from 01:49:50.2 to 01:49:51.5 UT, just at about 4 minutes after the onset of the flare (01:46 UT), and about 6 minutes before the GOES flare peak (01:56 UT)."
 The duration of the ZP structure lasts for about 1.3 s. It is composed with 3 stripes in distorted sinusoidal wave shape arraved ou the longitudinal direction., The duration of the ZP structure lasts for about 1.3 s. It is composed with 3 stripes in distorted sinusoidal wave shape arrayed on the longitudinal direction.
 The whole sinusoidal wave shape drifts slowly to the low frequency. aud the drifting rate is about -300 MIITz/s. The frequency bandwidth of the zebra stripes is about LO ~ 60 MIIz.," The whole sinusoidal wave shape drifts slowly to the low frequency, and the drifting rate is about -300 MHz/s. The frequency bandwidth of the zebra stripes is about 40 $\sim$ 60 MHz."
 The frequency separation between the adjacent zebra stripes is about SO ~ 120 ATTTz. and the relative frequency separation is Af/fz 1.17 L.76Cs dnercases slowly with respect to time.," The frequency separation between the adjacent zebra stripes is about 80 $\sim$ 120 MHz, and the relative frequency separation is $\Delta f/f\simeq$ 1.17 – 1.76, increases slowly with respect to time."
 The emission of the structure is stronelv left-lLanded circular polarization with polarization degree (defined as po!=S7«1005©. here R and L are the emission flux subtracted the backeround emission) close to10054.," The emission of the structure is strongly left-handed circular polarization with polarization degree (defined as $pol=\frac{R-L}{R+L}\times 100\%$, here R and L are the emission flux subtracted the background emission) close to."
.. Additionally. the ZP presents superfine structures: each zebra stripe consists of miillisecoud spikes which are practically vertical bright lines in the spectrogram with duration at the limit of the time resolution (5 ms) aud the bandwidth of about LO 60 MIIz.," Additionally, the ZP presents superfine structures: each zebra stripe consists of millisecond spikes which are practically vertical bright lines in the spectrogram with duration at the limit of the time resolution (5 ms) and the bandwidth of about 40 – 60 MHz."
 The right paucl of Fig., The right panel of Fig.
 2 is the temporal profile of the microwave flux at frequency of 6.80 Cz associated with the ZP1 structure., 2 is the temporal profile of the microwave flux at frequency of 6.80 GHz associated with the ZP1 structure.
 In this fleure. the averaged level of the ambient background. emissiou before aud after the ZP structure is also preseuted (the dashed line) which is about 535 sfu at left- aud right-handed circular polarization.," In this figure, the averaged level of the ambient background emission before and after the ZP structure is also presented (the dashed line) which is about 535 sfu at left- and right-handed circular polarization."
 The intensity of the zebra stripes have cuhancements of about 10 ~ SO sfu ith respect to the ambient emissions. which is much hieher than the instrmucut seusitivitv (here 2%Se is about 1 5 sfu at frequency of 6.20 ~ 7.20 κ).," The intensity of the zebra stripes have enhancements of about 40 $\sim$ 80 sfu with respect to the ambient emissions, which is much higher than the instrument sensitivity (here $2\% S_{\bigodot}$ is about 4 $\sim$ 5 sfu at frequency of 6.20 $\sim$ 7.20 GHz)."
 The positious of stripes ive marked with letters (A. D aud C. which represent respectively the stripes at high. middle. aud low frequencies).," The positions of stripes are marked with letters (A, B and C, which represent respectively the stripes at high, middle, and low frequencies)."
 At the same time. Fig.," At the same time, Fig."
 2 is also inplviug another feature: the cussion intensity on the zebra stripes decreases from the low frequency stripe to the high frequency stripes., 2 is also implying another feature: the emission intensity on the zebra stripes decreases from the low frequency stripe to the high frequency stripes.
 Fig.2 indicates that the zebra stripes of ZP1 are in distorted sinusoidal wave shapes., Fig.2 indicates that the zebra stripes of ZP1 are in distorted sinusoidal wave shapes.
 If we suppose that the pattern is a iuixture of a general ZP structure and a quasi-periodic pulsation (QPP). then it is easy. to nuderstand the distorted sinusoidal wave shapes.," If we suppose that the pattern is a mixture of a general ZP structure and a quasi-periodic pulsation (QPP), then it is easy to understand the distorted sinusoidal wave shapes."
 Fig., Fig.
 3 is the result of Fast Fourier Transtormation (FFT) analvsis on the microwave cussion at frequency of 6:80. GIIz. which indicates that the period of QPP is about 375 ms.," 3 is the result of Fast Fourier Transformation (FFT) analysis on the microwave emission at frequency of 6.80 GHz, which indicates that the period of QPP is about 375 ms."
 This QPP belongs to a very short-period pulsation (VSP). (, This QPP belongs to a very short-period pulsation (VSP). (
2) Zebra Pattern Structure at 2.60 ~ 2.75 Cllz The second ZP structure is a weakly oue (narked as ZP2 in Fig.,2) Zebra Pattern Structure at 2.60 $\sim$ 2.75 GHz The second ZP structure is a weakly one (marked as ZP2 in Fig.
 1) at frequency of 2.60 ~ 2.75 CGIIz in 02:01:19 ~ 02:01:21 UT on the left-hauded circular polarization spectrogram observed at SBBRS/IInairou (Fig., 1) at frequency of 2.60 $\sim$ 2.75 GHz in 02:01:19 $\sim$ 02:01:21 UT on the left-handed circular polarization spectrogram observed at SBRS/Huairou (Fig.
 D)., 4).
 There are 2 stripes iu this structure which cau be discriminated from the spectrogram., There are 2 stripes in this structure which can be discriminated from the spectrogram.
 The frequency separation between the adjaceut zebra stripes is GO 70 MITz. and the relative frequency separation is about Afify~ 2.25. 2.60%.," The frequency separation between the adjacent zebra stripes is 60 – 70 MHz, and the relative frequency separation is about $\Delta f/f_{0}\simeq$ 2.23 – 2.60."
. The ZP structure lasts for about 1.5 s. It is very weak. that the cussion flux at," The ZP structure lasts for about 1.5 s. It is very weak, that the emission flux at"
of stars 1n a group that came from a single satellite.,of stars in a group that came from a single satellite.
 If 7 is the set of all satellites and ./ is the set of all groups. then where #/;; is the number of stars from satellite / in group / and n;=YNnij 18 the total number of number of stars in that group.," If $I$ is the set of all satellites and $J$ is the set of all groups, then where $n_{ij}$ is the number of stars from satellite $i$ in group $j$ and $n_{.j}=\sum_i n_{ij}$ is the total number of number of stars in that group."
 The mean value of purity 2=57Puritvt/)/|J| is à good indicator of to what extent the recovered groups correspond to real physical associations., The mean value of purity $P=\sum {\rm Purity}(j)/|J|$ is a good indicator of to what extent the recovered groups correspond to real physical associations.
 As a crucial step towards determining how groups recovered from a photometric survey can be related to aceretion events in the Galaxy’s history we first tsolate the limitations imposed by the phase-mixing of debris. rather than observational constraints.," As a crucial step towards determining how groups recovered from a photometric survey can be related to accretion events in the Galaxy's history we first isolate the limitations imposed by the phase-mixing of debris, rather than observational constraints."
 In we characterize how the peak density of debris — which determines whether it might be detectable. whether by a group-finder as a group or by some other means — depends on the physical characteristics of an accretion event.," In we characterize how the peak density of debris — which determines whether it might be detectable, whether by a group-finder as a group or by some other means — depends on the physical characteristics of an accretion event."
 We then apply what we have learned to understand which progenitor satellites produce detectable groups in a full stellar halo 4.2)) and finally to outline the prospects for interpreting the full distribution of group characteristics in terms of accretion histories 4.3))., We then apply what we have learned to understand which progenitor satellites produce detectable groups in a full stellar halo ) and finally to outline the prospects for interpreting the full distribution of group characteristics in terms of accretion histories ).
 For the latter two sections we use the idealized data set SI (see 2.2.0)). consisting of synthetic surveys with MSTO stars as tracers but the distance to the stars being assumed to be known exactly.," For the latter two sections we use the idealized data set S1 (see ), consisting of synthetic surveys with MSTO stars as tracers but the distance to the stars being assumed to be known exactly."
 Note that we restrict our attention to the properties of groups in stellar halos that correspond to unbound structures because the number and distribution of the structures that are still bound (i.e.. the missing satellite problem of the ACDM paradigm) is still an issue of much debate.," Note that we restrict our attention to the properties of groups in stellar halos that correspond to unbound structures because the number and distribution of the structures that are still bound (i.e., the missing satellite problem of the $\Lambda$ CDM paradigm) is still an issue of much debate."
 In the simulations. the star formation recipes were tuned to match the observed distribution of the still surviving satellite galaxies in. the Milky Way known at the time of the publication of ?..," In the simulations, the star formation recipes were tuned to match the observed distribution of the still surviving satellite galaxies in the Milky Way known at the time of the publication of \cite{2005ApJ...635..931B}."
 Hence the number of satellites are not a prediction of the models but rather a boundary condition., Hence the number of satellites are not a prediction of the models but rather a boundary condition.
 The true prediction of the simulated models ts the number of unbound structures corresponding to the ancestral siblings of the present day classical dwarf spheroidal satellites that are formed by the dynamical interaction of these satellites with the potential of the host galaxy., The true prediction of the simulated models is the number of unbound structures corresponding to the ancestral siblings of the present day classical dwarf spheroidal satellites that are formed by the dynamical interaction of these satellites with the potential of the host galaxy.
 The peak density of an unbound structure evolves in time from an initial Όμωςmaximum state set by the structural properties of the satellite before disruption to a final minimum state where it 1s fully mixed in the entire volume roughly outlined by the orbit of the satellite., The peak density $\rho_{\rm peak}$ of an unbound structure evolves in time from an initial maximum state set by the structural properties of the satellite before disruption to a final minimum state where it is fully mixed in the entire volume roughly outlined by the orbit of the satellite.
 In the intermediate state. transient structures such as streams. shells and clouds are apparent.," In the intermediate state, transient structures such as streams, shells and clouds are apparent."
 Below. we individually analyze each simulated satellite debris (a total of 1515 satellites? debris from our eleven stellar halo models) in order to understand the dependence of peak density on the properties of its progenitor satellite during each of these stages.," Below, we individually analyze each simulated satellite debris (a total of 1515 satellites' debris from our eleven stellar halo models) in order to understand the dependence of peak density on the properties of its progenitor satellite during each of these stages."
 Initially the structure is bound and its peak density Ομως is observed to depend primarily upon the luminosity L of the progenitor satellite.," Initially the structure is bound and its peak density $\rho_{\rm
 peak}$ is observed to depend primarily upon the luminosity $L$ of the progenitor satellite."
" The exact dependence assumed in the simulations is given by ΌμωςxL"" and comes from fitting King profiles to observed dwarf galaxy data (?)..", The exact dependence assumed in the simulations is given by $\rho_{\rm peak} \propto L^{0.4}$ and comes from fitting King profiles to observed dwarf galaxy data \citep{1998ARAA..36..435M}.
" In the final. completely phase mixed state the ensemble of stars in the original satellite occupy a large volume Το, and have peak density of order L/V,,;,.."," In the final, completely phase mixed state the ensemble of stars in the original satellite occupy a large volume $V_{mix}$ and have peak density of order $L/V_{mix}$."
 In a spherical potential. Ἐν is expected to be inversely dependent on the circularity of the orbit € because the debris occupies the region between the apocenter and the pericenter of the orbit and an eccentric orbit (with lower €) explores a wider range m radius.," In a spherical potential, $V_{mix}$ is expected to be inversely dependent on the circularity of the orbit $\epsilon$ because the debris occupies the region between the apocenter and the pericenter of the orbit and an eccentric orbit (with lower $\epsilon$ ) explores a wider range in radius."
 In the case of an axisymmetric potential the debris also occupies the region generated by the precession of the orbit about the axis of the system., In the case of an axisymmetric potential the debris also occupies the region generated by the precession of the orbit about the axis of the system.
" Evolution in the intermediate state is governed by phase mixing. which occurs on a characteristic time scale T,,i, that depends on both the circularity ε and the mass of the satellite (??).. which in turn is correlated to luminosity L."," Evolution in the intermediate state is governed by phase mixing, which occurs on a characteristic time scale $T_{mix}$ that depends on both the circularity $\epsilon$ and the mass of the satellite \citep{1998ApJ...495..297J,1999MNRAS.307..495H}, which in turn is correlated to luminosity $L$ ."
 In this stage. the density has both spatial and temporal dependence.," In this stage, the density has both spatial and temporal dependence."
" Using the action angle formalism ? analytically deduced the general dependence as Όμωςκ.1Ger.) for spherical potentials and as picax for axis-symmetric potentials. where r is the distance1/G761,) of the peak density of a satellite debris from the center of the parent galaxy and £4, Is the time since accretion of the progenitor satellite."," Using the action angle formalism \cite{1999MNRAS.307..495H} analytically deduced the general dependence as $\rho_{\rm peak} \propto 1/(r^2 t_{\rm acc}^2)$ for spherical potentials and as $\rho_{\rm peak} \propto 1/(r^2 t_{\rm acc}^3)$ for axis-symmetric potentials, where $r$ is the distance of the peak density of a satellite debris from the center of the parent galaxy and $t_{\rm acc}$ is the time since accretion of the progenitor satellite."
 Rather than model these multiple stages and effects independently. we instead examine empirically to what extent the peak densities of debris in our simulations can be connected to progenitor properties with a single formula.," Rather than model these multiple stages and effects independently, we instead examine empirically to what extent the peak densities of debris in our simulations can be connected to progenitor properties with a single formula."
 Overall. the physical effects described above lead us to expect that pi; follows a 1i? radial dependence andis largest for objects that are luminous. recently accreted and are disrupting on mildly eccentric orbits.," Overall, the physical effects described above lead us to expect that $\rho_{\rm peak}$ follows a $1/r^2$ radial dependence andis largest for objects that are luminous, recently accreted and are disrupting on mildly eccentric orbits."
 We find that the relation cake provides the least scatter about the actual value pai measured for the simulated satellite debris 3)) and is a reasonable representation of trendsdiscussed earlier.," We find that the relation provides the least scatter about the actual value $\rho_{\rm
 peak}$ measured for the simulated satellite debris ) and is a reasonable representation of trendsdiscussed earlier."
" Note that in 3.. instead of associating p,4i simply with the"," Note that in , instead of associating $\rho_{\rm peak}$ simply with the"
velocily dispersion of the Volume as a whole.),velocity dispersion of the Volume as a whole.)
 There are more points on the plot at the larger distance. hence the distribution is more filly shown including the wings: but it is essentiallv (he same distribution.," There are more points on the plot at the larger distance, hence the distribution is more fully shown including the wings; but it is essentially the same distribution."
 Similarly. there is a diagonal line of points between about 1.5 and 2.5 Mpc. which could easily be taken as a sign of front side infall into Centaurus or Msi.," Similarly, there is a diagonal line of points between about 1.5 and 2.5 Mpc, which could easily be taken as a sign of front side infall into Centaurus or M81."
 But Chis is made up of galaxies in at least four wicely-separated parts of the sky. so altributing a coordinated motion to M81 or Centaurus is questionable: and anvwav the structure disappears when error bars are introduced (as in the right panel of Figure (1))).," But this is made up of galaxies in at least four widely-separated parts of the sky, so attributing a coordinated motion to M81 or Centaurus is questionable; and anyway the structure disappears when error bars are introduced (as in the right panel of Figure \ref{distance}) ))."
 po itis al least plausible (hat apparently clear signs of infall seen in other investigations are mistaken., So it is at least plausible that apparently clear signs of infall seen in other investigations are mistaken.
 We now turn (o other wavs of relating peculiar motions to luminous matter., We now turn to other ways of relating peculiar motions to luminous matter.
 Perhaps the assumption of dvnamical vouth implicit in the idea of infall is not right. al least close to the giants.," Perhaps the assumption of dynamical youth implicit in the idea of infall is not right, at least close to the giants."
 If these regions are dvnanmically old. (hat is if the groups are virialized. (hen there should be a much lareer velocity dispersion close to the eiant compared with farther out. (," If these regions are dynamically old, that is if the groups are virialized, then there should be a much larger velocity dispersion close to the giant compared with farther out. ("
It should also be true more generally that satellite galaxies deeper inside a potential well show higher speeds. as potential energy is converted (o kinetic.),"It should also be true more generally that satellite galaxies deeper inside a potential well show higher speeds, as potential energy is converted to kinetic.)"
 To look for (his we choose (wo zones. an inner zone out to 1 Mpe from the giant and an outer zone [rom 1 to 2 Mpe away.," To look for this we choose two zones, an inner zone out to 1 Mpc from the giant and an outer zone from 1 to 2 Mpc away."
 The resulating velocity clispersions are shown in Table (6))., The resulating velocity dispersions are shown in Table \ref{apple}) ).
 Included is an F-ratio calculation. which shows the probability of the outer sample being drawn from a significantly different. (smaller) parent distribution.," Included is an F-ratio calculation, which shows the probability of the outer sample being drawn from a significantly different (smaller) parent distribution."
 + 2i (bub.) «κ παpD((2/3).,"v_y v_z - d^2 ) + ], b_y b_z - d^2 ] (2/3)."
" Here. a= 3/Ah,. a=hk/h, and £=visfA."," Here, $\alpha = \beta / \A k_y$ , $a=k_z / k_y$ and $\xi = \nu k_y^2 / \A$."
 One can see that the correction to ihe magnetic stress is much larger than that to the Ievnolds stress (recall Chat £«&1]., One can see that the correction to the magnetic stress is much larger than that to the Reynolds stress (recall that $\xi \ll 1$ ).
 Comparing the zeroth order solution for the Reynolds stress and second order for the Maxwell sivess. we find that the former is dominant if 57a.τε77«| (giving us the result of the 2D case).," Comparing the zeroth order solution for the Reynolds stress and second order for the Maxwell stress, we find that the former is dominant if $\gamma^2 \alpha^{-2} \xi^{-2/3} \ll 1$ (giving us the result of the 2D case)."
 In the opposite case (στα7€.2%ss 1). (he transport of angular momentum is mainly controlled by the magnetic part of the stress by Alfvénn wave turbulence.," In the opposite case $\gamma^2 \alpha^{-2} \xi^{-2/3} \gg 1$ ), the transport of angular momentum is mainly controlled by the magnetic part of the stress by Alfvénn wave turbulence."
 In particular. the turbulent viscosity vp=ου)—(byb-))/A is now positive. reminiscent of the direct. cascade of enerev in MIID turbulence (Nim&Duhrulle2001).," In particular, the turbulent viscosity $\nu_T = (\langle v_y v_z \rangle - \langle b_y b_z \rangle)/ \A$ is now positive, reminiscent of the direct cascade of energy in MHD turbulence \citep{Kim01}."
. These results enable us to calculate the characteristic seale at which the transition from the Rossby wave turbulence (with negative viscosity) to the Alfvénn turbulence (positive viscosity) occurs. Ly=20ABI3v)! which corresponds to the Rhines (elerrvine to the scale which separates. in ordinary geostrophic turbulence. eddy turbulence from wave turbulence).," These results enable us to calculate the characteristic scale at which the transition from the Rossby wave turbulence (with negative viscosity) to the Alfvénn turbulence (positive viscosity) occurs, $L_R = 2 \pi (\A B_0^3 / \beta^3 \nu)^{1/4}$, which corresponds to the Rhines (referring to the scale which separates, in ordinary geostrophic turbulence, eddy turbulence from wave turbulence)."
 Our result shows that Lj~Ri5RU as we assunie a unit Prandtl number., Our result shows that $L_R \sim R_e^{1/4} \sim R_m^{1/4}$ as we assume a unit Prandtl number.
" However. it is easy to show that for 2v (as is relevant [or the tachocline) or equivalentlv [for P,<I. the first correction to the magnetic potential is the same as previously,"," However, it is easy to show that for $\eta \gg \nu$ (as is relevant for the tachocline) or equivalently for $P_m \ll 1$, the first correction to the magnetic potential is the same as previously."
 Consequently. (he magnetic stress remains the same. wilh ν΄ being replaced by i.," Consequently, the magnetic stress remains the same, with $\nu$ being replaced by $\eta$."
 Thus. for low Prandtl number. the crossover scale is Lj;=2=(ABI/Pubn»! ," Thus, for low Prandtl number, the crossover scale is $L_R = 2 \pi (\A B_0^3 / \beta^3 \eta)^{1/4} \sim R_m^{1/4}$."
This scaling is (he same as that of Diamondetal.(2006) obtained in a slightly different context (MIID turbulence on a 9 plane without a shear flow)., This scaling is the same as that of \citet{Diamond06} obtained in a slightly different context (MHD turbulence on a $\beta$ plane without a shear flow).
 In a similar way. we can calculate the turbulent. diffusivity app.," In a similar way, we can calculate the turbulent diffusivity $\eta_T$."
 To caleulate the first two leading order terms. we need to keep terms up to the third order in + for the magnetic potential (not shown here for brevitv).," To calculate the first two leading order terms, we need to keep terms up to the third order in $\gamma$ for the magnetic potential (not shown here for brevity)."
" Here again. only the leading order term survives when assuming o(b,)=oó(—h,) andthus the turbulent diffusivity remains the same as the one obtained in the hydrodynamical case [see Eq. (1))]."," Here again, only the leading order term survives when assuming $\phi(k_y) = \phi(-k_y)$ andthus the turbulent diffusivity remains the same as the one obtained in the hydrodynamical case [see Eq. \ref{TransportParticules}) )]."
 the companions prevent a fit from giving meaningful h and fy values (this ellect was not significant for objects labeled. 7) so the By values were derived individually by identifving à symmetry axis and computing the flux from the part of the galaxy on the other side of this axis [rom the companion ancl multiplving this Hux by two.,"“**” the companions prevent a fit from giving meaningful $h$ and $I_0$ values (this effect was not significant for objects labeled “*”), so the $B_T$ values were derived individually by identifying a symmetry axis and computing the flux from the part of the galaxy on the other side of this axis from the companion and multiplying this flux by two."
 Additionally. ¥CC 2062 was treated this way. since the light from this irregular galaxy peaks so far away from its centre.," Additionally, VCC 2062 was treated this way, since the light from this irregular galaxy peaks so far away from its centre."
 Since we will ultimatelv consider the LE for the entire magnitude rango 22«Alpll. it is important to consider how the RCS D magnitudes used at very bright magnitudes compare with the By magnitudes we clerive from the exponential fittingD method at fainter magnitudes.," Since we will ultimately consider the LF for the entire magnitude range $-22 < M_B < -11$, it is important to consider how the RC3 $B_T$ magnitudes used at very bright magnitudes compare with the $B_T$ magnitudes we derive from the exponential fitting method at fainter magnitudes."
5 We made the transition from using the dilferent. types of measurements at £y=15., We made the transition from using the different types of measurements at $B_T = 15$.
 For objects within 0.5 magnitude of this transition magnitude. the cdillerence A between the Dp values obtained from the B€3 and those we derive is A=0.60.10 mag. which is smaller than our expected uncertainty in the ο values we derive (see above) and much smaller than the binwidth (1 mag) we use to compute the LE.," For objects within 0.5 magnitude of this transition magnitude, the difference $\Delta$ between the $B_T$ values obtained from the RC3 and those we derive is $\Delta = 0.16 \pm 0.10$ mag, which is smaller than our expected uncertainty in the $B_T$ values we derive (see above) and much smaller than the binwidth (1 mag) we use to compute the LF."
 The HC€3 magnitudes will be more accurate for brighter galaxies. but even at the faintest limits that we use them. they therefore appear to be accurate enough for our purposes.," The RC3 magnitudes will be more accurate for brighter galaxies, but even at the faintest limits that we use them, they therefore appear to be accurate enough for our purposes."
 ]t is also instructive to compare our derived magnitudes with the magnitudes in the VCC for all chwarfs common o both catalogues (see Fig., It is also instructive to compare our derived magnitudes with the magnitudes in the VCC for all dwarfs common to both catalogues (see Fig.
 4)., 4).
 For Mg=16 to Mg1l we have 318 galaxies. 232 (73 %)) of which are VOC members.," For $M_B = -16$ to $M_B = -11$ we have 318 galaxies, 232 (73 ) of which are VCC members."
 Some are not (like (OC 7346) since they are not in he VCC fields. — in this study we observed fields bevond the northern extent of the VOC (see Fig., Some are not (like UGC 7346) since they are not in the VCC fields – in this study we observed fields beyond the northern extent of the VCC (see Fig.
 1) survey area., 1) survey area.
 Some are not VOC members (like the IBMSS objects) because heir surface brightnesses were too low to be included. in hat sample., Some are not VCC members (like the IBM88 objects) because their surface brightnesses were too low to be included in that sample.
 From Fig. 3. we see that the VCC magnitudes xighter than By=18 tend to be systematically brighter han ours bv about half a magnitude (a similar. albeit smaller. effect was noted by Bineeeli et al.," From Fig, 3, we see that the VCC magnitudes brighter than $B_T = 18$ tend to be systematically brighter than ours by about half a magnitude (a similar, albeit smaller, effect was noted by Binggeli et al."
 1995 who compared the VCX magnitudes with the By values derived w Ichikawa et al., 1995 who compared the VCC magnitudes with the $B_T$ values derived by Ichikawa et al.
 1986)., 1986).
 Painter than Dy=17 there is no systematic olfset between the VCC magnitudes anc ours (the mean cillerence is Br (present) Br (VOC) = 001 FE 0.04 mag)., Fainter than $B_T = 17$ there is no systematic offset between the VCC magnitudes and ours (the mean difference is $B_T$ (present) $-$ $B_T$ (VCC) = $-$ 0.01 $\pm$ 0.04 mag).
 > (11) the nominal probability £ of being a cluster meniber. as determined in Section 3 from a comparison with the background fields: (12) Galaxy rating according to the scheme described in Table 1: (13) absolute magnitude Alp=Dr31.15Ads. where the Galactic LExtinetion ely values (twpically 0.1 0.2 magnitudes) are adopted from NED and. based on the measurements of Schlegel et al. (," 1pt (11) the nominal probability $P$ of being a cluster member, as determined in Section 3 from a comparison with the background fields; 1pt (12) Galaxy rating according to the scheme described in Table 1; 1pt (13) absolute magnitude $M_B = B_T - 31.15 - A_B$, where the Galactic Extinction $A_B$ values (typically 0.1 – 0.2 magnitudes) are adopted from NED and based on the measurements of Schlegel et al. ("
1998).,1998).
 We adopt a distance mocdulus of 31.15 for the Virgo Cluster (Ponry ct al., We adopt a distance modulus of 31.15 for the Virgo Cluster (Tonry et al.
 2001)., 2001).
auplitude is quite unusual,amplitude is quite unusual.
 We can estinate the mean displacements aud velocities that are needed. to produce he observed density contrast., We can estimate the mean displacements and velocities that are needed to produce the observed density contrast.
 A density wave is created by cobereut oscillations of stars around their equilibrium orbits., A density wave is created by coherent oscillations of stars around their equilibrium orbits.
 Because the spiral is ightly wrapped. the largest contribution comes from the radial displacements. or. whose pliases are rapidly varving πιοΊος of the equilibrium position r.," Because the spiral is tightly wrapped, the largest contribution comes from the radial displacements, $\delta r$, whose phases are rapidly varying functions of the equilibrium position $r$."
 If we write the yhase as ορ{μοιInti)20|}. then a little bit of algebra shows that the change in surface density. ὃμ satisfies the relation when oZ1.," If we write the phase as $\exp\{i[\alpha \ln(r) - 2\theta]\}$, then a little bit of algebra shows that the change in surface density, $\delta \mu$ satisfies the relation when $\alpha \gg 1$."
 For our spiral a=P1201 9.4. which mcaus that when A~0.03. oe~0.003. a very small munber indecd!," For our spiral $\alpha = \frac{2}{\sin 12\fdg 1} = 9.5$ , which means that when $\frac{\delta \mu}{\mu} \simeq 0.03$, $\frac{\delta r}{r} \simeq 0.003$, a very small number indeed!"
 Tn a similar spirit we cau cstimate the velocities. óV. associated with such displacements.," In a similar spirit we can estimate the velocities, $\delta V$, associated with such displacements."
 The radial displacements will oscillate with a frequeney typically less than the natural racial frequeney 5. which in turn is roughly V2 times the mean angular rotation rate Q.," The radial displacements will oscillate with a frequency typically less than the natural radial frequency $\kappa$, which in turn is roughly $\sqrt{2}$ times the mean angular rotation rate $\Omega$."
" Hence n~0.005, or 0.3lo"," Hence $\frac{\delta V}{V_{\mathrm{c}}} \simeq 0.005$, or $0.3$."
 These siuall uunubers cau be increased by reducing the disk lieht. contribution and attributing it to a spheroid., These small numbers can be increased by reducing the disk light contribution and attributing it to a spheroid.
 Even a teu-fol reduction would produce fractional displacements of and velocities ofonlv ο1.," Even a ten-fold reduction would produce fractional displacements of $3\%$, and velocities of only $3$."
 Such low velocities are unlikely to lead to shocks in anv neutral eas that mav be present im the disk. aud therefore to iu star formation that would signal the preseuce of gas.," Such low velocities are unlikely to lead to shocks in any neutral gas that may be present in the disk, and therefore to any star formation that would signal the presence of gas."
 By reducing the disk light contribution we also reduce the importance of the spiral’s self-gravity aud increase the likelihood of a tidal origin., By reducing the disk light contribution we also reduce the importance of the spiral's self-gravity and increase the likelihood of a tidal origin.
 Fig.6 shows several possible perturbers., \ref{fig6} shows several possible perturbers.
 When the sclteravity becomes negligible. oue can rule out even a distant passage of a bie perturber like[J80.. since he tidal distortions will not propagate to the center.," When the self-gravity becomes negligible, one can rule out even a distant passage of a big perturber like, since the tidal distortions will not propagate to the center."
 That leaves close passages bv the faiuter objects. such as ιο two VOC dwarf galaxies. as possible cause of +the spira seen in IC3328.," That leaves close passages by the fainter objects, such as the two VCC dwarf galaxies, as possible cause of the spiral seen in IC3328."
 The other plausible scenario is that most of the lelt does come from the disk and we are seeing swing alplited noise (Toomre haluajs 1991)., The other plausible scenario is that most of the light does come from the disk and we are seeing swing amplified noise (Toomre Kalnajs 1991).
 The gain of the swing amplifier is uot large enough to amplify the WN stellar density fluctuations to the observed .QUÉ{ level. but a siall amount huupy gas could provide the necessary leading perturbations which then are amplified by a factor of 10230 as the shear transforms them iuto a trailing spirals (Toouwe 1981).," The gain of the swing amplifier is not large enough to amplify the $\sqrt{N}$ stellar density fluctuations to the observed $3\%$ level, but a small amount lumpy gas could provide the necessary leading perturbations which then are amplified by a factor of $10-30$ as the shear transforms them into a trailing spirals (Toomre 1981)."
 ITuchtineier Richter’s (1986) upper Int of =ο.LOCAL: for the content of IC3328 docs not rule out this possibility., Huchtmeier Richter's (1986) upper limit of $\approx 6\cdot 10^7 M_{\sun}$ for the content of IC3328 does not rule out this possibility.
 Dwarf galaxies come in two basic brands: dwiuf ellipticals (dEs) and dwarf mregublus5 (dhrs)., Dwarf galaxies come in two basic brands: dwarf ellipticals (dEs) and dwarf irregulars (dIrrs).
 The conuuou view is that dEs are spheroids and divs are disks., The common view is that dEs are spheroids and dIrrs are disks.
 Iu the case of nregulus. the disk nature is clearly indicated by the typical rotation pattern found in III radio data (at least for dbrs more bhpuuiuous than Ap= 12).," In the case of irregulars, the disk nature is clearly indicated by the typical rotation pattern found in HI radio data (at least for dIrrs more luminous than $M_B = -12$ )."
 Owine to the absence of gas and hencethe lack of an casily accessible kinematic tracer. the situation is less clear for," Owing to the absence of gas and hencethe lack of an easily accessible kinematic tracer, the situation is less clear for"
"where the scale height /1, in eq. (5))",where the scale height $h_\mathrm{e}$ in eq. \ref{eq-emp}) )
 is to 4/2., is to $h/2$.
 We can derivethe maximum of DMsin|/| and the seale height / from the distribution of DMπρι) of our sampleof 38 pulsars shown in Fig. {.., We can derivethe maximum of $\DM \sin|b|$ and the scale height $h$ from the distribution of $\DM \sin|b|(z)$ of our sampleof 38 pulsars shown in Fig. \ref{fig:1}.
 A two-parameter fit yielded μή=21.7€1.Sempe and 1=0.93+0.13kpe. giving nO)=fono0.023+0.004cm.," A two-parameter fit yielded $f_0
n_\mathrm{c0} h = 21.7\pm 1.5\cmcube\pc$ and $h = 0.93\pm 0.13\kpc$, giving $\nel(0) = f_0 n_\mathrm{c0} = 0.023\pm 0.004\cmcube$."
" As the maximum of EM,sin|| equals EM,sin ||. we have nz(0/2-fonih/2 (see Sect."," As the maximum of $\EM_\mathrm{p} \sin|b|$ equals $\EM_\mathrm{c} \sin|b|$ , we have $\nel^2(0)h/2 = f_0 n_\mathrm{c0}^2
h/2 = 2.8\pm 0.3\cmsix\pc$ (see Sect."
 2): hence nz(0)=0.0060 em. i4=0.26+0.03em™ and fo=0.09+ 0.02.," 2); hence $\nel^2(0) = 0.0060
\pm 0.0011\cmsix$ , $n_\mathrm{c0} = 0.26\pm 0.03\cmcube$ and $f_0 =
0.09\pm 0.02$ ."
that models of constant star formation with a canonical LAI could explain observations of the Galactic centre.,that models of constant star formation with a canonical IMF could explain observations of the Galactic centre.
 1n this work we use calculated mass loss rates due to stellar collisions as a method to constrain the PDME for main sequence stars in the Galactic centre., In this work we use calculated mass loss rates due to stellar collisions as a method to constrain the PDMF for main sequence stars in the Galactic centre.
 We construct a simple model to estimate the actual mass loss rate in the Galactic centre based on observed. x-ray emission., We construct a simple model to estimate the actual mass loss rate in the Galactic centre based on observed x-ray emission.
 PDMES hat predict mass loss rates from stellar collisions &reater han the observed. rate are precluded., PDMFs that predict mass loss rates from stellar collisions greater than the observed rate are precluded.
" This method allows us to place conservative constraints on the PDME, because we do not include the contribution to the mass loss rate rom stellar winds from massive evolved. stars (Baganolletal. 2003)."," This method allows us to place conservative constraints on the PDMF, because we do not include the contribution to the mass loss rate from stellar winds from massive evolved stars \citep{baganoff:2003}."
. Specifically. this method allows us to place a lower imit on the power-law slope and an upper limit on the minimum stellar mass of the PDME in the Galactic centre (sce 5)).," Specifically, this method allows us to place a lower limit on the power-law slope and an upper limit on the minimum stellar mass of the PDMF in the Galactic centre (see \ref{sec:constrain}) )."
 Inclusion of the mass loss rate from stellar winels (or other sources) could further constrain the PDAIF of the Galactic centre., Inclusion of the mass loss rate from stellar winds (or other sources) could further constrain the PDMF of the Galactic centre.
 We present novel. analytical models to calculate the amount of stellar mass lost due to stellar collisions between main sequence stars in 2. through 2.3.," We present novel, analytical models to calculate the amount of stellar mass lost due to stellar collisions between main sequence stars in \ref{sec:condition} through \ref{sec:mass_loss_direct}."
 In 3. we develop the formalism for caleulating collision rates in the Galactic center., In \ref{sec:coll_rates} we develop the formalism for calculating collision rates in the Galactic center.
 We utilize our caleulations of the mass loss per collision. and the collision rate as a function of Galactic racius to find the radial profile of the mass loss rate in 4..," We utilize our calculations of the mass loss per collision, and the collision rate as a function of Galactic radius to find the radial profile of the mass loss rate in \ref{sec:mass_loss_rates}."
 Since the amount of mass lost is dependent on the masses of the colliding stars. the mass loss rate in the Galactic centre is sensitive to the underlying PDME.," Since the amount of mass lost is dependent on the masses of the colliding stars, the mass loss rate in the Galactic centre is sensitive to the underlying PDMF."
 By comparing our calculations to mass loss rates obtained from the x-ray luminosity measured byChandra. in 5 we constrain the PDAIF of the Galactic centre.," By comparing our calculations to mass loss rates obtained from the x-ray luminosity measured by, in \ref{sec:constrain} we constrain the PDMF of the Galactic centre."
 We derive analytic solutions of the PDALF as a function of an adopted LME for cillerent star formation scenarios. which allows us to place constraints on the IME in 86.," We derive analytic solutions of the PDMF as a function of an adopted IMF for different star formation scenarios, which allows us to place constraints on the IMF in 6."
 In 7. we estimate the contribution to the mass loss rate from collisions involving red giant. (RC) stars.," In 7, we estimate the contribution to the mass loss rate from collisions involving red giant (RG) stars."
" ""Throughout this paper we refer to the star that loses material as the perturbed star. ancl the star that causes material to be lost as the perturber star."," Throughout this paper we refer to the star that loses material as the perturbed star, and the star that causes material to be lost as the perturber star."
 Quantities with the subscript or superscript spd” or “pro refer to the perturbed star and. perturber star respectively., Quantities with the subscript or superscript “pd” or “pr” refer to the perturbed star and perturber star respectively.
 We work in units where mass is measured in the mass of the perturbed star. My. distance in the radius of the perturbed star. rpg. velocity in the escape velocity of the perturbed star. pnt (=Y2OM fri). and time in ryíprn," We work in units where mass is measured in the mass of the perturbed star, $M_{pd}$ , distance in the radius of the perturbed star, $r_{pd}$ , velocity in the escape velocity of the perturbed star, $v_{esc}^{pd}$ $= \sqrt{2GM_{pd}/r_{pd}}$ ), and time in $r_{pd}/v_{esc}^{pd}$."
" We denote normalization by these quantities Cor the appropriate combination of these quantities) with a tilde: We refer to collisions in which brj;|ry as ""indirect? collisions. and. collisions in which bxrir|]rp as μασ collisions."," We denote normalization by these quantities (or the appropriate combination of these quantities) with a tilde: We refer to collisions in which $b > r_{pd}+r_{pr}$ as “indirect” collisions, and collisions in which $b\leq r_{pd}+r_{pr}$ as “direct” collisions."
 The impact parameter. b. is the distance ofclosest approach measured from the centers of both stars.," The impact parameter, b, is the distance of closest approach measured from the centers of both stars."
 We consider the condition for mass loss at à position. P. within the perturbed star to be that the Kick. velocity due to the encounter at 7 exceeds the escape velocity of the perturber star at P. 2;NP(P)PG).," We consider the condition for mass loss at a position, $\tilde r$ , within the perturbed star to be that the kick velocity due to the encounter at $\tilde r$ exceeds the escape velocity of the perturber star at $\tilde r$, $\Delta \tilde v (\tilde r) \geq \tilde{v}_{esc}(\tilde r)$."
 The escape velocity as à function of position within the perturbed star can be found from the initial kinetic and potential energies of a test particle al position r. where Adj;~ is. the mass interior+ at positionme P?=f and po is the density. profile of the star.," The escape velocity as a function of position within the perturbed star can be found from the initial kinetic and potential energies of a test particle at position $\tilde r$, where $\tilde M_{int}$ is the mass interior at position $\tilde {r}^{\prime}$ and $\tilde \rho$ is the density profile of the star."
 ‘To calculate the mass lost due to an indirect collision. we first calculate the kick velocity given to the perturbed star as à function of position within the star.," To calculate the mass lost due to an indirect collision, we first calculate the kick velocity given to the perturbed star as a function of position within the star."
 We work under the impulse approximation (Spitzer1958).. valid under the condition that the encounter time is much shorter than the characteristic crossing time of a constituent of the perturbed svsteor.," We work under the impulse approximation \citep{spitzer:1958}, valid under the condition that the encounter time is much shorter than the characteristic crossing time of a constituent of the perturbed system."
" Given a mass distribution for the perturbed system. ppd and a potential for the perturber system. 9. the kick velocity alter an encounter under the impulse approximation is given by Binney&Tremaine(2008):: Equation (3)) can be simplified by expanding the eracdient of the potential in à Taylor series. resulting in The expansion is. valic uncer the ""distant tide"" ⋜↧↓≻↓≻↓⋅∢⋟⇀∖↕↓↕↓⋜↧⇂↕⋖≱↓↥∖∖⋎↓↕⊲⊔↛↓↥⋠↓≱∖≱∖⋜⊔⊀↓≱∖∐⋖⋅∠⇂∖"," Given a mass distribution for the perturbed system, $\rho_{pd}$ and a potential for the perturber system, $\Phi$, the kick velocity after an encounter under the impulse approximation is given by \citet{binney:2008}: Equation \ref{eqn:dv_impulse}) ) can be simplified by expanding the gradient of the potential in a Taylor series, resulting in The expansion is valid under the “distant tide” approximation which is satisfied when $r_{pd} \ll b$."
∖⋎↓∐⊾⊔∣⋮↗⊔∣∣∖≮↾∶∶∣↗⊳↾↓∖↓↥⋖⋅ ↓≻⋜⊔⋅⋜⋯↓⋖⊾↿∢⊾↓⋅∣⋎⊓∣↕⊳∖↿↓↥⋖⊾↓⋅⋖⋅↓⋜∐⊲↓∖⇁⋖⋅≱∖↓≻⋖⋅⋯⇂∣⋡⋖⋅↿∖∖⊽⋖⋅⋖⋅⊔↿∐⋖⋅⊳∖↿⋜⊔⋅⊳∖ ∣↴↗∣∣∪⋡, The parameter $v_{rel}$ is the relativespeed betweenthe stars $v_{rel} \equiv |\vec{v}_{pd}-\vec{v}_{pr}|$ ).
∖∖⊽∢⋅⋜⊔⋅⋖⋅⊲↓⊔∩⊾↓⋅⋖⋅⊳∖↿⋖⋅∠⇂∐↕↿↓↥⋖⋅⊔↓⋜↧⋏∙≟⊔⊀↓↿⋯⇂∢⊾ of equation. (4)). which when normalized to the units that we have adopted for this paper is," Weare interested in the magnitude of equation \ref{eqn:dv_tide}) ), which when normalized to the units that we have adopted for this paper is"
excecding 70 Ik. Such Πιν mean that uou-detectious result in coustraimts on dust mass du these systems of typically 200 times the dust content of the solar svsteun. but down to ouly approximately four times this for OUL nearest tarect. GJ 699. at Las pe.,"exceeding 70 K. Such limits mean that non-detections result in constraints on dust mass in these systems of typically 200 times the dust content of the solar system, but down to only approximately four times this for our nearest target, GJ 699, at 1.8 pc."
 Iu additiou to the 500 primary targets of the survey. the laree FOV of SCUDA-2 will enable us to image conanions nui nuItiple systems. aud in some fields. to include other nearyw stars where they are preseut.," In addition to the 500 primary targets of the survey, the large FOV of SCUBA-2 will enable us to image companions in multiple systems, and in some fields, to include other nearby stars where they are present."
 IDu our target list. at least 210 stars are members of muliple systems of at least two members.," In our target list, at least 240 stars are members of multiple systems of at least two members."
 This lieh fracHon of multiple systems raises the total number of stars surveved to well over 750., This high fraction of multiple systems raises the total number of stars surveyed to well over 750.
" Very few of these svsOnus require more than oue SCUDA-2 sub-aurav. since evel wide binaries up o 1000 AU iu separation will have an aueular separation of«2""."," Very few of these systems require more than one SCUBA-2 sub-array, since even wide binaries up to 1000 AU in separation will have an angular separation of$< 2$."
 Di the case of biuaries. hei iuclusion in the survey is based on the spectral type ofthe primary: exeiuptiue the secondary frou the initial sauple eusures that no bias is introduced due to the probabiities of disks in binary svstenis or the weighting of secoudaries toward lower masses.," In the case of binaries, their inclusion in the survey is based on the spectral type of the primary; exempting the secondary from the initial sample ensures that no bias is introduced due to the probabilities of disks in binary systems or the weighting of secondaries toward lower masses."
 The sample-size for each of the spectral types was chosen so that we could statistically distinguish between detection rates of 5. 10. 25 aud when the primary dataset is divided iuto tlie| following sub-categories.," The sample-size for each of the spectral types was chosen so that we could statistically distinguish between detection rates of 5, 10, 25 and when the primary dataset is divided into the following sub-categories."
 The overall choice of sample size is driven by the aimi of differcutiating detecion rates among the listed sub-eroups., The overall choice of sample size is driven by the aim of differentiating detection rates among the listed sub-groups.
 The kev is to ive a sample size large chough that the sinallest. sub-group wih the lowest detection rate can be distinguished from rates in other sub-eroups., The key is to have a sample size large enough that the smallest sub-group with the lowest detection rate can be distinguished from rates in other sub-groups.
 The detection rate for X disks amoug Y stars is defined as Z=(NX|vX)/Y as thei ncertaiutyv is Polssouiau for small counts., The detection rate for X disks among Y stars is defined as $Z = (X +/- \sqrt{X}) / Y$ as the uncertainty is Poissonian for small counts.
 In the simplest iuterpretation. woe accept rates as cdifferime if the upper bound of one Z value is below the lower bound of ti6 conrparison Z value.," In the simplest interpretation, we accept rates as differing if the upper bound of one Z value is below the lower bound of the comparison Z value."
 Iu mractice this meaus the survey size is driven by the sanallest samples. that are one-euth of the conparisonu eroup.," In practice this means the survey size is driven by the smallest samples, that are one-tenth of the comparison group."
 The two cases are the planetary svstenis (— liu 10 Solaa-type svstenis) aud voung Suu-like stars (20 among 200 of the € (ond In stars)., The two cases are the planetary systems $\sim$ 1 in 10 Solar-type systems) and young Sun-like stars (20 among 200 of the G and K stars).
 Tere we cau distinguish observed detHon rates Z for any pair of intrinsic detection rates auoug Cad.," Here we can distinguish observed detection rates $Z$ for any pair of intrinsic detection rates among, and."
. The formal nininimn requie to do so is 163 stars. which is just met bv our snallest eroup: the approximately Ls y.ars of late-F o 1uid-K 4»o»teutial planet-hosts.," The formal minimum required to do so is 163 stars, which is just met by our smallest group: the approximately 180 stars of late-F to mid-K potential planet-hosts."
 Detection rates do no need to be compared withicere very spectral type., Detection rates do not need to be compared within every spectral type.
 For exiuuple. the effects of age ο ebris will be investigated for stars of Sun-like mass. iec.. types both € and Ix. rather than € aud Ix separately.," For example, the effects of age on debris will be investigated for stars of Sun-like mass, i.e., types both G and K, rather than G and K separately."
 I- fact. this approach is recpred to allow us to establiscor robust distinctions (1.6.. at the 2-3 o level) betwee etection rates as a function of spectral type aud. age. ue to the practical luits on saniple size.," In fact, this approach is required to allow us to establish robust distinctions (i.e., at the 2-3 $\sigma$ level) between detection rates as a function of spectral type and age, due to the practical limits on sample size."
 For instance. with 200 stars per smuupk a difference in detection rate can be identified at the 3 σ level.," For instance, with 200 stars per sample, a difference in detection rate can be identified at the 3 $\sigma$ level."
 Results from the far-IR surveys withJRAS.790 aud Opitzer show that these rates actually occur amongst," Results from the far-IR surveys with, and show that these rates actually occur amongst"
"anele: where ciffereutial distances iu the fIuid [rame are fouud by projection onto a local tetracl. i.e.. daUn=οLy,","angle: where differential distances in the fluid frame are found by projection onto a local tetrad, i.e., $dx^{(\mu)} = \hat e^{(\mu)}_{\nu} dx^\nu$."
 The factor of 2 enters because the disk has both a top and a bottom surface., The factor of 2 enters because the disk has both a top and a bottom surface.
 This surface brightuess is the closest approximation we cau readily make to a vertical integration through a thin disk., This surface brightness is the closest approximation we can readily make to a vertical integration through a thin disk.
 Because ThinHR has au aspect ratio of only 20.06. there is very little dillerence between a polar angle and a vertical integration.," Because ThinHR has an aspect ratio of only $\simeq 0.06$, there is very little difference between a polar angle and a vertical integration."
 The effective temperature is then where o is the Stefau-Boltzimaun constant., The effective temperature is then where $\sigma$ is the Stefan-Boltzmann constant.
 The shape of the integrated spectrum at infinity is uniquely determined by the simulation data. but the dimeusionless code data specify neither the units of photon energy nor of luminosity.," The shape of the integrated spectrum at infinity is uniquely determined by the simulation data, but the dimensionless code data specify neither the units of photon energy nor of luminosity."
 Attaching physical units to the results demands choosing two parameters., Attaching physical units to the results demands choosing two parameters.
 A particular choice of central mass À determines the units of leugth aud time: a particular value of the accretion rate AZ determines the unit of mass in the flukl., A particular choice of central mass $M$ determines the units of length and time; a particular value of the accretion rate $\dot M$ determines the unit of mass in the fluid.
 Between these two. both the unit of photon energy aud the unit of luminosity cau be determined.," Between these two, both the unit of photon energy and the unit of luminosity can be determined."
" The characteristic photon energy. (or temperature) "" ↜∖∢∙⋜≹↥↩↠∖∣↖∪∫↙∟⊔−⋝↓↓⋜↕⊳∖≺↵⊸∖↥↽≻≺↵∢∙↕↩≺⊔↕⋅∩∐⊔∙∩∖⇁≺↵↥∐↥∩∐⋜↕↥≺∐⊳∖↕⊆↕∐≺↵∩↕⋅⊽∖⊽⋅⊺∐", The characteristic photon energy (or temperature) scale is $\propto (\dot M/M^2)^{1/4}$ as expected from conventional disk theory.
≺↵⊔∐≺↵∑≟↓⋅⋜↕↕≺↲≺⇂≺↵∐∐⊳∖⊳∖↥∖⊽∐⊽∖⇁↥⊳∖⋅∩↥i> ⋅ ⋅ ⋅ ⋅ ⋅ ⋅ ⋅ ∢∙⋯⊔⋅⊳∖≺↵⋅∣↖⋀⊔⋅⊺≺," The integrated emissivity is, of course, $\propto \dot M$."
↵∢∙∐∐∎∢∙⋜↕↥≺⇂≺↵↕⋜↕∐⊳∖↕⋅≺↵∑≟⋜∐⋅≺∐∐∑≟↕∐≺↵⋃⋅⋜↕∐⊳∖↥⋜↕∏∩∐∣⋈↵↕∖∖↽≺↵≺↵∐∢∙∩⊓↲⋃∐∐⊳∖⋜↕∐≺⇂↥↽≻∐⊽∖⊽⊳∖↥∢∙⋜↕↥⋃∐∐⊳∖∢∙⋜↕∐ ∣⋈↲↥∎⋯∐≺⊔, Technical details regarding the translation between code units and physical units can be found in the Appendix.
∐↕∐≺↵⇀−∖↥↽≻↥↽≻≺↲∐≺∐⊸∖⋅ This approximate solution to the radiation trausfer problem is. iu some seuse. “truest” to the johwsies of the simulation.," This approximate solution to the radiation transfer problem is, in some sense, “truest"" to the physics of the simulation."
 Even thoughe we assume a thermal spectrum. we permit each cell to raciate isotropically in its (rame. aud the photous then travel without Πιονασος to iufinity.," Even though we assume a thermal spectrum, we permit each cell to radiate isotropically in its frame, and the photons then travel without hindrance to infinity."
 To compute the bolometric luminosity received at infinity. we use tlie cooling Buuction snapshots as input to the general relativistic ray-tracing engine described in ?..," To compute the bolometric luminosity received at infinity, we use the cooling function snapshots as input to the general relativistic ray-tracing engine described in \cite{NKH09}."
 We trace rays [rom au array ol observers at infinity at diflfereut polar aueles to the sinmulation volume aid then integrate over solid angle in order to determine the total luminosity., We trace rays from an array of observers at infinity at different polar angles to the simulation volume and then integrate over solid angle in order to determine the total luminosity.
 Photon geodesics are defined by aud the intensity arriving at infinity comes from integrating the invariant transfer equation, Photon geodesics are defined by and the intensity arriving at infinity comes from integrating the invariant transfer equation
similar to the effects of rotation noted by ?.. anc may not be useful for disentangling the effects of rotation and overshoot. ?.. ?.. 6)),"similar to the effects of rotation noted by \citet{me09}, and may not be useful for disentangling the effects of rotation and overshoot. \citet{briquet07}, \citet{briquet07}, \ref{HH})"
similar to the effects of rotation noted by ?.. anc may not be useful for disentangling the effects of rotation and overshoot. ?.. ?.. 6)).,"similar to the effects of rotation noted by \citet{me09}, and may not be useful for disentangling the effects of rotation and overshoot. \citet{briquet07}, \citet{briquet07}, \ref{HH})"
structure can cause additional structure iu P.,structure can cause additional structure in $P$.
 However. we have shown that this effect cannot be very inportaut iu these observations (Iaverkoru ct 220050).," However, we have shown that this effect cannot be very important in these observations (Haverkorn et 2003c)."
 The Westerbork Northern Sky Survey (WENSS. Reneclink et 11997) is a ow-frequencey radio survev that covers the whole «kv north of 6=30° aat 325 MIIz to a limiting flux density of approximately 18 wy (56) aud with a resolution of cosec à.," The Westerbork Northern Sky Survey (WENSS, Rengelink et 1997) is a low-frequency radio survey that covers the whole sky north of $\delta=30$ at 325 MHz to a limiting flux density of approximately 18 mJy $\sigma$ ) and with a resolution of $\times$ cosec $\delta$."
 Polarization data taken during the dav are ercatly affected lov solar radiation. which is detected iu sidelobes.," Polarization data taken during the day are greatly affected by solar radiation, which is detected in sidelobes."
 Furthermore. jonospheric Faraday rotation rapidly changes during sunrise and sunset. causing a reduction of the appareut polarized intensity. and adding much noise.," Furthermore, ionospheric Faraday rotation rapidly changes during sunrise and sunset, causing a reduction of the apparent polarized intensity, and adding much noise."
 Tlowever. tn mosaics observed Giiuost) chtively during nieht tine. the polarization data are of eood quality.," However, in mosaics observed (almost) entirely during night time, the polarization data are of good quality."
" As a result. we could nake a ""supernmosaic of a large region of 307«35 “which we will refer to as the WENSS polarization region."," As a result, we could make a “supermosaic” of a large region of $\sim 30\dg\times35\dg$, which we will refer to as the WENSS polarization region."
 Details on the observations and correction for ionospheric Faraday rotation are given in Schuitzeler et (n prep)., Details on the observations and correction for ionospheric Faraday rotation are given in Schnitzeler et (in prep).
 Fie., Fig.
 9. shows polarized intensity P in the WENSS polarization region. resampled in Calactic longitude aud latitude.," \ref{f9:wenss} shows polarized intensity $P$ in the WENSS polarization region, resampled in Galactic longitude and latitude."
 For this analysis. a Gaussian taper with value 0.25 at a baseline of 500m. was applied. which viclds a resolution of ~2.5'.," For this analysis, a Gaussian taper with value 0.25 at a baseline of 500m was applied, which yields a resolution of $\sim$."
 Iun the figure. P saturates at 35 aidvbeam. which coincides with a polarized brightuess temperature of ~ LF I to ~ 25 Ix. depending on declination.," In the figure, $P$ saturates at 35 mJy/beam, which coincides with a polarized brightness temperature of $\sim$ 17 K to $\sim$ 25 K, depending on declination."
 The average Pox2.6 mJv/beuu., The average $P \approx 2.6$ mJy/beam.
 Note that these observatious are taken in a single 4 MIIZ wide frequency band. so rotation measure data are nof available.," Note that these observations are taken in a single 5 MHz wide frequency band, so rotation measure data are not available."
 To quautify the structure iu the polarization maps. we calculate angular power spectra ο) as a function of nimltipole (.," To quantify the structure in the polarization maps, we calculate angular power spectra $PS (\ell)$ as a function of multipole $\ell$."
 Α iuultipole 6 is a measure of angular scales equivalent to wave umber. aud is defined as {51807ή. where 0 is the angular scale in degrees.," A multipole $\ell$ is a measure of angular scales equivalent to wave number, and is defined as $\ell \approx 180\dg /\theta$, where $\theta$ is the angular scale in degrees."
 The angular power spectrin P5 of a radiation field IX is the square of the Fourier transform of X: Γον=[F(N)P. where F enotes the Fouricr transforma.," The angular power spectrum $PS$ of a radiation field $X$ is the square of the Fourier transform of $X$: $PS_X (\ell) = |{\cal F}(X)|^2$, where ${\cal F}$ denotes the Fourier transform."
 Were. the observable (X ‘an be either Stokes Q. Stokes C. polarized intensity P or rotation mcasire RAL.," Here, the observable $X$ can be either Stokes $Q$, Stokes $U$ , polarized intensity $P$ or rotation measure $RM$."
 The power spectra were computed in two dimensions. and averaged over azimuth im radial iius.," The power spectra were computed in two dimensions, and averaged over azimuth in radial bins."
 The miultipole spectra iudex o. defined as ( “lis caleulated from a log-log fit to the power spectrum.," The multipole spectral index $\alpha$, defined as $PS_X (\ell) \propto \ell^{-\alpha}$ , is calculated from a log-log fit to the power spectrum."
 I1 the tapered data. multipoles with higher values of { are affected by the tapering. and in the uutapered data lugher multipoles are dominatedby noise.," In the tapered data, multipoles with higher values of $\ell$ are affected by the tapering, and in the untapered data higher multipoles are dominatedby noise."
 Multipoles with, Multipoles with
show the slopes in. aand aagainst Av mag for 19 GC's studied. by. (2009).. assuming following relations. We also examined the variations in. aand aagainst cllective temperatures of RD stars in individual GCs.,"show the slopes in and against $K$ mag for 19 GCs studied by \citet{vii}, assuming following relations, We also examined the variations in and against effective temperatures of RGB stars in individual GCs."
 Substantial slopes in the plots of. vversus effective temperature can be seen in many GC's. in sharp contrast to what Carrettactal.(2009). claimed (see their Figure 5).," Substantial slopes in the plots of versus effective temperature can be seen in many GCs, in sharp contrast to what \citet{vii} claimed (see their Figure 5)."
 Lt is bevond the scope of our work to explain such discrepancy between aand iin Carrettactal., It is beyond the scope of our work to explain such discrepancy between and in \citet{vii}.
(2009).. However. as a first approximation. the gradient in. wwith respect to nunayv indicate that their adopted: surface gravitics are systematically in error.," However, as a first approximation, the gradient in with respect to may indicate that their adopted surface gravities are systematically in error."
 For example. surface gravities of UGD stars in NGC 4590 could be systematically larger with increasing A mae.," For example, surface gravities of RGB stars in NGC 4590 could be systematically larger with increasing $K$ mag."
 Phe mocilication of surface gravity also ends to alfect. other input parameters. such as effective emperature and. turbulent: velocity.," The modification of surface gravity also tends to affect other input parameters, such as effective temperature and turbulent velocity."
 Therefore. the results oesented by Carrettaetal.(2009) may not be as accurate as they claimed withapparentfy very small measurement errors.," Therefore, the results presented by \citet{vii} may not be as accurate as they claimed with very small measurement errors."
 Furthermore. their. final results may not. be as 10mogeneous as they claimed. since the cliscrepaney between aand vvaries one cluster to the other and the break in the slope between vversus A mag or effective temperature can be seen in some GCs.," Furthermore, their final results may not be as homogeneous as they claimed, since the discrepancy between and varies one cluster to the other and the break in the slope between versus $K$ mag or effective temperature can be seen in some GCs."
 We note that if this discrepaney is entirely due to incorrect surface gravity and other parameters adopted. by Carrettaetal.(2009)... their derivations of other elemental abundances. such as oxveen and sodium. may not be reliable.," We note that if this discrepancy is entirely due to incorrect surface gravity and other parameters adopted by \citet{vii}, their derivations of other elemental abundances, such as oxygen and sodium, may not be reliable."
 Also importantly. if the limited number of available Fe HE lines is responsible for this discrepancy (seeCarrettaetal. 2006).. their O/Fe] and Na/Ee] ratios would. be in the same situation since only few lines are available for both elements.," Also importantly, if the limited number of available Fe II lines is responsible for this discrepancy \citep[see][]{n2808}, their [O/Fe] and [Na/Fe] ratios would be in the same situation since only few lines are available for both elements."
 We await their careful re-analysis of their data with keen interest., We await their careful re-analysis of their data with keen interest.
 As we mentioned above. Carrettaetal.(2009). derived OfFe] ratios based on Fe Land a substantial slope in the vversus Amag may imply that their O/Fe] ratios could be in error.," As we mentioned above, \citet{vii} derived [O/Fe] ratios based on Fe I and a substantial slope in the versus $K$mag may imply that their [O/Fe] ratios could be in error."
 ln Table 1.. eight GCs (NGC 6838. NGC 6121. NGC 288. NGC 6218. NGC 5904. NGC 3201. NGC 6809 and NGC 7078) have negligibly small gradients in their. aand vversus. A mag.," In Table \ref{tab:slope}, eight GCs (NGC 6838, NGC 6121, NGC 288, NGC 6218, NGC 5904, NGC 3201, NGC 6809 and NGC 7078) have negligibly small gradients in their and versus $K$ mag."
 Therefore. it is expected. that pofendiat systematic errors in surface gravity adopted by al.(2009) do not signilicantly allect differential rratios against A mag in these GC's.," Therefore, it is expected that systematic errors in surface gravity adopted by \citet{vii} do not significantly affect ratios against $K$ mag in these GCs."
 ln Figure 2.. we show oxvgen and sodium distributions against fyην For RGB stars in NGC 3201 and NGC62187.," In Figure \ref{fig:n3201}, we show oxygen and sodium distributions against $K-K_{bump}$ for RGB stars in NGC 3201 and NGC."
. As shown in the figure. differences in aancd delistributions between the upper ROB and the lower RGB stars arenotable?.," As shown in the figure, differences in and distributions between the upper RGB and the lower RGB stars are."
. The mean rratio for the lower RGB stars is larger than that for upper RGB stars in NGC 3201 and. NGC 0918. while the mean rratio for the lower ROB stars is smaller than that. for upper RGB stars.," The mean ratio for the lower RGB stars is larger than that for upper RGB stars in NGC 3201 and NGC 6218, while the mean ratio for the lower RGB stars is smaller than that for upper RGB stars."
 LE the results by Carrettaetal.(2009) are not greatly in error. are these variations in O/Fe] and ΝαΡΟ) ratios between the lower and the upper RGB stars a signature of the internal deep mixing episode. so that oxygen is depleted by the CNO evele while sodium isenrichedfrom the Νορ. ο] Να reaction in upper RGB stars of NGC 3201 and NGC 62187," If the results by \citet{vii} are not greatly in error, are these variations in [O/Fe] and [Na/Fe] ratios between the lower and the upper RGB stars a signature of the internal deep mixing episode, so that oxygen is depleted by the CNO cycle while sodium isenrichedfrom the $^{22}$ $p$ , $\gamma$ $^{23}$ Na reaction in upper RGB stars of NGC 3201 and NGC 6218?"
2006).,.
". In particular, the blue continuum arising from unobscured type-1 quasars could potentially bias the BM/BX selection."," In particular, the blue continuum arising from unobscured type-1 quasars could potentially bias the BM/BX selection."
" However the majority of the most IR-luminous AGNs in COSMOS have also been detected in the X-rays (Hasingeretal.,2007;Brusa2010)."," However the majority of the most IR-luminous AGNs in COSMOS have also been detected in the X-rays \citep{Hasinger:07,Brusa:10}."
. They have been removed from our initial sample and therefore they should not influence the various trends that we find in reffig:LF.., They have been removed from our initial sample and therefore they should not influence the various trends that we find in \\ref{fig:LF}.
" Contrary to the behavior characterizing the three frame UV/optical selections, the fraction of Optically-Faint IR-bright galaxies is clearly rising with IR luminosity."," Contrary to the behavior characterizing the three rest-frame UV/optical selections, the fraction of Optically-Faint IR-bright galaxies is clearly rising with IR luminosity."
" This is consistent with what we had qualitatively inferred from the comparison of their redshift distributions at P54,í,,> 0.08mmJy and Foaym>0.5 mmJy (see 33).", This is consistent with what we had qualitatively inferred from the comparison of their redshift distributions at $F_{24\mu m}>0.08$ mJy and $F_{24\mu m}>0.3$ mJy (see 3).
 In the previous section we argued that these sources are biased toward highly-obscured galaxies., In the previous section we argued that these sources are biased toward highly-obscured galaxies.
" Assuming that dust extinction must correlate with the observed 24j:m/optical flux ratio of galaxies at z~2, it could thus explain the trend that we find with ]uminosity."," Assuming that dust extinction must correlate with the observed $\mu$ m/optical flux ratio of galaxies at $z \sim 2$, it could thus explain the trend that we find with luminosity."
" Furthermore, this trend could partly originate from an increasing contribution of obscured AGNs to the mid-IR luminosities of galaxies as a function of 24m flux (e.g.,Brandetal.,2006;Menéndez-Delmestre 2009)."," Furthermore, this trend could partly originate from an increasing contribution of obscured AGNs to the mid-IR luminosities of galaxies as a function of $\mu$ m flux \citep[e.g.,][]{Brand:06,Menendez:09}."
". This rising contribution would boost their mid-IR emission with respect to their optical luminosity, hence leading to an rt—[24] color satisfying the selection criterion of this population."," This rising contribution would boost their mid-IR emission with respect to their optical luminosity, hence leading to an $r^+-[24]$ color satisfying the selection criterion of this population."
" Even though we did not account for the contribution of galaxies below our 24m flux limit we see on 77 that star-forming galaxies detected in the COSMOS mid-ΤΗ, survey contribute a substantial fraction of the total IR luminosity density at z2, in agreement with previous findings (e.g.,Caputietal2007;Rodighieroetal, 2010).."," Even though we did not account for the contribution of galaxies below our $\mu$ m flux limit we see on 7 that star-forming galaxies detected in the COSMOS mid-IR survey contribute a substantial fraction of the total IR luminosity density at $z \sim 2$, in agreement with previous findings \citep[e.g.,][]{Caputi:07,Rodi:10}."
 It implies that the incompleteness which affects the different color selection techniques previously described can result in non-negligible biases with respect to the global picture of galaxy evolution inferred from complete samples of high-redshift sources., It implies that the incompleteness which affects the different color selection techniques previously described can result in non-negligible biases with respect to the global picture of galaxy evolution inferred from complete samples of high-redshift sources.
 For example we find that the MIPS sources selected with the BM/BX, For example we find that the MIPS sources selected with the BM/BX
the wide variations in propertics flat plague stilies of sinall nunbers of galaxies.,the wide variations in properties that plague studies of small numbers of galaxies.
 This aualvsis makes use of normalised ποτ profiles. which we calculate Εςx each ealaxy based on the elliptical aperure plotometry that was used in the previous section to derive conceltration nidiees.," This analysis makes use of normalised light profiles, which we calculate for each galaxy based on the elliptical aperture photometry that was used in the previous section to derive concentration indices."
 However. we choose not to prescut these in the form of surfacebriglhtuess profiles. 1j uuits of maenitudes per square aresec. but rather use the fluxes in the elliptical annuli. uncorrected for the ixreasing area of the annulus as the apertures grow.," However, we choose not to present these in the form of surface brightness profiles, in units of magnitudes per square arcsec, but rather use the fluxes in the elliptical annuli, uncorrected for the increasing area of the annulus as the apertures grow."
 Thus je area under such a profile within a given range iu radit sis directv proportional to the amount of lieht. contributed to the total hunuimositv of he galaxy., Thus the area under such a profile within a given range in radius is directly proportional to the amount of light contributed to the total luminosity of the galaxy.
 The shape of jo resultiis xofile is mach more intuitively related to conceutratio1 indices than ds a surfacebrightness profile. )ecaise of this lnk between flux and area under the profie.," The shape of the resulting profile is much more intuitively related to concentration indices than is a surface brightness profile, because of this link between flux and area under the profile."
 For exauple. the ¢fective radius can be simply estimate YOU slch a plot as the radius which eveulv divides the area nn ert1e profile.," For example, the effective radius can be simply estimated from such a plot as the radius which evenly divides the area under the profile."
" It ls. LOCOSSALY to nornalise Much p»files )efore colmbining them iuto. for exaale, a lea1 preiile for all galaxies of a given type."," It is necessary to normalise such profiles before combining them into, for example, a mean profile for all galaxies of a given type."
 A siuiple averwe without normalisation will be doni iate]x the xielest σαANIC: anc if a racial scale in kpe is adoped. theu all galaxies will contribute to the centre of the nean profile. but oulv t10 largestaveest galaxiesealaxies fto 15 0ouer reelons. eiviug a disorted result that is hard to interpret.," A simple average without normalisation will be dominated by the brightest galaxies; and if a radial scale in kpc is adopted, then all galaxies will contribute to the centre of the mean profile, but only the largest galaxies to the outer regions, giving a distorted result that is hard to interpret."
 Thus we normalise the individual profiles obtaiue« from cach galaxy in two ware., Thus we normalise the individual profiles obtained from each galaxy in two ways.
 Firstly. the radial x‘ale is expressed iu units of the τοι isophotal ποπ]λαjeDi axis of the galaxy.," Firstly, the radial scale is expressed in units of the $r_{24}$ isophotal semi-major axis of the galaxy."
 Secolh. t1ο area under each profie ds normalised to uuity. eiviug cach ealaxy. bright or faint. equal weighting in the 10021 profie.," Secondly, the area under each profile is normalised to unity, giving each galaxy, bright or faint, equal weighting in the mean profile."
" This process was appied first to the R- bandduaee of cach ealaxy. aud then to the liuage using exactly the same positiois, shapes aud size of apertures. and scaling again by the R-baud isophotal size."," This process was applied first to the $R$ -band image of each galaxy, and then to the image using exactly the same positions, shapes and size of apertures, and scaling again by the $R$ -band isophotal size."
 AS a resiIt of the normalisation. the profiles have uo calibration in flux or magnitude units on the vertical axis: the profiles are simple shape functions. iudicaine the fraction of lel: residing iu a given racial range.," As a result of the normalisation, the profiles have no calibration in flux or magnitude units on the vertical axis; the profiles are simple shape functions, indicating the fraction of light residing in a given radial range."
" Iu he remainder of this section and in tho next. we look at the mean profiles produced bv this WOCCSS, aud examine tιο depeudeuce of profie shape on ealaxy ype. aud ou fae presence or absence of |ALS,"," In the remainder of this section and in the next, we look at the mean profiles produced by this process, and examine the dependence of profile shape on galaxy type, and on the presence or absence of bars."
" The mean xofiles or the individual T-tvpes couain betwee ιδ and GT profiCR, as listed in Table which also subdivides he ΠΡΟΣ for cach type according to bar classification."," The mean profiles for the individual $T$ -types contain between 8 and 67 profiles, as listed in Table \ref{tbl:ngal}, which also subdivides the numbers for each type according to bar classification."
 Some galaxies from the GS sample were omitted from his analvsis due to contamination by foreground stars or xoblenmis with varving backgrouud levels across naages., Some galaxies from the GS sample were omitted from this analysis due to contamination by foreground stars or problems with varying background levels across images.
 ence tje total ummber of galaxies listed in Table ls 313. out of the full GS sample of 327.," Hence the total number of galaxies listed in Table \ref{tbl:ngal}
 is 313, out of the full GS sample of 327."
 The mean profiles for allgalaxies as à function of ITubble T-type are shown in Fig. 3..," The mean profiles for allgalaxies as a function of Hubble $T$ -type are shown in Fig. \ref{fig:meanprof},"
 with the pprofiles ou the left and the A- baud ]xofiles ou the right., with the profiles on the left and the $R$ -band profiles on the right.
 The unuber of galaxies contributiug» to cach profile is indicated: it is imporaut to note that all 313 profiles of sufficieut cosmetic quality have been icluded. so there las been no exclusion of outhers with wlsual profile shapes.," The number of galaxies contributing to each profile is indicated; it is important to note that all 313 profiles of sufficient cosmetic quality have been included, so there has been no exclusion of outliers with unusual profile shapes."
 Dor] vars nudicate f1ο standard error of the mdividua scaled. profiles about he mean at eacli radial point., Error bars indicate the standard error of the individual scaled profiles about the mean at each radial point.
 The errors on the pprofiles are significantly lareer than those on the 2-1 MC profiles. originating from the clumpicr nature of the COLUSSION.," The errors on the profiles are significantly larger than those on the $R$ -band profiles, originating from the clumpier nature of the emission."
 The main point to note from the R-baud profiles iu Fie., The main point to note from the $R$ -band profiles in Fig.
 3 is the degree of correlation between mean profile shape aud Uhibble type. swuch is quite eratifvius eiven he subjective nature of galaxy classification.," \ref{fig:meanprof} is the degree of correlation between mean profile shape and Hubble type, which is quite gratifying given the subjective nature of galaxy classification."
" The mean wohes of the earliest. types (f-types 0 and 1) are very ceutyally concentrated. with the pcvas of the profiles. aud hence the largest couvibution to the total galaxy iTuninositv. coniug from ~O.1 ro,"," The mean profiles of the earliest types $T$ -types 0 and 1) are very centrally concentrated, with the peak of the profiles, and hence the largest contribution to the total galaxy luminosity, coming from $\sim$ 0.1 $r_{24}$."
 The one anomaly in an otherwise smooth sequence of A-baud profiles is for he Sh eaANICS = (T3): the addition:d liebt in this mean wohe close to the centre aud at 0.1 +οἱ will be discussed in Sect. [.., The one anomaly in an otherwise smooth sequence of $R$ -band profiles is for the Sb galaxies $T=$ 3); the additional light in this mean profile close to the centre and at 0.4 $r_{24}$ will be discussed in Sect. \ref{sec:bars}.
 For later types. the cussion svstcmatically HOVCS ouwards and broadens. with the peak sitting at 0. Poy or the S11 and hu ealaxies (P-tvpes 9 anc 10).," For later types, the emission systematically moves outwards and broadens, with the peak sitting at $\sim$ 0.4 $r_{24}$ for the Sm and Im galaxies $T$ -types 9 and 10)."
" This sincth progression in profile properties is shown in ""Table 2..which lists the peak aud effective radii for the Π-]uk Laxd luniea1 profiles for cach P-type. iu units of ro, (tle mean value of ro, for cach type is also eive iin kpc in the final columi. to enable approximate conversion of these values to physical units. although it should be noted that there is a large range of galaxy sizes presen at cach type)."," This smooth progression in profile properties is shown in Table \ref{tbl:reff},which lists the peak and effective radii for the $R$ -band and mean profiles for each $T$ -type, in units of $r_{24}$ (the mean value of $r_{24}$ for each type is also given in kpc in the final column, to enable approximate conversion of these values to physical units, although it should be noted that there is a large range of galaxy sizes present at each type)."
 The siue cata are plotted in Fie. L. , The same data are plotted in Fig. \ref{fig:reff}. .
Sizes in kpe are calculated using au effective Wahldle coustaut of 75 hans + |. corrected for Virgo iufall as explainec in Paper I., Sizes in kpc are calculated using an effective Hubble constant of 75 km $^{-1}$ $^{-1}$ corrected for Virgo infall as explained in Paper I.
somewhat naturally to the hypothesis that the souree of the accreted heavy elements was the interstellar medium (LSAL). encountered in sullicient. density a few to several times per Galactic orbit (on average) for the Solar neighborhoo: (Dupuisetal.1993a.b.1992).,"somewhat naturally to the hypothesis that the source of the accreted heavy elements was the interstellar medium (ISM), encountered in sufficient density a few to several times per Galactic orbit (on average) for the Solar neighborhood \citep{dup93a,dup93b,dup92}."
. Several factors have continually challenged: the ISA pollution scenario., Several factors have continually challenged the ISM pollution scenario.
 First. the typical dearth of hydrogen relative to calcium in DZ stars argues for the accretion of volatile depleted: material (Dufouretal.2007:WollEet.2002:Sionetal. 19902).," First, the typical dearth of hydrogen relative to calcium in DZ stars argues for the accretion of volatile depleted material \citep{duf07,wol02,sio90a}."
. Second. the firm establishment of the hyerogen-rich DAZ spectral class. with heavy elemen dilfusion timescales as short as a few davs. and Galactic positions far from known interstellar clouds: (Ixoesterοal.2005a:Zuckermanet 2003).," Second, the firm establishment of the hydrogen-rich DAZ spectral class, with heavy element diffusion timescales as short as a few days, and Galactic positions far from known interstellar clouds \citep{koe05a,zuc03}."
. Third. the increasing number of metal-contaminated white cdwarks with infrared excesses and closely orbiting rings of dusty ancl gaseous debris (μαιetal.2009:Jura20092:Cansickeοἱal. 2008).," Third, the increasing number of metal-contaminated white dwarfs with infrared excesses and closely orbiting rings of dusty and gaseous debris \citep{far09a,jur09a,gan08}."
. Thus. while there is firm observational evidence of pollution via circumstellar material. as vet there is none avoring the ISM.," Thus, while there is firm observational evidence of pollution via circumstellar material, as yet there is none favoring the ISM."
 Aannestadctal.(1993). provided an extensive spatial. kinematical. and caleium abundance analysis against which hey tested likely scenarios of interstellar accretion based on he best available data at that time. concluding interstellar acecretion could not explain the abundances in.avy of their 15 DZ white cwarls.," \citet{aan93} provided an extensive spatial, kinematical, and calcium abundance analysis against which they tested likely scenarios of interstellar accretion based on the best available data at that time, concluding interstellar accretion could not explain the abundances in of their 15 DZ white dwarfs."
 However. their landmark study. was imited to targets within SOppe for the most part. and to stars that were proper-motion selected.ancl hence biased oward higher space velocities (Sionctal.LOSS)..," However, their landmark study was limited to targets within pc for the most part, and to stars that were proper-motion selected,and hence biased toward higher space velocities \citep{sio88}."
. Ixilic&tedfield(2007)— examined: possible correlations between he current positions of 35 DAZ white chwarfs ancl models of the surrounding ISAT as inferred. from nearby. column density measurements. finding a lack of sullicicnt material o account for most. polluted stars.," \citet{kil07} examined possible correlations between the current positions of 35 DAZ white dwarfs and models of the surrounding ISM as inferred from nearby column density measurements, finding a lack of sufficient material to account for most polluted stars."
 This paper returns to the question of the origin of the DZ spectral class., This paper returns to the question of the origin of the DZ spectral class.
 Historicallv. the DZ stars outnumboered heir hyclrogen-rich counterparts stars by 25:1 (Dupuisοἱal.19936) until high resolution. spectroscopic searches with aree telescopes uncovered a commoensurate number of DAZ stars (Ixoesteretal.2005a:Zuckerman2003).," Historically, the DZ stars outnumbered their hydrogen-rich counterparts stars by 25:1 \citep{dup93b} until high resolution, spectroscopic searches with large telescopes uncovered a commensurate number of DAZ stars \citep{koe05a,zuc03}."
.. Tocdav. he ratio of known DZ to DAZ stars is greater than 4:1. due almost exclusively to the Sloan Digital Sky Survey (SDSS: Eisensteinetal. 2006)).," Today, the ratio of known DZ to DAZ stars is greater than 4:1, due almost exclusively to the Sloan Digital Sky Survey (SDSS; \citealt{eis06}) )."
 The ability to detect. mild levels of reavy element pollution in these stars. combined with the relative lack of circumstellar material at both the DZ and cool DAZ white dwarls (Farihietal.2009).. makes the origin of the DZ stars challenging to constrain.," The ability to detect mild levels of heavy element pollution in these stars, combined with the relative lack of circumstellar material at both the DZ and cool DAZ white dwarfs \citep{far09a}, makes the origin of the DZ stars challenging to constrain."
 It is argued here that the hundreds of newly identified. cool. helium-rich. white dwarfs from the SDSS provide an excellent. spatially ancl kinematically unbiased sample upon which to test the ISAL accretion hypothesis.," It is argued here that the hundreds of newly identified, cool, helium-rich white dwarfs from the SDSS provide an excellent, spatially and kinematically unbiased sample upon which to test the ISM accretion hypothesis."
 These stars form a large statistical sample that favors an alternative to the aceretion of Galactic eas and. dust: as a whole the SDSS DZ stars do not show the spatial. kinematical. or elemental abundance correlations expected. from. accretion episodes within the plane of the Galaxy.," These stars form a large statistical sample that favors an alternative to the accretion of Galactic gas and dust; as a whole the SDSS DZ stars do not show the spatial, kinematical, or elemental abundance correlations expected from accretion episodes within the plane of the Galaxy."
 Furthermore. the SDSS DZ and DC (helium-rich. featureless) stars appear to be similar populations of polluted and non-pollut stars broadly sharing all relevant characteristics save photosphericed calcium. and. perhaps hydrogen.," Furthermore, the SDSS DZ and DC (helium-rich, featureless) stars appear to be similar populations of polluted and non-polluted stars broadly sharing all relevant characteristics save photospheric calcium, and perhaps hydrogen."
 The SDSS fourth data release (DIU: Acdclman-AleCarthyetal.2006)) contains a prodigious number of new white cdwarfs. in the neighborhood. of 0000 (Eisensteinοἱal.2006).," The SDSS fourth data release (DR4; \citealt{ade06}) ) contains a prodigious number of new white dwarfs, in the neighborhood of 6000 \citep{eis06}."
. Ixcluding multiple entries for the same source. binaries or binary cancidates. and uncertain classifications. there are 95 DZ and 1432. DC stars within the DIA white dwarl catalog.," Excluding multiple entries for the same source, binaries or binary candidates, and uncertain classifications, there are 95 DZ and 142 DC stars within the DR4 white dwarf catalog."
 Under further scrutiny. Dufour.etal.(2007) has expanded the number of confirmed. unique DZ stars to 146. providing effective temperatures. calcium abuncdances. and photometric distancesunder the assumption of logg(cms7)]=8.0.," Under further scrutiny, \citet{duf07} has expanded the number of confirmed, unique DZ stars to 146, providing effective temperatures, calcium abundances, and photometric distancesunder the assumption of $\log\,[g\,({\rm 
cm\,s}^{-2})]=8.0$."
 Of these 288 cool. helium-rich. white chwarks in. DIA... only one was known previously: G1L1-54 or SDSS | 3832124 (Dufouretal.2007:Greenstein 1975).," Of these 288 cool, helium-rich white dwarfs in DR4, only one was known previously; G111-54 or SDSS $+$ 383212.4 \citep{duf07,gre75}."
. Perhaps surprisinglv. the current number of DC and DZ stars in the SDSS are nearly identical.," Perhaps surprisingly, the current number of DC and DZ stars in the SDSS are nearly identical."
 Phe DR4 catalog temperatures for helium-rich stars are generally unreliable near or below KIX. as their. helium atmosphere models do not extend. below this temperature (Eisensteinctal.2006).," The DR4 catalog temperatures for helium-rich stars are generally unreliable near or below K, as their helium atmosphere models do not extend below this temperature \citep{eis06}."
. Owing to this fact. the DC star spectroscopic and photometric data were fitted using the same technique and helium-rich. (metal-free) models as in Dufourctal.(2007)... vielding effective temperatures and Whotometric distances. again assuming logg=&.0.," Owing to this fact, the DC star spectroscopic and photometric data were fitted using the same technique and helium-rich (metal-free) models as in \citet{duf07}, yielding effective temperatures and photometric distances, again assuming $\log\,g=8.0$."
 All Whotometric cata were taken from the DIU white cwarl catalog (Eisensteinetal.2006)... with stellar parameters for he DZ stars from Dufouretal.(2007). and the DC moclel its performed for this work," All photometric data were taken from the DR4 white dwarf catalog \citep{eis06}, with stellar parameters for the DZ stars from \citet{duf07} and the DC model fits performed for this work."
 Spatial ancl kinematical data for these white diwarfs were obtained [rom the SDSS seventh data release (DIU: Abazajianetal. 2009))., Spatial and kinematical data for these white dwarfs were obtained from the SDSS seventh data release (DR7; \citealt{aba09}) ).
 This latest data release includes the USNO-SDSS derived proper motion catalog of Munnοἱal.(2004). and its recent amencments (Munnctal.2008).. with statistical errors around 3mmasvye+.," This latest data release includes the USNO-SDSS derived proper motion catalog of \citet{mun04} and its recent amendments \citep{mun08}, with statistical errors around $^{-1}$."
 For a handful of white dwarls. proper motions were not available in the DIU catalog: for these stars measurements were taken from the SuperCOSMOS Science Archive (SSA:Llamblyctal. 10111. the LSPAL catalog (Lepine&Shara2005).. or were calculated for this work based on two or more positions on archival photographic plates and SDSS images separated by approximately 50 vears.," For a handful of white dwarfs, proper motions were not available in the DR7 catalog; for these stars measurements were taken from the SuperCOSMOS Science Archive (SSA;\citealt{ham01}) ), the LSPM catalog \citep{lep05}, or were calculated for this work based on two or more positions on archival photographic plates and SDSS images separated by approximately 50 years."
 These proper motions are listed in ‘Table 1.., These proper motions are listed in Table \ref{tbl1}.
 The SDSS white cdwarfs. olfer a deeper snapshot of the Solar neighborhood and bevond., The SDSS white dwarfs offer a deeper snapshot of the Solar neighborhood and beyond.
 Owing to the g722 AB magnitude limit of the survey. and its general avoidance of the Galactic plane. these white cdwarls typically represent distant stars (tens 1o. hundreds ofppc) at. intermediate Galactic latitudes.," Owing to the $g>22$ AB magnitude limit of the survey, and its general avoidance of the Galactic plane, these white dwarfs typically represent distant stars (tens to hundreds pc) at intermediate Galactic latitudes."
 Hence the SDSS DZ stars are an excellent probe of the local SM. (LISM) both within and bevond the Local Bubble or Chimney (Welshetal.1999. 1994)... which extends roughly to ppe (Itedfield&Linsky 2008)..," Hence the SDSS DZ stars are an excellent probe of the local ISM (LISM) both within and beyond the Local Bubble or Chimney \citep{wel99,wel94}, , which extends roughly to pc \citep{red08}. ."
 A number of relevant quantities were calculated. [rom the DZ and. DC white chvarl spatial ancl kinematical data: tangential speed. Pi. height above (or below) the Galactic micl-plane ο]. space velocities CVM (assuming zero radial," A number of relevant quantities were calculated from the DZ and DC white dwarf spatial and kinematical data: tangential speed $v_{\rm tan}$ height above (or below) the Galactic mid-plane $|z|$ , space velocities $UVW$ (assuming zero radial"
heoretical models degenerate at the present epoch (Peacock 1991).,theoretical models degenerate at the present epoch \citep{pea97}.
. Until recently. the galaxy. correlation function had been hought to follow a power law Clotsuji&IxiharaLOGO:Pee-Mes1974:Gott&Turner 1979).," Until recently, the galaxy correlation function had been thought to follow a power law \citep{tot69,pee74,got79}."
. SubsequentIvy. a departure rom the power law on small scales (of the order 1 to a few Alpc/h) in the galaxy. correlation function was noticed. in xoneering survevs of the eighties. c.g. (CGuzzoetal.1991). and has now been fully confirmed by many large and deep ealaxy surveys.," Subsequently, a departure from the power law on small scales (of the order 1 to a few Mpc/h) in the galaxy correlation function was noticed in pioneering surveys of the eighties, e.g. \citep{guz91}, and has now been fully confirmed by many large and deep galaxy surveys."
 These consist of various surveys at low and intermediate redshifts (z < 1.5) such as the SDSS etal.2002:Zehavi 2004).. 2dEGIUS (Magliocchettilorciani 2003).. VVDS (LeFevereetal.2005b:Pollo 2006).. COMBO-17 (Phlepsetal.2006).. DEEP? (Coiletal.2006) and for Ivman break galaxies (LBGs) at high redshifts in the SNDS and GOODS surveys (Ouchietal.2005:Leeetal. 90060).," These consist of various surveys at low and intermediate redshifts (z $\le$ 1.5) such as the SDSS \citep{con02,zeh04}, , 2dFGRS \citep{mag03}, VVDS \citep{lef05b, pol06}, COMBO-17 \citep{phl06}, DEEP2 \citep{coi06} and for lyman break galaxies (LBGs) at high redshifts in the SXDS and GOODS surveys \citep{ouc05, lee06}."
.. Earliest measures showing a non-power law [for the galaxy correlation. function. were clillicull to. interpret as 1rey relied heavily upon the angular correlation. function. wf). at low redshifts (Connollyetal.2002).," Earliest measures showing a non-power law for the galaxy correlation function were difficult to interpret as they relied heavily upon the angular correlation function, $\omega(\theta)$, at low redshifts \citep{con02}."
. This meant an integration over a wide range of galaxy luminosities and redshifts as well as complicated correlations between the statistical errors (Zehavietal.2004)., This meant an integration over a wide range of galaxy luminosities and redshifts as well as complicated correlations between the statistical errors \citep{zeh04}.
. Phis was followed by measurements of the projected correlation function. wir). and w(@) in larger and deeper surveys.," This was followed by measurements of the projected correlation function, $w_p(r_p)$, and $\omega(\theta)$ in larger and deeper surveys."
" It has been seen that the deviation [rom a power law becomes more pronounced for bright galaxy samples with L 2 L, (Coilctal.Polloetal.2006). and for LBGs at high redshifts (Leeetal. 2006)."," It has been seen that the deviation from a power law becomes more pronounced for bright galaxy samples with $L$ $>$ $L_\star$ \citep{coi06,pol06} and for LBGs at high redshifts \citep{lee06}."
. Likewise. in SPII simulations the luminous and more strongly clustered. galaxies show a similar behaviour where the ‘kink’ is clearer than in the case of the full galaxy sample (Weinbergctal.2004).," Likewise, in SPH simulations the luminous and more strongly clustered galaxies show a similar behaviour where the 'kink' is clearer than in the case of the full galaxy sample \citep{wei04}."
. Interestingly cnough. the correlation function [or dark matter in N-body simulations is well known to be non-adherent to a power law (Jenkins et al.," Interestingly enough, the correlation function for dark matter in N-body simulations is well known to be non-adherent to a power law (Jenkins et al."
 1998. Ixaulf/mann et al.," 1998, Kauffmann et al."
 1999. Cooray Sheth 2002 and references therein)," 1999, Cooray Sheth 2002 and references therein)."
 The natural question that arises is how biased is the galaxy distribution with respect to the uncdoerlving matter distribution?, The natural question that arises is how biased is the galaxy distribution with respect to the underlying matter distribution?
 “Phe overall shape of the dark matter correlation function is mostly unallectecl as one goes to higher redshifts as seen in SPLL simulations (Weinbergetal.2004). and N-bocly simulations (Jenkinsetal.1998)., The overall shape of the dark matter correlation function is mostly unaffected as one goes to higher redshifts as seen in SPH simulations \citep{wei04} and N-body simulations \citep{jen98}.
. ‘This is different to what is seen for high-z galaxies (Ouchietal.2005:Lee2006:Zhengct 2007).. where the so-called. “break? is more prominent implying that the biasing of the galaxy distribution on Large scales. spatial exclusion of dark matter halos on small scales. alonewith a host ofother complex physical processes. such as dvnanical [riction. feedback. (roni supernovae. ram-pressure stripping etc.," This is different to what is seen for high-z galaxies \citep{ouc05, lee06, zhe07}, where the so-called 'break' is more prominent implying that the biasing of the galaxy distribution on large scales, spatial exclusion of dark matter halos on small scales, alongwith a host of other complex physical processes, such as dynamical friction, feedback from supernovae, ram-pressure stripping etc."
 conspire in a non-trivial way to produce cdillerences in the galaxy. correlation function at different redshilts., conspire in a non-trivial way to produce differences in the galaxy correlation function at different redshifts.
 The break in the power law can be physically interpreted in the language of the halo model (see Cooray Sheth 2002 for a detailed review) as the transition between two scales - small scales lving within the halo to those larger than the halo., The break in the power law can be physically interpreted in the language of the halo model (see Cooray Sheth 2002 for a detailed review) as the transition between two scales - small scales lying within the halo to those larger than the halo.
 Ht is only natural to use a halo-based prescription where galaxies form by the cooling of gas within dark matter halos (White&Rees1978).. which are bound. virialized. clumps of dark matter that are roughly 200 times the background density at that time (Gunn&Gott.1972).," It is only natural to use a halo-based prescription where galaxies form by the cooling of gas within dark matter halos \citep{whi78}, which are bound, virialized clumps of dark matter that are roughly 200 times the background density at that time \citep{gun72}."
. The galaxies occupy dark matter halos following a HOD (halo occupation distribution) model., The galaxies occupy dark matter halos following a HOD (halo occupation distribution) model.
 In turn the HOD fully describes the bias in the distribution of galaxies with respect to the underlying dark matter clistribution (Berlind&Wein-berg2002)., In turn the HOD fully describes the bias in the distribution of galaxies with respect to the underlying dark matter distribution \citep{ber02}.
. The motivation for LOD based models. arose when it was noticed. that the clustering of galaxies could be reproduced. by populating halos in semi-analytic mocdels with galaxies following a particular probability cistributioncitekau9., The motivation for HOD based models arose when it was noticed that the clustering of galaxies could be reproduced by populating halos in semi-analytic models with galaxies following a particular probability distribution.
ben00.ber02. Furthermore. it has been seen that without the help of a proper halo based. cleseription the strong clustering of red galaxies (at z 2 3) can be explained bv high and unrealistic (anvwhere in the range of τ-200 ealaxies per halo) occupation numbers to match the observed. number density and strong clustering of a small number of high mass halos (Zheng2004).," Furthermore, it has been seen that without the help of a proper halo based description the strong clustering of red galaxies (at z $\simeq$ 3) can be explained by high and unrealistic (anywhere in the range of 70-200 galaxies per halo) occupation numbers to match the observed number density and strong clustering of a small number of high mass halos \citep{zhe04}."
. Recently. several groups. have studied the galaxy correlation in light of the IIOD mocels (Magliocchetti&2006:Conrovetal.Zheng 2007).," Recently, several groups have studied the galaxy correlation in light of the HOD models \citep{mag03, van03, ham04, zeh05, ouc05, phl06, con06, zhe07}."
. Most. of hese works have mainly concentrated on obtaining best-it HOD parameters and. consequently the evolution. in he HOD and. information on the underlving dark matter distribution., Most of these works have mainly concentrated on obtaining best-fit HOD parameters and consequently the evolution in the HOD and information on the underlying dark matter distribution.
 ln some cases. data from cdillerent surveys waving dillerent selections were used to study the evolution (Conroyetal.2006:Zheng2007) The work in this »iper Complements and extends these analyses by studying he evolution in the WOD over a redshift range z z 0.1 - 1.3 for à variety of Iuminositv-limited. samples but always or data from the same survey. the VIAIOS VLT. Deep Survey (VVDS).," In some cases, data from different surveys having different selections were used to study the evolution \citep{con06, zhe07}
 The work in this paper complements and extends these analyses by studying the evolution in the HOD over a redshift range z $\approx$ 0.1 - 1.3 for a variety of luminosity-limited samples but always for data from the same survey, the VIMOS VLT Deep Survey (VVDS)."
 We also study two cifferent HOD moclels in order to obtain a better understanding on the degeneracies tween the various parameters., We also study two different HOD models in order to obtain a better understanding on the degeneracies between the various parameters.
 The paper is organized. as follows., The paper is organized as follows.
 ln section 2. we πο] describe the data. used. from the VVDS. survey., In section \ref{data} we briefly describe the data used from the VVDS survey.
 This section is followed by Section 3. giving an outline of he theoretical framework., This section is followed by Section \ref{model} giving an outline of the theoretical framework.
 Section 4 describes the fitting procedure alongwith the best-fit parameters obtained for the different models., Section \ref{results} describes the fitting procedure alongwith the best-fit parameters obtained for the different models.
 The estimates for the average halo mass and number of galaxies per halo (galaxy weighted) for the different: Buminositv-threshold. samples are. also. presented., The estimates for the average halo mass and number of galaxies per halo (galaxy weighted) for the different luminosity-threshold samples are also presented.
 Finally. Section 5. wraps up with a discussion and conclusion of the results.," Finally, Section \ref{disc} wraps up with a discussion and conclusion of the results."
" ""Throughout we will assume a flat. ACDAL model for which (Og. h.m.) = (0.3.0.7.0.9) at s=0."," Throughout we will assume a flat $\Lambda$ CDM model for which $\Omega_0,h,\sigma_8$ ) = (0.3,0.7,0.9) at $z=0$."
 Here Qo is the density in units of critical density today. P is the Llubble constant today in units of 100. km 1 and ax describes the rms Ductuations of the initial field evolved to the present time using linear theory ancl smoothed with a top-hat filter of radius 8 Alpesh.," Here $\Omega_0$ is the density in units of critical density today, $h$ is the Hubble constant today in units of 100 km $^{-1}$ $^{-1}$, and $\sigma_8$ describes the rms fluctuations of the initial field evolved to the present time using linear theory and smoothed with a top-hat filter of radius 8 $h$."
 ALL absolute magnitudes are in the AB svstem., All absolute magnitudes are in the AB system.
 The data used in this analysis comes from the First Epoch VVDS (VIMOS-VLI Deep Survey) 0226-04 “Deep” field (hereafter VVDS-Deep)., The data used in this analysis comes from the First Epoch VVDS (VIMOS-VLT Deep Survey) 0226-04 “Deep” field (hereafter VVDS-Deep).
 X complete description of the data. survey strategy. data reduction ancl primary goals can be founcl in Le Fevyreet al. (," A complete description of the data, survey strategy, data reduction and primary goals can be found in Le Fèvvreet al. ("
2005a).,2005a).
 Llere we will simply give a short description of the data used., Here we will simply give a short description of the data used.
 Carciofi.(2011).. (Leeetal.1," \citep[for reviews, see][]{por03,baa01}. \citep{oud11,kra11}. \citet{car11},"
991) (e.g..Cranmer2005).," \citep{lee91} \citep{riv98,nei02} \citep{and86,cra09} \citep[e.g.,][]{cra05}, \citep[see][for a detailed assessment]{owo06}."
. Be stars exhibit a wide range of variability. ranging from emission line profile variations (e.g.V/Rvariability.indica-ciofietal. 2009).. to photometric and polarimetric variability. indicating that frequently the disks are not fed at a constant rate.," Be stars exhibit a wide range of variability, ranging from emission line profile variations \citep[e.g, $V/R$ variability, indicative of density waves within the disk; ][]{ste09,car09}, to photometric and polarimetric variability, indicating that frequently the disks are not fed at a constant rate."
 In some cases. the disk can disappear for decades before being rebuilt (seereviewbyUnderhill&Doazan1982;Wis-niewskietal. 2010).," In some cases, the disk can disappear for decades before being rebuilt \citep[see review by][]{und82,wis10}."
. Oftentimes. this rebuilding phase occurs in à series of outbursts. during which the disk. apparently is fed for a short period of time followed by a period of quiescence. during which the nascent disk can partially reaccrete onto the star.," Oftentimes, this rebuilding phase occurs in a series of outbursts, during which the disk apparently is fed for a short period of time followed by a period of quiescence, during which the nascent disk can partially reaccrete onto the star."
 Similar outbursts can oceur while the disk is disappearing., Similar outbursts can occur while the disk is disappearing.
 Once such example is the Be star 28CCMa., Once such example is the Be star CMa.
 Over the past 40 years. CCMa has exhibited quasi-regular outbursts. every 8 years or so. when the star brightens by about half a magnitude in the V-band.," Over the past 40 years, CMa has exhibited quasi-regular outbursts, every 8 years or so, when the star brightens by about half a magnitude in the $V$ -band."
 This usually occurs. via multiple individual injections (e.g.. the rapid rises seen in Fig.," This usually occurs via multiple individual injections (e.g., the rapid rises seen in Fig."
 1). during which the star brightens at a rate of 0.02+ magd7!. due to free-bound and free-free radiation from the disk.," 1), during which the star brightens at a rate of $0.02\pm0.01\,\mathrm{mag\,d}^{-1}$ , due to free-bound and free-free radiation from the disk."
 The outbursts last for about two years (seeStefletmetricvariationsof28 CMa)..," The outbursts last for about two years \citep[see][for a detailed discussion of the long-term photometric variations of 28\,CMa]{ste03}."
 A particularly well-monitored outburst occurred from 2001 to 2004. and this outburst was followed by a long quiescent phase from 2004 to 2000.," A particularly well-monitored outburst occurred from 2001 to 2004, and this outburst was followed by a long quiescent phase from 2004 to 2009."
 Furthermore it appears that during this quiescent phase. very little additional material was injected into the disk (only a few minor events).," Furthermore it appears that during this quiescent phase, very little additional material was injected into the disk (only a few minor events)."
 Thus nature has provided us the perfect experiment to study how Be star disks dissipate., Thus nature has provided us the perfect experiment to study how Be star disks dissipate.
 If we accept as a premise that the viscous decretion disk model is correct (and that no other mechanisms participate in the disk mass budget). then the timescale of the post-outburst photometric decline ts controlled solely by the disk viscosity.," If we accept as a premise that the viscous decretion disk model is correct (and that no other mechanisms participate in the disk mass budget), then the timescale of the post-outburst photometric decline is controlled solely by the disk viscosity."
 In this paper. we model the V-band light curve with the goals of: 1) directly measuring the disk viscosity parameter. o. and 2) further testing the viscous decretion disk model by comparing its time-dependent predictions of the disk dissipation against the detailed shape of the photometric light curve.," In this paper, we model the $V$ -band light curve with the goals of: 1) directly measuring the disk viscosity parameter, $\alpha$, and 2) further testing the viscous decretion disk model by comparing its time-dependent predictions of the disk dissipation against the detailed shape of the photometric light curve."
A moclel is presented here in which a black hole grows with its surrounding stellar spheroid by the accretion of ιοί gas.,A model is presented here in which a black hole grows with its surrounding stellar spheroid by the accretion of hot gas.
 The gas is also cooling and forming a clistributed gaseous cold component within which stars slowly form., The gas is also cooling and forming a distributed gaseous cold component within which stars slowly form.
 In he central regions the cold component provides the recquired column density for absorption of the quasar radiation from he acereting black hole., In the central regions the cold component provides the required column density for absorption of the quasar radiation from the accreting black hole.
 The quasar is assumed to make a slow wind. which eventually becomes powerful enough to ow away the absorbing gas. probably due to à wind. and an unobscured quasar is seen (see Silk Rees 1998 and DIandford 1999 for a similar wind model).," The quasar is assumed to make a slow wind, which eventually becomes powerful enough to blow away the absorbing gas, probably due to a wind, and an unobscured quasar is seen (see Silk Rees 1998 and Blandford 1999 for a similar wind model)."
 Because the fuel supply for the black hole has also been ejected however. the quasar dies when its accretion disc is exhausted.," Because the fuel supply for the black hole has also been ejected however, the quasar dies when its accretion disc is exhausted."
 Ehe high obscuration phase occurs only during the growing phase and thus at redshifts greater than unity which have not been explored. vet in the hard. X-ray. band., The high obscuration phase occurs only during the growing phase and thus at redshifts greater than unity which have not been explored yet in the hard X-ray band.
 Reprocessing of the radiation by the absorbing gas into far infrared emission may make these objects detectable in the sub-millimetre band., Reprocessing of the radiation by the absorbing gas into far infrared emission may make these objects detectable in the sub-millimetre band.
 Alagorrian et al (1998) find from the demographies of nearby massive black holes that the mass of the hole Adpy is proportional to the mass of the surrounding spheroid A. or bulge: AdpHz0.005AZn.," Magorrian et al (1998) find from the demographics of nearby massive black holes that the mass of the hole $M_{\rm BH}$ is proportional to the mass of the surrounding spheroid $M_{\rm sph}$ or bulge; $M_{\rm
BH}\approx 0.005M_{\rm sph}$."
 There is considerable scatter about this relation of order X1 dex., There is considerable scatter about this relation of order $\pm1$ dex.
 This can be combined with the mass function of ealactic bulges. where most of he mass resides.," This can be combined with the mass function of galactic bulges, where most of the mass resides."
 Salucci et al (1999) use the Schechter ‘unetion luminosity function. of bulges to obtain the local mass function of black holes., Salucci et al (1999) use the Schechter function luminosity function of bulges to obtain the local mass function of black holes.
 Most of the mass lies in black voles o£ ineliviclual mass around the break in the function. be 35IOM...," Most of the mass lies in black holes of individual mass around the break in the function, i.e. $3\times 10^8\Msun$."
 Taking this as a typical mass. and assuming growth w accretion with a radiative elliciency of η=0.14.1. the rolometric Luminosity of the final object Ly=31023Oleres1 |i is at one and 0.01 times the Eddington imit Leag. respectively.," Taking this as a typical mass, and assuming growth by accretion with a radiative efficiency of $\eta=0.1 \eta_{-1}$, the bolometric luminosity of the final object $L_{\rm B}=3\times 10^{46}-3\times
10^{44}\ergps$ if it is at one and 0.01 times the Eddington limit $L_{\rm Edd}$, respectively."
 The Salpeter time for the growth of the black hole. Le. the mass doubling time. ἐς=310 GLeaa/Lai)noave. M we now assume that the initial mass of all massive black holes is less than that of the central black hole in our Galaxy. sav ~LO?M... and has grown to 3.OPAL.{ withinH 3. Civr.D which+ corresponds to z~κ2. then δές35.10(0 vrand the black hole must have grown at arate of at about 10 per cent of the Exldington rate or more.," The Salpeter time for the growth of the black hole, i.e. the mass doubling time, $t_{\rm s}=3\times 10^7
(L_{\rm Edd}/L_{\rm B})\eta_{0.1}\yr.$ If we now assume that the initial mass of all massive black holes is less than that of the central black hole in our Galaxy, say $\sim 10^6\Msun$, and has grown to $3\times 10^6\Msun$ within 3 Gyr, which corresponds to $z\sim2$, then $8t_{\rm s}<3\times 10^9\yr$ and the black hole must have grown at a rate of at about 10 per cent of the Eddington rate or more."
 Thus over the final 3.107vr. Lp>310Meres+.," Thus over the final $3\times 10^8\yr,$ $L_{\rm B}>3\times 10^{45}\ergps$."
 From the work of Elvis et al (1994) the observed 2.10 keV luminosity is about 3 per cent of the bolometric one for a quasar so we can conclude that £L(210Κον}>ores curing the growth of a typical black hole and we are dealing with powerful. obsceured quasars.," From the work of Elvis et al (1994) the observed 2–10 keV luminosity is about 3 per cent of the bolometric one for a quasar so we can conclude that $L(2-10\keV)>10^{44}\ergps$ during the growth of a typical black hole and we are dealing with powerful, obscured quasars."
 lt is clear that the objects making up the X-ray background. and the radiative growth phase of most of the mass in nearby black holes are not represented by any object observed so Bu.," It is clear that the objects making up the X-ray background, and the radiative growth phase of most of the mass in nearby black holes are not represented by any object observed so far."
 The spectrum of the N-arayv Background. requires. that most accretion power is absorbed., The spectrum of the X-ray Background requires that most accretion power is absorbed.
 This means that some absorption (sav Ny1077cm27) occurs in 90 per cent of objects ancl heavy absorption (Ngtotem7) in 3050 per cent of them.," This means that some absorption (say $N_{\rm
H}>10^{22}\psqcm$ ) occurs in 90 per cent of objects and heavy absorption $N_{\rm H}>10^{24}\psqcm$ ) in 30–50 per cent of them."
 10 per cent are unabsorbed and are the quasars identified in blue optical surveys or are bright in the soft. X-ray band of ROSAT., 10 per cent are unabsorbed and are the quasars identified in blue optical surveys or are bright in the soft X-ray band of ROSAT.
 So far these could. be distinct different populations of quasars or different phases in the erowth of all quasars., So far these could be distinct different populations of quasars or different phases in the growth of all quasars.
 Such high absorbed fractions mean that the covering fraction of the sky by high column density material seen from a erowing black hole must approach 4tasr., Such high absorbed fractions mean that the covering fraction of the sky by high column density material seen from a growing black hole must approach $4\pi \sr$.
 This cannot be provided by any thin dise or by the standard torus of unified models., This cannot be provided by any thin disc or by the standard torus of unified models.
 Phe absorbing matter is probably concentrated within the innermost [ον 100 kpe. or its total mass becomes high.," The absorbing matter is probably concentrated within the innermost few 100 kpc, or its total mass becomes high."
 Lt must. presumably be cold and. fairly. neutral. or it will not absorb the X-rays.," It must presumably be cold and fairly neutral, or it will not absorb the X-rays."
 At first sight this is at variance with it being space covering. especially if it must be so for a Cyr or more.," At first sight this is at variance with it being space covering, especially if it must be so for a Gyr or more."
 Such matter should collide and clissipate into a clisclike structure., Such matter should collide and dissipate into a disclike structure.
 Following our earlier work on a cireumnuclear starburst (Fabian ct al 1998). it is plausible that collisions do take place leading to dissipation but that some massive star formation occurs in the shocked and. cooled dense gas.," Following our earlier work on a circumnuclear starburst (Fabian et al 1998), it is plausible that collisions do take place leading to dissipation but that some massive star formation occurs in the shocked and cooled dense gas."
 The winds and supernovae [from those stars then supply the energy to keep the rest of the cold matter in a chaotic and space covering stato., The winds and supernovae from those stars then supply the energy to keep the rest of the cold matter in a chaotic and space covering state.
 A more detailed model can be developed by assuming that the black hole is growing at the same time as the galaxy does., A more detailed model can be developed by assuming that the black hole is growing at the same time as the galaxy does.
 Ehe stellar spheroid continues to grow by cooling of the eas heated by gravitational collapse of the protogalaxy., The stellar spheroid continues to grow by cooling of the gas heated by gravitational collapse of the protogalaxy.
 Lt is likely that the hot phase density while the galaxy continues to grow is the maximum possible. which means that the radiative cooling time of the gas equals the eravitationa infall time.," It is likely that the hot phase density while the galaxy continues to grow is the maximum possible, which means that the radiative cooling time of the gas equals the gravitational infall time."
 This condition has been studied in the contex of quasar fuelling and growth by Nulsen Fabian (1999)., This condition has been studied in the context of quasar fuelling and growth by Nulsen Fabian (1999).
 The gas aceretes in a Bondi Dow. probably forming a disc well within the Doncdi radius.," The gas accretes in a Bondi flow, probably forming a disc well within the Bondi radius."
 The accretion rate is such tha it can twpically be LO per cent of the Exldington value., The accretion rate is such that it can typically be 10 per cent of the Eddington value.
 ‘The situation as envisaged here is essentially a maxima cooling How and we can use the properties. of observec cooling flows. (Fabian 1994 and references therein) to indicate how the cooled gas is distributed and how it may cack to absorption., The situation as envisaged here is essentially a maximal cooling flow and we can use the properties of observed cooling flows (Fabian 1994 and references therein) to indicate how the cooled gas is distributed and how it may lead to absorption.
 X-rav observations of cluster. cooling ows show that the mass deposited by cooling is distribute with Ad(<r)xr. where r is the radius.," X-ray observations of cluster cooling flows show that the mass deposited by cooling is distributed with $\dot M(<r)\propto r$, where $r$ is the radius."
 The density istribution of the cooled gas is therefore pxor27., The density distribution of the cooled gas is therefore $\rho\propto r^{-2}$.
 Uf 1 gravitational potential of the galaxy is isothermal i remains so., If the gravitational potential of the galaxy is isothermal it remains so.
 X-ray absorption with column clensities of the der of 107em.7. is also observed in many cluster cooling ows., X-ray absorption with column densities of the order of $10^{21}\psqcm$ is also observed in many cluster cooling flows.
 Although not observed in other wavehands. it coulc represent very cold. clusty gas (Fabian et al 1994).," Although not observed in other wavebands, it could represent very cold, dusty gas (Fabian et al 1994)."
 lt is therefore. proposed. that. the. cooled. gas in »otogalaxies does not all rapidly form stars but that much of it forms long-lived cold dusty clouds., It is therefore proposed that the cooled gas in protogalaxies does not all rapidly form stars but that much of it forms long-lived cold dusty clouds.
 As discussed above.," As discussed above,"
" As we will show, this method almost perfectly reproduces the CDA computations while the computation time required is much smaller."," As we will show, this method almost perfectly reproduces the CDA computations while the computation time required is much smaller."
 The method we propose is based on the idea of simulating the entire aggregate as a sphere with an effective refractive index., The method we propose is based on the idea of simulating the entire aggregate as a sphere with an effective refractive index.
 There are several methods in the literature to compute the effective refractive index under certain assumptions., There are several methods in the literature to compute the effective refractive index under certain assumptions.
 Recently have analyzed several of these effective medium theories for the optical properties of porous particles in the visual part of the spectrum., Recently \citet{2007ApOpt..46.4065V} have analyzed several of these effective medium theories for the optical properties of porous particles in the visual part of the spectrum.
 In our case the, In our case the
spectroscopic survey of the stellar radial velocities of M-type giant stars whose population membership in the bulge is well established.,spectroscopic survey of the stellar radial velocities of M-type giant stars whose population membership in the bulge is well established.
 These giants most of the 2 jum radiation whose box-shaped light providedistribution motivates bar models., These giants provide most of the 2 $\mu$ m radiation whose box-shaped light distribution motivates bar models.
 emphasizes measurements in two strips at latitude b2—4? and b2—8? and at longitude -10? «/ -10?., emphasizes measurements in two strips at latitude $b=-4^\circ$ and $b=-8^\circ$ and at longitude $-10^\circ$ $< l <$ $+10^\circ$.
 A the minor axis (/= 0?) has also been observed.," A strip along the minor axis $l
\equiv 0^\circ$ ) has also been observed."
" We stripuse alongnearly 5,000 stellar radial velocities in this report."," We use nearly 5,000 stellar radial velocities in this report."
" A preliminary analysis of data found strong cylindrical rotation (Howardetal.2009) consistent with an edge-on, bar-like pseudobulge, although a precise fit of a bar model to the data was not available."," A preliminary analysis of data found strong cylindrical rotation \citep{how_etal_09} consistent with an edge-on, bar-like pseudobulge, although a precise fit of a bar model to the data was not available."
 This success leads us here to construct a full evolutionary N-body model that we can fit to the radial velocity data., This success leads us here to construct a full evolutionary $N$ -body model that we can fit to the radial velocity data.
 We use cylindrical particle-mesh code (Sellwood&Val-luri1997;aShen&Sellwood2004) to build fully self-consistent N-body galaxies.," We use a cylindrical particle-mesh code \citep{sel_val_97, she_sel_04} to build fully self-consistent $N$ -body galaxies."
 It is well suited to study the evolution of disk galaxies: we model the disk with at least 1 million to provide particle resolution near the center where particlesthe density is highhigh., It is well suited to study the evolution of disk galaxies: we model the disk with at least 1 million particles to provide high particle resolution near the center where the density is high.
" We try to construct the simplest self-consistent N-body models that fit the data, avoiding contrived models with too many free parameters."," We try to construct the simplest self-consistent $N$ -body models that fit the data, avoiding contrived models with too many free parameters."
" Initially, they contained only an unbarred disk and a dark halo."," Initially, they contained only an unbarred disk and a dark halo."
 The profile of the Galactic halo is poorly constrained observationally; we adopt a rigid pseudo-isothermal halo potentialD)., The profile of the Galactic halo is poorly constrained observationally; we adopt a rigid pseudo-isothermal halo potential.
" Here V, — 250 km s is the asymptotic circular-orbit rotation velocity at infinity, and R,=15 kpc is the core radius inside which the potential is effectively constant."," Here $\Vc$ $\sim$ 250 km $^{-1}$ is the asymptotic circular-orbit rotation velocity at infinity, and $\Rc=15$ kpc is the core radius inside which the potential is effectively constant."
 This halo gives a nearly flat rotation curve of ~ 220 km s! between 5 to 20 kpc., This halo gives a nearly flat rotation curve of $\sim$ 220 km $^{-1}$ between 5 to 20 kpc.
 A simple halo form allows us to run many simulations quickly; this is important for a parameter search such as the present one., A simple halo form allows us to run many simulations quickly; this is important for a parameter search such as the present one.
" A rigid halo also omits dynamical friction on the bar, but the central density of the cored halo we adopt is low enough so that friction will be very mild (e.g.Debattista&Sellwood2000;Athanassoula&Misiriotis 2002)."," A rigid halo also omits dynamical friction on the bar, but the central density of the cored halo we adopt is low enough so that friction will be very mild \citep[e.g.][]{deb_sel_00,ath_mis_02}."
". More importantly, we are mainly interested in the bulge, which is embedded well interior to Κο."," More importantly, we are mainly interested in the bulge, which is embedded well interior to $\Rc$."
 So the exact profile of the dark halo at large radii is not critical., So the exact profile of the dark halo at large radii is not critical.
 We will explore more sophisticated halos in a future study., We will explore more sophisticated halos in a future study.
" In our models, a bar develops self-consistently from the initially unbarred, thin disk."," In our models, a bar develops self-consistently from the initially unbarred, thin disk."
" Bar formation enhances the radial streaming motions of disk particles, so the radial velocity dispersion quickly grows much bigger than the vertical one."," Bar formation enhances the radial streaming motions of disk particles, so the radial velocity dispersion quickly grows much bigger than the vertical one."
 Consequently the disk buckles vertically out of the plane like a fire hose; this is the well known buckling or corrugation instability (Toomre1966;Combesetal.1990;Raha 1991).," Consequently the disk buckles vertically out of the plane like a fire hose; this is the well known buckling or corrugation instability \citep{toomre_66,com_etal_90,rah_etal_91}."
. It raises the vertical velocity dispersion and increases the bar's thickness., It raises the vertical velocity dispersion and increases the bar's thickness.
 This happens on a short dynamical timescale and saturates in a few hundred million years., This happens on a short dynamical timescale and saturates in a few hundred million years.
 The central part of the buckled bar is elevated well above the disk mid-plane and resembles the peanut morphology of many bulges including the one in our Galaxy Rahaetal. 1991).," The central part of the buckled bar is elevated well above the disk mid-plane and resembles the peanut morphology of many bulges including the one in our Galaxy \citep{com_etal_90,rah_etal_91}."
". Out of a large set of N-body models, we find the one that best matches ourBRAVA kinematic data after suitable mass "," Out of a large set of $N$ -body models, we find the one that best matches our kinematic data after suitable mass scaling."
"The barred disk evolved from thin exponential disk that scaling.contains M,=4.25x abouta 55 of the total mass at the truncation radius (5 10/094...scale-lengths)."," The barred disk evolved from a thin exponential disk that contains $M_{\rm d}=4.25\times 10^{10}\Ms$ , about 55 of the total mass at the truncation radius (5 scale-lengths)."
 Thescale- and scale-height of the initial disk are ~ 1.9 kpc and, Thescale-length and scale-height of the initial disk are $\sim$ 1.9 kpc and
((Alereghetti et al.,(Mereghetti et al.
 2000)., 2000).
 Furthermore. 21-010 continu and several molecular line maps of the region (Nim νου. 2002) show that 0.3 is rather separated from the primary componcuts of W31 O8and 0.1).," Furthermore, 21-cm continuum and several molecular line maps of the region (Kim Koo 2002) show that $-$ 0.3 is rather separated from the primary components of W31 $-$ 0.3 and $-$ 0.1)."
 Thus. as of this writing. there is uo definitive plivsical linkage between aanucd the major componcuts of W31.," Thus, as of this writing, there is no definitive physical linkage between and the major components of W31."
 Toecther. these imply that 0.3 mar lie at a different distance along," Together, these imply that $-$ 0.3 may lie at a different distance along"
The X-ray Spectrn. fransjent behaviour. and lack of an optical aud infrared couuterpar place strong constraints on the nature ofIO.,"The X-ray spectrum, transient behaviour, and lack of an optical and infrared counterpart place strong constraints on the nature of."
 Nou-transicnt objects such as isolatec cooling neutron stars aud rotation powered pulsars with moderate magnetic field streugths can inuuediatcly be ruled out. leaving behind objects that are known to exhibit large scale variability.," Non-transient objects such as isolated cooling neutron stars and rotation powered pulsars with moderate magnetic field strengths can immediately be ruled out, leaving behind objects that are known to exhibit large scale variability."
 Although the CFUT optical imaging was uo taken at the same time as wwas detected in A-ravs. we would still expect to detect an optical counterpart cousisteut with the A-rav source position if wwere a stars Ὁν variable (CV) or an Active Calactic Nucleus (AGN).," Although the CFHT optical imaging was not taken at the same time as was detected in X-rays, we would still expect to detect an optical counterpart consistent with the X-ray source position if were a star, a cataclysmic variable (CV) or an Active Galactic Nucleus (AGN)."
 The X-ray detection and deep g-baud and ‘-baucd limits indicate N-ray to optical Hux ratios of Py Εν = A2OO and Fy /F; > 9δ00 for10., The X-ray detection and deep $g$ -band and $i$ -band limits indicate X-ray to optical flux ratios of $_X$ $_g$ $\geq$ 8200 and $_X$ $_i$ $\geq$ 5800 for.
". Even if the source just outsice the cerror circle were the counterpart the flux ratios would be Py /E,  1300 aud Fy /F; ~ 310.", Even if the source just outside the error circle were the counterpart the flux ratios would be $_X$ $_g$ $\sim$ 1300 and $_X$ $_i$ $\sim$ 340.
 These are far in excess of the ratios observed for ACN and stars. which typically have X-ray to optical flux ratios ofFy f/Fo < 10 and < 0.01. respectively (e.g.Mainierietal.2002).," These are far in excess of the ratios observed for AGN and stars, which typically have X-ray to optical flux ratios of $_X$ $_O$ $<$ 10 and $<$ 0.01, respectively \citep[e.g.][]{mai02}."
. It is also inconsistent with CVs which typically have Fy το x 01 (I&uulkersetal.2006)., It is also inconsistent with CVs which typically have $_X$ $_O$ $\leq$ 0.1 \citep{kuu06}.
. Iheh levels of dust extinction could explain he lack of an opticalcounterpart”...however the xositiou of ~22° out of the Galactic plane puts it iu a reeion of low dust extinction with E(B-V) = 1.062 mae (Schlegeletal.1998).," High levels of dust extinction could explain the lack of an optical,however the position of $\sim$ $^\circ$ out of the Galactic plane puts it in a region of low dust extinction with E(B-V) = 0.062 mag \citep{sch98}."
". The weighted average Galactic neutral lvdrogen absorption iu he direction of aas ineasured by the Leiden/Arecutine/Boun (LAB) survey of Galactic IIT is 7.6 « 1078 ? (Nalberlactal.2005).. consistent with the Ny values determuned through the XN-rav spectral fitting (Le. 8 .— 35 « 107"" 7. Tables 1 2j) with some additioual intrinsic absorption!."," The weighted average Galactic neutral hydrogen absorption in the direction of as measured by the Leiden/Argentine/Bonn (LAB) survey of Galactic HI is 7.6 $\times$ $^{20}$ $^{-2}$ \citep{kal05}, consistent with the $_H$ values determined through the X-ray spectral fitting (i.e. 8 – 35 $\times$ $^{20}$ $^{-2}$, Tables \ref{bhxrb} \ref{specpar}) ) with some additional intrinsic ."
. Tf, If
 LO° 105 , $^6$ $^8$ 
"If black hole sclfvcenlation iu the radio-loud phase is mediated bv cosnüc rays produced in the jet core. we have shown that the coupliug efficiency 64, Is indeed independent (for cosmic ravs scattering with Alfvéunu waves self-senerated via the streaming instability. see C27))) or weakly dependent (for resonant diffusion by magnetic turbulence with a Ikolinogorov. spectrum. seo (22))) on the black hole In our cosuic rav-diiven feedback model we should then expect that AtiAvg,X77 ov. if Ay, is a constant with black hole mass. that the duration of the radio-loud phase should scale as μιx» σε.","If black hole self-regulation in the radio-loud phase is mediated by cosmic rays produced in the jet core, we have shown that the coupling efficiency $\epsilon_\ditto{CR}$ is indeed independent (for cosmic rays scattering with Alfvénn waves self-generated via the streaming instability, see ) or weakly dependent (for resonant diffusion by magnetic turbulence with a Kolmogorov spectrum, see ) on the black hole In our cosmic ray-driven feedback model we should then expect that $\Delta t_{_{\rm RP}}\,\Lambda_{_{\rm Edd}}\propto\sigma_\star^2$ or, if $\Lambda_\ditto{Edd}$ is a constant with black hole mass, that the duration of the radio-loud phase should scale as $\Delta t_\ditto{RP}\propto\sigma_\star^2$ ."
" As discussed iu δν, a reasonable upper limit for Af, may be given by the cosmic ray diffusion timescale fill within the host ealaxy."," As discussed in \ref{sec:propag}, a reasonable upper limit for $\Delta t_\ditto{RP}$ may be given by the cosmic ray diffusion timescale $t_\ditto{diff,\dg}$ within the host galaxy."
" In fact. if the jet is powered bv black hole accretion, AGN radio activity would be terminated when the interstellar cosmic ravs lave diffused up to the scales where most of the gas resides. and started to push it outward. thus preventing further accretion."," In fact, if the jet is powered by black hole accretion, AGN radio activity would be terminated when the interstellar cosmic rays have diffused up to the scales where most of the gas resides, and started to push it outward, thus preventing further accretion."
" Tuterestingly. if f, aud A,/A, do uot depeud ou black hole mass. shows that £44,ME.x2σε. the saine scaling that Af, should have in order to satisfy with a scale-indepeudent coupling efficiency."," Interestingly, if $f_g$ and $\lambda_\ditto{CR}/\lambda_\ditto{T}$ do not depend on black hole mass, shows that $t_\ditto{diff,\dg}\propto\sigma_\star^2$, the same scaling that $\Delta t_\ditto{RP}$ should have in order to satisfy with a scale-independent coupling efficiency."
 Note that for Af.~fuu. suggests that the duration of the radio-loud phase for the most massive ealaxies iav be comparable to the Salpeter time iu eq. (128)).," Note that for $\Delta t_\ditto{RP}\sim t_\ditto{diff,\dg}$, suggests that the duration of the radio-loud phase for the most massive galaxies may be comparable to the Salpeter time in eq. \ref{e: Salpeter}) ),"
 whereas in relatively simall galaxies it should last less than the quasar pliase., whereas in relatively small galaxies it should last less than the quasar phase.
 In other words. for low-mass galaxies the black hole seltresulation would be confined to an episodic radio-loud epoch much shorter than the tine required to accrue the black hole mass.," In other words, for low-mass galaxies the black hole self-regulation would be confined to an episodic radio-loud epoch much shorter than the time required to accrue the black hole mass."
 We propose ᾱ selfresulation imiechanisni for supermassive black holes aud them —host bulees —that operates in actively-accreting svstenis with powerful radio-loud activity., We propose a self-regulation mechanism for supermassive black holes and their host bulges that operates in actively-accreting systems with powerful radio-loud activity.
 In the core of relativistic radio jets. internal shocks resulting from the dissipative interaction of shells of matter. intermitteuth ejected from the central engine. couvert a fraction of the ordered kinetic cherey of the jet flow iuto thermal form.," In the core of relativistic radio jets, internal shocks resulting from the dissipative interaction of shells of matter, intermittently ejected from the central engine, convert a fraction of the ordered kinetic energy of the jet flow into thermal form."
 If a sufücieut nuuber of the protous (or “cosmic rays] heated and randonuüzed at internal shocks escape the radio core into the interstellar mediun of the host galaxy. they nay profoundly affect the evolution of the ealaxv and its black hole.," If a sufficient number of the protons (or “cosmic rays”) heated and randomized at internal shocks escape the radio core into the interstellar medium of the host galaxy, they may profoundly affect the evolution of the galaxy and its black hole."
 To quantify them effect on the wdrostatic balance of the ealactic gas we define an Eddington limit iu cosmic ravs for the host galaxy., To quantify their effect on the hydrostatic balance of the galactic gas we define an Eddington limit in cosmic rays for the host galaxy.
 For a shenomenologicallyanotivated jet model. we show that diving powerful radio-loud phases the power iu 1-100 GeV interstellar cosiic rays is large enough to break the cosmic rav Eddinetou nuit for the host ealaxy. so that balance of the galactic gaseous component is lost and a cosmic rav-driven wind develops. that removes the galactic eas.," For a phenomenologically-motivated jet model, we show that during powerful radio-loud phases the power in 1-100 GeV interstellar cosmic rays is large enough to break the cosmic ray Eddington limit for the host galaxy, so that balance of the galactic gaseous component is lost and a cosmic ray-driven wind develops, that removes the galactic gas."
 Ifthis super-Eddiugtou cosmic rav activity lasts lone enough. the time-inteerated cosmic rav input into the interstellar medi may excecd the binding euergv of the eas. aud the whole ealaxws gaseous phase will become uubound.," If this super-Eddington cosmic ray activity lasts long enough, the time-integrated cosmic ray input into the interstellar medium may exceed the binding energy of the gas, and the whole galaxy's gaseous phase will become unbound."
" Iu domme so. any fuel for further black hole growth aud star formation will be removed. thus affecting the combined evolution of the black hole aud its stellar bulec. as nupled bv the Af,AM, aud Af,σι relations."," In doing so, any fuel for further black hole growth and star formation will be removed, thus affecting the combined evolution of the black hole and its stellar bulge, as implied by the $M_\bullet-M_\star$ and $M_\bullet-\sigma_\star$ relations."
 report the detection of (1000kins 1j (110 Aspe) massive outflows of hydrogeu in several powerful radio galaxies.," report the detection of $\sim1000\unit{km\,s^{-1}}$ ) $\sim1-10\unit{kpc}$ ) massive outflows of hydrogen in several powerful radio galaxies."
 They claim that the outflows are spatially associated with radio knots extended along the jet. but this may just result frou the fact that II I absorption cau be traced ouly in the preseuce of a bright background radio contimuun.," They claim that the outflows are spatially associated with radio knots extended along the jet, but this may just result from the fact that H I absorption can be traced only in the presence of a bright background radio continuum."
 Fast large-scale outflows of low-ionization species with nearlv morphology have been detected in the powerful radio source MRC 1138-262 bv?., Fast large-scale outflows of low-ionization species with nearly morphology have been detected in the powerful radio source MRC 1138-262 by.
" The fact that the outflow is “cold” (atomic or weakly ionized) aud ""fast? (super-virial) iuplies that it is the result of catastrophic momentum. rather than euerev. exchauge with the interstellar medi."," The fact that the outflow is “cold” (atomic or weakly ionized) and “fast” (super-virial) implies that it is the result of catastrophic momentum, rather than energy, exchange with the interstellar medium."
 Iu other words. if is possible that some sort of Eddiuston Huit is being broken.," In other words, it is possible that some sort of Eddington limit is being broken."
" Among the 1iomieutuu-driven feedback mechauisuis. conclude that neither radiation-powoered quasar winds nor direct coupling between the jet aud the interstellar medium in the ""dentist drill”modol(?7) seem sufficieut"," Among the momentum-driven feedback mechanisms, conclude that neither radiation-powered quasar winds nor direct coupling between the jet and the interstellar medium in the “dentist drill”model seem sufficient"
Botl the far infrared. luminosity and the Ha Iuminosity are thought to be good racers of star formation (e.g. Young et al.,Both the far infrared luminosity and the $\alpha$ luminosity are thought to be good tracers of star formation (e.g. Young et al.
 1996)., 1996).
 The eloba Lyyr/Lye ratio lor NGC. E110. uncor'ected for extinction. is 67. soniewhat lower that typical for star-forming galaxies (e.g.. Devereux Young 1990: Yo ige al.," The global $_{FIR}$ $_{H\alpha}$ ratio for NGC 4410, uncorrected for extinction, is 67, somewhat lower than typical for star-forming galaxies (e.g., Devereux Young 1990; Young et al."
 1996)., 1996).
 However. usiie only he observed Ha luminosity from young stars gives Lian/bg ~ 100 — 230. consistent witli ratios [οιud for star forming galaxies (Devereux Young 1990: Yo1& el al.," However, using only the observed $\alpha$ luminosity from young stars gives $_{FIR}$ $_{H\alpha}$ $\sim$ 100 $-$ 230, consistent with ratios found for star forming galaxies (Devereux Young 1990; Young et al."
 1996. after correcting for clile‘ent definitious of [ar-iufrared. luminosity).," 1996, after correcting for different definitions of far-infrared luminosity)."
 This similarity witl star-formiug galaxies swweests tha lassive stars cloninuate the heating «of dust in this system ra1her than the active nuclei or slocks., This similarity with star-forming galaxies suggests that massive stars dominate the heating of dust in this system rather than the active nuclei or shocks.
 It. also unplies that the Ha extinction is uot high. cousistett with what we find οι1 the opticaT spectroscopy.," It also implies that the $\alpha$ extinction is not high, consistent with what we find from the optical spectroscopy."
" Obse'vatious in the 2.6 mm CO (1 — 0) line show that GC LOA is rich in molecular gas. with My, = 3.9 x 10? Abe (Smith 2000). asstming the standard Galactic {σου conversion factor."," Observations in the 2.6 mm CO (1 $-$ 0) line show that NGC 4410A is rich in molecular gas, with $_{H_2}$ = 3.9 $\times$ $^9$ $\sun$ (Smith 2000), assuming the standard Galactic $_{CO}$ $_{H_2}$ conversion factor."
 Compajug the Ha luminosity [ro1) yOu18oO stars to this mass gives Ly Νο = 0.001 — 0.000 ΕΝ:. at the low end of the rane[n]e for noanal spirals aud. irregulars (Young et al.," Comparing the $\alpha$ luminosity from young stars to this mass gives $_{H\alpha}$ $_{H_2}$ = 0.004 $-$ 0.009 $\sun$ $\sun$, at the low end of the range for normal spirals and irregulars (Young et al."
 1996)., 1996).
 Thus NGC 111Q is not forming stars at a partielarly high rate compared to the available amount of fuel for star formation. relative to noqual splals.," Thus NGC 4410 is not forming stars at a particularly high rate compared to the available amount of fuel for star formation, relative to normal spirals."
 The ο... ratio for NGC LL0A+B is ouly  1 Ls /M7 (init. 2000). also iuplvitD>oO a moderate star formation rate relative to the avallable molecular gas. at the low ed of the raiee for normal galaxies (Young et al.," The $_{FIR}$ $_{H_2}$ ratio for NGC 4410A+B is only $\sim$ 1 $\sun$ $\sun$ (Smith 2000), also implying a moderate star formation rate relative to the available molecular gas, at the low end of the range for normal galaxies (Young et al."
 1996)., 1996).
 Δι present. CO fluxes of only a haudful of radiο gaaxies lave been published to date (Mirabel et al.," At present, CO fluxes of only a handful of radio galaxies have been published to date (Mirabel et al."
 1989: Mazzarella et al., 1989; Mazzarella et al.
 1993: Evans e al., 1993; Evans et al.
 1999a.b: Li ret al.," 1999a,b; Lin et al."
 2000)., 2000).
" Many of these galaxies have Lere/ Mua ratios higher than that of NGC"" [410AB. by a factor of ~10.", Many of these galaxies have $_{FIR}$ $_{H_2}$ ratios higher than that of NGC 4410A+B by a factor of $\sim$ 10.
 The Àazzarella et al. (, The Mazzarella et al. (
1993) galaxies. however. were selectec ontje basis of heir ar-iufrared brightuesses. aud have luminosities au order of magnitucle higher than that of NGC LLLOA. so are 100 an unbiasecl comparison saliple for NGC 110 in terms of LiyeMgr.,"1993) galaxies, however, were selected on the basis of their far-infrared brightnesses, and have far-infrared luminosities an order of magnitude higher than that of NGC 4410A, so are not an unbiased comparison sample for NGC 4410 in terms of $_{FIR}$ $_{H_2}$."
 hi the raclio-selected sumον of Lin et al. (, In the radio-selected survey of Lin et al. (
2000). ouly twὉ of the 19 surveyed radio galaxies we'e cetected above the threshold ~10ML.,"2000), only two of the 19 surveyed radio galaxies were detected above the threshold $\sim 10^{8}~\rm{M}_\odot$."
 Compared to tat sample. NGC [110 lias abundant molecuar gas.," Compared to that sample, NGC 4410 has abundant molecular gas."
 The CO lie width in NCC ΕΠΑ is quite broad. with FWHMNM ~ 600 km  (Sinit1 2000).," The CO line width in NGC 4410A is quite broad, with FWHM $\sim$ 600 km $^{-1}$ (Smith 2000)."
 Based on a coiparisou of the ostical and CO velocities (Smith 2000). about 60% of his CO eiission may be associated with the H HI regions. while the reinainder arises from the nortweslert filamentary strcture.," Based on a comparison of the optical and CO velocities (Smith 2000), about $\%$ of this CO emission may be associated with the H II regions, while the remainder arises from the northwestern filamentary structure."
" Thus for tve eastern aud western portions of the ring. Ly /Mgy, e» 0.007 Γ.Ν: and ~ 0.003 ΕΝ:. Jesyectively."," Thus for the eastern and western portions of the ring, $_{H\alpha}$ $_{H_2}$ $\sim$ 0.007 $\sun$ $\sun$ and $\sim$ 0.003 $\sun$ $\sun$, respectively."
 The value for the western portion should be coisiclerec an upper liuit in comparing witli values derived for star fortline regious. as uch of the ionization may be due to shocks rather than young stars.," The value for the western portion should be considered an upper limit in comparing with values derived for star forming regions, as much of the ionization may be due to shocks rather than young stars."
 Iu contrast. the X-ray. emissiou [rom NGC £110 is provably dominated by the ΔΝ aud at extended thermal source unrelated to the star formation.," In contrast, the X-ray emission from NGC 4410 is probably dominated by the AGN and an extended thermal source unrelated to the star formation."
" The X-ray emission from NCC 111 Qi the ROSAT High Resolution Imager (HRI aps has been 'esolved into two compouents: a polit source with Ly = 2.7 x 104 eres tassociated with tre nucleus of NCC. L1I0A. aud an exteiclec (10"") halo with Ly = 1.3 x 10H erg sL οἱset towards Ixuo #2 (Tschókke et al."," The X-ray emission from NGC 4410 in the ROSAT High Resolution Imager (HRI) maps has been resolved into two components: a point source with $_{\rm X}$ = 2.7 $\times$ $^{41}$ erg $^{-1}$ associated with the nucleus of NGC 4410A, and an extended $''$ ) halo with $_{\rm X}$ = 1.3 $\times$ $^{41}$ erg $^{-1}$, offset towards Knot $\#$ 2 (Tschökke et al."
 1999)., 1999).
" Assuning the halo originates from the star foring regions. ii vas Lx /Ly, = 2.0."," Assuming the halo originates from the star forming regions, it has $_{\rm X}$ $_{H\alpha}$ = 2.0."
" For the center of GC IH0À. Ly, = L2 x 101"" erg 1 (Tables 1 and 3)). πο Lx flaw = 6.1."," For the center of NGC 4410A, $_{H\alpha}$ = 4.2 $\times$ $^{40}$ erg $^{-1}$ (Tables \ref{Table1} and \ref{Table3}) ), so $_{\rm X}$ $_{H\alpha}$ = 6.4."
 Pérrez-Olea Colina (1996), Pérrez-Olea Colina (1996)
(LOG sources detected in the 0.5-2.0 keV band down to a fux limit of 3.8« 1° ere 7 Lelunann ct al..,"(106 sources detected in the 0.5-2.0 keV band down to a flux limit of $\times$ $^{-16}$ erg $^{-2}$ $^{-1}$, Lehmann et al.,"
 2002) eives 13 reliable racio/X-ray associations., 2002) gives 13 reliable radio/X-ray associations.
 The percentages of the racdio/N-ray associations are — 21 per ceut of the racio sample aud ~ 12 per ceut of the NAIAL/X- salple., The percentages of the radio/X-ray associations are $\sim$ 21 per cent of the radio sample and $\sim$ 12 per cent of the XMM/X-ray sample.
" From an analysis of colour diagram (V-I versus LIX’) we divided the optical counterparts in two classes,", From an analysis of a colour–colour diagram (V-I versus $^{\prime}$ ) we divided the optical counterparts in two classes.
 The οἱjects in the first class are likely to be high redshittthe(z = jbjects0.5) carlytype. passively evolving galaxies. while ol in the second class can be both latetype. starformune galaxies at all redshifts aud carlytype ealaxies at low redshift (z < 0.5).," The objects in the first class are likely to be high redshift (z $\ge$ 0.5) early–type, passively evolving galaxies, while the objects in the second class can be both late--type, star–forming galaxies at all redshifts and early–type galaxies at low redshift (z $\le$ 0.5)."
 This separation in two classes appears to have a plivsical basis. related to different radio emission mechanisns. as further supported v the fact these two classes of objects lave a quite differeut distribution iu the radio fluxoptical iiaguitude diagram.," This separation in two classes appears to have a physical basis, related to different radio emission mechanisms, as further supported by the fact these two classes of objects have a quite different distribution in the radio flux–optical magnitude diagram."
 Iu particular. all the objects at large radio-to-optical ratios Geom 1000) have colours typical of passively evolving ealaxies at relatively high redshift.," In particular, all the objects at large radio-to-optical ratios $\it R$ $\ge$ 1000) have colours typical of passively evolving galaxies at relatively high redshift."
 Ou the basis of this plot we suegest that also the five radio sources without optical identification are likely to belong to this class., On the basis of this plot we suggest that also the five radio sources without optical identification are likely to belong to this class.
 No object of this class; iustead. appears iu this plot at OW values of R CR x 30). consisteut with findings (see. for exanuple. Camppioui et al.," No object of this class, instead, appears in this plot at low values of $\it R$ $\it R$ $\le$ 30), consistent with previous findings (see, for example, Gruppioni et al.,"
 1999) that previeneoljects with these values of R iare mainly ideutifie with starforming ealaxies., 1999) that objects with these values of R are mainly identified with star–forming galaxies.
 Frou this analysis based ou colours and the radio-to-optical ratios we conclude tha at least of the radio sources m a 5 GIIz selected sample with limiting fix S6Cl= 0.05 πι] is associated to carlytype galaxies., From this analysis based on colours and the radio-to-optical ratios we conclude that at least of the radio sources in a 5 GHz selected sample with limiting flux $_{\rm 6~cm}\geq$ 0.05 mJy is associated to early–type galaxies.
 A simular conclusion was reached by Cauppioni et al. (, A similar conclusion was reached by Gruppioni et al. (
1999) fora 1.1 Gz selected sample Solon2 0.2 11Jv). for which a substantial fraction of spectroscopic ideutifications was available.,"1999) for a 1.4 GHz selected sample $_{\rm 21cm}\geq$ 0.2 mJy), for which a substantial fraction of spectroscopic identifications was available."
 Oi data also show a significant correlation between LR colour and both I magnitude and radio-to-optica ratio ΠΠ. with the redder galaxies being associated with optically fainter radio sources aud with higher radio-to-optical ratio.," Our data also show a significant correlation between $^{\prime}$ colour and both I magnitude and radio-to-optical ratio $\it R$, with the redder galaxies being associated with optically fainter radio sources and with higher radio-to-optical ratio."
 Iu particular. amost all the counterparts of he radio sources in the magnitude range 22.5«1«21.5 are red oljects with LIN’ > 3 and a lieh fraction (6/10) of hese red objec sare EROs with LK’ > L.," In particular, almost all the counterparts of the radio sources in the magnitude range $<$ $<$ 24.5 are red objects with $^{\prime}$ $>$ 3 and a high fraction (6/10) of these red objects are EROs with $^{\prime}$ $>$ 4."
 By combining our radio data with existiug ISO data we conclude tha uost of our six racioselectec EROs sources are likely no | bo associated with dusty starburst ealaxies. but rather with carl-type Oogalaxies. hostiugc» an ACN respousible of he radio activity.," By combining our radio data with existing ISO data we conclude that most of our six radio–selected EROs sources are likely not to be associated with dusty starburst galaxies, but rather with early-type galaxies, hosting an AGN responsible of the radio activity."
 The six radio selected. EROs represent only ~ of the optically selected. EROs prescut in the field., The six radio selected EROs represent only $\sim$ of the optically selected EROs present in the field.
 If heir radio bhunuinositv is indeed a sign of ACN activity. he stuall fraction of radio detectious sugeests that the optically/near-infrared selected EROs population contains a relatively small fraction of active ACN. in agreement with what sueeested by recent deep ταν SHPVOYS (Toruschemeicr et al.," If their radio luminosity is indeed a sign of AGN activity, the small fraction of radio detections suggests that the optically/near-infrared selected EROs population contains a relatively small fraction of active AGN, in agreement with what suggested by recent deep X-ray surveys (Hornschemeier et al.,"
 2001)., 2001).
 The very high percentage of optical ideutifications aud he availability of near infrared and deep N-rav data make his suuple one of the best samples of racjo selected sources available at αν flux density level and well suited ο study in detail the nature of faiut 6 cmi radio sources., The very high percentage of optical identifications and the availability of near infrared and deep X-ray data make this sample one of the best samples of radio selected sources available at $\mu$ Jy flux density level and well suited to study in detail the nature of faint 6 cm radio sources.
 With this aim. we already obtained spectroscopic data for about one third of the radio sample. while a xograumie o obtain spectroscopic data for all the sources is ongoing.," With this aim, we already obtained spectroscopic data for about one third of the radio sample, while a programme to obtain spectroscopic data for all the sources is ongoing."
 The results of the optical spectroscopic classification and a 1n0re detailed comparison with the NMM source list will © presented in a forthcoming paper., The results of the optical spectroscopic classification and a more detailed comparison with the XMM source list will be presented in a forthcoming paper.
features a star formation rate density anc UV. emissivity that is higher than that of the fiducial model at z>3.,features a star formation rate density and UV emissivity that is higher than that of the fiducial model at $z>3$.
 The star-formation history in this model peaks at ο=5. and is Consistent with some of the higher determinations ol star-formation rate density above redshift 3.," The star-formation history in this model peaks at $z=5$, and is consistent with some of the higher determinations of star-formation rate density above redshift 3."
 3eCause this model provides maximum stellar emission. at. high redshift. and produces a stellar mass density exceeding most observations. this should be considered à somewhat extreme model resulting in maximal background. levels.," Because this model provides maximum stellar emission at high redshift, and produces a stellar mass density exceeding most observations, this should be considered a somewhat extreme model resulting in maximal background levels."
 This mocel converges to the fiducial model at. redshifts lower than 3. and opacities for gamma-ray sources closer than this are not significantly cillerent [rom the fiducial case.," This model converges to the fiducial model at redshifts lower than 3, and opacities for gamma-ray sources closer than this are not significantly different from the fiducial case."
 For our purposes the performance of the LAT instrument onferme can be described with a relatively small number of parameters., For our purposes the performance of the LAT instrument on can be described with a relatively small number of parameters.
 We take the LAT to have an cllective area of 9000 em? up to an energv of 300 GeV. While this may be an overestimate of the effective area at the highest energies. the number of photons seen at these energies is likely to be negligible.," We take the LAT to have an effective area of 9000 $^2$ up to an energy of 300 GeV. While this may be an overestimate of the effective area at the highest energies, the number of photons seen at these energies is likely to be negligible."
 Phe integrated Geld of view lorfermi is found to be approximately ~20500 em? se: we therefore assume a Ποιά of view 2050079000.z2.28 sr., The integrated field of view for is found to be approximately $\sim 20500$ $^2$ sr; we therefore assume a field of view $20500/9000 \approx 2.28$ sr.
 Lt is conservatively: assumed in our analysis that will be in survey mocde at all times. and that triggered rotations to view GRBs will not significantly raise the number of high-cncrey photons eathered.," It is conservatively assumed in our analysis that will be in survey mode at all times, and that triggered rotations to view GRBs will not significantly raise the number of high-energy photons gathered."
 The observations of GRBs by LACTs such as ALAGIC are highly sensitive to the capabilities of the instrument., The observations of GRBs by IACTs such as MAGIC are highly sensitive to the capabilities of the instrument.
 While these telescopes have much larger cllective collection areas than space-based instruments such as the LAL. other constraints such as the energy threshold. duty. evele. and time to respond to an alert must be taken into account in the analysis.," While these telescopes have much larger effective collection areas than space-based instruments such as the LAT, other constraints such as the energy threshold, duty cycle, and time to respond to an alert must be taken into account in the analysis."
 The much larger elfective area of these telescopes compared toferm is compensated by the relatively small probability that any sinele event will be observable., The much larger effective area of these telescopes compared to is compensated by the relatively small probability that any single event will be observable.
 The ollective area of LAC’Ps is a strong function. of energy in the subTeV regime. as the combined. elfects of the trigecr elliciency ancl the reconstruction cllicleney of Iow-enerev showers in the analysis leads to a sharp decrease in elfective area near threshold.," The effective area of IACTs is a strong function of energy in the sub–TeV regime, as the combined effects of the trigger efficiency and the reconstruction efficiency of low-energy showers in the analysis leads to a sharp decrease in effective area near threshold."
 The potential of detecting GRBs is strongly allected by the low energy. capabilities of the instrument. due to the rapidly increasing opacity of 10 universe to gamma ravs above a couple hundred: GeV for all but the closest. bursts.," The potential of detecting GRBs is strongly affected by the low energy capabilities of the instrument, due to the rapidly increasing opacity of the universe to gamma rays above a couple hundred GeV for all but the closest bursts."
" For. ALAGIC. we have used je ""after cuts’ form for the elfective area as a function of energv. [rom Albertctal.(2008b)... which is. defined for angles near zenith."," For MAGIC, we have used the `after cuts' form for the effective area as a function of energy from \citet{albert08a}, which is defined for angles near zenith."
 For observations away from zenith. the telescope loses sensitivity at. low energies due to the increasing amount of atmosphere through which the particle shower is being observed.," For observations away from zenith, the telescope loses sensitivity at low energies due to the increasing amount of atmosphere through which the particle shower is being observed."
 To estimate the change in the collective area function as a function of angle [rom zenith 8. we use the following simple approximation: where This expression is based. on a fit to the energy. threshold results at. lower angles from zenith (6%40 deg) presented in Firpo (2007).," To estimate the change in the effective area function as a function of angle from zenith $\theta$, we use the following simple approximation: where This expression is based on a fit to the energy threshold results at lower angles from zenith $\theta \lesssim 40$ deg) presented in \citet{firpoPHDT}. ."
". Additionally. we implement in our analysis an absolute cut at energy. Z5,;,. below which the sensitivity is zero and no gamma rays will be observed."," Additionally, we implement in our analysis an absolute cut at energy $E_{min}$, below which the sensitivity is zero and no gamma rays will be observed."
 We will perform our analysis for values of 50 and LOO GeV in Lyi). to allow for uncertaintv in low-cnereyv sensitivity ancl data reconstruction.," We will perform our analysis for values of 50 and 100 GeV in $E_{min}$, to allow for uncertainty in low-energy sensitivity and data reconstruction."
 Note that this parameter is distinct. from the instrument encrey threshold in that it is an absolute cut and is not dependent on the spectral properties of the source., Note that this parameter is distinct from the instrument energy threshold in that it is an absolute cut and is not dependent on the spectral properties of the source.
 Recent estimates of the instrument sensitivity for stereo. observations (Colinetal.2000) suggest that our more optimistic cut of 50 GeV may be reasonable for future observations., Recent estimates of the instrument sensitivity for stereo observations \citep{colin09} suggest that our more optimistic cut of 50 GeV may be reasonable for future observations.
 AX realistic estimate of the instrument duty. evele is critical for our analvsis., A realistic estimate of the instrument duty cycle is critical for our analysis.
 As outlined in Bastierietal.(2005).. there are several requirements for the operation of ALAGILC. including distance of the sun from zenith (>LOS deg). à minimum angular distance of the moon from the observation field (>30 deg). and humidity and wine requirements.," As outlined in \citet{bastieri05}, there are several requirements for the operation of MAGIC, including distance of the sun from zenith $>108$ deg), a minimum angular distance of the moon from the observation field $>30$ deg), and humidity and wind requirements."
 For the duty evele of the instrument. we use the standard value of 10 per cent. and note that this is supported by the fraction of GRBs (9.2 per cent) which were observed during follow-up in 42 months of observations (Ciarczarczvketal.2008)... and is similar to reports from observations with other ΑςἘν.," For the duty cycle of the instrument, we use the standard value of 10 per cent, and note that this is supported by the fraction of GRBs (9.2 per cent) which were observed during follow-up in 42 months of observations \citep{garczarczyk08}, and is similar to reports from observations with other IACTs."
 A major challenge for ground-based attempts to detect CRBs has been the highly transitory nature of the emission., A major challenge for ground-based attempts to detect GRBs has been the highly transitory nature of the emission.
 tesponding to an event requires the minimization of the individual components that contribute to the total delay ime. including the time for the detecting satellite to confirm a burst in progress. the time to transmit this information. and then the time for the ground-based telescope to slew to he coordinates and begin taking data.," Responding to an event requires the minimization of the individual components that contribute to the total delay time, including the time for the detecting satellite to confirm a burst in progress, the time to transmit this information, and then the time for the ground-based telescope to slew to the coordinates and begin taking data."
 The first quantity depends on the satellite response time and strength of the ourst., The first quantity depends on the satellite response time and strength of the burst.
 Fvpical numbers for the BAT are «15 sec to confirm and transmit à coordinate with a precision of a ew arcminutes., Typical numbers for the BAT are $<15$ sec to confirm and transmit a coordinate with a precision of a few arcminutes.
 Communications of these coordinates. are received by the 1οἘν in real time (~2 sec. Basticrietal. 2005)) over the GRB Coordinate Network(GCN)?.," Communications of these coordinates are received by the IACTs in real time $\sim 2$ sec, \citealp{bastieri05}) ) over the GRB Coordinate Network."
. Llowever. the final step of repositioning a telescope the size of MAGIC on the time scale required to view prompt CRB emission. presents à major engineering challenge.," However, the final step of repositioning a telescope the size of MAGIC on the time scale required to view prompt GRB emission presents a major engineering challenge."
 Ensuring »rsonnel safetv is. also a major practical concern in minimizing response time. which requires that the telescope » able to reposition without warning at any time during operation.," Ensuring personnel safety is also a major practical concern in minimizing response time, which requires that the telescope be able to reposition without warning at any time during operation."
" For our analvsis. we assume a delay time Lien, in wonmpt observations. which incorporates all three times iscussecl above."," For our analysis, we assume a delay time $T_{delay}$ in prompt observations, which incorporates all three times discussed above."
 We assume a typical report. time of 15 seconcls for the burst alert to reach the instrument. and. in 10 case of NLAGIC. a 30 see slew time to move to the target uxi begin observations.," We assume a typical report time of 15 seconds for the burst alert to reach the instrument, and, in the case of MAGIC, a 30 sec slew time to move to the target and begin observations."
 The total of 45 seconds is about equal to the lowest time. 43 sec. reported. in. Ciarczarczvketal.(2008). for all of NLACUIC. GRB responses to date. and is rerefore optimistic.," The total of 45 seconds is about equal to the lowest time, 43 sec, reported in \citet{garczarczyk08} for all of MAGIC GRB responses to date, and is therefore optimistic."
 To compute the [lux in the prompt phase of a GRB that can be seen by. eround-based observations after this delay time we use the {ου variable in BAT flux reported for our sample of bursts. which is the time elapsed. between the arrival of the first 5 and first 95 percent ofthe total GRB Iluence in the BAT band.," To compute the flux in the prompt phase of a GRB that can be seen by ground-based observations after this delay time we use the $T_{90}$ variable in BAT flux reported for our sample of bursts, which is the time elapsed between the arrival of the first 5 and first 95 percent ofthe total GRB fluence in the BAT band."
 The Iluence of the burst is modified by a factor, The fluence of the burst is modified by a factor
cosmnologies (Adelhergerefal 1998)). a result verified by N-body simulations (Wechslere£el1998: Isvavtsov&EKlvpiu 1998)).,"cosmologies \cite{adelberger98ap}) ), a result verified by $N$ -body simulations \cite{wechsler98ap}; \cite{kravtsov98ap}) )."
" This fact is reflected in the results for hieh-redslüft clustering of semi-anualvtic ealaxy formation models, which also find that galaxies at 2~3 are highly biased (Ikauffinann.Nusser. 1997: IWauffinaunefaf 1999: Somerville.Primack&Faber1998:: Baughe£a£.1998))."," This fact is reflected in the results for high-redshift clustering of semi-analytic galaxy formation models, which also find that galaxies at $z\sim 3$ are highly biased \cite{kauffmann97ap}; \cite{kauffmann98ap}; \cite{somerville98ap}; \cite{baugh98ap}) )."
 Firthermore. hydrodynamic simulations using simple criteria for galaxy formation (Evrard.Sunuuers.&Davis 1991: Katz.Heruquist.&Weinhere 1998: Cen&Ostriker1998a)) find that the distribution of ealaxies at 2=3 is indeed biased with respect to the mass distribution. aud that the clustering strength of ealaxies depends only weakly ou redshift.," Furthermore, hydrodynamic simulations using simple criteria for galaxy formation \cite{evrard94ap}; \cite{katz98ap}; \cite{cen98ap}) ) find that the distribution of galaxies at $z=3$ is indeed biased with respect to the mass distribution, and that the clustering strength of galaxies depends only weakly on redshift."
 While the issue of 5462) aud ry(2) is the one which observatious are currently best-suited to address. perhaps other statistics can coustrain the nature of galaxy formation more powerfully.," While the issue of $b_g(z)$ and $r_g(z)$ is the one which observations are currently best-suited to address, perhaps other statistics can constrain the nature of galaxy formation more powerfully."
 After all. the fundamental prediction of theories for galaxy formation. such as ours aud such as the scmi-analytic models mentioned above. is the location of star formation as a function of time.," After all, the fundamental prediction of theories for galaxy formation, such as ours and such as the semi-analytic models mentioned above, is the location of star formation as a function of time."
 Important information about ealaxy formation mav be lost if one examines only the inteeral quantities. which have been affected by the entire listory of galaxy formation couvolved with their subsequent eravitational evolution.," Important information about galaxy formation may be lost if one examines only the integral quantities, which have been affected by the entire history of galaxy formation convolved with their subsequent gravitational evolution."
 Therefore. we focus here on the evolutiou of the large-scale clustering of galaxy forination. defined as the formation of the ealaxvs constitucut stars that is. the evolution of 5h. andr...," Therefore, we focus here on the evolution of the large-scale clustering of galaxy , defined as the formation of the galaxy's constituent stars — that is, the evolution of $b_\ast$ and $r_\ast$."
 Now is au opportune time to investigate such questions. since large redshift aud aneular surveys which can examine this aud related properties of ealaxies are munuiment.," Now is an opportune time to investigate such questions, since large redshift and angular surveys which can examine this and related properties of galaxies are imminent."
 Iu this paper. we measure the clustering of galaxy formation in the bydrocdvuamical simulations of (1998a).. which we describe iu Section 1..," In this paper, we measure the clustering of galaxy formation in the hydrodynamical simulations of \cite{cen98a}, which we describe in Section \ref{simulations}."
 Tn Section 1.. we cousider the properties of galaxy ornmation in these simulations as a function of time. studving the first two effects discussed above aud he redshitt dependence of 5. aud r..," In Section \ref{timeformation}, we consider the properties of galaxy formation in these simulations as a function of time, studying the first two effects discussed above and the redshift dependence of $b_\ast$ and $r_\ast$."
 In Section 1.. we discuss the eravitational debiasing. the third effect ueutioued above.," In Section \ref{burst}, we discuss the gravitational debiasing, the third effect mentioned above."
 In Section 1.. we discuss observable effects of the trends in galaxy formation found in the sinulatious.," In Section \ref{observables}, we discuss observable effects of the trends in galaxy formation found in the simulations."
 Expecially iniportaut is the strong evolution of the cross-correlation between the star-formation weighted deusity field aud the galaxy density field: this evolution should be observable with surveys such as he Sloan Digital Skv Survey (SDSS: Cunu&Weinberg 19953)., Especially important is the strong evolution of the cross-correlation between the star-formation weighted density field and the galaxy density field; this evolution should be observable with surveys such as the Sloan Digital Sky Survey (SDSS; \cite{gunn95ap}) ).
 In addition. the evolution of the location of ealaxy formation in the universe nav be related to the observed Butcher-Ocuiler effect and may affect the nass-to-light ratios of clusters.," In addition, the evolution of the location of galaxy formation in the universe may be related to the observed Butcher-Oemler effect and may affect the mass-to-light ratios of clusters."
 We conclude in Section 1.., We conclude in Section \ref{conclusions}.
 The work of Ostriker&Steinhardt(1995) motivated the choice of a flat cold dark matter cosmology for the simulation used in this paper. with O4=0.37. O4=0.63. and O5=0.019.," The work of \cite{ostriker95a} motivated the choice of a flat cold dark matter cosmology for the simulation used in this paper, with $\Omega_0 = 0.37$, $\Omega_\Lambda = 0.63$, and $\Omega_b=0.049$."
 Recent observations of high redshift supernovae have lent support to the picture of a flat. low-density universe (Perlmuttere£alL997: CGarnavichcfaf 1998)).," Recent observations of high redshift supernovae have lent support to the picture of a flat, low-density universe \cite{perlmutter97ap}; \cite{garnavich98ap}) )."
 Creat uncertaiuty remains. however. and future work along the lines of this paper will need to address differeut cosimologics.," Great uncertainty remains, however, and future work along the lines of this paper will need to address different cosmologies."
 The Hubble coustant was set to Ly=100 5 lan ! |l with 5h=0.7.," The Hubble constant was set to $H_0 = 100$ $h$ km $^{-1}$ $^{-1}$, with $h=0.7$."
 The primordial perturbatious were adiabatic and raudom phase. with a power Προς slope of η=0.95 and amplitudes such that ox=0.5 for the dark matter at 2= 0. at which time the age of the universe is 12.7. Cyrs.," The primordial perturbations were adiabatic and random phase, with a power spectrum slope of $n = 0.95$ and amplitudes such that $\sigma_8=0.8$ for the dark matter at $z=0$ , at which time the age of the universe is 12.7 Gyrs."
 We use a periodic box 100 2.1 Mpe on a side. with 512° erid colls and 256° dark matter particles.," We use a periodic box 100 $h^{-1}$ Mpc on a side, with $512^3$ grid cells and $256^3$ dark matter particles."
 Thus. the dark matter mass resolution isabout 5«10? 5.1 AL. aud the erid cell size is ~ 0.2 51 Mpe.," Thus, the dark matter mass resolution isabout $5\times
10^9 $ $h^{-1}$ $M_\odot$ and the grid cell size is $\sim$ 0.2 $h^{-1}$ Mpc."
 The smallest smoothing leugth we consideris à l5+ Mpe radius top lat. which is considerably larger than a cell size.," The smallest smoothing length we consideris a 1 $h^{-1}$ Mpc radius top hat, which is considerably larger than a cell size."
 Ou these scales and leger. the relevant exavitational and," On these scales and larger, the relevant gravitational and"
the horizontal flow component of this form of gappy convection (Scharmer et al.,the horizontal flow component of this form of gappy convection (Scharmer et al.
 2008)., 2008).
 The striation pattern was already clear from observations made with the SST in 2003 (unpublished)., The striation pattern was already clear from observations made with the SST in 2003 (unpublished).
" In the following we analyze these observations, which show the overturning motions in great detail at a resolution approaching 0/11."," In the following we analyze these observations, which show the overturning motions in great detail at a resolution approaching 1."
 We compare them with overturning flows seen in the granulation around small isolated magnetic structures (pores)., We compare them with overturning flows seen in the granulation around small isolated magnetic structures (pores).
" In both cases, substructure in the form of striation is seen."," In both cases, substructure in the form of striation is seen."
 The observational properties of this striation are discussed., The observational properties of this striation are discussed.
" Its structure is interpreted as a corrugation (fluting) of the boundary between the penumbral magnetic field surrounding the gap and the flow inside it, and compared with the results from recent 3-D radiative MHD simulations."," Its structure is interpreted as a corrugation (fluting) of the boundary between the penumbral magnetic field surrounding the gap and the flow inside it, and compared with the results from recent 3-D radiative MHD simulations."
 Observations with the Hinode telescope reported by Ichimoto et al. (, Observations with the Hinode telescope reported by Ichimoto et al. (
2007) also show the striation pattern.,2007) also show the striation pattern.
" These authors did not give a clear interpretation, but nevertheless suggested that it is due to ‘twisting motions’ of the magnetic field, a view that also dominates later references."," These authors did not give a clear interpretation, but nevertheless suggested that it is due to `twisting motions' of the magnetic field, a view that also dominates later references."
 A more explicit interpretation was made by Zakharov et al. (, A more explicit interpretation was made by Zakharov et al. (
2008) using observations with the Swedish 1-m Solar Telescope (SST).,2008) using observations with the Swedish 1-m Solar Telescope (SST).
" They show that a twisting field explanation conflicts directly with the observed motion of the pattern, which is systematically away from the solar limb."," They show that a twisting field explanation conflicts directly with the observed motion of the pattern, which is systematically away from the solar limb."
" Whichever way a magnetic field might be twisting, it does not know about the solar limb seen by the observer."," Whichever way a magnetic field might be twisting, it does not know about the solar limb seen by the observer."
 They interpret the pattern sensibly in terms of a convective flow., They interpret the pattern sensibly in terms of a convective flow.
" As we shall argue below, a convection pattern is indeed indicated by the observations, but not in the form of the magnetic ‘rolls’ proposed by these authors."," As we shall argue below, a convection pattern is indeed indicated by the observations, but not in the form of the magnetic `rolls' proposed by these authors."
" Instead, the flow must be in the form of an at best weakly magnetic, overturning form of convection more akin to that seen in the granulation."," Instead, the flow must be in the form of an at best weakly magnetic, overturning form of convection more akin to that seen in the granulation."
 shows a part of the penumbra of a large symmetric spot., \\ref{stria} shows a part of the penumbra of a large symmetric spot.
" The field is from a time series of images taken on 2 May 2003 of the spot in AR351 (heliographic coordinates N8 E65, cos@z0.53)."," The field is from a time series of images taken on 2 May 2003 of the spot in AR351 (heliographic coordinates N8 E65, $\cos\theta\approx 0.53$ )."
 Image data were collected through a 430.5 nm G-band interference filter with a MegaPlus 1.6 1534 x 1024-pixel camera at 13 ms exposure., Image data were collected through a 430.5 nm G-band interference filter with a MegaPlus 1.6 1534 x 1024-pixel camera at 13 ms exposure.
" The camera was equipped with a phase diversity beamsplitter, so that half the detector recorded a conventional in-focus image, while the other half recorded a simultaneous but intentionally defocused image of the same field of view."," The camera was equipped with a phase diversity beamsplitter, so that half the detector recorded a conventional in-focus image, while the other half recorded a simultaneous but intentionally defocused image of the same field of view."
" The SST Adaptive Optics system was running and frame selection was used, storing the best 3 frames in 20-sec intervals."," The SST Adaptive Optics system was running and frame selection was used, storing the best 3 frames in 20-sec intervals."
" T'he residual seeing effects were further reduced by use of Joint Phase Diverse Speckle (Lóffdahl 2002) on the selected images, resulting in a sequence of restored images with an average cadence of 20 sec."," The residual seeing effects were further reduced by use of Joint Phase Diverse Speckle (Löffdahl 2002) on the selected images, resulting in a sequence of restored images with an average cadence of 20 sec."
 These images were corrected for image rotation and anisoplanatic warping., These images were corrected for image rotation and anisoplanatic warping.
 Well-defined bright filaments show a clear striation: dark lanes at angles of 10-45? to the filament axis., Well-defined bright filaments show a clear striation: dark lanes at angles of $\degr$ to the filament axis.
" Bright filaments are most clearly defined as individual structures in the inner part of the penumbra, and the striation is also seen most clearly there."," Bright filaments are most clearly defined as individual structures in the inner part of the penumbra, and the striation is also seen most clearly there."
" The field shown in refstria is centered on a wide, well resolved bright filament."," The field shown in \\ref{stria} is centered on a wide, well resolved bright filament."
" The time-dependence of this striation is illustrated as a ‘time slice’ in 2,, showing the intensity as a function of time along a slit taken perpendicular to the filaments."," The time-dependence of this striation is illustrated as a `time slice' in \ref{perpslit}, showing the intensity as a function of time along a slit taken perpendicular to the filaments."
 The substructure seen crossing the filaments in this image corresponds to a proper motion of about 2 km/s. 'The motion of the striation can also be seen by taking a slit along a single filament., The substructure seen crossing the filaments in this image corresponds to a proper motion of about 2 km/s. The motion of the striation can also be seen by taking a slit along a single filament.
 This is illustrated by the time slice shown in refparrslice.., This is illustrated by the time slice shown in \\ref{parrslice}.
" On the right (umbral) side of the image, it shows the well-know inward penetration of filament heads."," On the right (umbral) side of the image, it shows the well-know inward penetration of filament heads."
" The striation, seen most clearly from x=1.5 to x=3”, shows an outward motion (away from the umbra)."," The striation, seen most clearly from $x=1.5$ to $x=3$, shows an outward motion (away from the umbra)."
 This is also striking as a visual impression in a movie of the time series., This is also striking as a visual impression in a movie of the time series.
 The apparent speed ranges from 0.7 to 3 km/s. The pattern seen in 2 is also seen in the data from the Hinode satellite in Ichimoto et al. (, The apparent speed ranges from 0.7 to 3 km/s. The pattern seen in \ref{perpslit} is also seen in the data from the Hinode satellite in Ichimoto et al. (
2007).,2007).
 The systematic motion away from the limb at all azimuthal angles and disk center positions excludes the interpretation in terms of ‘twisting motions’ by Ichimoto et al. (, The systematic motion away from the limb at all azimuthal angles and disk center positions excludes the interpretation in terms of `twisting motions' by Ichimoto et al. (
2007) and by later authors.,2007) and by later authors.
 As pointed out in Zakharov et al. (, As pointed out in Zakharov et al. (
"2008) the systematic flow away from the limb, instead, shows it to be a convective flow.","2008) the systematic flow away from the limb, instead, shows it to be a convective flow."
 These authors discuss it in terms of a ‘convective roll’., These authors discuss it in terms of a `convective roll'.
" As we shall argue presently (see also Paper LII), a simpler and more convincing interpretation is in terms of an"," As we shall argue presently (see also Paper I,II), a simpler and more convincing interpretation is in terms of an"
"aand pphotoionization rates [yr and ΓἨοπ, the column density ratio, for a hydrogen to helium number density ratio of 12.9, is where the last form is accurate to 5 per cent.","and photoionization rates $\Gamma_{\rm HI}$ and $\Gamma_{\rm HeII}$, the column density ratio, for a hydrogen to helium number density ratio of 12.9, is where the last form is accurate to 5 per cent."
 for 104<10* K (??)..," for $10^4<T<5\times10^4$ K \citep{MM94, 2009RvMP...81.1405M}."
 The optical depth ratio is thus a measure of the fluctuations in the ratio of the radiation field as well as the velocity-broadening of the absorption lines., The optical depth ratio is thus a measure of the fluctuations in the ratio of the radiation field as well as the velocity-broadening of the absorption lines.
 'The distribution of the optical depth ratio per pixel is shown in Fig., The distribution of the optical depth ratio per pixel is shown in Fig.
 17 for 2«z3., \ref{fig:fluxdist_HeH_EnzoRT} for $2<z<3$.
" The distributions show a wide range of values, with the peak declining from 105 at to 75 at z=2."," The distributions show a wide range of values, with the peak declining from 105 at $z=3$ to 75 at $z=2$."
" The FWHM of the distributions is about 30, changing little with redshift."," The FWHM of the distributions is about 30, changing little with redshift."
 The variations arise entirely from the radiative transfer of the incident radiation through an inhomogeneous medium following helium reionization., The variations arise entirely from the radiative transfer of the incident radiation through an inhomogeneous medium following helium reionization.
 Allowing for the Poisson fluctuations from the finite number of sources within an attenuation volume for Herr--ionizing A," Allowing for the Poisson fluctuations from the finite number of sources within an attenuation volume for -ionizing photons will further enhance the fluctuations \citep{1998AJ....115.2206F, BHVC06, 2009RvMP...81.1405M,
 2010ApJ...714..355F}."
"bsorption lines were fit in the sspectrum at z=2.5, and matched to the ffeatures, requiring a line-centre velocity difference smaller than 10kms! for a successful match."," Absorption lines were fit in the spectrum at $z=2.5$, and matched to the features, requiring a line-centre velocity difference smaller than $10\kms$ for a successful match."
 The relation between the matching line centre optical depths is shown in Fig. 18.., The relation between the matching line centre optical depths is shown in Fig. \ref{fig:tau0HI_tau0HeII_z2p5}.
" While the deblending of absorption features will introduce scatter into the relation, a clear correlation is apparent in the upper left panel corresponding to To,Herr/To,H1ὦ100."," While the deblending of absorption features will introduce scatter into the relation, a clear correlation is apparent in the upper left panel corresponding to $\tau_{0, {\rm HeII}}/ \tau_{0, {\rm HI}}\lsim100$."
" The ratio is found to decrease towards increasing 70,H1 (lower left panel)."," The ratio is found to decrease towards increasing $\tau_{0, {\rm HI}}$ (lower left panel)."
 This arises as the ffeatures become saturated., This arises as the features become saturated.
" A similar trend was found for simulated spectra by ? for absorption lines in the redshift range 2.30ὦz< 2.75, who argued statistical comparisons must be restricted to systems with τοηι«0.1."," A similar trend was found for simulated spectra by \citet{2006A&A...455...91F}
 for absorption lines in the redshift range $2.30\lsim z\lsim 2.75$ , who argued statistical comparisons must be restricted to systems with $\tau_{0, {\rm HI}}<0.1$."
 'The distribution of column density ratios η=Ngen/Nui in Fig.," The distribution of column density ratios $\eta=N_{\rm HeII}/ N_{\rm
 HI}$ in Fig."
 18 shows a similar decreasing trend towards increasing Λι (top right panel)., \ref{fig:tau0HI_tau0HeII_z2p5} shows a similar decreasing trend towards increasing $N_{\rm HI}$ (top right panel).
 A broad distribution of 7 values for all the systems is found (lower right panel)., A broad distribution of $\eta$ values for all the systems is found (lower right panel).
" Restricting the systems to those with 0.01«Tour<0.1, however, results in a far sharper distribution centred at η~130."," Restricting the systems to those with $0.01<\tau_{0, {\rm HI}}<0.1$, however, results in a far sharper distribution centred at $\eta\simeq130$."
" The distributions are very similar to those measured (?),, with the distribution for systems restricted to 0.01«τοι«0.1 peaking at η<100 and somewhat broader than the distribution found from the RT simulation."," The distributions are very similar to those measured \citep{2006A&A...455...91F}, with the distribution for systems restricted to $0.01<\tau_{0, {\rm HI}}<0.1$ peaking at $\eta\lsim100$ and somewhat broader than the distribution found from the RT simulation."
 The peak in the distribution from the simulation is fixed by the normalisations imposed on the effective aand ooptical depths., The peak in the distribution from the simulation is fixed by the normalisations imposed on the effective and optical depths.
" The agreement in the widths of the distribution, however, is a product of radiative transfer within the simulation."," The agreement in the widths of the distribution, however, is a product of radiative transfer within the simulation."
" While the measured mean and breadth of the distribution may be reproduced by the QSO luminosity function, taking into account the Poisson fluctuations in QSO numbers within a Her--ionizing photon attenuation volume in a homogeneous IGM (?),, the results found here demonstrate that a large contribution to the breadth in the distribution is expected to arise from variations in the iionizing background resulting from the effect of inhomogeneities in the density field of the IGM on the radiation transfer of the impinging radiation field."," While the measured mean and breadth of the distribution may be reproduced by the QSO luminosity function, taking into account the Poisson fluctuations in QSO numbers within a -ionizing photon attenuation volume in a homogeneous IGM \citep{2009RvMP...81.1405M}, the results found here demonstrate that a large contribution to the breadth in the distribution is expected to arise from variations in the ionizing background resulting from the effect of inhomogeneities in the density field of the IGM on the radiation transfer of the impinging radiation field."
 The peak and FWHM in the distribution of 7 at z=2.5 from the simulation correspond to an effective V—Γπι/Γποιc 3007190.," The peak and FWHM in the distribution of $\eta$ at $z=2.5$ from the simulation correspond to an effective $\Psi=\Gamma_{\rm
 HI}/\Gamma_{\rm HeII} \simeq300^{+150}_{-100}$ ."
" For a metagalactic spectral shape fuvMG, the ratio would be ΓΗΙ/ΤΓΠοῃ~4!Μα, "," For a metagalactic spectral shape $f_\nu\sim\nu^{-\alpha_{\rm MG}}$, the ratio would be $\Gamma_{\rm HI}/\Gamma_{\rm HeII}\sim4^{1+\alpha_{\rm MG}}$."
"The match to the effective Ψ corresponds to effective spectral indices of 2.9<ama3.5, much larger than the assumed source spectral index 0.5."," The match to the effective $\Psi$ corresponds to effective spectral indices of $2.9<\alpha_{\rm MG}<3.5$, much larger than the assumed source spectral index $0.5$."
 The ratio corresponds to the expected softening of the incident spectrum due to radiative transfer through the IGM (??)..," The ratio corresponds to the expected softening of the incident spectrum due to radiative transfer through the IGM \citep{MM94, HM96}."
 The required rescaling of V imposed on the simulations to match the measured effective aand ooptical depths corresponds to a pre-rescaled ambient radiation field with an average ama~—0.25 betweenthe aand LLyman edges arising exclusively from the radiative transfer of the input sourcespectrum through the simulation volume., The required rescaling of $\Psi$ imposed on the simulations to match the measured effective and optical depths corresponds to a pre-rescaled ambient radiation field with an average $\alpha_{\rm MG}\simeq-0.25$ betweenthe and Lyman edges arising exclusively from the radiative transfer of the input sourcespectrum through the simulation volume.
 This corresponds to the hardening effect of the IGM on, This corresponds to the hardening effect of the IGM on
 Opt Opt Qnun «ταν = 160mm = 211Lhuin Opt 5pt Observations of stellar orbits in the proximity of the enigmatic radio source Ser À* located a very ανασα. ceiter of our galaxy have established that it is a { x10 ack hole (Reid&Brunthaler2001:Chezetal.2008:Cillessen2009).,"= 0pt = 0pt = 0mm = -17mm = 160mm = 244mm 0pt 5pt Observations of stellar orbits in the proximity of the enigmatic radio source Sgr A* located at the very dynamical center of our galaxy have established that it is a 4 $\times 10^6$ black hole \citep{reid04,ghez08,gillessen09}."
. The emission : A* consists of €uasi-steady aud variable components (e.g.Horsteinοἱal.2007:Docds-Ecle.2010:Yusel-Zadehetal. 2010).," The emission from Sgr A* consists of quasi-steady and variable components \citep[e.g.][]{hornstein07,dodds10,zadeh10}."
. The variable component is now detected in aluost all waveleuegt uxds aud is used o estimate the physical quantities of the gas low. the radiation mechanis erent wavelenetL bands aud the scale leneths at which variable emission o»erates.," The variable component is now detected in almost all wavelength bands and is used to estimate the physical quantities of the gas flow, the radiation mechanism in different wavelength bands and the scale lengths at which variable emission operates."
 One of t wzline aspect of he flaringOm activity of Sere A* is the physical 1ature of the variable eiissio11 iferent waveleugis and how its spectrum evolves with time., One of the puzzling aspect of the flaring activity of Sgr A* is the physical nature of the variable emission in different wavelengths and how its spectrum evolves with time.
 Fu‘(hermore. it is not clear wheher fla‘es are cleeply embedded in the accretion disk of Ser A* or are physically distiict from the regio Wwlere the bulk of he etission a‘ises (Yusel-ZadehetECharteal.2008:Trapet20! LQ).," Furthermore, it is not clear whether flares are deeply embedded in the accretion disk of Sgr A* or are physically distinct from the region where the bulk of the emission arises \citep{zadeh06a,zadeh08,zadeh09,marrone08,eckart08,trap10}."
 pare enmission at 1ear-IR. (IIR) wavelengths is due to optically thin svuchrot‘oO eimnission aud the slbium [requences js couside'ecd o be the dividiug line between optically ti emission at high frequencies (1ear-IBR). aud opically thick emission at lower frequencies (subir1//adio).," Flare emission at near-IR (IR) wavelengths is due to optically thin synchrotron emission and the submm frequency is considered to be the dividing line between optically thin emission at high frequencies (near-IR), and optically thick emission at lower frequencies (submm/radio)."
 To examiue ue possibility tliat there is a casal association between radio/submun variabiity aud IB. [lariug ‘livily. wesucied in detail he ight curves of flare emission [rom optically tulFR and thin plasla ‘here t1ere was sufficient overlap between IB/X-ray and submm/radio waveetlis.," To examine the possibility that there is a causal association between radio/submm variability and IR flaring activity, we studied in detail the light curves of flare emission from optically thick and thin plasma where there was sufficient overlap between IR/X-ray and submm/radio wavelengths."
 We searcled — ‘signatures of optical deptl effect against the background quiex'ent enmilssiol a tlie earliest pliase of the evolutiou of flares., We searched for signatures of optical depth effect against the background quiescent emission at the earliest phase of the evolution of flares.
"z=10, maintaining an almost constant value of a=—2.0.","$z=10$, maintaining an almost constant value of $\alpha=-2.0$."
 'This value is slightly larger than the one derived from the data (see Sec. 1)), This value is slightly larger than the one derived from the data (see Sec. \ref{introduction}) )
 at z—56 but perfectly consistent with the recent determination of a for the z=7—8 LF obtained by Bouwens et al. (, at $z=5-6$ but perfectly consistent with the recent determination of $\alpha$ for the $z=7-8$ LF obtained by Bouwens et al. (
2010b).,2010b).
" This behavior is most likely produced by a combination of the halo mass function evolution, feedback effects and evolving stellar populations."," This behavior is most likely produced by a combination of the halo mass function evolution, feedback effects and evolving stellar populations."
 Clearly the faint-end slope will be better constrained by forthcoming facilities such as the JWST; the above successful test of our model allows us to make reliable predictions for surveys that will be performed with such instruments., Clearly the faint-end slope will be better constrained by forthcoming facilities such as the JWST; the above successful test of our model allows us to make reliable predictions for surveys that will be performed with such instruments.
" For a deep exposure of 10° s, JWST is expected to reach a photometric sensitivity of about 1 nJy (10σ) allowing an investigation of the faint-end of the LF predicted by our simulations at least up to z=9 (see Fig. 1))."," For a deep exposure of $10^6$ s, JWST is expected to reach a photometric sensitivity of about $1$ nJy $10\sigma$ ) allowing an investigation of the faint-end of the LF predicted by our simulations at least up to $z=9$ (see Fig. \ref{fig:LF}) )."
" This will be particularly exciting because it will very likely unveil the physical properties of these objects which are now thought to be the main reionization sources, as we will discuss in detail in Sec. ??.."," This will be particularly exciting because it will very likely unveil the physical properties of these objects which are now thought to be the main reionization sources, as we will discuss in detail in Sec. \ref{Rei_sources}."
" We want to stress here that our simulation allow us to resolve properly (i.e. with more than 200 DM particles) the galaxy population that will be observed by JWST, being their total halo masses >10°Mo."," We want to stress here that our simulation allow us to resolve properly (i.e. with more than 200 DM particles) the galaxy population that will be observed by JWST, being their total halo masses $\ge 10^9\;\Msun$."
" This, toghether with the carefull threatment of the chemical feedback, makes our predictions particularly robust."," This, toghether with the carefull threatment of the chemical feedback, makes our predictions particularly robust."
 To give a visual idea of the high-z (z>7.5) galaxy population in a typical JWST field (of size 5.3 arcmin) we have produced the maps shown in Fig., To give a visual idea of the $z$ $z > 7.5$ ) galaxy population in a typical JWST field (of size 5.3 arcmin) we have produced the maps shown in Fig.
 2 which demonstrate the wealth of objects that JWST will be able to detect up to z211., \ref{maps} which demonstrate the wealth of objects that JWST will be able to detect up to $z \simgt 11$.
" The maps are a 2D cut of a 3D image produced by stacking all the available simulation snapshots between z~7.6—11.6, and filling in the gaps between snapshots by randomly rotating the orientation axes in all the three-dimensions simultaneously."," The maps are a 2D cut of a 3D image produced by stacking all the available simulation snapshots between $z \sim 7.6-11.6$, and filling in the gaps between snapshots by randomly rotating the orientation axes in all the three-dimensions simultaneously."
" As it has emerged from the overview of the observational results given in the Introduction, dust seems to have only a minor effect on shaping the SED of faint high-redshift galaxies."," As it has emerged from the overview of the observational results given in the Introduction, dust seems to have only a minor effect on shaping the SED of faint high-redshift galaxies."
" This is easily understood, as we will see later on, as a result of the relatively low stellar masses and metallicities characterizing these objects."," This is easily understood, as we will see later on, as a result of the relatively low stellar masses and metallicities characterizing these objects."
 Dust can however become more important in luminous Lyman Break Galaxies (LBGs) and Lyman Alpha Emitters (LAEs) as many studies have shown (Lai et al., Dust can however become more important in luminous Lyman Break Galaxies (LBGs) and Lyman Alpha Emitters (LAEs) as many studies have shown (Lai et al.
 2007; Atek et al., 2007; Atek et al.
" 2008; Nagamine, Zhang Hernquist 2008; Dayal et al."," 2008; Nagamine, Zhang Hernquist 2008; Dayal et al."
" 2008, 2009, 2010, 2011; Finkelstein et al."," 2008, 2009, 2010, 2011; Finkelstein et al."
 2009)., 2009).
 The dust enrichment of galaxies can be followed self-consistently by post-processing our simulation through the model introduced by Dayal et al. (, The dust enrichment of galaxies can be followed self-consistently by post-processing our simulation through the model introduced by Dayal et al. (
"2010), using which, we can quantitatively check the previous statement.","2010), using which, we can quantitatively check the previous statement."
" In brief, the model assumes that dust is produced by supernovae (we neglect here the contribution of AGB stars, though see Valiante et al. ("," In brief, the model assumes that dust is produced by supernovae (we neglect here the contribution of AGB stars, though see Valiante et al. ("
2009) who found a significant contribution,2009) who found a significant contribution
to condensate.,to condensate.
 In Fig., In Fig.
 5. for A=0.999.," 5, for $\lambda=0.999$."
 we demonstrate the asviuptotic ποα istributio1., we demonstrate the asymptotic wealth distribution.
 It clearly displays a power-law with expoucut z1.5. The siutlation Was carriec out for N=100. f=10° iterations and averaged over 310° initia onditious.," It clearly displays a power-law with exponent $\simeq 1.5$, The simulation was carried out for $N=100$, $t= 10^6$ iterations and averaged over $\times 10^3$ initial conditions."
 This model has a tunable parameter A aud only for values cose to nity we oserve a power law behavior., This model has a tunable parameter $\lambda$ and only for values close to unity we observe a power law behavior.
 Ilowever. evervone las a different appetite for risk.," However, everyone has a different appetite for risk."
 We attempt a inode oeji which Αν have a quenched raudon distribution., We attempt a model in which $\lambda_i$ 's have a quenched random distribution.
 We consider a uniforiu istributio iof A As in previous case. we find the saturation time £f. at which Iuaxiumui wealth saurates.," We consider a uniform distribution of $\lambda$ 's. As in previous case, we find the saturation time $t_c$ at which maximum wealth saturates."
 Iu Fig., In Fig.
 6. we have potted £N) as a function of N for this model.," 6, we have plotted $t_c(N)$ as a function of $N$ for this model."
 T1ο saturation time f.0N) scaes with number of agents as LIN)zmaNT with a~210 and bz1.561., The saturation time $t_c(N)$ scales with number of agents as $t_c(N)\simeq a N^b$ with $a\simeq 240$ and $b\simeq 1.561$.
 It is interesting to note that the steady state of wealth distribution has a power-law tail with »z2.0 as that shown in Fie., It is interesting to note that the steady state of wealth distribution has a power-law tail with $\nu\simeq 2.0$ as that shown in Fig.
 7., 7.
 Wowever. in the region correspondiug to low wealth. the wealth distribution is found to be exponential.," However, in the region corresponding to low wealth, the wealth distribution is found to be exponential."
 The inset of Fig., The inset of Fig.
 7 show that behavior., 7 show that behavior.
 Tu this strategy Pareo exponent is 21.0., In this strategy Pareto exponent is $\simeq 1.0$.
 This strategy seen to be more realistic as compared to ποια A previously discussed aud it also gives expoucuts which are colparable ο realislc Case., This strategy seem to be more realistic as compared to model A previously discussed and it also gives exponents which are comparable to realistic case.
 Iu model C. the transaction is carried on depeudiug on the ageuκ choice who chooses to lay wit LYS or TF model with certain probability.," In model C, the transaction is carried on depending on the agent's choice who chooses to play with YS or TF model with certain probability."
 Now some transactions wil follow YS model aud asset exchauge iu rest of the transactions will be decided by TE model., Now some transactions will follow YS model and asset exchange in rest of the transactions will be decided by TF model.
 We define the saturation time tC) by lookiug at wealth of richest agent as done in previous cases;, We define the saturation time $t_c(N)$ by looking at wealth of richest agent as done in previous cases.
 We observe that with ax16.52 and b~1.56 in this case. (, We observe that $t_c(N)\simeq a N^b $ with $a\simeq 16.82$ and $b\simeq 1.56$ in this case. (
See Fig.,See Fig.
 8.), 8.)
 The asvauptotic distribution of wealth shown in Fie., The asymptotic distribution of wealth shown in Fig.
 9., 9.
 It has a power-law tail with expoucut vox3.7., It has a power-law tail with exponent $\nu\simeq 3.7$.
 In this case Pareto expoucut is z2.7 which is comparable to one observed in societies like [20] We have checked that the power law is a better fit than lognormal for all the cases discussed in the paper in several wavs., In this case Pareto exponent is $\simeq 2.7$ which is comparable to one observed in societies like \cite{italy} We have checked that the power law is a better fit than lognormal for all the cases discussed in the paper in several ways.
 We have checked it visually., We have checked it visually.
 We have checked the eooduess offit by finding \7/DoF aud Π for που models by fitting it a power law functional form aud lognormal fit., We have checked the goodness of fit by finding $\chi^2/DoF$ and $R^2$ for three models by fitting it a power law functional form and lognormal fit.
 The values are given i Table 1., The values are given in Table 1.
" It is ceur tha ReD values are higherB aud very close to unity+ for. fit which shows that this mode is relevant for higher fraction of data and v/DoF2 values are lower for. a power law fit which. shows that error is. πια]er in this fit in al Cases,", It is clear that $R^2$ values are higher and very close to unity for power-law fit which shows that this model is relevant for higher fraction of data and $\chi^2/DoF$ values are lower for a power law fit which shows that error is smaller in this fit in all cases.
 For models B and € where Pareο exponent is more than oue. we lave also," For models B and C where Pareto exponent is more than one, we have also"
"not be spatially sufficiently resolved and are not considered further, leaving 16748 field lines (60% of all) to be studied.","not be spatially sufficiently resolved and are not considered further, leaving 16748 field lines $60 \%$ of all) to be studied."
 Figure 2aa shows a randomly chosen fraction of 2% of these loops and Fig., Figure \ref{fig2}a a shows a randomly chosen fraction of $2\%$ of these loops and Fig.
 2bb a histogram of the loop height distribution for all resolved loops., \ref{fig2}b b a histogram of the loop height distribution for all resolved loops.
 The average loop height is 1.24+2.45 Mm!., The average loop height is $1.24 \pm 2.45$ Mm.
. The number of small loops (below about 1 Mm) is significantly larger than that of higher loops., The number of small loops (below about 1 Mm) is significantly larger than that of higher loops.
" 51% of the resolved loops close within phototospheric heights («500km) and 2996 in the chromosphere (500—2500km). Finally 20% of the loops reach into the corona (>2500km), but half of these long coronal loops do not close within the Sunrise FOV."," $51 \%$ of the resolved loops close within phototospheric heights $(<500 \, {\rm km})$ and $29 \%$ in the chromosphere $(500-2500 \, {\rm km})$ Finally $20 \%$ of the loops reach into the corona $(>2500 \, {\rm km})$ , but half of these long coronal loops do not close within the Sunrise FOV."
 We also investigated how using a linear force-free model instead of a potential field model influences the loop height statistics., We also investigated how using a linear force-free model instead of a potential field model influences the loop height statistics.
" These models contain the force-free parameter a, which cannot be deduced from the available data."," These models contain the force-free parameter $\alpha$, which cannot be deduced from the available data."
 Consequently one unique linear force-free model cannot be computed., Consequently one unique linear force-free model cannot be computed.
" To investigate the influence of linear electric currents, we carried out computations with aL=+4, respectively, where L=37 Mm is the size of the employed portion of the magnetogram (see dotted and dashed lines in Fig."," To investigate the influence of linear electric currents, we carried out computations with $\alpha L= \pm 4$, respectively, where $L=37$ Mm is the size of the employed portion of the magnetogram (see dotted and dashed lines in Fig."
" 2bb, respectively)."," \ref{fig2}b b, respectively)."
 4 is already quite a large value for the normalized force-free parameter al and close to the mathematical allowed maximum value of |aL|=V2z (seediscussioninSeehaferforce-free models)..," $4$ is already quite a large value for the normalized force-free parameter $\alpha L$ and close to the mathematical allowed maximum value of $|\alpha L|= \sqrt{2} \pi$ \citep[see discussion in][for a
suitable rank of $\alpha$-values for linear force-free
models]{seehafer78}."
 As visible in Fig., As visible in Fig.
 2bb the loop height statistics are almost identical to those obtained from potential fields., \ref{fig2}b b the loop height statistics are almost identical to those obtained from potential fields.
 Small differences occur only for field lines higher than about 5Mm., Small differences occur only for field lines higher than about 5Mm.
" For our study, which mainly concentrates on lower chromospheric loops, it is therefore justified to consider only potential fields."," For our study, which mainly concentrates on lower chromospheric loops, it is therefore justified to consider only potential fields."
" Figure 2cc contains a scatter plot of loop height vs. magnetic field strengths at the leading foot point, i.e. the foot point with the largest field strength."," Figure \ref{fig2}c c contains a scatter plot of loop height vs. magnetic field strengths at the leading foot point, i.e. the foot point with the largest field strength."
" It shows that loops of any height can begin from weak fields, but only long loops start from strong-field foot points."," It shows that loops of any height can begin from weak fields, but only long loops start from strong-field foot points."
This means that,This means that
The optical light curve of IRAS 13224-3809 is consistent with the short time-scale variability of the source being less than | per cent.,The optical light curve of IRAS 13224-3809 is consistent with the short time-scale variability of the source being less than 1 per cent.
 We ure unable to comment on the wavelength dependence of any optical variability., We are unable to comment on the wavelength dependence of any optical variability.
 Our aperture. however. also collects flux from the underlying galaxy which dilutes the signal from the nucleus.," Our aperture, however, also collects flux from the underlying galaxy which dilutes the signal from the nucleus."
 As may be seen from Fig., As may be seen from Fig.
 | the three main components of the flux received within a radius of 10 pixels. namely the contribution of the background. the underlying galaxy and the nucleus may be separated.," 1 the three main components of the flux received within a radius of 10 pixels, namely the contribution of the background, the underlying galaxy and the nucleus may be separated."
 If the nucleus is assumed to be a Gaussian point source and we (over-) estimate the contribution from the galaxy then the variability of the nucleus may be constrained., If the nucleus is assumed to be a Gaussian point source and we (over-) estimate the contribution from the galaxy then the variability of the nucleus may be constrained.
 It is found that a | per cent variation of the entire source may correspond to a ~2 per cent variation in the nucleus., It is found that a 1 per cent variation of the entire source may correspond to a $\sim 2$ per cent variation in the nucleus.
" The continuum optical spectrum of IRAS 13224-3809 (Boller et al 1993) has a photon index of approximately Εν,~0 across the optical waveband.", The continuum optical spectrum of IRAS 13224-3809 (Boller et al 1993) has a photon index of approximately $\Gamma_{\rm o}\sim 0$ across the optical waveband.
 The V band magnitude of the nucleus is 15.2. corresponding to a vi. flux in the V band of 1.5«10.++ erg 7s tat5800À..," The V band magnitude of the nucleus is 15.2, corresponding to a $\nu F_\nu$ flux in the V band of $1.5\times10^{-11}$ erg $^{-2}$ $^{-1}$ at."
 Simultaneous X-ray data are not available but if the 30 day ROSAT light curve of Boller et al (1997) is assumed to be typical of IRAS 13224-3809 it is extremely unlikely that the X-ray flux did no at least double during the period of our observations., Simultaneous X-ray data are not available but if the 30 day ROSAT light curve of Boller et al (1997) is assumed to be typical of IRAS 13224-3809 it is extremely unlikely that the X-ray flux did not at least double during the period of our observations.
 This ligh curve is shown in Fig., This light curve is shown in Fig.
 6 and there is no set of three consecutive days during which the X-ray flux did not change considerably., 6 and there is no set of three consecutive days during which the X-ray flux did not change considerably.
 The only period of relative “quiescence” is beween 8-IO days. and even during this the X-ray flux at least de»ubles.," The only period of relative `quiescence' is between 8–10 days, and even during this the X-ray flux at least doubles."
 A more recen HRI monitoring campaign has contirmed tje continued existence of such extreme variability (private communication)., A more recent HRI monitoring campaign has confirmed the continued existence of such extreme variability (private communication).
" The X-ray photon index of IRAS 13224-3809 was observed to be very high. Py~4. and the average HRI count rate was counts * (Boller et al 1997) corresponding to a (4/7, flux of 2.9.10 lenem τς tat 1.25 keV. Optical emission can originate from the soft X-ray emission region. or very close to it. by several means."," The X-ray photon index of IRAS 13224-3809 was observed to be very high, $\Gamma_{\rm x}\sim 4$, and the average HRI count rate was $4.7\times10^{-3}$ count $^{-1}$ (Boller et al 1997) corresponding to a $\nu F_\nu$ flux of $2.9\times10^{-13}$ erg $^{-2}$ $^{-1}$ at 1.25 keV. Optical emission can originate from the soft X-ray emission region, or very close to it, by several means."
 It may be emitted by the same electrons that produce the X-rays. by cyclotron or synchrotron emission or Comptonization (perhaps of photons produced by the eyelo-synchrotron process: see e.g. Di Matteo. Celotti Fabian 1997).," It may be emitted by the same electrons that produce the X-rays, by cyclotron or synchrotron emission or Comptonization (perhaps of photons produced by the cyclo-synchrotron process; see e.g. Di Matteo, Celotti Fabian 1997)."
 Comptonization is however unlikely to make a large contribution to the observed flux since the steep X- spectrum implies a low Compton y-parameter which gives a similarly steep optical spectrum., Comptonization is however unlikely to make a large contribution to the observed flux since the steep X-ray spectrum implies a low Compton $y$ -parameter which gives a similarly steep optical spectrum.
 Optical emission may also originate from reprocessing of the X-ray flux. and therefore would have a thermal spectrum (probably from a range of temperatures).," Optical emission may also originate from reprocessing of the X-ray flux, and therefore would have a thermal spectrum (probably from a range of temperatures)."
 Finally it may be triggered by whatever mechanism causes the ray variability but not share the same emitting electrons as the rays., Finally it may be triggered by whatever mechanism causes the X-ray variability but not share the same emitting electrons as the X-rays.
 The X-ray flux is expected to have a least varied by a factor of two during our observations., The X-ray flux is expected to have at least varied by a factor of two during our observations.
 Let us assume that the X-ray flux doubled at some point and consider the implications of the lack of optical variability., Let us assume that the X-ray flux doubled at some point and consider the implications of the lack of optical variability.
 A lack of variability in the optical band can be explained in several ways., A lack of variability in the optical band can be explained in several ways.
 It may simply be due to most. more than 98 per cent. of the optical emission emerges from a region unconnected with that producing the soft X-ray flux.," It may simply be due to most, more than 98 per cent, of the optical emission emerges from a region unconnected with that producing the soft X-ray flux."
 Note however that the average soft X-ray flux (below | keV) is comparable to the optical flux and, Note however that the average soft X-ray flux (below 1 keV) is comparable to the optical flux and
is surrounded by a compact structure which would generate an extinction Ay of about 90 magnitudes if the dust-to-gas ratio in that structure were similar to that of the interstellar medium.,is surrounded by a compact structure which would generate an extinction $A_V$ of about 90 magnitudes if the dust-to-gas ratio in that structure were similar to that of the interstellar medium.
" They interpret that structure as a nearly edge-on, compact (Rk ~ 2-3 R.) circumstellar disk."," They interpret that structure as a nearly edge-on, compact $R$ $\sim$ 2–3 $R_\odot$ ) circumstellar disk."
 It is noteworthy that EC 95a and T Tau Sa are similar in several respects., It is noteworthy that EC 95a and T Tau Sa are similar in several respects.
" Both are most likely young Herbig AeBe stars orbited by a low-mass companion (T Tau Sb, and EC 95b, respectively) located at about 10 AU."," Both are most likely young Herbig AeBe stars orbited by a low-mass companion (T Tau Sb, and EC 95b, respectively) located at about 10 AU."
" Thus, in both cases, the circumstellar disks are expected to be üdally truncated to about 2-3 AU."," Thus, in both cases, the circumstellar disks are expected to be tidally truncated to about 2–3 AU."
" The rovibrational CO lines in EC 95 are somewhat fainter than those in T Tau Sa, and the K-band veiling of EC 95 (ry — 0.1: Doppmann et 22005) is also smaller than that ofT Tau Sa (ry ~ 2; Duchénne et 22003)."," The rovibrational CO lines in EC 95 are somewhat fainter than those in T Tau Sa, and the K-band veiling of EC 95 $r_K$ $\sim$ 0.1; Doppmann et 2005) is also smaller than that of T Tau Sa $r_K$ $\sim$ 2; Duchênne et 2003)."
" Thus, the amount ofcircumstellar material around EC 95a must be significantly smaller than that around T Tau Sa."," Thus, the amount of circumstellar material around EC 95a must be significantly smaller than that around T Tau Sa."
" Still. à scaled-down version of the compact disk surrounded T Tau Sa, consistent with the fainter CO lines and smaller veiling of EC 95, might well provide the 20 magnitudes of exUnction required to reconcile the infrared and X-ray exünction determinations."," Still, a scaled-down version of the compact disk surrounded T Tau Sa, consistent with the fainter CO lines and smaller veiling of EC 95, might well provide the 20 magnitudes of extinction required to reconcile the infrared and X-ray extinction determinations."
 We mentioned in the introduction that SVS 4 is believed to be one of the densest sub-clusters in the Milky Way., We mentioned in the introduction that SVS 4 is believed to be one of the densest sub-clusters in the Milky Way.
" Given the determination of the distance to SVS 4 given here, it is useful to re-examine this issue."," Given the determination of the distance to SVS 4 given here, it is useful to re-examine this issue."
 Eiroa Casali (1989) idenüfied 11 young stellar objects in an angular diameter of ~ 50°., Eiroa Casali (1989) identified 11 young stellar objects in an angular diameter of $\sim$ 50”.
" Assuming a typical mass of 1 M. for each object, they obtained stellar mass densities between 4x10? and 10° M.  (the broad range reflected the uncertainty on the distance to Serpens)."," Assuming a typical mass of 1 $_\odot$ for each object, they obtained stellar mass densities between $4\times10^3$ and $10^{5}$ $_\odot$ $^{-3}$ (the broad range reflected the uncertainty on the distance to Serpens)."
" The 11 objects in SVS 4 are now known to be embedded in à common faint nebulosity covering an area of ~ 35""x35"" (Eiroa et al.", The 11 objects in SVS 4 are now known to be embedded in a common faint nebulosity covering an area of $\sim$ $''~\times$ $''$ (Eiroa et al.
 2008)., 2008).
 This implies a diameter of ~ 0.07 pe for the SVS 4 group., This implies a diameter of $\sim$ 0.07 pc for the SVS 4 group.
" The mass of EC 95 is about 5 M., so if we assume the other 10 objects in SVS 4 to be 1 M. stars, we obtain a total mass of 15 M. for the sub-cluster."," The mass of EC 95 is about 5 $M_\odot$, so if we assume the other 10 objects in SVS 4 to be 1 $M_\odot$ stars, we obtain a total mass of 15 $M_\odot$ for the sub-cluster."
 This yields a stellar mass density of about 8x 10 M. ., This yields a stellar mass density of about $\times$ $^4$ $_\odot$ $^{-3}$.
" This is consistent with the estimates of Eiroa Casali (1989), and confirms that SVS 4 is one of the densest concentration of pre main sequence stars known in the Galaxy."," This is consistent with the estimates of Eiroa Casali (1989), and confirms that SVS 4 is one of the densest concentration of pre main sequence stars known in the Galaxy."
" As we menuoned earlier, the distance to Serpens has been a matter of controversy during the past several decades."," As we mentioned earlier, the distance to Serpens has been a matter of controversy during the past several decades."
" In 22, we compile the distance estimates obtained by different authors over the years."," In 2, we compile the distance estimates obtained by different authors over the years."
" All these esumates are indirect, and most of them rely on spectroscopic parallaxes either to estimate the distance directly to Serpens members, or to measure extncüon as a function of distance [or objects in the direction of Serpens."," All these estimates are indirect, and most of them rely on spectroscopic parallaxes either to estimate the distance directly to Serpens members, or to measure extinction as a function of distance for objects in the direction of Serpens."
" Interesüngly, several of those studies are based on the same samples of stars (or at least on strongly overlapping samples). but reach widely different conclusions."," Interestingly, several of those studies are based on the same samples of stars (or at least on strongly overlapping samples), but reach widely different conclusions."
 This is the consequence of (wo main sources of uncertainty., This is the consequence of two main sources of uncertainty.
" The first one is the assignment of a spectral type to the stars considered, which sets their absolute magnitudes."," The first one is the assignment of a spectral type to the stars considered, which sets their absolute magnitudes."
 The second is the determination of the exüncton factor Ry which is required to translate reddening into exüncton., The second is the determination of the extinction factor $R_V$ which is required to translate reddening into extinction.
 There has been some discussion about the appropriate value, There has been some discussion about the appropriate value
We selected NLTT 11748 (LP413—40) as a white dwarf candidate from the revised NLTT catalogue (Salim&Gould.2003) using the reduced proper motion diagram (Salim&Gould.2003:Kawkaetal.. 2004).,"We selected NLTT 11748 $-$ 40) as a white dwarf candidate from the revised NLTT catalogue \citep{sal2003} using the reduced proper motion diagram \citep{sal2002,kaw2004}."
. Luyten(1977) originally listed NLTT 11748 in his white dwarf catalogue. and. more recently the star has also been recovered in the LSPM catalogue as LSPM J0345+1748 (Lépine&Shara.2005).," \citet{luy1977} originally listed NLTT 11748 in his white dwarf catalogue, and, more recently the star has also been recovered in the LSPM catalogue as LSPM $+$ 1748 \citep{lep2005}."
 The rNLTT and LSPM catalogue remeasured the proper motions independently., The rNLTT and LSPM catalogue remeasured the proper motions independently.
 We list the weighted average of these measurements in Table 2.., We list the weighted average of these measurements in Table \ref{tbl_prop}.
 We also measured the radial velocity of the star using the Balmer line series (Ha to Hy) and the corresponding heliocentric velocity is provided in Table 2.., We also measured the radial velocity of the star using the Balmer line series $\alpha$ to $\gamma$ ) and the corresponding heliocentric velocity is provided in Table \ref{tbl_prop}.
 The photometric distance estimate toward NLTT 11748 ts possibly effected by extinction in the line of sight., The photometric distance estimate toward NLTT 11748 is possibly effected by extinction in the line of sight.
 NLTT 11748 is near the Taurus-Auriga. molecular cloud. which is a well known star formation region located at a distance of ~140 pe (e.g..Li&Hu.1998).," NLTT 11748 is near the Taurus-Auriga molecular cloud, which is a well known star formation region located at a distance of $\sim 140$ pc \citep[e.g.,][]{li1998}."
. The extinction in the line of sight toward NLTT 11748 based on themaps of Schlegeletal.(1998). is E(B—Vy= 0.2828. which corresponds to a large extinction in the ultraviolet.," The extinction in the line of sight toward NLTT 11748 based on themaps of \citet{sch1998} is $E(B-V) = 0.2828$ , which corresponds to a large extinction in the ultraviolet."
 Assuming Ry=3.2 and using the relations from Cardellietal.(1989).. we estimate the extinction in the NUV and FUV bands to be Anuy=2.6 and Apuy=2.3. respectively.," Assuming $R_V = 3.2$ and using the relations from \citet{car1989}, we estimate the extinction in the NUV and FUV bands to be $A_{\rm NUV} = 2.6$ and $A_{\rm FUV} = 2.3$, respectively."
 However. we do not expect such a large extinction because the object is still well within the plane.," However, we do not expect such a large extinction because the object is still well within the plane."
 The extinction in the infrared should be very small. so we calculated a distance of 199+40 pe using the / magnitude.," The extinction in the infrared should be very small, so we calculated a distance of $199\pm40$ pc using the $J$ magnitude."
 Using the effective temperature and surface gravity from the spectral fit. we calculated the absolute magnitude in the FUV. NUV. V. J. and H bands.," Using the effective temperature and surface gravity from the spectral fit, we calculated the absolute magnitude in the $FUV$, $NUV$, $V$, $J$, and $H$ bands."
 The expected NUV—J for the white dwarfis 1.90. but the observed NUV—7 colour is 2.83640.090.," The expected $NUV - J$ for the white dwarf is $1.90$, but the observed $NUV - J$ colour is $2.836\pm0.090$."
 This corresponds to an extinction of Aygy=0.935 in the NUV and E(B—V)=0.10., This corresponds to an extinction of $A_{NUV} = 0.935$ in the NUV and $E(B-V) \approx 0.10$.
 Using the synthetic spectrum of the white dwarf at the derived temperature and surface gravity. we calculated a synthetic spectrum taking extinction into account where E(B—V)=0.10 and Ry=3.2.," Using the synthetic spectrum of the white dwarf at the derived temperature and surface gravity, we calculated a synthetic spectrum taking extinction into account where $E(B-V) = 0.10$ and $R_V = 3.2$."
 Figure 3 shows the observed energy distribution compared to a synthetic spectrum with and without extinction., Figure \ref{fig_sed} shows the observed energy distribution compared to a synthetic spectrum with and without extinction.
 The corresponding model magnitudes including extinction are given in Table 1.., The corresponding model magnitudes including extinction are given in Table \ref{tbl_phot}. .
 Based on the proper motion. distance. and velocity measurements. we calculated the kinematics (Table. 2) following Johnson&Soderblom(1987) and assuming that the solar motion relative to the local standard of rest (LSR) is (Us.V4.Wa)=(10.1.4.0.6.7) as determined by Hoge (2005).," Based on the proper motion, distance, and velocity measurements, we calculated the kinematics (Table \ref{tbl_prop}) ) following \citet{joh1987} and assuming that the solar motion relative to the local standard of rest (LSR) is $(U_\odot, V_\odot, W_\odot) = (10.1, 4.0, 6.7)$ as determined by \citet{hog2005}."
. Figure + shows the U and V velocities of NLTT 11748 assuming a radial velocity of 415 km s., Figure \ref{fig_UV} shows the $U$ and $V$ velocities of NLTT 11748 assuming a radial velocity of $415$ km $^{-1}$.
" However. NLTT 11748 probably resides in a close binary. and the correct systemic velocity still needs to bemeasured""."," However, NLTT 11748 probably resides in a close binary, and the correct systemic velocity still needs to be."
.. We illustrate the possible variations in UV velocities assuming a range of values from 0 to 600 km s! for the systemic velocity., We illustrate the possible variations in $UV$ velocities assuming a range of values from 0 to 600 km $^{-1}$ for the systemic velocity.
 NLTT 11748 bears a striking resemblance with the other high-velocity ELM white dwarf LP 400-22 2000). ," NLTT 11748 bears a striking resemblance with the other high-velocity ELM white dwarf LP 400-22 \citep{kil2009,ven2009}. ."
Given that the two field ELM white dwarfs (LP 400-22. SDSS JO917+46) that were checked for radial velocity variations were found to be in à binary system. it is highlyprobable that NLTT 11748 isalso in à binary system.," Given that the two field ELM white dwarfs (LP 400-22, SDSS $+$ 46) that were checked for radial velocity variations were found to be in a binary system, it is highlyprobable that NLTT 11748 isalso in a binary system."
 In the two confirmed, In the two confirmed
Ou 23 January 1999 detected990123.. one of the strongest recorded gamuna-ray bursts (CRBs) (cise 1999: Piro 1999)).,"	On 23 January 1999 detected, one of the strongest recorded gamma-ray bursts (GRBs) (Heise \cite{H99}; Piro \cite{P99}) )."
 The optical transicut (OT) associated with this GRD reached a peak maeuitude of Vo=as.95 CAkorlof Melxav 1999)) ouly LF seconds after the initial detection., The optical transient (OT) associated with this GRB reached a peak magnitude of $V = 8.95$ (Akerlof McKay \cite{AM99}) ) only 47 seconds after the initial detection.
 Spectroscopic observations with the Nordic Optical Telescope (IIjorth et citeITA99:: Andersen et citeA C993). and the Keck Telescope Uxelson et al. 1999:," Spectroscopic observations with the Nordic Optical Telescope (Hjorth et \\cite{HA99}; Andersen et \\cite{AC99}) ), and the Keck Telescope (Kelson et al. \cite{KI99};"
 Iuli et citelk D909)). showed stroug metallic absorption limes with redshifts of 2=1.600. suggesting that the OT is located either in. or behind. an absorbing svsteni atf i=1.6.," Kulkarni et \\cite{KD99}) ), showed strong metallic absorption lines with redshifts of $z=1.600$, suggesting that the OT is located either in, or behind, an absorbing system at $z = 1.6$ ."
 Andersen et ((1999}) placed an upper limit of ou the redshift of -2.05 based on ultraviolet. photometiy aud the absence of Lviiui-o. forest lues., Andersen et \cite{AC99}) ) placed an upper limit of on the redshift of $z < 2.05$ based on ultraviolet photometry and the absence of $\alpha$ forest lines.
" For a Frieden cosmolosv with Πρ=65 km 120,2 0.2. aud A=0 a redshift of +=1.6 corresponds to 1]982 Mpc aud a distance modulus of µ.=15.39."," For a Friedman cosmology with $H_0 = 65$ km $^{-1}$, $\Omega_0 = 0.2$, and $\Lambda
= 0$ a redshift of $z = 1.6$ corresponds to $11\,982$ Mpc and a distance modulus of $\mu = 45.39$."
 This means that the OT reache a peak absolute V-baix magnitude of Ay~ 836.51. which eives au intrinsic luminosity of Ly~2<1... approximately 105 times brighter than a Type I supernova.," This means that the OT reached a peak absolute $V$ -band magnitude of $M_V \sim -36.5$ , which gives an intrinsic luminosity of $L_V \sim 2 \times
10^{16} L_{\sun}$, approximately $10^7$ times brighter than a Type I supernova."
 For an isotropic burst the implied enerev release approached 3. 1.5<10°! eve; using BATSE fluence (IXippeu 1999)).," For an isotropic burst the implied energy release approached $3$ $4.5 \times 10^{54}$ erg, using BATSE fluence (Kippen \cite{K99}) )."
 There are wo broad families of models for CRBs. both of which are related to the cud stages of the lives of mnassive stars. and both of which predict that GRBs should trace the star-formation rate in the Universe WWijers et citeWB9s)}).," 	There are two broad families of models for GRBs, both of which are related to the end stages of the lives of massive stars, and both of which predict that GRBs should trace the star-formation rate in the Universe Wijers et \\cite{WB98}) )."
" The first family is the exploding object family. which inclides the ""αλλος supernova’ ποσο of Woosley (1993)) aud the “hypernova” modol of Paczvüsski (1998))."," The first family is the exploding object family, which includes the “failed supernova” model of Woosley \cite{W93}) ) and the “hypernova” model of Paczyńsski \cite{P98}) )."
 These models predict that the progenitors of GRBs are short-lived objects with low space velocities. so GRBs should be found within ~0.5 kpe of the star-forming regions where the progenitors formed.," These models predict that the progenitors of GRBs are short-lived objects with low space velocities, so GRBs should be found within $\sim 0.5$ kpc of the star-forming regions where the progenitors formed."
 The second family of models is the binary progenitor family. which inchides the binary neutron star (NS-NS) model of Naravan et al. (1992))," The second family of models is the binary progenitor family, which includes the binary neutron star (NS-NS) model of Narayan et al. \cite{NP92}) )"
 and the mereer of a black hole and a neutron star (BS-NS) inodel of Paczvüsski (1991))., and the merger of a black hole and a neutron star (BS-NS) model of Paczyńsski \cite{P91}) ).
 These models predict that GRD progenitors will have high space velocities. due to two supernova explosions having occurred in the progenitors binary svstem. so the CRB could be found a significant distance from the star-forming region where the progenitor was formed.," These models predict that GRB progenitors will have high space velocities, due to two supernova explosions having occurred in the progenitor's binary system, so the GRB could be found a significant distance from the star-forming region where the progenitor was formed."
" OTs have eeu associated withseveral CRBs. aud most of these OTs have been located within 1"" of afaint ealaxy."," OTs have been associated withseveral GRBs, and most of these OTs have been located within $1\arcsec$ of afaint galaxy."
 Redshifts have been measured for four of these candidate host galaxies and all are located at cosinological distances., Redshifts have been measured for four of these candidate host galaxies and all are located at cosmological distances.
mostly maenetically active stars observed by the high resolution ταν spectrometers. Drake&Testa(2005) found Ne/O abundance ratios to be consistently higher by a factor of 2.7 than the recommended solar value of the time. Ne/O=0.15 bynumber!.,"mostly magnetically active stars observed by the high resolution X-ray spectrometers, \citet{Drake.Testa:05} found Ne/O abundance ratios to be consistently higher by a factor of 2.7 than the recommended solar value of the time, Ne/O=0.15 by."
. The latter value is traceable to the assessment of Anders&Grevesse(19, The latter value is traceable to the assessment of \citet{Anders.Grevesse:89}.
89).. Drake&Testa(2005) suggested that (he same ratio might also be appropriate for the Sun aud help solve the “Solar Model Problem., \citet{Drake.Testa:05} suggested that the same ratio might also be appropriate for the Sun and help solve the “Solar Model Problem”.
 Models emploving a recently aclvancecl solar chemical composition based on 3-D non-LTE lvdrodvnamic photospheric modelling (Asplundetal.2005) lecl to predictions of the depth of the convection zone. helium abundance. density ancl sound speed in serious disagreement wilh helioseismoloey measurements (Dasu&Antia2004:Bahealletal.2005a).," Models employing a recently advanced solar chemical composition based on 3-D non-LTE hydrodynamic photospheric modelling \citep{Asplund.etal:05} led to predictions of the depth of the convection zone, helium abundance, density and sound speed in serious disagreement with helioseismology measurements \citep{Basu.Antia:04,Bahcall.etal:05}."
. The Asplundetal.(2005) mixture contained less of the elements C. N. O and Ne that are important for the opacitw of the solar interior by 25-35 ecompared to earlier assessments (e.g.Anders&Grevesse1980;Satval19," The \citet{Asplund.etal:05}
 mixture contained less of the elements C, N, O and Ne that are important for the opacity of the solar interior by 25-35 compared to earlier assessments \citep[e.g.][]{Anders.Grevesse:89,Grevesse.Sauval:98}."
93).. Antia&Basu(2005) and Dahealletal.(2005b) suggested the uncertain solar Ne abundance might be raised (ο compensate., \citet{Antia.Basu:05} and \citet{Bahcall.etal:05c} suggested the uncertain solar Ne abundance might be raised to compensate.
 The estimated Ne/O ratio required was in agreement wilh that Found in stellar coronae by Drake&Testa(2005)., The estimated Ne/O ratio required was in agreement with that found in stellar coronae by \citet{Drake.Testa:05}.
. More detailed investigationse of helioseismologye. models now suggestex that raisinge the Ne abundance alone cannot [ully reconcile models and observations (Delahave& Pinsonneault2006:Linetal.2007:Delahave 2010).," More detailed investigations of helioseismology models now suggest that raising the Ne abundance alone cannot fully reconcile models and observations \citep{Delahaye.Pinsonneault:06,Lin.etal:07,Delahaye.etal:10}."
. Nevertheless. the Ne abundance still remains an important individual ingredient. not onlv lor the Sun but also for nucleosynthesis ancl galactic chemical evolution.," Nevertheless, the Ne abundance still remains an important individual ingredient, not only for the Sun but also for nucleosynthesis and galactic chemical evolution."
 Ne exhibits no lines in the visible light spectra of Iate-tvpe stars and is not retained in meteorites., Ne exhibits no lines in the visible light spectra of late-type stars and is not retained in meteorites.
 Consequently. the solar abundance is based largely on transition region ancl coronal lines. and energetic particle measurements. supplemented with local cosmic estimates (e.g.Mever1935;Anders&Grevesse1989:Asplundοἱal.2009).," Consequently, the solar abundance is based largely on transition region and coronal lines, and energetic particle measurements, supplemented with local cosmic estimates \citep[e.g.\ ][]{Meyer:85,Anders.Grevesse:89,Asplund.etal:09}."
. Two studies of transition region EUV and coronal X-ray lines find consistency with the ratio Ne/O-—0.15 (Youngal.2005) and conclude that (his ratio represents (hat of the bulk of the Sun.," Two studies of transition region EUV and coronal X-ray lines find consistency with the ratio Ne/O=0.15 \citep{Young:05,Schmelz.etal:05b} and conclude that this ratio represents that of the bulk of the Sun."
 llowever. based on UV lines observed during a flare Landietal.(2007) obtained an absolute abundance Ne/IL a factor of 1.9 lareer than the Asplundetal.(2005). value.," However, based on UV lines observed during a flare \citet{Landi.etal:07} obtained an absolute abundance Ne/H a factor of 1.9 larger than the \citet{Asplund.etal:05} value."
 Subsecuently. the latter authors have revised their recommended Ne anc O abundances back up towarel older values by 0.13 and 0.07 dex. respectively (Aspluncletal.2009).. raising their favoured absolute logarithmic Ne abundance to Ne/II—7.9740.1 and Ne/O ratio to 0.11just within formal agreement with the Landietal.(2007) Ne abundance. but still somewhat short of," Subsequently, the latter authors have revised their recommended Ne and O abundances back up toward older values by 0.13 and 0.07 dex, respectively \citep{Asplund.etal:09}, raising their favoured absolute logarithmic Ne abundance to $=7.97\pm 0.1$ and Ne/O ratio to 0.17—just within formal agreement with the \citet{Landi.etal:07} Ne abundance, but still somewhat short of"
developed by ? was used. and readers are referred there for information on the method and its testing.,"developed by \cite{Watson2009} was used, and readers are referred there for information on the method and its testing."
 This is an automated system for detecting and tracking sunspots through large datasets and also records physical parameters of the sunspots detected., This is an automated system for detecting and tracking sunspots through large datasets and also records physical parameters of the sunspots detected.
 It involves using techniques from the field of morphological image processing to detect the outer boundaries of sunspot penumbrae., It involves using techniques from the field of morphological image processing to detect the outer boundaries of sunspot penumbrae.
 This 1s achieved by means of the top-hat transform which allows us to remove any limb-darkening profile from the data and to perform the detections in one step., This is achieved by means of the top-hat transform which allows us to remove any limb-darkening profile from the data and to perform the detections in one step.
 In addition to the method given in ? the code had to be developed further to separate the umbra and penumbra of spots as we would be looking at the magnetic fields present in the umbra., In addition to the method given in \citet{Watson2009} the code had to be developed further to separate the umbra and penumbra of spots as we would be looking at the magnetic fields present in the umbra.
 When visually inspecting the data there is a clear intensity difference between the umbra and penumbra in sunspots., When visually inspecting the data there is a clear intensity difference between the umbra and penumbra in sunspots.
 This difference is due to the magnetic structure of susnpots., This difference is due to the magnetic structure of susnpots.
 The umbra has a higher density of magnetic flux which inhibits convection more than in the penumbral region., The umbra has a higher density of magnetic flux which inhibits convection more than in the penumbral region.
 This causes the umbra to be cooler and therefore appear darker., This causes the umbra to be cooler and therefore appear darker.
 However. as sunspots move towards the limb both the umbra and penumbra are limb-darkened.," However, as sunspots move towards the limb both the umbra and penumbra are limb-darkened."
 For this reason. we cannot use a single threshold value to define the outer edge of the umbra.," For this reason, we cannot use a single threshold value to define the outer edge of the umbra."
 The algorithm we use removes all limb darkening effects at the same time as sunspot detection. greatly increasing speed as these two steps are carried out together.," The algorithm we use removes all limb darkening effects at the same time as sunspot detection, greatly increasing speed as these two steps are carried out together."
 This problem has been approached by other authors using different techniques. for example the inflection point method of ?.. the cumulative histogram method of ?.. the fuzzy logic approach of ?.. and the morphological approach of ?..," This problem has been approached by other authors using different techniques, for example the inflection point method of \cite{1997SoPh..171..303S}, the cumulative histogram method of \cite{1997SoPh..175..197P}, , the fuzzy logic approach of \cite{2009SoPh..260...21F}, and the morphological approach of \cite{2005SoPh..228..377Z}."
 Our method begins with the sunspots (which includes umbrae and penumbrae) detected by STARA. and then produces a histogram of sunspot pixel intensities for each spot.," Our method begins with the sunspots (which includes umbrae and penumbrae) detected by STARA, and then produces a histogram of sunspot pixel intensities for each spot."
 This clusters in two peaks. the local minimum between which corresponds to the intensity value at the edge of the umbra.," This clusters in two peaks, the local minimum between which corresponds to the intensity value at the edge of the umbra."
 A similar. histogram-based approach was implemented by ? who then used concepts from fuzzy logic to assign membership to umbra or penumbra: they showed that particularly the pixel membership of the penumbra can vary significantly (tens of percent) depending on a parameter known as the membership function. but this is apparently less of a problem for low-resolution data. in which brightness variations within the penumbra are smeared out.," A similar histogram-based approach was implemented by \cite{2009SoPh..260...21F} who then used concepts from fuzzy logic to assign membership to umbra or penumbra; they showed that particularly the pixel membership of the penumbra can vary significantly (tens of percent) depending on a parameter known as the membership function, but this is apparently less of a problem for low-resolution data, in which brightness variations within the penumbra are smeared out."
 We have not adopted such a method. but have instead identified the local minimum for each sunspot's histogram. and created a mask for umbral pixels.," We have not adopted such a method, but have instead identified the local minimum for each sunspot's histogram, and created a mask for umbral pixels."
 We normally find that the umbra region of sunspots has an MDI pixel value of less than 7000 - 8000., We normally find that the umbra region of sunspots has an MDI pixel value of less than 7000 - 8000.
 However. our algorithm does have the benefit of being applied consistently across the entire data series. and being able to deal with the varying intensity across the solar disk due to limb darkening which eases the problems of sunspot detection and area estimation that occur if a straightforward intensity threshold is used.," However, our algorithm does have the benefit of being applied consistently across the entire data series, and being able to deal with the varying intensity across the solar disk due to limb darkening which eases the problems of sunspot detection and area estimation that occur if a straightforward intensity threshold is used."
 The data used in this study are taken from the MDI instrument (?) on the SOHO spacecraft., The data used in this study are taken from the MDI instrument \citep{Scherrer1995} on the SOHO spacecraft.
 We use the level 1.8 continuum data as well as the level 1.8 magnetograms to analyse magnetic fields present in the spots., We use the level 1.8 continuum data as well as the level 1.8 magnetograms to analyse magnetic fields present in the spots.
 Our dataset uses [5 years of data and we analyse daily measurements taken at OOOOUT when co-temporal continuum images and magnetograms are recorded., Our dataset uses 15 years of data and we analyse daily measurements taken at 0000UT when co-temporal continuum images and magnetograms are recorded.
 The STARA code takes around 24 hours to process the approximately 5000 days of data available to generate the sunspot catalogue used in this article and holds 300084 separate sunspot detections., The STARA code takes around 24 hours to process the approximately 5000 days of data available to generate the sunspot catalogue used in this article and holds 084 separate sunspot detections.
 The same sunspot will be detected in many different images and tracked from image to image allowing them to be associated with one another., The same sunspot will be detected in many different images and tracked from image to image allowing them to be associated with one another.
" The physical parameters obtained from this analysis are the sunspot total area and ""centre of mass! location. number and area of umbrae; mean. maximum and minimum magnetic fields in the umbrae and penumbra: total and excess flux in the umbrae and penumbra and the information relating to the observation itself such as time and instrument used."," The physical parameters obtained from this analysis are the sunspot total area and `centre of mass' location, number and area of umbrae; mean, maximum and minimum magnetic fields in the umbrae and penumbra; total and excess flux in the umbrae and penumbra and the information relating to the observation itself such as time and instrument used."
 The trend of sunspot number throughout a solar cycle is well documented and generally rises rapidly at the start of a solar cycle before a slower decrease towards the end of the cycle., The trend of sunspot number throughout a solar cycle is well documented and generally rises rapidly at the start of a solar cycle before a slower decrease towards the end of the cycle.
 The Solar Influences Data Center (SIDC. be/sunspot-data/)) keeps records on the sunspot index and so we compare the results of our detections with the findings of the SIDC as an initial test.," The Solar Influences Data Center (SIDC, ) keeps records on the sunspot index and so we compare the results of our detections with the findings of the SIDC as an initial test."
 It must be noted that both indicators are not measuring the same thing as the international sunspot number recorded by the SIDC weights the sunspots seen in groups so that it becomes a stronger proxy for solar activity whereas STARA only gives us the raw number of observed sunspots., It must be noted that both indicators are not measuring the same thing as the international sunspot number recorded by the SIDC weights the sunspots seen in groups so that it becomes a stronger proxy for solar activity whereas STARA only gives us the raw number of observed sunspots.
 However. it is beneficial to see if the same trends are present.," However, it is beneficial to see if the same trends are present."
 The data used here are the smoothed monthly sunspot number (?) and so our daily measurements have been treated in the same way to give a fair comparison., The data used here are the smoothed monthly sunspot number \citep{SIDC} and so our daily measurements have been treated in the same way to give a fair comparison.
 In Fig., In Fig.
 | we can see that both curves share several features., \ref{fig:emergences} we can see that both curves share several features.
 The STARA output has been scaled up to the same level as the International Sunspot Number around 2001 - 2003 when sunspot count rates were higher and the general trends are more important here than absolute values due to the differences in counting methods (this scaling is permissible due to the somewhat arbitrary factors present in the SIDC sunspot numbers - see Equation |..), The STARA output has been scaled up to the same level as the International Sunspot Number around 2001 - 2003 when sunspot count rates were higher and the general trends are more important here than absolute values due to the differences in counting methods (this scaling is permissible due to the somewhat arbitrary factors present in the SIDC sunspot numbers - see Equation \ref{eq:1}. .)
 We see that both datasets exhibit the same patterns of increasing and decreasing at the same time and the agreement Is very good in the declining phase of the cycle., We see that both datasets exhibit the same patterns of increasing and decreasing at the same time and the agreement is very good in the declining phase of the cycle.
 This also continues into cycle 24 shown at the right hand side of the plot with both curves rising at the same time and we will continue to track the agreement of these further into the next cycle., This also continues into cycle 24 shown at the right hand side of the plot with both curves rising at the same time and we will continue to track the agreement of these further into the next cycle.
 The SIDC data (?).. shown as a dashed line on the plot has a smooth rise up to the first maximum sometime in the year 2000 and falls before reaching a second maximum in 2002.," The SIDC data \citep{Clette2007}, shown as a dashed line on the plot has a smooth rise up to the first maximum sometime in the year 2000 and falls before reaching a second maximum in 2002."
" This ""double maximum' feature. separated by the ""Gnevyshev gap’ (?) is also seen in the STARA output although the first maximum is weaker when compared to the second. in contrast with the SIDC data in which the first maximum is larger than the second."," This `double maximum' feature, separated by the `Gnevyshev gap' \citep{1967SoPh....1..107G} is also seen in the STARA output although the first maximum is weaker when compared to the second, in contrast with the SIDC data in which the first maximum is larger than the second."
 However. both sets of data scale well with one another after this second maximum with very little deviation and this continues from 2002 up to the current day.," However, both sets of data scale well with one another after this second maximum with very little deviation and this continues from 2002 up to the current day."
 The differences in the first peak. and indeed in the rise beforethat are most likely due to the method of counting," The differences in the first peak, and indeed in the rise beforethat are most likely due to the method of counting"
atmosphere around the core (like in the isothermal case). should have as initial conditions a core surrounded by a static atmosphere extending all the way up to the Roche lobe radius.,"atmosphere around the core (like in the isothermal case), should have as initial conditions a core surrounded by a static atmosphere extending all the way up to the Roche lobe radius."
 In such a case. there will be no cold accretion How around the protoplanet.," In such a case, there will be no cold accretion flow around the protoplanet."
 It has been argued that the orbital migration of low mass cores depends very significantly on the torque exerted by the gas within the Roche lobe of the protoplanet (D'Angelo et al., It has been argued that the orbital migration of low mass cores depends very significantly on the torque exerted by the gas within the Roche lobe of the protoplanet (D'Angelo et al.
 2003. Crida et al.," 2003, Crida et al."
 2009)., 2009).
 The results presented above indicate that this is probably not the case. as (he protoplanet itself fills in its Roche lobe.," The results presented above indicate that this is probably not the case, as the protoplanet itself fills in its Roche lobe."
 In (his paper. we have considered the evolution of a protoplanet which accretes planetesimadls from the protoplanetary disc.," In this paper, we have considered the evolution of a protoplanet which accretes planetesimals from the protoplanetary disc."
 Such an accretion may be reduced significantly before (he gas in (he nebula has disappeared. with the consequence that the core reaches the critical mass and the atmosphere has (to contract to provide the energy radiated at its surface.," Such an accretion may be reduced significantly before the gas in the nebula has disappeared, with the consequence that the core reaches the critical mass and the atmosphere has to contract to provide the energy radiated at its surface."
 In that case. ils evolution is (hat calculated by Papaloizou," In that case, its evolution is that calculated by Papaloizou"
dense cores). but that they would agree on the effect of the bulk motion.,"dense cores), but that they would agree on the effect of the bulk motion."
 Regarclless of which fit we use. Figure 1. shows that the relative motion has a large ellect on the minimum circular velocity.," Regardless of which fit we use, Figure \ref{Fig:Vc} shows that the relative motion has a large effect on the minimum circular velocity."
 The implications for the minimum cooling mass as a function of redshift are also shown in Figure 1. (bottom panel)., The implications for the minimum cooling mass as a function of redshift are also shown in Figure \ref{Fig:Vc} (bottom panel).
 In a patch with no relative motion. the mass drops rapidly with recshift. since at higher redshift the eas density is higher and a given halo mass heats the infalling gas to a higher virial temperature.," In a patch with no relative motion, the mass drops rapidly with redshift, since at higher redshift the gas density is higher and a given halo mass heats the infalling gas to a higher virial temperature."
 However. in a region at the root-mean-square value of ei; the higher bulk velocity at jiegh redshift implies that a higher halo mass is needed. for cllicient molecular cooling.," However, in a region at the root-mean-square value of $v\bc$ the higher bulk velocity at high redshift implies that a higher halo mass is needed for efficient molecular cooling."
 In particular. at. redshift 20 a ouch with ei=0 will form stars in 3.6101 M. halos. while a patch with the root-mean-square value of ei; has a mininiun cooling mass of 6.010° M. according to the optimal fit. or a range of (4.87.3)10 M. from the other its.," In particular, at redshift 20 a patch with $v\bc = 0$ will form stars in $3.6 \times 10^5$ $_\odot$ halos, while a patch with the root-mean-square value of $v\bc$ has a minimum cooling mass of $6.0 \times 10^5$ $_\odot$ according to the optimal fit, or a range of $(4.8 - 7.3) \times 10^5$ $_\odot$ from the other fits."
 A s=60 these numbers become 7.2.103 M... 7.0.107 AL. cand (4.1.10.3).10° AL... respectively.," At $z=60$ these numbers become $7.2 \times 10^4$ $_\odot$, $7.0 \times 10^5$ $_\odot$, and $(4.1 - 10.3) \times 10^5$ $_\odot$, respectively."
 In patches with ow bulk velocity we expect stars to form earlier. since the ialos with lower masses are more abundant and form earlier in the hierarchical picture of structure formation.," In patches with low bulk velocity we expect stars to form earlier, since the halos with lower masses are more abundant and form earlier in the hierarchical picture of structure formation."
 This is the xisis of the discussion that follows., This is the basis of the discussion that follows.
 The population of eas-Lilled halos at high redshift. divides naturally into two major categories., The population of gas-filled halos at high redshift divides naturally into two major categories.
 The first category consists of large halos in which the eas can cool (via molecular hvdrogen cooling): these are presumed. to be the sites of formation of the first. stars. ancl are obviously most important. since the stellar radiation is in principle observable. and it also produces feedback on the intergalactic medium and on other nearby sites of star formation.," The first category consists of large halos in which the gas can cool (via molecular hydrogen cooling); these are presumed to be the sites of formation of the first stars, and are obviously most important since the stellar radiation is in principle observable, and it also produces feedback on the intergalactic medium and on other nearby sites of star formation."
 Also interesting. though. is the second category. namely the smaller halos (*minihalos”) in which the eas accumulates to roughly virial density ancl vet cannot cool.," Also interesting, though, is the second category, namely the smaller halos (“minihalos”) in which the gas accumulates to roughly virial density and yet cannot cool."
 The latter may allect the epoch of reionization by acting as à sink for ionizing photons (e.g..Llaiman.Abel&Macdau2001:Ciardiοἱal.2005) and may generate a 21-em signal [rom collisional excitation of (e.g.Hievetal.2003:Furlanetto&Oh 2006).," The latter may affect the epoch of reionization by acting as a sink for ionizing photons \citep[e.g.,][]{Haiman01, BL02, Iliev05,
Ciardi05} and may generate a 21-cm signal from collisional excitation of \citep[e.g.,][]{Iliev03, FO06}."
. In this subsection we apply the result. we founcl for the minimum cooling mass to find the recdshilt evolution of the eas fraction in these two categories., In this subsection we apply the result we found for the minimum cooling mass to find the redshift evolution of the gas fraction in these two categories.
 In the following subsections we explore the probability distribution function (PDE) of the gas fraction. beginning with its dependence on the bulk. velocity.," In the following subsections we explore the probability distribution function (PDF) of the gas fraction, beginning with its dependence on the bulk velocity."
 In addition. though. in each patch of coherent velocity the mean density is slightly. different. varving as a result. of random density fluctuations on scales larger than the patch size.," In addition, though, in each patch of coherent velocity the mean density is slightly different, varying as a result of random density fluctuations on scales larger than the patch size."
 We thus also study the full PDE as determined by the joint dependence of the gas fraction in alos on the bulk velocity and the local overdensity in cach xuch., We thus also study the full PDF as determined by the joint dependence of the gas fraction in halos on the bulk velocity and the local overdensity in each patch.
" Following Tseliakhovich.Barkana&LHirata(2010). we ind the fraction of the barvon density contained in halos with mass larger than the minimunm cooling mass AM UM where py is the mean matter density today. απαλ is the comoving abundance of halos of mass M. fi,=O,/(,, is the mean cosmic barvon fraction and. f,CA) is the fraction of the total halo mass which is in the form of gas."," Following \citet{Tseliakhovich:2010b} we find the fraction of the baryon density contained in halos with mass larger than the minimum cooling mass $M\cool$ where $ \bar\rho_0$ is the mean matter density today, $dn/dM$ is the comoving abundance of halos of mass $M$, $f_b\equiv
\Omega_b/\Omega_m$ is the mean cosmic baryon fraction and $f_g(M)$ is the fraction of the total halo mass which is in the form of gas."
 The eas Tractions {ΑΔΗ} depend on the filtering mass. which measures the scale at which the barvon [uctuations ciller substantially [rom those in the dark matter.," The gas fractions $f_g(M)$ depend on the filtering mass, which measures the scale at which the baryon fluctuations differ substantially from those in the dark matter."
 In each patch. the filtering mass depends on the bulk velocity. ancl thus so do the gas fractions.," In each patch, the filtering mass depends on the bulk velocity, and thus so do the gas fractions."
 Since the barvons contribute to the total power spectrum. the halo abundance dn/dÀI (which depends on Huctuations in the total matter density) varies as well with eyo.," Since the baryons contribute to the total power spectrum, the halo abundance $dn/dM$ (which depends on fluctuations in the total matter density) varies as well with $v\bc$ ."
 We use the halo mass function of Tormen (1999)., We use the halo mass function of \citet{Shetht:1999}.
. We refer the reader to Darkana&Lirata(2010). for the full details., We refer the reader to \citet{Tseliakhovich:2010b} for the full details.
 We begin by recalculating some of the results. of Tseliakhovich.Barkana&Hirata(2010).., We begin by recalculating some of the results of \citet{Tseliakhovich:2010b}.
 We show in Figure 2. the redshift evolution of the globally averaged eas fraction in star-forming halos or in gas minihalos., We show in Figure \ref{Fig:Vc2} the redshift evolution of the globally averaged gas fraction in star-forming halos or in gas minihalos.
 Compared. with Figure δ of Tseliakhovich.Barkana&LHi-rata (2010).. our gas fractions are substantially lower. e... the gas fraction in halos above the minimum cooling mass is lower by a factor of ~3 at redshift z=20.," Compared with Figure 8 of \citet{Tseliakhovich:2010b}, our gas fractions are substantially lower, e.g., the gas fraction in halos above the minimum cooling mass is lower by a factor of $\sim 3$ at redshift $z = 20$."
" This is due to our higher A, and lower power spectrum. normalization (sce Section 6 for a Cull discussion of our cdillerences with previous papers).", This is due to our higher $M\cool$ and lower power spectrum normalization (see Section 6 for a full discussion of our differences with previous papers).
 Note that the gas fraction in halos above the minimum cooling mass is proportional to the stellar mass density. assuming a fixed star formation ellicieney. (averaged over each 3 Alpe patch).," Note that the gas fraction in halos above the minimum cooling mass is proportional to the stellar mass density, assuming a fixed star formation efficiency (averaged over each 3 Mpc patch)."
 In general. the importance of the relative velocities increases with redshift.," In general, the importance of the relative velocities increases with redshift."
 Comparing the two categories ofhalos. we find that the relative suppression of theminihalos is larger than that of the star-forming halos at low recshilt:," Comparing the two categories ofhalos, we find that the relative suppression of theminihalos is larger than that of the star-forming halos at low redshift;"
0.05. ενω=0:37c0.010. we discuss tlis feature further iu Section 7.,", $\gamma_{spiral}=0.37\pm0.05$, we discuss this feature further in Section 7."
 Abraham et al. (, Abraham et al. (
1996) aud Driver et al. (,1996) and Driver et al. (
1998) studied the morphological προς. counts for all galaxies with yp <26.0 detected in the TDF-N. Our results are in agreement with theirs: uuuber counts are dominated by late type galaxies and carly type sources show a flat curve.,1998) studied the morphological number counts for all galaxies with $_{AB}<$ 26.0 detected in the HDF-N. Our results are in agreement with theirs: number counts are dominated by late type galaxies and early type sources show a flat curve.
 Results for late tvpe galaxies are compatible with a strong evolution iu nuuber or with a non uceligible fraction of sources being at moderate redshift. as suggested also by the analysis of colors.," Results for late type galaxies are compatible with a strong evolution in number or with a non negligible fraction of sources being at moderate redshift, as suggested also by the analysis of colors."
 This result fits in Alavzke ct al. (, This result fits in Marzke et al. (
"1997) Ποιος about the LF of galaxies as a function of their color: the slope of the LE. faint cud is steeper at low luniunosities for thebluest galaxies (Mj,στ19.1. oco12 for galaxies bluer than V=0.1).","1997) findings about the LF of galaxies as a function of their color: the slope of the LF faint end is steeper at low luminosities for thebluest galaxies $M^*_B\approx -19.4$, $\alpha<-1.7$ for galaxies bluer than $-$ V=0.4)."
 Actually by inteerating this LE. within 2= 0.5. the munber deusitv," Actually by integrating this LF, within $z=0.5$ , the number density"
(diagonaliziug) of this matrix for high fs are ucarly impossible.,(diagonalizing) of this matrix for high $\ell$ s are nearly impossible.
 Diagonaliziug a matrix is more effücient when elements are close to the diagonal., Diagonalizing a matrix is more efficient when off-diagonal elements are close to the diagonal.
 However. the (TT.EE.BB) veprescutation of the matrix in figure AT gives rise to clements spread around the full matrix.," However, the (TT,EE,BB) representation of the matrix in figure \ref{fig:covmatold} gives rise to elements spread around the full matrix."
 Typically. Cholesky-factorization for such a TT-EE-DD representation breaks down for A444;=Gl due to the dense structure of the upper-triaugular decomposed. Lauatzix.," Typically, Cholesky-factorization for such a TT-EE-BB representation breaks down for $l_{\textrm{max}}=64$ due to the dense structure of the upper-triangular decomposed L-matrix."
 To overcome this problem. we operate with a differcut representation of the (TT.EE.DD)anatrix.," To overcome this problem, we operate with a different representation of the (TT,EE,BB)-matrix."
" Instead of building the matrix as presented in equation À2.. we choose a different way of expressing the matrix: with corresponding e/,,5 As the EE and BB correlations share the same structure as the stand-alone TT. the complete covariance matrix will in this represcutation resenible the original three-bauded covariance matrix."," Instead of building the matrix as presented in equation \ref{eq:covariancematrix1}, we choose a different way of expressing the matrix: with corresponding $a_{\ell m}$ s As the EE and BB correlations share the same structure as the stand-alone TT, the complete covariance matrix will in this representation resemble the original three-banded covariance matrix."
 The clements are now munich closer to the diagonal. solving the problem of ineficient diagonalizine.," The elements are now much closer to the diagonal, solving the problem of inefficient diagonalizing."
 The matrix representation is depicted in figure Ατιν, The matrix representation is depicted in figure \ref{fig:covmatold}.
 Note that this representation is ouly used when iultipliug the matrices with vectors aud performing Cholesky decompositions., Note that this representation is only used when multiplying the matrices with vectors and performing Cholesky decompositions.
" Withiug the rest of the framework. the 0;,,5 are treated as in equation Al.."," Withing the rest of the framework, the $a_{\ell m}$ s are treated as in equation \ref{eq:alms}. ."
"far-apart energy of the pair Beyond this point, the pair is no longer tracked.","far-apart energy of the pair Beyond this point, the pair is no longer tracked."
 4., 4.
" Find a new partner for proton 7, and repeat the process for 7 and its new partner."," Find a new partner for proton $i$, and repeat the process for $i$ and its new partner."
" Since a proton pair's departure from their closest point to a distance Ay is symmetric to an approach from R; to r,. we use the information obtained from the departure to examine the dynamic effect of the plasma on the screening enhancement for interacting proton pairs as they approach a separation at which nuclear reactions are possible."," Since a proton pair's departure from their closest point to a distance $R_f$ is symmetric to an approach from $R_f$ to $r_c$, we use the information obtained from the departure to examine the dynamic effect of the plasma on the screening enhancement for interacting proton pairs as they approach a separation at which nuclear reactions are possible."
 The dynamic information of cach proton and its partner are recorded at cach time step., The dynamic information of each proton and its partner are recorded at each time step.
 First we examine the distribution of the distance of closest approach of proton pairs., First we examine the distribution of the distance of closest approach of proton pairs.
 Fig., Fig.
 2 shows the number of pairs that achieved each closest distance r.., \ref{plottwo} shows the number of pairs that achieved each closest distance $r_c$.
" We see that most pairs are able to get at least as close as the average distance of neighboring protons «r,>=0.256a,.", We see that most pairs are able to get at least as close as the average distance of neighboring protons $<r_p>= 0.256 a_B$.
" As seen in Fig. 3., "," As seen in Fig. \ref{plotthree}, ,"
pairs of protons with higher far-apart kinetic energy £kinetic are more likely to get closer to each other than pairs with lower far-apart kinetic energy., pairs of protons with higher far-apart kinetic energy $E^f_{\rm{kinetic}}$ are more likely to get closer to each other than pairs with lower far-apart kinetic energy.
 The screening energy is the energy exchanged between a pair of protons and the surrounding plasma., The screening energy is the energy exchanged between a pair of protons and the surrounding plasma.
 It represents the combined effect of the neiehboring protons and electrons on the interacting pair., It represents the combined effect of the neighboring protons and electrons on the interacting pair.
 Salpeter's mean-field theory gives an expression for the average screened Coulomb potential that includes the energy exchanged between interacting protons and the plasma during their approach Salpeter's theory averages the effect of the kinetic energy of each particle so that the screening energy only depends on distance., Salpeter's mean-field theory gives an expression for the average screened Coulomb potential that includes the energy exchanged between interacting protons and the plasma during their approach Salpeter's theory averages the effect of the kinetic energy of each particle so that the screening energy only depends on distance.
" However, our simulations show that the screening energy also depends on the kinetic energy of the protons."," However, our simulations show that the screening energy also depends on the kinetic energy of the protons."
" Protons with hieh Kinetic energy tend to eain less energy from the plasma during the interaction, while protons with low kinetic energy eain more energy from the plasma."," Protons with high kinetic energy tend to gain less energy from the plasma during the interaction, while protons with low kinetic energy gain more energy from the plasma."
" Shaviv&(2000,2001) refer to this as the dynamic effect."," \citet{Shaviv_2000, Shaviv_2001} refer to this as the dynamic effect."
" For the purpose of comparison, we define the screening energy in the simulations to be the difference in the energy of the pair when they are far apart and when theyare close together."," For the purpose of comparison, we define the screening energy in the simulations to be the difference in the energy of the pair when they are far apart and when theyare close together."
"by magnetic fields @ can be computed as the integral of the distributions in Fig. 7,,","by magnetic fields $\alpha$ 	can be computed as the integral of the distributions in 	Fig. \ref{fig4},"
" i.e., Considering the full subfield, the value of a for the major and the minor components is 3.1 and 1.4 %,, respectively."," i.e., Considering the full subfield, 	the value of $\alpha$ for the major and the minor components 	is $3.1$ and $1.4$ , respectively."
" In other words, 95.5 of the analyzed photosphere is field free or, more precisely, it has field strengths much smaller than the ones inferred from inversion, so that the inversion code cannot distinguish them from zero."," In other words, $95.5$ 	of the analyzed photosphere is field free or, more precisely, 	it has field strengths much smaller than the ones inferred 	from inversion, so that the inversion code cannot distinguish them from zero."
 Other fractions corresponding to the various components in the FOV are listed in Table 1.., Other fractions corresponding to the various components in the FOV are listed in Table \ref{tab}.
" The first moment of the magnetic field strength distributions also provides an estimate of the unsigned flux density (B) (see,e.g.,DomínguezCerdefiaetal.2006b),, and of the average field strength, The values of (B) and B for the various components in the FOV are listed in Table 1.."," The first moment of the magnetic field strength distributions also provides an estimate of the unsigned flux density $\langle B\rangle$ \citep[see, e.g., ][]{DomC06b}, and of the average field strength, The values of $\langle B\rangle$ and $\overline{B}$ for the various components in the FOV are listed in Table \ref{tab}."
" As the table shows, (B)—66 G considering network and IN all together."," As the table shows, $\langle B\rangle = 66$ G considering network and IN all together."
" This flux density is so large partly because it reflects the fraction of the 29.52""x31.70"" subfield having the largest polarization signals.", This flux density is so large partly because it reflects the fraction of the $29.52''\times31.70''$ subfield having the largest polarization signals.
" Such fraction f is about 0.29, therefore, if one considers the full subfield, the fraction of photosphere occupied by the magnetic fields inferred from our inversion is af, and the unsigned flux density (B)f."," Such fraction $f$ is about $0.29$, therefore, if one considers the full subfield, the fraction of photosphere occupied by the magnetic fields inferred from our inversion is $\alpha f$, and the unsigned flux density $\langle B\rangle f$."
 The values of the filling factor and unsigned flux densities corresponding to this other normalization are also included in Table 1 within parentheses., The values of the filling factor and unsigned flux densities corresponding to this other normalization are also included in Table \ref{tab} within parentheses.
" Then the unsigned flux density decreases to about 19 G. This still large value is put into context in 5,, but it is important to realize that it implicitly assumes the (1—f) non-inverted subfield to have no magnetic fields."," Then the unsigned flux density decreases to about $19$ G. This still large value is put into context in \ref{disc}, but it is important to realize that it implicitly assumes the $(1-f)$ non-inverted subfield to have no magnetic fields."
 We have also tried a crude separation between granules and intergranules., We have also tried a crude separation between granules and intergranules.
" A simple thresholding criterion was used, so that pixels brighter than the mean intensity are granules, and vice-versa."," A simple thresholding criterion was used, so that pixels brighter than the mean intensity are granules, and vice-versa."
" We find that 75 of the magnetized plasma is in intergranules, and this fraction increases with increasing field strength — 85 of the plasma having B>1.5 kG is in intergranules."," We find that $75$ of the magnetized plasma is in intergranules, and this fraction increases with increasing field strength – $85$ of the plasma having $B>1.5$ kG is in intergranules."
 These figures are meant to be rough estimates since our simple criterion is insufficient for an accurate separation between granules and intergranules., These figures are meant to be rough estimates since our simple criterion is insufficient for an accurate separation between granules and intergranules.
 The ubiquity of Stokes profiles with asymmetries implies that even SOT/SP resolution is insufficient to resolve the structure of the quiet Sun magnetic fields (seeSánchezAlmeidaetal.1996)., 	The ubiquity of Stokes profiles with asymmetries 	implies that even SOT/SP resolution 	is insufficient to resolve the structure of the 	quiet Sun magnetic fields \citep[see][]{SanA96}.
".. This unresolved structure has been neglected in the analyses carried out so far, and we present the first attempt to account for it when interpreting spectra."," 	This unresolved structure has been neglected in the analyses carried out 	so far, and we present the first attempt to account 	for it when interpreting spectra."
" We employ the MISMA inversion code by SánchezAlmeida(1997),, which produces very good fits."," 	We employ the MISMA inversion code by \citet{SanA97}, 	which produces very good fits."
" The MISMA description is able to reproduce the whole range of profile shapesobserved by SOT/SP in the quiet Sun, which are often extremely asymmetric."," 	The MISMA description is able to reproduce the whole range of profile shapesobserved by SOT/SP in the quiet Sun, which are often extremely asymmetric."
" This was known for Advanced Stokes Polarimeter (ASP) spectra (Sánchez 2000),, but SOT/SP has three"," This 	was known for Advanced Stokes Polarimeter (ASP) spectra \citep{SanALit00}, , but SOT/SP has three"
(LindandDlandford1985).,\citep{lin85}.
". This equation derives [rom an assumed powerlaw energy distribution for the relativistic electrons. .N(E)=Np"". where the radio spectral index α—(n1)/2 and E is the energy of the electrons in units of πο."," This equation derives from an assumed powerlaw energy distribution for the relativistic electrons, $ N(E)= N_{\Gamma}E^{-n}$, where the radio spectral index $\alpha = (n-1)/2$ and $E$ is the energy of the electrons in units of $m_{e}c^{2}$."
" The radiative transfer equation was solved in GinzburgandSvrovatskii(1969). to vield the following parametric form for the observed flux density. 5,. from the SSA source. where Ris the radius of the spherical region in the rest Irae of the plasma and 5, is a normalization factor."," The radiative transfer equation was solved in \citet{gin69} to yield the following parametric form for the observed flux density, $S_{\nu}$, from the SSA source, where R is the radius of the spherical region in the rest frame of the plasma and $S_{o}$ is a normalization factor."
 In the spherical. homogeneous approximation. one can make a simple ⊔⋯↥⊔∐↲⋟∖⊽∪⋯⋅≺∢≼↲↕⋟∖⊽⋟∖⊽↕↽≻↥∐↲↕⋅↕≺∢≀↕↴↥≀↧↴∐≼⇂∐∪∐↓∪≸↽↔↴≼↲∐≼↲∪∏⋟∖⊽⊔∐↲∐⊔∐↲↕⋅≼↲≀↧↴↕⋅≼↲⊔∐⋅≼↲≼↲∏∐↳↽∐⋯∖⊽∐⋟∖⇁↕∐≼↲≺⇂∏≀↕↴∐∪∐⋖⋡⊑↽⊰↕⋝⋅.," In the spherical, homogeneous approximation, one can make a simple parameterization of the SSA attenuation coefficient, $\mu(SSA)=\mu_{o}\nu_{o}^{(-2.5 +
\alpha)}$."
 . − ∙− ↴⊽⊾⊏⊇∡⊽⋡∃ ↕↽≻≀↧↴↕⋅≀↧↴∐∐↲↥≼↲↕⋅↕∠≀↕↴↥↕∪∐∪↓⊔∐↲⊳↔⊲⊳↔⊲∶∖≀↧↴⊔≼↲∐∏≀↕↴⊔∪∐≺∢⋯↲∐↓≺∢↕≼↲∐↥⋅∕∣⋖↜∣⇆∣⇆⇀∣⊔∶∕∣∩↗∣∕↭↗⊔⋅∐∪∐≼↲≀↧⊔∖⊽⋟∖⊽∏∐∐↲⋟∖⊽ ∫↘⋝∕∣∩⋅∪≀↧↴∐≼⇂∣⇆⊽∩⋅≀↧↴∐≼⇂∣⇆⊽∣∕↕⋝∖⊽≼⇂≼↲∥↲↕⋅∐↓↕∐≼↲≺⇂∣↽≻⋡∖↽∪∣↽≻⊳∖⇁≼↲," If one assumes that the source is spherical and homogeneous then there are three unknowns in equation (3), $R\mu_{o}$, $\alpha$ and $S_{o}$, and $S_{\nu}$ is determined by observation."
↕⋅∖↽≀↧↴∐∪∐⋅ Figure 1 shows the simultaneous VLA data from December 6. 1993.," Figure 1 shows the simultaneous VLA data from December 6, 1993."
 The data seems to indicate three local maxima in the spectrum (one between 1.4 Gllz and 5 Gllz. one near 8.3 GlIlIz and one near 22 Gllz).," The data seems to indicate three local maxima in the spectrum (one between 1.4 GHz and 5 GHz, one near 8.3 GHz and one near 22 GHz)."
 It should be noted (that the SSA function is not an arbitrary shape., It should be noted that the SSA function is not an arbitrary shape.
 For example. it has precisely one maximum.," For example, it has precisely one maximum."
 Thus. even though each SSA [unetion has three parameters. there is no enarantee. in general. that two SSA components can fit six or less points exactly.," Thus, even though each SSA function has three parameters, there is no guarantee, in general, that two SSA components can fit six or less points exactly."
 Two SSA components can create al most (wo local maxima in specirun., Two SSA components can create at most two local maxima in spectrum.
 Thus a minimum of three components seems indicated by (he simultaneous subset of the data., Thus a minimum of three components seems indicated by the simultaneous subset of the data.
 The comparison between a (wo component fit and a three component fit is given in Figure |., The comparison between a two component fit and a three component fit is given in Figure 1.
 The fits are made by letting all the SAA component parameters vary. (hen minimizing the sum of the squares of the residuals between the simultaneous cata. and the models.," The fits are made by letting all the SAA component parameters vary, then minimizing the sum of the squares of the residuals between the simultaneous data and the models."
 The deviation of the (wo component model ancl Che data is bevoncd what is expected based on the statistical uncertainty in the data., The deviation of the two component model and the data is beyond what is expected based on the statistical uncertainty in the data.
" Each VLA data point can be treated as an independent statistical variable. δν, 7=1.2.3.4.5 that is normally distributed with a mean and standard. deviation given bv the flux density value and the flux density error in Table 1l. respectively."," Each VLA data point can be treated as an independent statistical variable, $X_{i}$, $i=1,2,3,4,5$ that is normally distributed with a mean and standard deviation given by the flux density value and the flux density error in Table 1, respectively."
 The (wo pints that are more than lo from the mean only have about a chance of occurring by random chance and indicate the inability to fit the local maximum in that region., The two pints that are more than $\sigma$ from the mean only have about a chance of occurring by random chance and indicate the inability to fit the local maximum in that region.
 Consider the null hypothesis that the five data points are fit bv the model., Consider the null hypothesis that the five data points are fit by the model.
 Each data point provides an independent statistical variable ancl probability., Each data point provides an independent statistical variable and probability.
 Therefore. comparing the models in Fieure 1 to the data implies that the probability. of rejecting the null hypothesis is 0.99 for a (wo component fit aad 0.512 for a three component fit.," Therefore, comparing the models in Figure 1 to the data implies that the probability of rejecting the null hypothesis is 0.99 for a two component fit and 0.512 for a three component fit."
 Thus. the three component fit cannot be rejected on a statistical basis. but the two component fit is rejected with a very high statistical signifieauce.," Thus, the three component fit cannot be rejected on a statistical basis, but the two component fit is rejected with a very high statistical significance."
 Restricting the discussion io the simultaneous VLA data removes (wo uncertainties for the fits of the December 6 data in Table 1., Restricting the discussion to the simultaneous VLA data removes two uncertainties for the fits of the December 6 data in Table 1.
 First of all. the data is simultaneous. so there is no issue of variability within," First of all, the data is simultaneous, so there is no issue of variability within"
"Canuna-Rav Bursts (CRBs) ire attributed to a release of ⋜↧↖⇁↸∖↥⋅⋅↖↽↕⋜∐⋅∶↴⋁↸∖⋜⋯↕∪∏∐↑∪↕⋟↸∖∐↸∖↥⋅∶↴⋁⋅↖↽≺∿↕∩⊽⊳⊥ ""ES10 eres) into. a sinall region of space (S100 kin) over a short period of . ↽ ↑∐⊔↸∖⋖∿↓∩↓∩−↴∖↴↕∪↥⋅↕∪∐∶↴∙⊾≼∶↕⊰↕≧↴∖↴⋜↧∐≼⋞⊰↽⋅≻⋟ ⊇↴∖↴↕≯∪↥⋅↴∖↴∐∪↥⋅↑ ≼∶↕⊰↕≧","Gamma-Ray Bursts (GRBs) are attributed to a release of a very large amount of energy $\sim 10^{51}-10^{52}$ ergs) into a small region of space $\lesssim 100$ km) over a short period of time $\sim
10-10^2$ s for long GRBs and $\lesssim 2$ s for short GRBs)."
↴∖↴⋝∙↽∕∏∐∖↴∖↴↸∖↸∖∐↸∖↥⋅∶↴∙⊾↸∖↑↕↸⊳↸∖↖⇁↸∖∐↑↴∖↴∐⋜↕↖⇁↸∖↑↖↖⇁∪↸⊳∐⋜∐⋅⋜↧↸⊳↑↸∖↥⋅↕↴∖↴↑↕↸⊳ ↥⋅⋜∥∐⋜↧↑↕↖↽↸∖↴∖↴↕∶↴∙⊾∐⋜↧↑⋃⋅↸∖↴∖↴∶⋖↕⋝↑∐↸∖↻↥⋅∪∐∏≻↑↷↴≓↥⋅, These energetic events have two characteristic radiative signatures: (i) the prompt $\gamma$ -ray and (ii) the afterglow emission.
⋜↕⋅↖↽⋜⋯≼⊔∏⋟↑⊔∖ ⋜↧↕⋟↑↸∖↥⋅∶↴∙⊾↕∪↖↖⇁↸∖∐∐↴∖↴↴∖↴↕∪∐∙↽∕∏∐∖≼∐∖↑↸∖↸⊳↑↕∪↕↕∪↕⋟↕∐∶↴⋁↕↸∖∐↸∖↥⋅∶↴∙⊾⋅↖⇁↻∐∪↑∪∐↴∖↴ ⋖←↽⋅∖↙∖∖↕⋀∖↕↸∖∖⊽⋝ ↕∐∏≻∐↸∖↴∖↴↴∖↴≺∏∐⋅↸⊳↸∖↴∖↴∪↕⋟↥⋅⋜∥∐⋜↧↑↕∪∐⋯∪↖↽↕∐∶↴∙⊾⋜↧↑ ↥⋅↸∖↕⋜↧↑↕↖↽↕↴∖↴↑↕↸⊳↴∖↴⋉∖↸∖≼↴∖↴↖↖↽↕↑∐∫⇀∪↥⋅↸∖∐↑∑↕≯⋜↧↸⊳↑," The detection of high energy photons $\epsilon_{\gamma} \gg 1$ MeV) implies sources of radiation moving at relativistic speeds with Lorentz factors $\Gamma$ exceeding 100 \citep{fenimore93,
lithwick01,Piran99}."
∪↥⋅↴∖↴↕⇁↸∖↘↸⊳↸∖↸∖≼∐∐∶↴⋁↕ ⋖∎⋅↗∎⋅↗∎⋅↗⋝∙∙ ↖↖⊽∐∐↸∖⋯⋜⋯⋅↖⇁↕↴∖∷∖↴⋯∖↴∖↴↸⊳∪∐↸⊳↸∖↥⋅↕∐∐∶↴⋁↑↕∐∖↻↥⋅∪∐∩↑↸∖∐∏↴∖∷∖↴↕∪↕ ⋜∐⋅↸∖↴∖↴↑↕∐∪↻↸∖∐∙↑∐∖⋜↧↕⋟↑↸∖↥⋅∶↴∙⊾↕∪↖↖⇁∙↕∙↸∖∙↑∐↸∖↕∪↖∏∖↥⋅↸∖∐↸∖↥⋅∶↴∙⊾⋅↖⇁↕∪↓≼↴⋁ ⋜↧↴∖↴↑↕∐∶↴∙⊾↸∖↕⊔↕↴∖∷∖↴↕∪∐⋅↕↴∖↴↴⋝↸∖∐↸∖↖↽↸∖≺↧↑∪⋜∐⋅↕↴∖↴↸∖↕≯↥⋅∪⊔↑∐↸∖↕∐↑↸∖↥⋅⋜↧↸⊳↑↕∪↕ of the relativistic cjecta with the ambicut uatter and can be adequately described by the so-callec “standard nodel (???)..," While many issues concerning the prompt emission are still open, the afterglow, i.e. the lower energy long lasting emission, is believed to arise from the interaction of the relativistic ejecta with the ambient matter and can be adequately described by the so-called `standard' model \citep{rees_mesz92, paczynski_rhoads93,
mesz_rees97}."
 According to this. the rclativistic blast wave xoduced. from. the explosion can cucreize the external uediunm. ic. accelerate electrons (and possibly protons) to Hel eucreies and generate magnetic fields.," According to this, the relativistic blast wave produced from the explosion can energize the external medium, i.e. accelerate electrons (and possibly protons) to high energies and generate magnetic fields."
 The relativistic clectrous radiate by svuchrotroun and inverse Compton radiation which is esseutially the observed afterglow cnussion., The relativistic electrons radiate by synchrotron and inverse Compton radiation which is essentially the observed afterglow emission.
 However. in order to calculate the radiated. photon spectra. one needs a detailed prescription of the electron distribution function and of the maeguetic field.," However, in order to calculate the radiated photon spectra, one needs a detailed prescription of the electron distribution function and of the magnetic field."
 This is usually done by postulating that the electrons rave a power law distribution between a minima (5j) and a maxiuun (spas) cutoff with an overall energy density content which is aa fixed fraction (usually denoted * €.) of the total post-shoclk internal energyv deusity ( while au analogous argument can be nade for the nagnetic field euergv which takes a fraction ep of Ü., This is usually done by postulating that the electrons have a power law distribution between a minimum $\gammamin$ ) and a maximum $\gammamax$ ) cutoff with an overall energy density content which is a a fixed fraction (usually denoted by $\ee$ ) of the total post-shock internal energy density $U$ while an analogous argument can be made for the magnetic field energy which takes a fraction $\eB$ of $U$.
 A sijenificaut amount of work has beeu performed bv nuiuiv researchers in calculating the multiwavelcueth spectra aud ight curves of GBD afterelows either based directly on the above prescription (222222) oy using different variations erm.," A significant amount of work has been performed by many researchers in calculating the multiwavelength spectra and light curves of GRB afterglows either based directly on the above prescription \citep{dermchiang98, sariPiran98,
panai_mesz98,wijers_galama99, dermbottcher00, panai_kumar00,
granotsari02} or using different variations \citep{granot_kumar06,
fan_piran06, panaietal06, zhang06, nousek06}."
 Tn the present paper we focus ou the effects that a low sanax will bringo on the multiwavelenetho spectra aud helt curvesο... of the afterelow emissiou., In the present paper we focus on the effects that a low $\gammamax$ will bring on the multiwavelength spectra and light curves of the afterglow emission.
" This has not been treated thus far as it is maplicitlv assumed that 5,444 ds very large aud its radiative siguature does not contribute to any observable baud.", This has not been treated thus far as it is implicitly assumed that $\gammamax$ is very large and its radiative signature does not contribute to any observable band.
 However. ifit has à low value. then a break nieht appear successively in various energy bauds of diminishing energy as the svuchrotron component gives eradually its position to the SSC' one.," However, if it has a low value, then a break might appear successively in various energy bands of diminishing energy as the synchrotron component gives gradually its position to the SSC one."
 This will produce igh curves which are not auy more pure power laws but wave more complicated shapes., This will produce light curves which are not any more pure power laws but have more complicated shapes.
 The paper is structured as follows., The paper is structured as follows.
 Is 822 we describe he principles of the model aud discuss. in a qualitative uv. some of the results.," Is 2 we describe the principles of the model and discuss, in a qualitative way, some of the results."
" In 8323 we quautify the above acl we derive some analytical relations between the itial xuaneters,. which. when satisfied. will produce cliffercut vpes of N-raw light curves."," In 3 we quantify the above and we derive some analytical relations between the initial parameters, which, when satisfied, will produce different types of X-ray light curves."
 In ll we make a tentative connection of our results to observations., In 4 we make a tentative connection of our results to observations.
 Finally in 855 we conclude aud give a brief discussion of the basic poiuts of the present work., Finally in 5 we conclude and give a brief discussion of the basic points of the present work.
 The general framework of the model we preseut here is based on the standard CRB afterelow ποσο]. albeit with some modifications regarding mainly he approach to the Xivsical problems (?? henceforth PMO9).," The general framework of the model we present here is based on the standard GRB afterglow model, albeit with some modifications regarding mainly the approach to the physical problem \cite{fanetal08, me09} – henceforth PM09)."
 For the sake of completeness we repeat here its basic premises: as the Relativistic Blast Wave (RBW) usually associated with CRB afterelows is expanding. it accelerates by sole unspecified mechanisin electrous of the eircuustellar οςτα to hieh cnereies.," For the sake of completeness we repeat here its basic premises: as the Relativistic Blast Wave (RBW) usually associated with GRB afterglows is expanding, it accelerates by some unspecified mechanism electrons of the circumstellar medium to high energies."
 These are assumed to be injected did the shock front ina volue of radius R containiug a taneled maeuetic field D. The particles suffer radiative and adiabatie losses. evolving with radius.," These are assumed to be injected behind the shock front in a volume of radius R containing a tangled magnetic field B. The particles suffer radiative and adiabatic losses, evolving with radius."
 At the same nue they enüt radiation. by svuchrotron aud ilverse Compton radiation., At the same time they emit radiation by synchrotron and inverse Compton radiation.
" Therefore. at each radius there is a coupling between clectrous aud photons. iu the sense hat the electron distributiou function determines the Xioton spectrum and. at the same time. the photons determine the electron distribution functiou through inverse Comptou losses and. possibly, pair reinjection."," Therefore, at each radius there is a coupling between electrons and photons, in the sense that the electron distribution function determines the photon spectrum and, at the same time, the photons determine the electron distribution function through inverse Compton losses and, possibly, pair reinjection."
 The usual procedure of approaching the problem is, The usual procedure of approaching the problem is
metallicities obtaimed by DP for two nain clouds (at c — (lans 1 bande=— d3]an 3) differ nearly by two orders of maguitude: [N/T] = 0.9 aud 2.7. respectively (in ourcase [N/I| &2.2 for the whole systeni).,"metallicities obtained by DP for two main clouds (at $v$ = 0 km $^{-1}$ and $v = -43$ km $^{-1}$ ) differ nearly by two orders of magnitude: [X/H] = –0.9 and –2.7, respectively (in our case [X/H] $\simeq -2.2$ for the whole system)."
 As shown in Paper L the Voiet fitting uav in general vield correct cohunn densities when applied to musaturated Hines. but the mean C aud. rence. the ionization corrections may not be unambiguous.," As shown in Paper I, the Voigt fitting may in general yield correct column densities when applied to unsaturated lines, but the mean $U$ and, hence, the ionization corrections may not be unambiguous."
 Therefore. the conclusion made by DP that the za;%1.92 system contains a region of iuteuse star-formation activity may not be well justified since this result is model depende-t.," Therefore, the conclusion made by DP that the $z_{\rm abs} \simeq 1.92$ system contains `a region of intense star-formation activity' may not be well justified since this result is model dependent."
 The values of the average gas deusitv ay and kinetic temperature Tj. aud the cloud thickness £ estimated iu ΟΥ nodel (see Table 1) are typical for the La-a svstenis discussed in the literature (6.5. Cüallougo Petitjeau 1991: Viegas et al.," The values of the average gas density $n_0$ and kinetic temperature $T_{\rm kin}$, and the cloud thickness $L$ estimated in our model (see Table 1) are typical for the $\alpha$ systems discussed in the literature (e.g., Giallongo Petitjean 1994; Viegas et al."
 1999: Prochaska Durles 1999: Chen et al.," 1999; Prochaska Burles 1999; Chen et al.,"
 1998: Chen et al., 1998; Chen et al.
 2001)., 2001).
 Low iuetallicitv. for the whole system (X/I] < 2.0) aud its dimension of 20 kpc nuplv that this svstem can originate in a ealactic halo or ina large scale structure object., Low metallicity for the whole system ([X/H] $< -2.0$ ) and its dimension of 20 kpc imply that this system can originate in a galactic halo or in a large scale structure object.
 This svstem exhibits a plenty of metal lines in differcut ionization stages., This system exhibits a plenty of metal lines in different ionization stages.
 The metal profiles are not very couples. and extend over the velocity range from 100 kins 4 to 100 lan +., The metal profiles are not very complex and extend over the velocity range from $-100$ km $^{-1}$ to 100 km $^{-1}$.
 Results obtained with the MCI are presented in Table 1 and shown iu Fies., Results obtained with the MCI are presented in Table 1 and shown in Figs.
 1 and 5., 4 and 5.
 As in the previous system. most absorption features can be well described with wniform metallicities aud a common TAI spectrim.," As in the previous system, most absorption features can be well described with uniform metallicities and a common HM spectrum."
 The La-a profile is contaminated by the forest absorption in the due and red wines aud therefore the Ίντα absorption feature was not involved iu the analysis., The $\alpha$ profile is contaminated by the forest absorption in the blue and red wings and therefore the $\alpha$ absorption feature was not involved in the analysis.
 The profiles of A2796 and A1670 were coni]tec ater using the derived velocity and density distributions., The profiles of $\lambda2796$ and $\lambda1670$ were computed later using the derived velocity and density distributions.
 A2796 1s contaminated by a telluric line aud this explains the difference between the computed and observed profiles., $\lambda2796$ is contaminated by a telluric line and this explains the difference between the computed and observed profiles.
 The svuthetie and observed profiles of ALGTO show unuch more pronounced discrepancy., The synthetic and observed profiles of $\lambda1670$ show much more pronounced discrepancy.
 Fractional jonisation curves for and were computed with CLOUDY., Fractional ionisation curves for and were computed with CLOUDY.
 These curves allowed us to fit the doublet quite well with the Al abundance similar to that obtained or the other metals., These curves allowed us to fit the doublet quite well with the Al abundance similar to that obtained for the other metals.
 However. when the profile was iuclided in the fitting. the Al aetallicity differed o» order of maguitude from the other metals.," However, when the profile was included in the fitting, the Al metallicity differed by order of magnitude from the other metals."
 Besides it was impossible to fit adequately the doublet., Besides it was impossible to fit adequately the doublet.
 Similar behaviow of Al was reported also by DP who roted hat ‘the recombination coefficients used to compute he aluminum ionisation equilibrium jin CLOUDY] are xobablyv questionable’., Similar behaviour of Al was reported also by DP who noted that `the recombination coefficients used to compute the aluminium ionisation equilibrium [in CLOUDY] are probably questionable'.
 Coluun densities derived by DP coincide well (within 155€4)) with that obtained in our procedure except for the saturated SiHLAL206 lune for which the Voigt fitting eave nearly 2 times ower value., Column densities derived by DP coincide well (within ) with that obtained in our procedure except for the saturated $\lambda1206$ line for which the Voigt fitting gave nearly 2 times lower value.
 The abuudauces estimated in DP scatter again frou component to component. but nevertheless hey conclude that ‘the gas in this system is Likely of quite Heh metallicity (larger than 0.1 solar).," The abundances estimated in DP scatter again from component to component, but nevertheless they conclude that `the gas in this system is likely of quite high metallicity (larger than 0.1 solar)'."
 Similar to the Voigt fitting. the MCT also delivered or this svstem high metal abuudancies: one third solar or carbon and silicon and nearly two times lower for itrogeon. magnesia aud aluminium.," Similar to the Voigt fitting, the MCI also delivered for this system high metal abundancies: one third solar for carbon and silicon and nearly two times lower for nitrogen, magnesium and aluminium."
 Takine iuto account his result aud a compact climension (25 kpc. see Table 1) of the absorbing region we come to the same couclusiou as Prochaska Durles (1999) did: the svstem at 2=1.91 can hardly be a large scale structure object (like a filamieut or a wall) and should be related to a galactic svsteim (av be a region of intense star formation).," Taking into account this result and a compact dimension $\simeq 5$ kpc, see Table 1) of the absorbing region we come to the same conclusion as Prochaska Burles (1999) did: the system at $z = 1.94$ can hardly be a large scale structure object (like a filament or a wall) and should be related to a galactic system (may be a region of intense star formation)."
 This is the uost interesting svsteni from the family of the absorbers at =1.9 toward J2233606., This is the most interesting system from the family of the absorbers at $z = 1.9$ toward J2233–606.
 The metal üσαie profiles show a rather couples structure extending OVor the velocity range of about 700 kin 1, The metal line profiles show a rather complex structure extending over the velocity range of about 700 km $^{-1}$.
 Some of these profiles are severcly bleuded that hampers the unique Voigt profile deconvolution (c.g. DP assumed. 17 coniponenuts to describe metal profiles)., Some of these profiles are severely blended that hampers the unique Voigt profile deconvolution (e.g. DP assumed 17 components to describe metal profiles).
 The MCT code turned out to be much more robust Hid was able to recover the self-cousisteut Lue profiles even under such uufavourable conditions., The MCI code turned out to be much more robust and was able to recover the self-consistent line profiles even under such unfavourable conditions.
 The physical parameters which the AICI delivered for the +=LAT svete together with the underlying velocity and density distributions are prescuted in Table 1 aud in Figs., The physical parameters which the MCI delivered for the $z = 1.87$ system together with the underlying velocity and density distributions are presented in Table 1 and in Figs.
 6 aud 7., 6 and 7.
 Tt is seen from Fie., It is seen from Fig.
 6 that like iu the previous two svstenis all lines ave well described with a single parameter set. uuiforim moetallicities aid a common DAL UV background.," 6 that like in the previous two systems all lines are well described with a single parameter set, uniform metallicities and a common HM UV background."
 The blue wing of the La-a line is coutaminated by the forest absorption as is clearly seen from the La.) aud, The blue wing of the $\alpha$ line is contaminated by the forest absorption as is clearly seen from the $\beta$ and
"Stars can experience a core-collapse supernova only if they are born with a mass much higher than the Sun, even though a chemically evolved core of order only one solar mass is required to ienite the advanced burning stages.","Stars can experience a core-collapse supernova only if they are born with a mass much higher than the Sun, even though a chemically evolved core of order only one solar mass is required to ignite the advanced burning stages."
 The culprit is mass loss severe enough to strip off the stellar envelope before a sufficiently massive processed core develops., The culprit is mass loss severe enough to strip off the stellar envelope before a sufficiently massive processed core develops.
" Although there is some mass loss even in shell hydrogen burning giants, the vast majority in intermediate mass stars occurs in the presence of thermal pulsations involving interactions between hydrogen and helium burning shells: we refer to this as the thermally pulsing AGB (TP-AGB) phase."," Although there is some mass loss even in shell hydrogen burning giants, the vast majority in intermediate mass stars occurs in the presence of thermal pulsations involving interactions between hydrogen and helium burning shells; we refer to this as the thermally pulsing AGB (TP-AGB) phase."
" Our understanding of even the basic properties of the TP-AGB phase, such as the lifetime or light emitted, is limited because their severe mass loss is difficult to constrain observationally or predict theoretically."," Our understanding of even the basic properties of the TP-AGB phase, such as the lifetime or light emitted, is limited because their severe mass loss is difficult to constrain observationally or predict theoretically."
 The uncertainties in TP-AGB evolution have profound consequences for stellar population studies., The uncertainties in TP-AGB evolution have profound consequences for stellar population studies.
 Stellar interiors models can accurately predict evolutionary properties up to the onset of, Stellar interiors models can accurately predict evolutionary properties up to the onset of
showed however that the effect of including quadratic terms is marginal (as was suspected). aud does not account for the observed discrepancy.,"showed however that the effect of including quadratic terms is marginal (as was suspected), and does not account for the observed discrepancy."
 Considering that our current result -which iucludes re-processed data from ALP- agrees quite well with meastwements by other groups. we conclude that ALP’s results might be affected by an unidentified systematic error in Decl.," Considering that our current result -which includes re-processed data from ALP- agrees quite well with measurements by other groups, we conclude that ALP's results might be affected by an unidentified systematic error in Decl."
 Since in the present work we used ALP's unmocilied (x.y) coordinates. we believe that this error could have originated iu the processing of the Decl.," Since in the present work we used ALP's unmodified (x,y) coordinates, we believe that this error could have originated in the processing of the Decl."
 PM instead of the coordinates themselves., PM instead of the coordinates themselves.
 It should be pointed out that the UCAC2-based result from Momany ZageiaMD (2005). which is consistent. with that from ALP. is currently also cousidered to be allected by au as yet unidentified systematic error (Momauy Zageia.MOD 2005: Wallivayalil et al..," It should be pointed out that the UCAC2-based result from Momany Zaggia (2005), which is consistent with that from ALP, is currently also considered to be affected by an as yet unidentified systematic error (Momany Zaggia, 2005; Kallivayalil et al.,"
 2005)., 2005).
 We would finally like to note that our new result for field QO159-6127 is consistent with PAM., We would finally like to note that our new result for field Q0459-6427 is consistent with PAM.
 Lyudeu-Bell Lyudeu-Bell (1995). have proposed that the LMC. together with the SMC. Draco auc Ursa Minor. anc possibly Carina and Sculptor. define a stream of galaxies witli similar orbits around our galaxy.," Lynden-Bell Lynden-Bell (1995), have proposed that the LMC, together with the SMC, Draco and Ursa Minor, and possibly Carina and Sculptor, define a stream of galaxies with similar orbits around our galaxy."
 Their moclels predict a PM lor each of inember of the stream. which cau be compared to their measured PAL (ο evaluate the reality of the stream.," Their models predict a PM for each of member of the stream, which can be compared to their measured PM to evaluate the reality of the stream."
 For the LMC they predict PM components of [p cosó.q15] = [41.5.0] mas |. giving a total PM of jp = 41.5 mas +. with a position angle of 0 = 90°.," For the LMC they predict PM components of $\mu_{\alpha}$ $\delta$ $\mu_{\delta}$ ] = [+1.5,0] mas $^{-1}$, giving a total PM of $\mu$ = +1.5 mas $^{-1}$ , with a position angle of $\theta$ = $^\circ$."
 A comparison of this prediction with our result [jj = (42.040.1) mas 1g [2£3.1)7]. shows that our measured. values of ji and ? are. respectively. 5.16 and 8.96 away [rom the predicted values.," A comparison of this prediction with our result $\mu$ = $+$ $\pm$ 0.1) mas $^{-1}$ , $\theta$ = $\pm$ $^\circ$ ], shows that our measured values of $\mu$ and $\theta$ are, respectively, $\sigma$ and $\sigma$ away from the predicted values."
 This result iudicates that the LMC does not seem to be a member of the above stream (it is worth meutioniug that Piatek et al. (, This result indicates that the LMC does not seem to be a member of the above stream (it is worth mentioning that Piatek et al. (
2005). using HSTdata. have concluded that Ursa Minor is not a member ofthis,"2005), using HSTdata, have concluded that Ursa Minor is not a member ofthis"
"astrophysical implications of (he expression (29)) we shall consider separately (he (wo limiting cases. corresponding {ο viscous aid non-viscous neutron matter. respectively,","astrophysical implications of the expression \ref
{cond1}) ) we shall consider separately the two limiting cases, corresponding to viscous and non-viscous neutron matter, respectively."
 In the case of ideal neutron matter. the condition of the formation of a quark bubble is eiven by Eq. (30)).," In the case of ideal neutron matter, the condition of the formation of a quark bubble is given by Eq. \ref{cond2}) )."
 This equation determines al which temperature the phase (Gransition takes place., This equation determines at which temperature the phase transition takes place.
" For example. in (he case of the Walecka mean field equation of state ).. with A=15758 and 5= 4.951994).. the temperature 7; necessary lor (he formation of a stable quark bubble in a second (A/=1I s). in a volume AV, with radius AR,=1 km GM=tr/3 km). and at the center of the neutron star. is around. T.=12—13 MeV. for o=(85 MeV)""."," For example, in the case of the Walecka mean field equation of state , with $K=15758$ and $%
\gamma =4.9 5, the temperature $T_{c}$ necessary for the formation of a stable quark bubble in a second $\Delta t=1$ s), in a volume $\Delta V_{0}$ with radius $\Delta R_{0}=1$ km $\Delta V_{0}=4\pi /3$ $%
^{3} ), and at the center of the neutron star, is around $T_{c}=12-13$ MeV, for $\sigma =(85$ $)^{3}$."
" Here we have also assumed the standard values for the quark matter chemical potential and for the bag constant. jp,=qj280 MeV and B=60 *1998).."," Here we have also assumed the standard values for the quark matter chemical potential and for the bag constant, $\mu
_{u}=\mu _{d}=280$ MeV and $B=60$ $^{-3}$."
 Since the temperature of a newborn neutron star is around 10 MeV.1994)... this value of the surface tension seenis (o rule out. for the given equation of state. the possibility of formation of quark nuclei in the neutron matter.," Since the temperature of a newborn neutron star is around $10$ MeV, this value of the surface tension seems to rule out, for the given equation of state, the possibility of formation of quark nuclei in the neutron matter."
 ILowever. a small variation in the value of o. 6=(75 MeV)? reduces ihe temperature for quark nuclei formation to 7;.=4—6 MeV. a value which does not exclude the possibilitv that quark iuclei formation can be imitiated in the early stages of the Smaller values of o.," However, a small variation in the value of $\sigma $, $\sigma =(75$ $%
)^{3} reduces the temperature for quark nuclei formation to $T_{c}=4-6$ MeV, a value which does not exclude the possibility that quark nuclei formation can be initiated in the early stages of the neutron star evolution."
" of the order of o=(50 MeV)"".evenneutronstarevolution. can lower more (he phase (ransilion temperature. making it possible even for cold neutron stars."," Smaller values of $\sigma $, of the order of $\sigma =(50$ }$, can lower even more the phase transition temperature, making it possible even for cold neutron stars."
 The (ransilion temperature is relatively insensitive to the details of the equation of state of the dense neutron matter., The transition temperature is relatively insensitive to the details of the equation of state of the dense neutron matter.
 If the temperature at which the phase Gransition is initiated is strongly dependent by the microscopic model adopted [ον describing the neutron matter properties. (he total transition lime of the neutron star to a quark star. given by. Eq. (35)).," If the temperature at which the phase transition is initiated is strongly dependent by the microscopic model adopted for describing the neutron matter properties, the total transition time of the neutron star to a quark star, given by Eq. \ref{conv}) ),"
 is much less dependent on the details of the equation of state or the exact numerical values of (he surface tension or curvature of the bubble., is much less dependent on the details of the equation of state or the exact numerical values of the surface tension or curvature of the bubble.
 Generally. for à neutron star with radius 2=12—15 km. the conversion (ime is verv long. of the order of 105 vears for σ22 (75 MeV)*.," Generally, for a neutron star with radius $R=12-15$ km, the conversion time is very long, of the order of $10^{6}$ years for $%
\sigma \approx (7 5 $)^{3}$."
" For σ&(50 MeV)? or even lower. the conversion (ime could have values of the order of LO"" vears. which are still of the same order of magnitude as the lifetime of a neutron star (10° vears)."," For $\sigma \approx (50$ $)^{3}$ or even lower, the conversion time could have values of the order of $10^{7}$ years, which are still of the same order of magnitude as the lifetime of a neutron star $10^{7}$ years)."
 Hence. in (he present approach. which is limited to the consideration of linear effects only in the hvdrodynamie description of the first order phase (ransiGions for an ideal neutron Πιτ. the erowth of the quark nuclei is verv slow.," Hence, in the present approach, which is limited to the consideration of linear effects only in the hydrodynamic description of the first order phase transitions for an ideal neutron fluid, the growth of the quark nuclei is very slow."
 Therefore the results of the present analysis suggest (hat even quark nuclei could appear inside the neutron star al a very early stage. the transition to a pure quark phase ends only in the final stage of its evolution.," Therefore the results of the present analysis suggest that even quark nuclei could appear inside the neutron star at a very early stage, the transition to a pure quark phase ends only in the final stage of its evolution."
 If the neutron matter-quark matter phase transition is mainly driven by viscousprocesses. like in the Csernai-Ixapusta scenario. (hen the time necessary for a neutron star lo convert {ο a quark star follows from Eqs. (28))," If the neutron matter-quark matter phase transition is mainly driven by viscousprocesses, like in the Csernai-Kapusta scenario, then the time necessary for a neutron star to convert to a quark star follows from Eqs. \ref{k1}) )"
 and (34))., and \ref{tconv}) ).
 In order to numerically estimate the conversion, In order to numerically estimate the conversion
to the separation.,to the separation.
 Closure phases are derived from the baseline phases as a simple summation of the three phases corresponding to a closed triangle of holes., Closure phases are derived from the baseline phases as a simple summation of the three phases corresponding to a closed triangle of holes.
" Companion searches are performed by fitting the various binary parameters (r, a, and 6) to the measured closure phases, so the detection limit obtained from any given dataset is directly related to the signal-to-noise ratio of the closure phase."," Companion searches are performed by fitting the various binary parameters $r$ , $\alpha$ , and $\delta$ ) to the measured closure phases, so the detection limit obtained from any given dataset is directly related to the signal-to-noise ratio of the closure phase."
" A Monte-Carlo simulation was performed to test the theoretical detection limit, as a function of the angular separation and closure phase measurement error bars."," A Monte-Carlo simulation was performed to test the theoretical detection limit, as a function of the angular separation and closure phase measurement error bars."
 The result is presented in Fig. 4.., The result is presented in Fig. \ref{Monte}.
" 'The detection limit is almost constant for any separation larger than A/D. The Monte-Carlo simulation shows that the dynamic range can be approximated as (as a function of c(C.P) in degrees) Detections below A/D can still be significant, but the contrast achieved decreases rapidly with the separation."," The detection limit is almost constant for any separation larger than $\lambda$ /D. The Monte-Carlo simulation shows that the dynamic range can be approximated as (as a function of $\sigma(CP)$ in degrees) Detections below $\lambda$ /D can still be significant, but the contrast achieved decreases rapidly with the separation."
" At half A/D, limits are worse by a factor of two, while at A/D the degradation in performance is a factor 15."," At half $\lambda$ /D, limits are worse by a factor of two, while at $\lambda$ /D the degradation in performance is a factor 15."
" For separations beyond the formal diffraction limit, any fitting process becomes highly model dependent and often strong degeneracies between fitted parameters can appear."," For separations beyond the formal diffraction limit, any fitting process becomes highly model dependent and often strong degeneracies between fitted parameters can appear."
 A typical example at small separations is the covariance between the flux ratio and separation: it becomes difficult to measure these two parameters independently., A typical example at small separations is the covariance between the flux ratio and separation: it becomes difficult to measure these two parameters independently.
" 1135549 was initially observed for the purpose of a PSF reference star, but it turned out to be a binary system with a contrast ratio around and a separation close to 180 mas."," 135549 was initially observed for the purpose of a PSF reference star, but it turned out to be a binary system with a contrast ratio around and a separation close to 180 mas."
" We present the data in this paper for the sake of introducing the technique of detecting faint companions with SAM illustrated by strong, clear systematic signals."," We present the data in this paper for the sake of introducing the technique of detecting faint companions with SAM illustrated by strong, clear systematic signals."
 The free parameters of Eq. (1)), The free parameters of Eq. \ref{eq}) )
" are the binary coordinates o and 6, along with the flux ratio r."," are the binary coordinates $\alpha$ and $\delta$ , along with the flux ratio $r$ ."
 A map of the reduced x? is derived as a function of a and 6 (Fig. 5))., A map of the reduced $\chi^2$ is derived as a function of $\alpha$ and $\delta$ (Fig. \ref{plot_map}) ).
 It indicates the presence of a companion with a strong likelihood at a separation of zz180 mmas to the southeast of the star., It indicates the presence of a companion with a strong likelihood at a separation of $\approx 180$ mas to the southeast of the star.
 The best-fit model is compared to the observed data in the two panels of Fig. 6.., The best-fit model is compared to the observed data in the two panels of Fig. \ref{plot_res}.
 The lefthand panel shows the observed versus model-predicted closure phase data., The lefthand panel shows the observed versus model-predicted closure phase data.
 The final reduced x? is equal to 1.85 while the average residual is 0.28 degree over a total of 35x15=525 closure phase measurements., The final reduced $\chi^2$ is equal to 1.85 while the average residual is 0.28 degree over a total of $35\times15=525$ closure phase measurements.
" According to Eq. (2)),"," According to Eq. \ref{dyn}) ),"
 this dataset therefore has the capability of a detection limit of 7.5x107 detection)., this dataset therefore has the capability of a detection limit of $7.5\times10^{-4}$ $\sigma$ detection).
 The righthand panel shows phases recovered from the closure phases., The righthand panel shows phases recovered from the closure phases.
 This is done by inverting the phase to closure phase matrix., This is done by inverting the phase to closure phase matrix.
" The matrix is not fully invertible; for example, tip-tilt is not constrained by the closure phase."," The matrix is not fully invertible; for example, tip-tilt is not constrained by the closure phase."
" Therefore, unknown parameters of the inversion were taken to match the model."," Therefore, unknown parameters of the inversion were taken to match the model."
 The best fit model is over-plotted as a solid line., The best fit model is over-plotted as a solid line.
 The first 15 lines of Table 2 results from the processing of each SCI-CAL acquisition pair separately., The first 15 lines of Table \ref{result_bin} results from the processing of each SCI-CAL acquisition pair separately.
" An independent model is fitted at each epoch, and the detection can be seen to be rotating in the detector frame (but fixed on the sky)."," An independent model is fitted at each epoch, and the detection can be seen to be rotating in the detector frame (but fixed on the sky)."
 A strong indication that our detection is indeed real lies in the rotation in the detector frame following the parallactic angle evolution., A strong indication that our detection is indeed real lies in the rotation in the detector frame following the parallactic angle evolution.
" Because the telescope is azimuthal, it is unlikely that a calibration error (including optical aberrations) rotate in such a motion."," Because the telescope is azimuthal, it is unlikely that a calibration error (including optical aberrations) rotate in such a motion."
 This validation could not have been made with an equatorial telescope., This validation could not have been made with an equatorial telescope.
 The last line of Table 2 gives the best-fitting binary parameters., The last line of Table \ref{result_bin} gives the best-fitting binary parameters.
 Separation and dynamic ranges are obtained with high statistical precision., Separation and dynamic ranges are obtained with high statistical precision.
" The precision of the closure phase measurement over time is highlighted in Fig. 7,,"," The precision of the closure phase measurement over time is highlighted in Fig. \ref{fit},"
 in which each panel depicts a different closure phase., in which each panel depicts a different closure phase.
 The solid line corresponding to the model can be seen to almost perfectly match the evolution of the closure phase as a function of time., The solid line corresponding to the model can be seen to almost perfectly match the evolution of the closure phase as a function of time.
" As stated by Eq. 2,"," As stated by Eq. \ref{dyn},"
 the dynamic range is a pure function of the error on the phase., the dynamic range is a pure function of the error on the phase.
" The binary system was observed during good atmospheric conditions (0.7"" seeing and a coherent time of 8mms) with a relatively fast frame time of 200mms.", The binary system was observed during good atmospheric conditions (0.7” seeing and a coherent time of ms) with a relatively fast frame time of ms.
" Phases recovered from the interferograms show a median frame-to-frame error of 17 degrees, equivalent to a 180-nm piston jitter at our observing wavelength."," Phases recovered from the interferograms show a median frame-to-frame error of 17 degrees, equivalent to a 180-nm piston jitter at our observing wavelength."
" On the other hand, the frame-to-frame closure phase error is only 2.8 degrees."," On the other hand, the frame-to-frame closure phase error is only 2.8 degrees."
 This is a clear indication that most of the 180-nm rms error is due to the residual of atmospheric perturbation not corrected by the AO system., This is a clear indication that most of the 180-nm rms error is due to the residual of atmospheric perturbation not corrected by the AO system.
" Assuming that these residuals were caused by photon noise or readout noise, the remaining 2.8 degreesframe-error should be uncorrelated."," Assuming that these residuals were caused by photon noise or readout noise, the remaining 2.8 degreesframe-to-frameerror should be uncorrelated."
" In this case, thefinal erroron a whole acquisition sequence (8 dither of 200 frames) should improve as the square root of the"," In this case, thefinal erroron a whole acquisition sequence (8 dither of 200 frames) should improve as the square root of the"
properties.,properties.
" The A. 0,4, test can be useful to investigate this point.", The $R$ $\alpha_\mathit{bol}$ test can be useful to investigate this point.
 Again. it is important to define a solid measure of L5 for the fainter objects.," Again, it is important to define a solid measure of $L_\mathit{IR}$ for the fainter objects."
" Concerning local sources. the ιν density fi; is known to be a good indicator of the total IR luminosity. so (hat eL, (60 jim) is sometimes adopted to obtain an alternative definition of the ULIRG class."," Concerning local sources, the flux density $f_{60}$ is known to be a good indicator of the total IR luminosity, so that $\nu L_\nu$ (60 $\mu$ m) is sometimes adopted to obtain an alternative definition of the ULIRG class."
 By estimating Lg in this wav in our sample we indeed [ind an excellent agreement (within 0.1 dex) with equation (1))., By estimating $L_\mathit{IR}$ in this way in our sample we indeed find an excellent agreement (within 0.1 dex) with equation \ref{e1}) ).
 Ilence the rest-frame luminosity al ~50 pam ds a reliable proxy of Lyy αἱ 2κ0.35., Hence the rest-frame luminosity at $\sim$ 50 $\mu$ m is a reliable proxy of $L_\mathit{IR}$ at $z < 0.35$.
 We can argue how far backward this fiducial point can be shifted in wavelength., We can argue how far backward this fiducial point can be shifted in wavelength.
 In Paper II we have discussed the enhancement of the AGN over SD brightness ratio at 5.8 yan due to the hot dust component: we have also shown how rapidly it declines and completely vanishes al 30 jan. In this view. the 70 jan fluxes provided bv the Multiband. Imaging Photometer onboard (MIDPS: Rieke et al.," In Paper II we have discussed the enhancement of the AGN over SB brightness ratio at 5–8 $\mu$ m due to the hot dust component; we have also shown how rapidly it declines and completely vanishes at $\sim$ 30 $\mu$ m. In this view, the 70 $\mu$ m fluxes provided by the Multiband Imaging Photometer onboard (MIPS; Rieke et al."
 2004) allow to assess the luminosity of Ih ealaxies al 2~1 regardless of their power source., 2004) allow to assess the luminosity of IR galaxies at $z \sim 1$ regardless of their power source.
 Our basic assumption is Chat (he local toolbox (AGN/SB templates and bolometric corrections. extinction law. 30 jan rest-frame to IR ratio) can be used as a fair approximation in the analysis of the distant populations.," Our basic assumption is that the local toolbox (AGN/SB templates and bolometric corrections, extinction law, 30 $\mu$ m rest-frame to IR ratio) can be used as a fair approximation in the analysis of the distant populations."
 We have first applied a filler with narrow gaussian profile to 120 sources in our main sample al z«0.2. in order to derive their monochromatic 30 jan resi-[rame Duminositv and check the correlation with the total energv output. which turis out to be verv good.," We have first applied a filter with narrow gaussian profile to 120 sources in our main sample at $z < 0.2$, in order to derive their monochromatic 30 $\mu$ m rest-frame luminosity and check the correlation with the total energy output, which turns out to be very good."
" A near proportionality holds. and the best fit is obtained [or Fy,xανα "," A near proportionality holds, and the best fit is obtained for $F_\mathit{IR} \propto (f_{30\mu m})^{0.93}$."
We have then searched in the literature for faint IR. galaxies whose AIIPS fey (or fey) τιν allows a good sampling of the 309 som rest-frame band. ancl found 52 ULIBRG-like systems ab 0.5</\2L5.," We have then searched in the literature for faint IR galaxies whose MIPS $f_{70}$ (or $f_{60}$ ) flux allows a good sampling of the 30 $\mu$ m rest-frame band, and found 52 ULIRG-like systems at $0.5 < z < 1.5$."
 The 30 jan resi-frame luminosities have been extrapolated from the measured fhixes by assuming a power-law trend for the continuum. whose spectral index Pann=3.87z0.82 is the average over the local sources.," The 30 $\mu$ m rest-frame luminosities have been extrapolated from the measured fluxes by assuming a power-law trend for the continuum, whose spectral index $\Gamma_{30\mu m} = 3.37 \pm 0.82$ is the average over the local sources."
 The selection of the 52 objects is very heterogeneous. and the lack of completeness precludes the study of the AGN detection rate and global contribution.," The selection of the 52 objects is very heterogeneous, and the lack of completeness precludes the study of the AGN detection rate and global contribution."
 On the other hand. the great variety is a desirable feature for a control sample.," On the other hand, the great variety is a desirable feature for a control sample."
 A [few known hyperluminous LR galaxies (ILIRGs. Ly;>LOM L.) and quasar-like objects are also included: by applying our deconvolution model we are able to obtain an adequate fit in every case.," A few known hyperluminous IR galaxies (HLIRGs, $L_\mathit{IR}>10^{13} L_\odot$ ) and quasar-like objects are also included: by applying our deconvolution model we are able to obtain an adequate fit in every case."
 The properties of the sources. (heir parent samples ancl the results of our analvsis are summarized in Table 3..," The properties of the sources, their parent samples and the results of our analysis are summarized in Table \ref{t3}."
" The location in the 2 0,5, diagram of the high-redshift IR galaxies is shown in Fig. 1 (", The location in the $R$ $\alpha_\mathit{bol}$ diagram of the high-redshift IR galaxies is shown in Fig. \ref{hz}( (
fa).,a).
 Their distribution is in good agreement with the best fit [or ULIRGs., Their distribution is in good agreement with the best fit for ULIRGs.
 We stress that (his, We stress that this
feedback processes.,feedback processes.
 We will discuss AMve further in refsecimzr.. when we will relate it to the metallicity of the infalling eas.," We will discuss $\dot{M}_{\rm recyc}$ further in \\ref{sec:mzr}, when we will relate it to the metallicity of the infalling gas."
 his leaves the preventive feedback. parameter C. which we now consider.," This leaves the preventive feedback parameter $\zeta$, which we now consider."
 There are a number of sources of preventive feedback. each with its own dependence on halo mass and redshift. the products of which comprise the total ¢: Alultiplving these terms together. we obtain ολ{πμ} that is schematically illustrated. in Figure. 1..," There are a number of sources of preventive feedback, each with its own dependence on halo mass and redshift, the products of which comprise the total $\zeta$: Multiplying these terms together, we obtain $\zeta(M_{\rm halo})$ that is schematically illustrated in Figure \ref{fig:zeta}."
 lt is smal ab low and high halo masses owing to photo-suppression and quenching respectively. and. approaches unity a intermediate halo masses which is where vigorous star formation can occur.," It is small at low and high halo masses owing to photo-suppression and quenching respectively, and approaches unity at intermediate halo masses which is where vigorous star formation can occur."
 This depiction is intended only to illustrate broad. trends. as the actual values ancl functiona forms of the various ¢ terms are at best only qualitatively known.," This depiction is intended only to illustrate broad trends, as the actual values and functional forms of the various $\zeta$ terms are at best only qualitatively known."
 The general shape of this curve with a cutolf in accretion at low and high masses is long known2006).. although the quantitative masses for the cutolls are debate to this day2011).," The general shape of this curve with a cutoff in accretion at low and high masses is long known, although the quantitative masses for the cutoffs are debated to this day."
. There παν be other sources of energetic preventive feedback that retard accretion such as cosmic rays. stellar winds. quasar outfows. local photo-ionisation. and magnetic fields. but their importance for global galaxy. evolution has not vet been firmly. established.," There may be other sources of energetic preventive feedback that retard accretion such as cosmic rays, stellar winds, quasar outflows, local photo-ionisation, and magnetic fields, but their importance for global galaxy evolution has not yet been firmly established."
" There may also be subtle ""amplification"" ellects by which two or more preventive nmechanisnis serve to strengthen. each other beyond. their individual impact)."," There may also be subtle “amplification"" effects by which two or more preventive mechanisms serve to strengthen each other beyond their individual impact."
. Hence the list of individual C's above is intended to illustrate the sort of physical processes contributing to preventive feedback. and how they might manifest in the overall shape of ¢.," Hence the list of individual $\zeta$ 's above is intended to illustrate the sort of physical processes contributing to preventive feedback, and how they might manifest in the overall shape of $\zeta$."
 Phe actual trend of ¢(Adi.) may involve more complicated and subtle effects than described here., The actual trend of $\zeta(M_{\rm halo})$ may involve more complicated and subtle effects than described here.
 Figure 2. illustrates the impact of these various inflow ancl feedback terms on the specific star formation rate, Figure \ref{fig:ssfrhalo} illustrates the impact of these various inflow and feedback terms on the specific star formation rate
"where e.⋅↗ c, and (""7; are related to the proper velocity. of the fluid. according. to and W is the ratio between the Lorentz factors. at a given physical poiut. in the two coordinate svstenis. eiven bv The mapping between the two coordinate svstenis necessarilv breaks down at the horizon.","where $v^{r}$ , $v^{\theta}$ , and $v^{\phi}$ are related to the proper velocity of the fluid according to and $\Psi$ is the ratio between the Lorentz factors, at a given physical point, in the two coordinate systems, given by The mapping between the two coordinate systems necessarily breaks down at the horizon."
 Ience conrparisous will be restricted to some exterior domain., Hence comparisons will be restricted to some exterior domain.
 Similar expressions eive themap., Similar expressions give the.
 Following the general approach laid out iu Bauvuls et al. (, Following the general approach laid out in Banyuls et al. (
1997). we now restrict the domain of integration of the hydrodynamic equations to the equatorial plane of the Ίνα spacetime. 6=7/2.,"1997), we now restrict the domain of integration of the hydrodynamic equations to the equatorial plane of the Kerr spacetime, $\theta=\pi/2$."
" In so doing they adopt the balance law form where à is. again. the lapse fiction of the spacetime. defined as à?=ας, "," In so doing they adopt the balance law form where $\alpha$ is, again, the lapse function of the spacetime, defined as $\alpha^2\equiv -1/g^{tt}$."
In Eq. (31)), In Eq. \ref{system}) )
" the vector of is defined as where p and 5 are. respectively, the rest-inass deusity (uot to be confused with the geometrical factor 0) and the specific internal cnerey. related to the pressure p via an equation of state which we close to be that of an ideal eas with * being the constant adiabatic mdex."," the vector of is defined as where $\rho$ and $\varepsilon$ are, respectively, the rest-mass density (not to be confused with the geometrical factor $\varrho$ ) and the specific internal energy, related to the pressure $p$ via an equation of state which we chose to be that of an ideal gas with $\gamma$ being the constant adiabatic index."
 Ou the other haud. the vector of unknowns (evolved quantities) in Eq. (31))," On the other hand, the vector of unknowns (evolved quantities) in Eq. \ref{system}) )"
" is The explicit relations between the two sets of variables. U and w. are with. WW. ⇁⋅being the Lorentz factor.. IV-—adf(1..07).»2V7, with04 οἳ»=«οi."," is The explicit relations between the two sets of variables, ${\bf U}$ and ${\bf w}$ are with $W$ being the Lorentz factor, $W\equiv\alpha u^t =
(1-v^2)^{-1/2}$, with $v^2=g_{ij}v^i v^j$."
 The specific EMform of. the fluxes. F'. and the source terms.S. are given in the Appendix for the two differcut represeutatious of the Ίνα nctric we use.," The specific form of the fluxes, ${\bf F}^{i}$, and the source terms, are given in the Appendix for the two different representations of the Kerr metric we use."
 Further details about the equations can be found in Dauvuls et al (1997)., Further details about the equations can be found in Banyuls et al (1997).
 We solve system (31)) on a discrete numerical erid., We solve system \ref{system}) ) on a discrete numerical grid.
 To this cud. we take advantage of the explicit hyperbolicity of the system in order to build up a linearized Riemanuu solver.," To this end, we take advantage of the explicit hyperbolicity of the system in order to build up a linearized Riemann solver."
 Schemes using approximate Ricmmann solvers are based ou the characteristic information contained ithesvstem of equations., Schemes using approximate Riemann solvers are based on the characteristic information contained inthesystem of equations.
 With this strategy. physical discoutinuities appearing in the solution. e.g. shock waves.are treated cousisteutlv," With this strategy, physical discontinuities appearing in the solution, e.g., shock waves,are treated consistently"
We have presented. detailed models of the inulti-wavelength astrometric displaceiient that SIAL Lite will observe due to eravityv darkening aud stellar spots using he code.,We have presented detailed models of the multi-wavelength astrometric displacement that SIM Lite will observe due to gravity darkening and stellar spots using the code.
 We find that SIM Lite observations. (specially when combined with other teciniques. will © able to determine the absolute iucliuation. gravity (larkeuiug exponent. aud 3-dimensional orietation of the l'otational axis for fast aud slow rotating eiut stars. aud ast-rotating miadu-sequeuce stars.," We find that SIM Lite observations, especially when combined with other techniques, will be able to determine the absolute inclination, gravity darkening exponent, and 3-dimensional orientation of the rotational axis for fast and slow rotating giant stars, and fast-rotating main-sequence stars."
 This technique will be especially useful iu probing binary star axd exoplanct ormation and evolution. as well as the plwsics of star ormune regions.," This technique will be especially useful in probing binary star and exoplanet formation and evolution, as well as the physics of star forming regions."
 Direct observational determination of he gravity darkening exponent has direct ayplications iu oth stellar and exoplauct astroplivsics., Direct observational determination of the gravity darkening exponent has direct applications in both stellar and exoplanet astrophysics.
" This technique is also relatively inexpensive iu terms of SIM Lite observing nue. ax one need only to observe a eve""n star once. as opposed to binary stars ane planets. which require constant monitoring over an cutire orbit."," This technique is also relatively inexpensive in terms of SIM Lite observing time, as one need only to observe a given star once, as opposed to binary stars and planets, which require constant monitoring over an entire orbit."
 Ἡ should be joted that this effect should be aken ito account when constructing the SIM. Lite astroinetrie τοῖςvenee frame. such that fast-rotatiugo ogiauts should be excluded so as not to produce a wavelenetl-dependent astrometric refercuce faic.," It should be noted that this effect should be taken into account when constructing the SIM Lite astrometric reference frame, such that fast-rotating giants should be excluded so as not to produce a wavelength-dependent astrometric reference fame."
 We also have preseuted models of star spots on suele stars. and find that SIAL Lite should be able to discern their location. temperature. aud size.," We also have presented models of star spots on single stars, and find that SIM Lite should be able to discern their location, temperature, and size."
 Combined with other techniques. this will provide ereat insight into stellar differential rotation. magnetic cvcles aud underlying cvasdnos. and magnetic interaction iu close Dinarics.," Combined with other techniques, this will provide great insight into stellar differential rotation, magnetic cycles and underlying dynamos, and magnetic interaction in close binaries."
 Fr:mu this modehue. it should especially be noted that uultiovaveleusgth astrometry is a kev tool in the lout for extrasolar planets. cither by ruling out false signals created by spots. or siuuply removing extra astrometric jitter introduced by spots.," From this modeling, it should especially be noted that multi-wavelength astrometry is a key tool in the hunt for extrasolar planets, either by ruling out false signals created by spots, or simply removing extra astrometric jitter introduced by spots."
 Thus. it remains critical that SIMD Lite niüntaius a inuultivaveleusth astrometric capability iu its final design.," Thus, it remains critical that SIM Lite maintains a multi-wavelength astrometric capability in its final design."
 This work was sponsored in part by a SIAL Scicuce Study (PI. D. Celine} from the National Aeronautics and Space Adiuinistration through a contract with the Jet Propulsion Laboratory. California Lhustitute of Technology.," This work was sponsored in part by a SIM Science Study (PI: D. Gelino) from the National Aeronautics and Space Administration through a contract with the Jet Propulsion Laboratory, California Institute of Technology."
 J.L.C. acknowledges additional support from a New Mexico Space Grant Consortium Fellowship., J.L.C. acknowledges additional support from a New Mexico Space Grant Consortium Fellowship.
 We thank the referee for conuneuts which ercatly helped to miprove this manuscript. particularly du making the presentation of our ideas much clearer.," We thank the referee for comments which greatly helped to improve this manuscript, particularly in making the presentation of our ideas much clearer."
"to ~100F,,.",to $\sim 100 F_{\mathrm{av}}$.
" For the same reasons, the analysis on the sample of GPs corrected for refractive scintillation did not give any meaningful results."," For the same reasons, the analysis on the sample of GPs corrected for refractive scintillation did not give any meaningful results."
 No change in the Crab GP generation rate was found on timescales from 10 to ss around y-ray photons and with any possible lag within +40 mmin with respect to y-ray photons., No change in the Crab GP generation rate was found on timescales from 10 to s around $\gamma$ -ray photons and with any possible lag within $\pm 40$ min with respect to $\gamma$ -ray photons.
" With probability, the high energy flux of the Crab pulsar during GPs is less than 4 times the average γ- pulsed flux for the on-pulse (MP--IP--bridge between them) phase window."," With probability, the high energy flux of the Crab pulsar during GPs is less than 4 times the average $\gamma$ -ray pulsed flux for the on-pulse (MP+IP+bridge between them) phase window."
" For IP GPs only, the upper limit on y-ray flux in the IP phase window is 12 times the average pulsed flux."," For IP GPs only, the upper limit on $\gamma$ -ray flux in the IP phase window is 12 times the average pulsed flux."
" If we consider the uncertainty due to discrete numbers of matches between photons and GPs, the upper limits are 3—5.5 times the average pulsed flux for the pulsed emission window, and 8—16 for the IP window."," If we consider the uncertainty due to discrete numbers of matches between photons and GPs, the upper limits are $-$ 5.5 times the average pulsed flux for the pulsed emission window, and $-$ 16 for the IP window."
 A few explanations may be offered for the non-detection of GP-y-photon correlations., A few explanations may be offered for the non-detection of $\gamma$ -photon correlations.
" The most natural is that production of GPs depends on non-stationary changing coherence conditions, which vary by a large degree even for similar magnetospheric particle densities."," The most natural is that production of GPs depends on non-stationary changing coherence conditions, which vary by a large degree even for similar magnetospheric particle densities."
" Another possibility is that beaming in radio and at high energies are somewhat different, so that simultaneous GPs and y-ray photons are emitted in different directions."," Another possibility is that beaming in radio and at high energies are somewhat different, so that simultaneous GPs and $\gamma$ -ray photons are emitted in different directions."
" Overall, our results suggest that enhanced pair creation is not a dominant factor for GP occurrence, at least for high frequency IP GPs."," Overall, our results suggest that enhanced pair creation is not a dominant factor for GP occurrence, at least for high frequency IP GPs."
" However, our flux increase estimations are not on the level of a few percent, as in the work of? "," However, our flux increase estimations are not on the level of a few percent, as in the work of \cite{shearer2003} "
ensemble average at spatially averaged Πο density p=τfdxp(x).,ensemble average at spatially averaged $\H2$ density $\overline{\rho}=\frac{1}{V}\int{}dx\rho(x)$.
" The star formation rate density p. is a random field and its value for a given point in space is defined by equation (1)), i.e., where p= and we assume that p. is proportional to the He density, i.e., Targ is a constant."," The star formation rate density $\dot{\rho}_*$ is a random field and its value for a given point in space is defined by equation \ref{eq:SFR}) ), i.e., where $\rho\equiv{}\rho_\H2$ and we assume that $\dot{\rho}_*$ is proportional to the $\H2$ density, i.e., $\tau_\SFR$ is a constant."
" Here X is a Gaussian random field pg,with zero mean and unit variance.", Here $X$ is a Gaussian random field with zero mean and unit variance.
" From equations (B1)) and (B2)), it follows σι=σ/1n10 on the scatter insertion scale lmin, i.e., without spatial averaging."," From equations \ref{eq:sigSFR}) ) and \ref{eq:SFRAlt}) ), it follows that $\sigma_l=\sigma/\ln{}10$ on the scatter insertion scale $l_{\rm min}$ , i.e., without spatial averaging."
" The reason being that, in this case, that0.=px, p= p, (logiop)=logjop and thus Equation (B1)) can be solved if all cells have the same density po and o<1."," The reason being that, in this case, $\overline{\dot{\rho}_*}=\dot{\rho}_*$, $\overline{\rho}=\rho$ , $\langle \log_{10}\rho \rangle=\log_{10}\rho$ and thus Equation \ref{eq:sigSFR}) ) can be solved if all cells have the same density $\rho_0$ and $\sigma{}\ll{}1$."
" In this case epe?*=epoe™, logy)e?*~oX, and If X is spatially uncorrelated then (X(z)X(y))=ó(x—y), (x)—1/V, and a«V95,"," In this case $\overline{\epsilon \rho\,e^{\sigma{}X}}=\epsilon \rho_0 \overline{e^{\sigma{}X}}$, $\log_{10}\overline{e^{\sigma{}X}}\approx{}\sigma\overline{X}$, and If $X$ is spatially uncorrelated then $\langle X(x) X(y) \rangle=\delta(x-y)$, $\langle \overline{X}^2 \rangle = 1/V$, and $\sigma_{l}\propto{}V^{-0.5}$."
" However, the scaling of σι may be different if X is spatially correlated, σ>>1, or if the density pdf is not sharply peaked."," However, the scaling of $\sigma_{l}$ may be different if X is spatially correlated, $\sigma\gg{}1$, or if the density pdf is not sharply peaked."
" We can testwhether thelatter, i.e., the finite width of a log-normal density distribution, provides a quantitative"," We can testwhether thelatter, i.e., the finite width of a log-normal density distribution, provides a quantitative"
we calculate the deceleration. jerk and snap parameters using type Ia supernovae and gamma-ray bursts.,"we calculate the deceleration, jerk and snap parameters using type Ia supernovae and gamma-ray bursts."
 GRBs can extend the redshift - distance relation up to high redshifts., GRBs can extend the redshift - distance relation up to high redshifts.
 We find that the deceleration and jerk parameters in the ΞξWo. GCG. Cardassian expansion and f(R) models are almost the same in the Io confidence level.," We find that the deceleration and jerk parameters in the $w=w_0$, GCG, Cardassian expansion and f(R) models are almost the same in the $1\sigma$ confidence level."
 So these models can not be discriminated using the present value of the statefinder pair., So these models can not be discriminated using the present value of the statefinder pair.
 We find that the dark energy models and modified gravity models could be distinguished between by the snap parameter., We find that the dark energy models and modified gravity models could be distinguished between by the snap parameter.
 Using the model-independent constraints on. cosmographic parameters. we find the ACDM model is consistent. with the current data.," Using the model-independent constraints on cosmographic parameters, we find the $\Lambda$ CDM model is consistent with the current data."
 This work is supported by the National Natural. Science Foundation of China (grants 10233010. 10221001. and 10873009) and the National Basic Research Program. of China (973 program) No.," This work is supported by the National Natural Science Foundation of China (grants 10233010, 10221001 and 10873009) and the National Basic Research Program of China (973 program) No."
 2007CB815404., 2007CB815404.
" F. Y. Wang was also supported by the Jiangsu Project Innovation for PhD Candidates (CX07B-039z),", F. Y. Wang was also supported by the Jiangsu Project Innovation for PhD Candidates (CX07B-039z).
The calculation method and other assumptions are the same as described in Umeda & Nomoto (2002. 2005) and Izutani et al. (,"The calculation method and other assumptions are the same as described in Umeda $\&$ Nomoto (2002, 2005) and Izutani et al. ("
2009) except. [or the size of the nuclear reaction networks.,2009) except for the size of the nuclear reaction networks.
" In (his paper. we adopt the Pop III progenitors as in Umeda & Nomoto (2002. 2005) and apply the model with = 15 M. and E; = 1 (normal SN model) and Af = 25 Af. and Ls, = 20 (118 model)."," In this paper, we adopt the Pop III progenitors as in Umeda $\&$ Nomoto (2002, 2005) and apply the model with $M$ = 15 $M_\odot$ and $E_{51}$ = 1 (normal SN model) and $M$ = 25 $M_\odot$ and $E_{51}$ = 20 (HN model)."
 Detailed nucleosvnthesis is calculated as a postprocessing alter the hydrodynamical ealeulation with a simple a-network., Detailed nucleosynthesis is calculated as a postprocessing after the hydrodynamical calculation with a simple $\alpha$ -network.
 The isotopes included in the post-process calculations for Y; « 0.51 are 809 species up to P1 Pd. and the ones for Y. > 0.51 are 652 species up tot’ Pd. (," The isotopes included in the post-process calculations for $Y_{\rm e}$ $<$ 0.51 are 809 species up to $^{121}$ Pd, and the ones for $Y_{\rm e}$ $\ge$ 0.51 are 652 species up to $^{112}$ Pd. ("
for detail. see Izutani. Umeda & Yoshida (2010) in preparation).,"for detail, see Izutani, Umeda $\&$ Yoshida (2010) in preparation)."
 As for the LIN model. we consider a model for which the density in the complete Si-burning region is artificially reduced to 1/3 (or entropy is enhanced by a [actor of 3) of the original as well as (he original model.," As for the HN model, we consider a model for which the density in the complete Si-burning region is artificially reduced to 1/3 (or entropy is enhanced by a factor of 3) of the original as well as the original model."
 The IEN model with the entropy enhanced by a factor of 3 fits the observation [rom Ca to Zn except Cr as shown in Figure 1., The HN model with the entropy enhanced by a factor of 3 fits the observation from Ca to Zn except Cr as shown in Figure 1.
" We call (his model “LLNS model (s/A, ~ 50). and call the IN model with the original entropy ΗΝ model (s/h), ~ 15)."," We call this model “HN3 model” $s$ $k_{\rm b}$ $\sim$ 50), and call the HN model with the original entropy “HN model” $s$ $k_{\rm b}$ $\sim$ 15)."
" As lor the normal SN model. the densitwv for the postprocessing is multiplied bv factors olL (s/h ~ 5). L/3 (s/h, ~ 15). 1/10 (s/f, e 40). ancl 1/35 (51 ~ 150)."," As for the normal SN model, the density for the postprocessing is multiplied by factors of 1 $s$ $k_{\rm b}$ $\sim$ 5), 1/3 $s$ $k_{\rm b}$ $\sim$ 15), 1/10 $s$ $k_{\rm b}$ $\sim$ 40), and 1/35 $s$ $k_{\rm b}$ $\sim$ 150)."
 This modification of density minucs hot-bubbles in niulti-dimensional simulations., This modification of density mimics hot-bubbles in multi-dimensional simulations.
 We obtain the final vields bv adding matter above M4 and below M4., We obtain the final yields by adding matter above $M_{\rm cut}$ and below $M_{\rm cut}$.
 We set Mq = 1.50 AZ. for the normal SN model and 1.88 M. for the IIN3S model., We set $M_{\rm cut}$ = 1.50 $M_\odot$ for the normal SN model and 1.88 $M_\odot$ for the HN3 model.
" As a result. the normal SN model ejects 0.07 11. of ""Ni. which is similar to SNI987À. and (the IIN3 model ejects 0.51 M. of Nit"," As a result, the normal SN model ejects 0.07 $M_\odot$ of $^{56}$ Ni, which is similar to SN1987A, and the HN3 model ejects 0.51 $M_\odot$ of $^{56}$ Ni."
 In Figure 1. the abundance patterns [rom $i (o Iu are compared with those of EMP stars.," In Figure 1, the abundance patterns from Si to Ru are compared with those of EMP stars."
 As for Co and Zn. the IIN3 model reproduces the observations well. while the normal SN model does not as described in the figure caption.," As for Co and Zn, the HN3 model reproduces the observations well, while the normal SN model does not as described in the figure caption."
 As for Sr. Y. and Zr. both the IIN3 model with neutron-rich hot-bubble matter (s/Aj ~ 15) amd the normal SN model with neutron-rich hot-bubble matter reproduce the observations well.," As for Sr, Y, and Zr, both the HN3 model with neutron-rich hot-bubble matter $s$ $k_{\rm b}$ $\sim$ 15) and the normal SN model with neutron-rich hot-bubble matter reproduce the observations well."
" ILowever. as for the heavier elements. Mo. Ru. and Rh. only the normal SN model with s/f), ~ 150 neutron-rich matter reach the observations."," However, as for the heavier elements, Mo, Ru, and Rh, only the normal SN model with $s$ $k_{\rm b}$ $\sim$ 150 neutron-rich hot-bubble matter reach the observations."
of discrete pre-calculated evolutionary. (racks: either interpolating between (hem. or using parametrized fit lormulae.,"of discrete pre-calculated evolutionary tracks: either interpolating between them, or using parametrized fit formulae."
 Clearly. (is procedure is incapable of dealing with ‘non-canonical’ stars. (he outcome of collisions and mergers.," Clearly, this procedure is incapable of dealing with `non-canonical' stars, the outcome of collisions and mergers."
 In (his paper we thus present a new evolutionary code (that we have developed with this ai in mind., In this paper we thus present a new evolutionary code that we have developed with this aim in mind.
 The outline of the code and method of solution are presented in the next section. Section 2: the input physics is described in some detail in Section 3.. ancl results of representative calculations are discussed in Section 4..," The outline of the code and method of solution are presented in the next section, Section \ref{code}; the input physics is described in some detail in Section \ref{input-phys}, and results of representative calculations are discussed in Section \ref{res}."
 The equations (hat govern the evolution of a star are those of continuity. hyvdrostatic equilibrium. energv transfer (radiative or convective). energv balance. and composition balance: In these equations.mass n and time / are the independent. variables.," The equations that govern the evolution of a star are those of continuity, hydrostatic equilibrium, energy transfer (radiative or convective), energy balance, and composition balance: In these equations,mass $m$ and time $t$ are the independent variables."
 The dependent ones are radius r. density p. temperature Z. and the number fractions Yj. related to the mass fractions JN; bv Y;—N;/Aj. where τς is the / th atomic mass.," The dependent ones are radius $r$, density $\rho$ , temperature $T$, and the number fractions $Y_j$ , related to the mass fractions $X_j$ by $Y_j=X_j/A_j$, where $A_j$ is the $j$ 'th atomic mass."
 The particle flux Fy of the jth species is assumed to be diffusive (proportional to the abundance gradient. of the j th species). determined by the diffusion coellicient σι.," The particle flux $F_j$ of the $j$ 'th species is assumed to be diffusive (proportional to the abundance gradient of the $j$ 'th species), determined by the diffusion coefficient $\sigma_j$ ."
 We regard (p.T.Y) as the basic thermodynamic variables.," We regard $(\rho,T,Y)$ as the basic thermodynamic variables."
 They determine. through the equation of state. the pressure p(p.T.Y) and the specific enerev (p.T.Y). as well as the opacity &(p.T. 3). energy production rate q(p.T.Y)and nuclear energyrates τρ.T.Y) (via an imported list oftables and formulae).," They determine, through the equation of state, the pressure $p(\rho,T,Y)$ and the specific energy $u(\rho,T,Y)$, as well as the opacity $\kappa(\rho,T,Y)$ energy production rate $q(\rho,T,Y)$and nuclear energyrates $R_j(\rho,T,Y)$ (via an imported list oftables and formulae)."
" The temperature ""gradient V(r.L.im.p.T.Y) and"," The temperature `gradient' $\nabla(r,L,m,\rho,T,Y)$ and"
outward and transports energy along twin jets that form naturally.,outward and transports energy along twin jets that form naturally.
 The results are in general agreement with an earlier proposal made by Punsly Coroniti (1990)., The results are in general agreement with an earlier proposal made by Punsly Coroniti (1990).
 In. addition. Ixoide el al. (," In addition, Koide et al. ("
2002) have shown that the energy extraction mechanism is related (o the Penrose process.,2002) have shown that the energy extraction mechanism is related to the Penrose process.
 As the field lines in the equatorial region are pulled forward. the local plasma on the field line acquires a negative energy.," As the field lines in the equatorial region are pulled forward, the local plasma on the field line acquires a negative energy."
 This plasma falls into the BIL while a torsional Allven wave is driven outward along the field lines. generating an outward Povuting flux and at the salle (ime increasing the enerev of plasma oulside the ergosphere.," This plasma falls into the BH, while a torsional Alfven wave is driven outward along the field lines, generating an outward Poynting flux and at the same time increasing the energy of plasma outside the ergosphere."
 The authors call it theprocess since there is a close analogy between this mechanism and Penrose's original idea., The authors call it the since there is a close analogy between this mechanism and Penrose's original idea.
 The main difference is (hat the negative energy particles close to the DII and the positive enerey particles farther oul communicate via magnetic field lines. which act as energv conduits.," The main difference is that the negative energy particles close to the BH and the positive energy particles farther out communicate via magnetic field lines, which act as energy conduits."
 semenov. Dvaclechkin Punsly (2004) have recently. used a simplified set of equations to follow the evolution of individual magnetic flux tubes in the vieinity of a rotating DII.," Semenov, Dyadechkin Punsly (2004) have recently used a simplified set of equations to follow the evolution of individual magnetic flux tubes in the vicinity of a rotating BH."
 Alovieó and Movie6 from their work show how a rotating DII with a magnetized plasma around it naturally develops a configuration with infalline negative enerev plasma in the equatorial plane and outgoing coiled magnetic fields (hat carry away energy in (win jets., Movie5 and Movie6 from their work show how a rotating BH with a magnetized plasma around it naturally develops a configuration with infalling negative energy plasma in the equatorial plane and outgoing coiled magnetic fields that carry away energy in twin jets.
 Aluch more work is obviously needed. but it appears that the connection between rotating DlIIs and jets is beginning to be understood.," Much more work is obviously needed, but it appears that the connection between rotating BHs and jets is beginning to be understood."
 Note that the field topology in the above-cited work is different from the traditional picture in which field lines connect the DII horizon to the accretion disk., Note that the field topology in the above-cited work is different from the traditional picture in which field lines connect the BH horizon to the accretion disk.
 The traditional field configuration seems to produce evacuated fuunels with outflowing funnel walls rather than bona fide jets (de Villiers. Hawley [ντος 2003: Melxiunnev Gammie 2004).," The traditional field configuration seems to produce evacuated funnels with outflowing funnel walls rather than bona fide jets (de Villiers, Hawley Krolik 2003; McKinney Gammie 2004)."
 The simulations of IXoide et al. (, The simulations of Koide et al. (
2002) and Semenov et al. (,2002) and Semenov et al. (
2004) start out with field lines going Tom the ereosphere to infinity along the spin axis.,2004) start out with field lines going from the ergosphere to infinity along the spin axis.
 Such a field configuration can be achieved iiurallv as a result. of magnetized plasma falling into the DII (Naravan. Igumenshceliev Abramowiez 2003).," Such a field configuration can be achieved naturally as a result of magnetized plasma falling into the BH (Narayan, Igumenshchev Abramowicz 2003)."
 In fact. not only does this field configuration enable energy extraction rom a rolating hole. it may also enhance the energy. extracted [rom gas accreting onto anon-rolating hole. 1.6.. it can give a larger value of 77 than the standard value of 0.057 shown in Fig.," In fact, not only does this field configuration enable energy extraction from a rotating hole, it may also enhance the energy extracted from gas accreting onto a hole, i.e., it can give a larger value of $\eta$ than the standard value of 0.057 shown in Fig."
 2 (Naravan et al., 2 (Narayan et al.
 2003)., 2003).
 Astrophysicists are contributing in important wavs to the study of BIls., Astrophysicists are contributing in important ways to the study of BHs.
 Thev are discovering increasing numbers of DII candidates. they are measuring (he fundamental parameters ol the BIIs. they are finding strong evidence that DII candidates have event horizons. ancl they are beginning to carry out additional tests of relativity in the regime of strong gravity.," They are discovering increasing numbers of BH candidates, they are measuring the fundamental parameters of the BHs, they are finding strong evidence that BH candidates have event horizons, and they are beginning to carry out additional tests of relativity in the regime of strong gravity."
In this section we model the spectral energy distribution of 7 intergalactic star forming regions using a photospheric+nebular model and a dust model.,In this section we model the spectral energy distribution of 7 intergalactic star forming regions using a photospheric+nebular model and a dust model.
 We divide the spectral energy distribution modeling into two steps., We divide the spectral energy distribution modeling into two steps.
 In the first we use to fit the continuum and line emission stepfrom the stars and the gas using observations ranging from the ultraviolet to the near-infrared., In the first step we use to fit the continuum and line emission from the stars and the gas using observations ranging from the ultraviolet to the near-infrared.
 The output of this fit provides information on the excitation and intrinsic interstellar radiation field., The output of this fit provides information on the excitation and intrinsic interstellar radiation field.
 In the second step we use those estimates to fit the mid-infrared observations using a dust model., In the second step we use those estimates to fit the mid-infrared observations using a dust model.
 In the process we also estimate the energy absorbed and re-emitted by the dust at longer wavelengths., In the process we also estimate the energy absorbed and re-emitted by the dust at longer wavelengths.
 We combine observations of emission lines with the spectral spectroscopicevolution code to model star forming regions in collision debris., We combine spectroscopic observations of emission lines with the spectral evolution code to model star forming regions in collision debris.
 We use the? IMF from 0.1 Μο to 100 Μο as a baseline for this study unless specified otherwise., We use the IMF from 0.1 $_\sun$ to 100 $_\sun$ as a baseline for this study unless specified otherwise.
 It is known that gaseous disks in parent galaxies are often more extended than the stellar ones., It is known that gaseous disks in parent galaxies are often more extended than the stellar ones.
" As a result, depending on the parameters of the interaction, the collision debris may be devoid of any older pre-existing evolved stellar population, or on the contrary be dominated by old stars."," As a result, depending on the parameters of the interaction, the collision debris may be devoid of any older pre-existing evolved stellar population, or on the contrary be dominated by old stars."
 We model this using a two burst star formation history., We model this using a two burst star formation history.
" The first burst models the disk of a spiral galaxy, that is the evolved stellar population that is during an interaction: it assumes a 12x10? years old strippedexponentially declining burst which has consumed of its original gas reservoir."," The first burst models the disk of a spiral galaxy, that is the evolved stellar population that is stripped during an interaction: it assumes a $12\times10^9$ years old exponentially declining burst which has consumed of its original gas reservoir."
 On top of this evolved population we add a second burst which models the star formation episode taking place in the collision debris after 12x10? years., On top of this evolved population we add a second burst which models the star formation episode taking place in the collision debris after $12\times10^9$ years.
" As the photometric properties change very rapidly during the first few million years the burst is modeled by an exponentially decreasing SFR: SFR(t)=axexp(-t/T), where a determines the intensity of the burst and τ is the timescale."," As the photometric properties change very rapidly during the first few million years the burst is modeled by an exponentially decreasing SFR: $SFR(t)=a\times\exp\left(-t/\tau\right)$, where $a$ determines the intensity of the burst and $\tau$ is the timescale."
" We follow this second burst from t=10 years to t=2x10° years, t being here the time elapsed since the beginning of the star formation episode in collision debris."," We follow this second burst from $t=10^6$ years to $t=2\times10^9$ years, $t$ being here the time elapsed since the beginning of the star formation episode in collision debris."
" For this population, τ varies between 10° years and 10x10? years to range of possible star formation histories from nearly spaninstantaneousalarge to nearly constant."," For this population, $\tau$ varies between $10^6$ years and $10\times10^9$ years to span a large range of possible star formation histories from nearly instantaneous to nearly constant."
 The quality of the determination of the relative contribution of both components will ascertain the accuracy of the determination of the other parameters., The quality of the determination of the relative contribution of both components will ascertain the accuracy of the determination of the other parameters.
" The relative importance of the two populations is determined by the ratio r, the stellar mass produced by the second burst divided the mass of the stars produced by the first burst (the byevolved population), at a given time { for a given timescale r."," The relative importance of the two populations is determined by the ratio $r$, the stellar mass produced by the second burst divided by the mass of the stars produced by the first burst (the evolved population), at a given time $t$ for a given timescale $\tau$."
 We span a range of ratios from logr=—2.75 to logr=2 by steps of 0.25., We span a range of ratios from $\log r=-2.75$ to $\log r=2$ by steps of $0.25$.
 The parameter can be seen as a burst strength., The parameter $r$ can be seen as a burst strength.
 A limitation is the r modeling of TP-AGB (thermally possiblepulsating asymptotic giant imprecisebranch) stars by., A possible limitation is the imprecise modeling of TP-AGB (thermally pulsating asymptotic giant branch) stars by.
.? has shown that they are dominant for populations aged between 300x10° and 2x10? years., has shown that they are dominant for populations aged between $300\times10^6$ and $2\times10^9$ years.
" For our specific case however, the older population is much older and is dominated by red giant branch stars and the younger populations best fit is never older than 300x10° years old as we will show later."," For our specific case however, the older population is much older and is dominated by red giant branch stars and the younger populations best fit is never older than $300\times10^6$ years old as we will show later."
 The nebular emission accounts for a significant fraction of the flux in the broad optical bands from young star forming (?)., The nebular emission accounts for a significant fraction of the flux in the broad optical bands from young star forming regions.
 To improve the standard nebular modeling of regionsemission lines by — which assumes a given set of ratios between emission lines — we use attenuation- emission lines obtained from spectra., To improve the standard nebular modeling of emission lines by – which assumes a given set of ratios between emission lines – we use attenuation-corrected emission lines obtained from spectra.
 To calculate, To calculate
There has been renewed interest in the kinematics of dwarf galaxies for two reasons. (,There has been renewed interest in the kinematics of dwarf galaxies for two reasons. (
1) There appeurs to be a conflict vetween licrarclical ealaxy formation model predictions for. and observational determinations of. the dark matter halo density profile of galaxies.,"1) There appears to be a conflict between hierarchical galaxy formation model predictions for, and observational determinations of, the dark matter halo density profile of galaxies."
 Iu particular. these models predict a cuspy ceutral halo. which many observers fiud iucousisteut with observations (e.g. Weldrakeetal. 2003)).," In particular, these models predict a cuspy central halo, which many observers find inconsistent with observations (e.g. \cite{weldrake03}) )."
 While the models predict cuspy halos for all galaxies. dwarf galaxies are believed ο be best suited for tests of the models because these ealaxies are iu general held to be dominated bv dark uatter. even in the innermost regions.," While the models predict cuspy halos for all galaxies, dwarf galaxies are believed to be best suited for tests of the models because these galaxies are in general held to be dominated by dark matter, even in the innermost regions."
 Du larecr galaxies. uncertainties m the stellar mass to light ratio unit one's ability to determine deusitv profile of the dark matter ido. (," In larger galaxies, uncertainties in the stellar mass to light ratio limit one's ability to determine density profile of the dark matter halo. ("
2) It appears that faint dwuf galaxies deviate sieuificautlv from the Tull-Fisher relationship as defined w bright galaxies.,2) It appears that faint dwarf galaxies deviate significantly from the Tully-Fisher relationship as defined by bright galaxies.
 While the TF relation is generally studied using the incliation corrected profile width (usually obtained from sinele dish observations). it is nuclear whether this is au appropriate measure fo use for the faimtest dwarf ealaxics. where the rotational velocities are comparable to the velocity dispersion.," While the TF relation is generally studied using the inclination corrected profile width (usually obtained from single dish observations), it is unclear whether this is an appropriate measure to use for the faintest dwarf galaxies, where the rotational velocities are comparable to the velocity dispersion."
 For such ealaxies. the vasvuunetric drift” corrected rotational velocity is probably a better measure.," For such galaxies, the “asymmetric drift” corrected rotational velocity is probably a better measure."
 However. estimating this quantity requires TT svuthesis imaging.," However, estimating this quantity requires HI synthesis imaging."
 Further. for the faint dwarf nreeulars. even the inclination may be difficult to estimate from the optical dmage aud may only be obtainable from the UT distribution and kineniatics.," Further, for the faint dwarf irregulars, even the inclination may be difficult to estimate from the optical image and may only be obtainable from the HI distribution and kinematics."
 Unfortunately. the ummber of faint dwarf galaxies with sutiicicntly hieh qualitv ITE svuthesis images available is quite simall.," Unfortunately, the number of faint dwarf galaxies with sufficiently high quality HI synthesis images available is quite small."
 We discuss here CAIRT IIT observations of two dwarf galaxies. Πο 250 aud KIE98 251. in light of both these issues.," We discuss here GMRT HI observations of two dwarf galaxies, KK98 250 and KK98 251, in light of both these issues."
 Iu discussing the TF relation. we combine our data for IKIX98 250 and 1ος 251 with our earlier data for still fainter galaxies aud data for brighter ealaxies available frou the literature.," In discussing the TF relation, we combine our data for KK98 250 and KK98 251 with our earlier data for still fainter galaxies and data for brighter galaxies available from the literature."
 Both KlX08 250 aud Wh9Os 251 WCTC identified as conipanious to the eiut spiral ealaxy NGC 6916 by Waracheutseva&IKaracheutsey(L998)}.. aud have also been observed in TT at Effelsbere (IIuchtiueier et al.," Both KK98 250 and KK98 251 were identified as companions to the giant spiral galaxy NGC 6946 by \cite{karachentseva98}, and have also been observed in HI at Effelsberg (Huchtmeier et al."
 1997) aud at the DRAO (Pisano&Wilcots20003)., 1997) and at the DRAO \cite{pisano00}) ).
 The data presented here are of much higher scusitivity aud resolution than the DRAO data., The data presented here are of much higher sensitivity and resolution than the DRAO data.
 Distance estimates to Ίος 250 aud ΠΙΟΣ 251 vary from 5.3 Alpe (Sharina ct al., Distance estimates to KK98 250 and KK98 251 vary from 5.3 Mpc (Sharina et al.
 1997. Ixaracheutsev et al.," 1997, Karachentsev et al."
 2000) to 8.2 Alpe (Sharina et al., 2000) to 8.2 Mpc (Sharina et al.
 1997)., 1997).
 Because of the proximity of the two galaxies to each other and to NCC 6916. on the sky as well as in the velocity. we feel that it is ikelv that all these ealaxics beloug to the same group viz.," Because of the proximity of the two galaxies to each other and to NGC 6946, on the sky as well as in the velocity, we feel that it is likely that all these galaxies belong to the same group viz."
 the NGC 6916 eroup., the NGC 6946 group.
 Ueuce. in this paper we take the mean distance to the group (5.6 Mpc. estimated from the brightest stars in eight imienibers of the eroup) as the distance to both the galaxies (Huchtiueier et al.," Hence, in this paper we take the mean distance to the group (5.6 Mpc, estimated from the brightest stars in eight members of the group) as the distance to both the galaxies (Huchtmeier et al."
 2000)., 2000).
 At this clistance. the absolute blue magnitude for ΠΙΟΣ 250 aud KI&98 251 are Mp~ 123.72 and 11514 respectively.," At this distance, the absolute blue magnitude for KK98 250 and KK98 251 are $\rm{M_B}\sim -$ 13.72 and $-$ 14.54 respectively."
 The rest of the paper is divided as follows., The rest of the paper is divided as follows.
 The GCAIRT observations are detailed in Sect. 2..," The GMRT observations are detailed in Sect. \ref{sec:obs},"
 while the results are presented and discussed in Sect. 3.., while the results are presented and discussed in Sect. \ref{sec:res}.
 CCD images of II&98 250 aud KIx95 251 in the Dessell I aud V filters were obtained on 297 and 307 May 2003. using the TFOSC (Timalavan Faint Object Spectrograph Camera}. at the 2.0 10 Wimalavan Clhiaudra Telescope.," CCD images of KK98 250 and KK98 251 in the Bessell I and V filters were obtained on $^{\rm{th}}$ and $^{\rm{th}}$ May 2003, using the HFOSC (Himalayan Faint Object Spectrograph Camera), at the 2.0 m Himalayan Chandra Telescope."
 The camera has a field of view of 10 «10. with a scale of 0.3. /pixel.," The camera has a field of view of $'\times$ $'$, with a scale of $^{''}$ /pixel."
 Both the galaxies were covered in the same pointing., Both the galaxies were covered in the same pointing.
 The total exposure time on the target was LO mun in Dand 50 ain in the V baud., The total exposure time on the target was 40 min in I and 50 min in the V band.
 The PFWIIM secing of the co-added images was ~LT ., The FWHM seeing of the co-added images was $\sim1.7^{''}$.
Since the fist nieht was not photometric. standard fields (from: Landolt (1983))) were observed ouly ou the secoud night.," Since the first night was not photometric, standard fields (from \cite{landolt83}) ) were observed only on the second night."
 Debiasing. flat-ficldine aud cosmic rav filtering were done in the usual manner. using standard IRAF routines.," Debiasing, flat-fielding and cosmic ray filtering were done in the usual manner, using standard IRAF routines."
 The exposures taken on the 3077 were calibrated and thou added after aliguiieut., The exposures taken on the $^{\rm{th}}$ were calibrated and then added after alignment.
 This combined frame was used to, This combined frame was used to
the clusters.,the clusters.
 Thus we can conclude that the clusters under study are dynamically relaxed., Thus we can conclude that the clusters under study are dynamically relaxed.
 Lt may be due to the result o£ dynamical evolution or imprint of star formation processes or both., It may be due to the result of dynamical evolution or imprint of star formation processes or both.
 In this paper we have presented CCD CDV photometry [or the stars in the fields of open clusters Basel - and NGC 7067 [or the first time., In this paper we have presented CCD $UBVRI$ photometry for the stars in the fields of open clusters Basel 4 and NGC 7067 for the first time.
 Using present CCD cata in combination with 2MLASS clata we derive the following results: (i) Using the (€D) versus (213 colour-colour diagram. we determine the cluster. metallicity Z ~ 0.008 ancl 0.02 for. Basel 4 and NGC 7067. respectively.," Using present CCD data in combination with 2MASS data we derive the following results: (i) Using the $(U-B)$ versus $(B-V)$ colour-colour diagram, we determine the cluster metallicity Z $\sim$ 0.008 and 0.02 for Basel 4 and NGC 7067 respectively."
 The corresponding mean value of £02V) = 0.45+0.05 mag anc 0.75z:0.05 mag respectively., The corresponding mean value of $E(B-V)$ = $\pm$ 0.05 mag and $\pm$ 0.05 mag respectively.
 Interstellar extinction law has, Interstellar extinction law has
interesting differences and similarities.,interesting differences and similarities.
" In general, the higher the abundance of a gas component is, the smaller the number density or the higher the temperature is."," In general, the higher the abundance of a gas component is, the smaller the number density or the higher the temperature is."
" Except for the component with the highest abundance, the other components all have similar pressure profiles."," Except for the component with the highest abundance, the other components all have similar pressure profiles."
" This means that the bulk of the gas is approximately in a pressure equilibrium, although the temperature and density distributions may still be highly structured."," This means that the bulk of the gas is approximately in a pressure equilibrium, although the temperature and density distributions may still be highly structured."
 Similar phenomena are also reported for simulations of the SN dominated interstellar medium in the Galaxy (deAvillez&Breitschwerdt2005;Lowetal.2005;Joung&Mac2006) and in starburst regions (Joung," Similar phenomena are also reported for simulations of the SN dominated interstellar medium in the Galaxy \citep{Avillez05,MacLow05,Joung06}
 and in starburst regions \citep{Joung08}."
etal., Fig.
2008).. Fig. 5 demonstrates that gas with higher iron abundance generally has a larger net outward radial velocity., \ref{F:ZFeGas_vr} demonstrates that gas with higher iron abundance generally has a larger net outward radial velocity.
" The velocity of the gas component with Zp.>100 fluctuates greatly, because it represents young SNRs which may or may not have been fully accelerated by the buoyancy (see 4.2)."," The velocity of the gas component with $Z_{\rm Fe}>100$ fluctuates greatly, because it represents young SNRs which may or may not have been fully accelerated by the buoyancy (see 4.2)."
" The gas component with 10<Zp.100, representing fully-accelerated and evolved SNRs with moderate dilution, has the largest velocity."," The gas component with $10<Z_{\rm Fe} \leq 100$, representing fully-accelerated and evolved SNRs with moderate dilution, has the largest velocity."
 Fig., Fig.
 6 shows the iron mass fractions of the individual gas components in the different abundance ranges., \ref{F:ironfractions} shows the iron mass fractions of the individual gas components in the different abundance ranges.
 The fractions in iron-rich components decrease considerably faster with increasing radius in the subsonic case than in the supersonic one., The fractions in iron-rich components decrease considerably faster with increasing radius in the subsonic case than in the supersonic one.
" This is because the outflow velocity is much slower in the former case, leading to more dilution and mixing of the iron ejecta before moving to large galactic radii."," This is because the outflow velocity is much slower in the former case, leading to more dilution and mixing of the iron ejecta before moving to large galactic radii."
irradiated brown clwarl-like secondary star of this svstem.,irradiated brown dwarf-like secondary star of this system.
 Their data showed that in theDVRIE bandpasses. the light curves were dominated by emission Irom the white dwarf ancl some residual evelotvon emission. while in the JE and A-bands. the light curves exhibited lavee amplitude (AIL = 0.8 mag) variations that repeated once per orbital evele.," Their data showed that in the bandpasses, the light curves were dominated by emission from the white dwarf and some residual cyclotron emission, while in the - and -bands, the light curves exhibited large amplitude $\Delta$ K = 0.8 mag) variations that repeated once per orbital cycle."
 ILarrison el al., Harrison et al.
 were able to model (hese multi-wavelength data with an extremely cool (Tg< 1.000 Ix) secondary star that is being irradiated by the primary.," were able to model these multi-wavelength data with an extremely cool $_{\rm eff} \leq$ 1,000 K) secondary star that is being irradiated by the primary."
 In this model. the irradiated hemisphere of the secondary. star is heated (ο T2 1.500 Ix. I£ this model were correct. then (he spectroscopic appearance of the secondary star should show dramatic changes over an orbital evele. oscillating between a state where CO absorption is strong. as in the L-dwarls. to one where methane might dominate. as seen in T-dwarls.," In this model, the irradiated hemisphere of the secondary star is heated to $_{\rm eff} \approx$ 1,500 K. If this model were correct, then the spectroscopic appearance of the secondary star should show dramatic changes over an orbital cycle, oscillating between a state where CO absorption is strong, as in the L-dwarfs, to one where methane might dominate, as seen in T-dwarfs."
 We have obtained new phase-resolved and A-bancl spectra of EF Evi using NIRSPEC on IKECI IL. and K-band spectra using NIRI on Gemini-North. (o test the irradiated brown dwar! moclel.," We have obtained new phase-resolved and -band spectra of EF Eri using NIRSPEC on KECK II, and -band spectra using NIRI on Gemini-North, to test the irradiated brown dwarf model."
 EF Eri was observed continuously from 7:21 until 9:00 UT on 24 December 2002 using on Gemini-North in queue scheduled mode., EF Eri was observed continuously from 7:21 until 9:00 UT on 24 December 2002 using on Gemini-North in queue scheduled mode.
 For these observations (he F/G camera. a four pixel slit ancl the were used to produce A-band spectra wilh a resolution of Rox 780.," For these observations the F/6 camera, a four pixel slit and the were used to produce band spectra with a resolution of R $\approx$ 780."
 Individual spectra. will exposure times of 120 s. were obtained al five dilferent positions along the slit.," Individual spectra, with exposure times of 120 s, were obtained at five different positions along the slit."
 Alter each exposure. the telescope was offset by G along the slit. and a new exposure started.," After each exposure, the telescope was offset by 6"" along the slit, and a new exposure started."
 In (his wav forty (vo individual spectra were obtained in 100 minules. covering of an orbital evele for EF Exi.," In this way forty two individual spectra were obtained in 100 minutes, covering of an orbital cycle for EF Eri."
 In addition to this set of continuous spectra. four additional spectra were obtained in (he period between 6:10 and 6:19 UT.," In addition to this set of continuous spectra, four additional spectra were obtained in the period between 6:10 and 6:19 UT."
AGN luminosity and the NLR metallicity is independent of the galaxy mass-metallcity relation.,AGN luminosity and the NLR metallicity is independent of the galaxy mass-metallcity relation.
 Whatever the scenario. our study strongly suggests that HzRGs completed their major chemical evolution in the very high-z universe. te.. z>4.," Whatever the scenario, our study strongly suggests that HzRGs completed their major chemical evolution in the very high-z universe, i.e., $z > 4$."
 If the minimum time-scale for a significant enrichment of carbon (~0.5 Gyr: e.g.. Matteucci 2008) 1s taken into account. the major epoch of the star formation in HzRGs may have occurred at z»5.," If the minimum time-scale for a significant enrichment of carbon $\sim0.5$ Gyr; e.g., Matteucci 2008) is taken into account, the major epoch of the star formation in HzRGs may have occurred at $z > 5$."
 Note that our conclusions are in contrast with some of the earlier studies mentioned in section |., Note that our conclusions are in contrast with some of the earlier studies mentioned in section 1.
" The most significant difference between previous works and our own is the adopted ""Slagnostics. 1e.. we do use emission to investigate the NLR metallicity."," The most significant difference between previous works and our own is the adopted diagnostics, i.e., we do use emission to investigate the NLR metallicity."
 While the emission is too weak in most HzRGs to be measured accurately. all of the lines used by us(Civ.Het. and i) are so strong that we can measure their fluxes rather easily.," While the emission is too weak in most HzRGs to be measured accurately, all of the lines used by us, and ]) are so strong that we can measure their fluxes rather easily."
 Therefore. we can investigate the chemical evolution of HzRGs without relying on upper-limit data., Therefore we can investigate the chemical evolution of HzRGs without relying on upper-limit data.
 Another serious problem in using the emission in metallicity studies is that the flux ratios including are sensitive not only to the NLR metallicity but also to. the ionization parameter as discussed in the following section., Another serious problem in using the emission in metallicity studies is that the flux ratios including are sensitive not only to the NLR metallicity but also to the ionization parameter as discussed in the following section.
 Given the high S/N of our spectra. the emission line is securely detected in 5 out of 9 objects.," Given the high S/N of our spectra, the emission line is securely detected in 5 out of 9 objects."
 Nitrogen is a secondary element and thus its abundance ts expected to be proportional to the square of the metallicity [1.e.. N/O « O/H. or equivalently. N/H « (O/HY].," Nitrogen is a secondary element and thus its abundance is expected to be proportional to the square of the metallicity [i.e., N/O $\propto$ O/H, or equivalently, N/H $\propto$ $\mathrm{(O/H)}^2$ ]."
 Therefore. the line is widely regarded as a good metallicity indicator. at least for BLRs (e.g.. Hamann Ferland 1992. 1993; Hamann et al.," Therefore, the line is widely regarded as a good metallicity indicator, at least for BLRs (e.g., Hamann Ferland 1992, 1993; Hamann et al."
 2002: Dietrich et al., 2002; Dietrich et al.
 2002. 2003: Nagao et al.," 2002, 2003; Nagao et al."
 2006c)., 2006c).
 Accordingly the line is sometimes regarded as a metallicity indicator also for NLRs (e.g.. van Ojik et al.," Accordingly the line is sometimes regarded as a metallicity indicator also for NLRs (e.g., van Ojik et al."
 1994: De Breuck et al., 1994; De Breuck et al.
 2000; Vernet et al., 2000; Vernet et al.
 2001: Overzier et al., 2001; Overzier et al.
 2001: Humphrey et al., 2001; Humphrey et al.
 2005)., 2008).
 We thus examine the and flux ratios as a function of redshift. to infer possible redshift evolution of the NLR metallicity independently of the diagnostic diagram of versusm[/Civ.," We thus examine the and flux ratios as a function of redshift, to infer possible redshift evolution of the NLR metallicity independently of the diagnostic diagram of versus."
 These two flux ratios are plotted as a function of redshift in Figure 9., These two flux ratios are plotted as a function of redshift in Figure 9.
 Our new data at z> clearly show the existence of HZRGs with high and flux ratios even at ς>2.7. which seems consistent to the idea that there is no significant evolution in the NLR metallicity of HZRGs as suggested by the and flux ratios.," Our new data at $z > 2.7$ clearly show the existence of HzRGs with high and flux ratios even at $z > 2.7$, which seems consistent to the idea that there is no significant evolution in the NLR metallicity of HzRGs as suggested by the and flux ratios."
 Figure 9 shows. however. the presence of a large number of HzRGs with à veryweak emission.," Figure 9 shows, however, the presence of a large number of HzRGs with a veryweak emission."
 In particular. some HzRGs show «0.1 andtv <0.1 while others show ~| and ~].," In particular, some HzRGs show $< 0.1$ and $< 0.1$ while others show $\sim 1$ and $\sim 1$."
 With the models presented by Vernet et al. (, With the models presented by Vernet et al. (
2001). these flux ratios would correspond to Ζνικ€O.A4Z. for v-weak HzRGs and Zug4Z« for v-strong HzROGs.,"2001), these flux ratios would correspond to $Z_\mathrm{NLR} < 0.4 Z_\mathrm{\odot}$ for -weak HzRGs and $Z_\mathrm{NLR} \sim 4 Z_\mathrm{\odot}$ for }-strong HzRGs."
 This would suggest that the NLR metallicity of HzRGs varies from object to object in a very wide range (>| dex)., This would suggest that the NLR metallicity of HzRGs varies from object to object in a very wide range $> 1$ dex).
 This seems inconsistent with the results obtainec through the analysis of the and emission-line ratios (Figure 3)., This seems inconsistent with the results obtained through the analysis of the and emission-line ratios (Figure 3).
 To investigate the origin of this apparent contradiction. in Figure 10 we plot the observational data oi the versus diagram and compare them with our model calculations.," To investigate the origin of this apparent contradiction, in Figure 10 we plot the observational data on the versus diagram and compare them with our model calculations."
 The distribution of the observational data on this diagram seems consistent with the area spannec by model grids., The distribution of the observational data on this diagram seems consistent with the area spanned by model grids.
 The near horizontal lines of constant-U 1 Figure 10 suggest that the flux ratio is sensitive to the ronization parameter but insensitive to the NLR metallicity., The near horizontal lines of $U$ in Figure 10 suggest that the flux ratio is sensitive to the ionization parameter but insensitive to the NLR metallicity.
 The ratio appears to depend on both U and Zwigs., The ratio appears to depend on both $U$ and $Z_\mathrm{NLR}$.
 I other words. the increase of the ronization parameter for a give metallicity results in high and flux ratios. while the increase of the metallicity for a given ionizatio parameter results in high but nearly constant flux ratios.," In other words, the increase of the ionization parameter for a given metallicity results in high and flux ratios, while the increase of the metallicity for a given ionization parameter results in high but nearly constant flux ratios."
 This suggests that the flux ratio is sensitive to the ionization parameter rather than to the NLR metallicity., This suggests that the flux ratio is sensitive to the ionization parameter rather than to the NLR metallicity.
 This is demonstrated more clearly in Figure 11. where we show the versus diagram.," This is demonstrated more clearly in Figure 11, where we show the versus diagram."
 The narrowcoverage of the model grids in this diagram suggests that both, The narrowcoverage of the model grids in this diagram suggests that both
sstucdies of selected NLS1s. Bolleretal.(1997) and Brancltetal.(1999). hypothesized that relativistic beamine by orbiting asvimnnetries in the inner accretion disk must be partly responsible for their variabilitv. which sometimes implies an efficiency greater (han that possible for an isotropically emitting accreting source.,"studies of selected NLS1s, \cite{bh97} and \cite{bbf99}
 hypothesized that relativistic beaming by orbiting asymmetries in the inner accretion disk must be partly responsible for their variability, which sometimes implies an efficiency greater than that possible for an isotropically emitting accreting source."
 If so. there should be at least some evidence for periodicitv in their light curves if the emitting structures retain (heir coherence for several orbits.," If so, there should be at least some evidence for periodicity in their light curves if the emitting structures retain their coherence for several orbits."
 It is that sort of quasi-periodic signal that we may be seeing in (hese long ]ieht curves., It is that sort of quasi-periodic signal that we may be seeing in these long light curves.
 However. we are süll [ar from developing a physical theory. for what creates such structures.," However, we are still far from developing a physical theory for what creates such structures."
 It appears that theory. will continue to lag behind observation in (his field until ancl unless some more detailed pattern to the observations emerges., It appears that theory will continue to lag behind observation in this field until and unless some more detailed pattern to the observations emerges.
 The foundation of the supermassive black hole model for AGNs rests squarely on their rapil X-rav variability. which establishes the compact nature of these liminous objects bevond a reasonable doubt.," The foundation of the supermassive black hole model for AGNs rests squarely on their rapid X-ray variability, which establishes the compact nature of these luminous objects beyond a reasonable doubt."
 Although AGN variability has been for various purposes. we have barely begun to address andwhy AGNs varv.," Although AGN variability has been for various purposes, we have barely begun to address and AGNs vary."
 The munerous “reverberation mapping canmpeigns have not. as a by-product. shed much lis£ht on this question.," The numerous “reverberation mapping” campaigns have not, as a by-product, shed much light on this question."
 Simultaneous. multiwavelength monitoring has cast a pall over (he whole enterprise. because it has become apparent that there is no simple relation among the lieht curves in (he various bands. e.g.. the studies of NGC! 7469 (Nandraetal.1993). and NGC 3516 (Edelsonetal.2000).," Simultaneous, multiwavelength monitoring has cast a pall over the whole enterprise, because it has become apparent that there is no simple relation among the light curves in the various bands, e.g., the studies of NGC 7469 \citep{n98} and NGC 3516 \citep{e00}."
 In NGC 5548. (here is some evidence (hat the underlving variability process is displaved in ihe EUV (Chiangetal.2000).," In NGC 5548, there is some evidence that the underlying variability process is displayed in the EUV \citep{crb00}."
. This might be true generally for Sevfert nuclei in which the peak of the spectral energy. distribution is in the EUV. but we still don't know how to interpret that variabilitv.," This might be true generally for Seyfert nuclei in which the peak of the spectral energy distribution is in the EUV, but we still don't know how to interpret that variability."
 A fundamental obstacle to understanding is the apparent lack of a characteristic (nme scale. a period or quasi-period that could be interpreted in terms of a ονπασα] or other effect related to the size of the putative inner accretion disk aud (he mass of the black hole.," A fundamental obstacle to understanding is the apparent lack of a characteristic time scale, a period or quasi-period that could be interpreted in terms of a dynamical or other effect related to the size of the putative inner accretion disk and the mass of the black hole."
 It is ciffieult to interpret the aperiodic variability (hat is characteristic of all accretion-powered objects including AGNs., It is difficult to interpret the aperiodic variability that is characteristic of all accretion-powered objects including AGNs.
 But if X-ray periods or quasi-periods could be found in just a few AGNs. they would provide a quantitative reference point for theory. just as QPOs in Galactic X-ray transients do lor stellar-mass black holes.," But if X-ray periods or quasi-periods could be found in just a few AGNs, they would provide a quantitative reference point for theory, just as QPOs in Galactic X-ray transients do for stellar-mass black holes."
 The lack of any confirmed X-ray periods in AGNs might be taken as evidence that they don't exist., The lack of any confirmed X-ray periods in AGNs might be taken as evidence that they don't exist.
 On (the other hand. the tentative detections reported here suggest that we may not vel have observed sufficiently on the required time scales.," On the other hand, the tentative detections reported here suggest that we may not yet have observed sufficiently on the required time scales."
 All three oobservalions (hat were 20 davs or more in duration show some indication of periodic behavior., All three observations that were 20 days or more in duration show some indication of periodic behavior.
 Evidence is mounting that a dedicated search for periodic phenomena in Sevfert, Evidence is mounting that a dedicated search for periodic phenomena in Seyfert
count rate and an average photon-iudex can distort substautialle the 6 relation.,count rate and an average photon-index can distort substantially the 6 relation.
 For each iudividual orbit. source counts were collected from a circular region of 12 arcsec radius. ceutered on the source position.," For each individual orbit, source counts were collected from a circular region of 12 arcsec radius, centered on the source position."
 Local backeround data were taken from nearby regions. separately for the PN. MOS1L aud MOS2 detectors where no source is found in the full exposure image.," Local background data were taken from nearby regions, separately for the PN, MOS1 and MOS2 detectors where no source is found in the full exposure image."
 The area of the background region had 200 arcsec radius., The area of the background region had 20 arcsec radius.
" The spectral data from individual exposures were sunuued up for the source and backeround. respectively. and the background subtraction was made assunuüues the conunuon scaling factor for the ποος""hbackgrouud geometrical areas."," The spectral data from individual exposures were summed up for the source and background, respectively, and the background subtraction was made assuming the common scaling factor for the source/background geometrical areas."
 Both PN aud MOS spectra are extracted in the 0.5-8 keV energy area., Both PN and MOS spectra are extracted in the 0.5-8 keV energy area.
 MOS1 and MOS2 spectra are sumed using the FTOOLS task., MOS1 and MOS2 spectra are summed using the FTOOLS task.
 Response and effective area. files were computed by averaging the iudividual files usine the FTOOLSADDRMEF audADDARF tasks., Response and effective area files were computed by averaging the individual files using the FTOOLS and tasks.
 We constructed SED to obtain an idea on the domiuaut powering mechanisin (ACN or star-formation) im the imid-IR part of the spectrmu., We constructed SED to obtain an idea on the dominant powering mechanism (AGN or star-formation) in the mid-IR part of the spectrum.
 The SEDs are also used iu order to obtain au accurate estimate of the Gian infrared taninositics., The SEDs are also used in order to obtain an accurate estimate of the $\rm 6\mu m$ infrared luminosities.
 We used a 7 nuunuisation techuiqne to fit the optimum conibinatious of host galaxy and pure AGN SED templates to the broad-band spectra of our objects., We used a $\chi^2$ minimisation technique to fit the optimum combinations of host galaxy and pure AGN SED templates to the broad-band spectra of our objects.
 The starburst templates we use. are those from he SWIRE template library of Pollettaetal.(2007).. which is a compilation of observed SED of nearby galaxies (M82. NGCGUO00.. Arp220. TRAS20517. IRAS22191).," The starburst templates we use, are those from the SWIRE template library of \citet{Polletta2007}, which is a compilation of observed SED of nearby galaxies (M82, NGC6090, Arp220, IRAS20547, IRAS22491)."
 Iu addition. we use the teiiplates of Chary&Elbaz(2001).," In addition, we use the templates of \citet{Chary2001}."
. The AGN templates we use come frou Silva.Maioliuo&Cranato (2001).. who combine nuclear SEDs of Sevterts with a rauge of absorption columus.," The AGN templates we use come from \citet*{Silva2004}, who combine nuclear SEDs of Seyferts with a range of absorption columns."
 We also coustruct a sample of AGN templates using the tvpe-1 QSO SED from Richardsetal.(2006) and applvine dust absorption following Rosenthal.Bertoldi&Drapatz(2000).. in order to account for the absorption feature.," We also construct a sample of AGN templates using the type-1 QSO SED from \citet{Richards2006} and applying dust absorption following \citet*{Rosenthal2000}, in order to account for the absorption feature."
 Iu Fig. 1..," In Fig. \ref{SED},"
 we show au example SED of Comptou-thick sources in both the high aud low- Universe., we show an example SED of Compton-thick sources in both the high and low-redshift Universe.
 We present the 6 diagram for the uearby sources In Fie. 2.., We present the 6 diagram for the nearby sources in Fig. \ref{brightman-lxl6}.
 The N-xav huuinosity is uncorrected for obscuration while the 6jan luminosity iucludes both the torus aud the star-formation component (see discussion)., The X-ray luminosity is uncorrected for obscuration while the $\rm \mu m$ luminosity includes both the torus and the star-formation component (see discussion).
 The lines denote various Compton-thick ACN areas., The lines denote various Compton-thick AGN areas.
 These have been derived in the following fashion., These have been derived in the following fashion.
 We adopt the average 6 Iluiiuosity depeudent ACN relation of Fioreetal./Lo(2009) derived in the COSMOS ποια., We adopt the average 6 luminosity dependent AGN relation of \citet{Fiore2009} derived in the COSMOS field.
 We then scale down this relation by a factor of 30., We then scale down this relation by a factor of 30.
 This is an approximation to the average reflected nubmositv relative to the iuntriusic cnussion from the vackside of the torus in the kkeV. band in reflection-dominated ACN (c.g.Comastri.2001)., This is an approximation to the average reflected luminosity relative to the intrinsic emission from the backside of the torus in the keV band in reflection-dominated AGN \citep[e.g.][]{Comastri2004}.
". Alternatively, is the relation between the absorbed ancl unabsorbed LOkkeV bhuuinosty for a column density of 1.5«107! and a power spectrum o| [=Ls."," Alternatively, is the relation between the absorbed and unabsorbed keV luminosity for a column density of $1.5\times10^{24}$ and a power-law spectrum of $\Gamma=1.8$."
 We show an alternative scenario. where the relation of (2009) is scaled for oulv reflected eiission (Vienalietal..2001).," We show an alternative scenario, where the relation of \citet{Fiore2009} is scaled for only reflected emission \citep{Vignali2004}."
. We also give the 6 relation derived by Lutzctal.(2001). iu the local Universe. together with ifs associate lo uncertainty scalec own again bv (dashed area).," We also give the 6 relation derived by \citet{Lutz2004} in the local Universe, together with its associated $1\sigma$ uncertainty scaled down again by (dashed area)."
 Throughout the paper. we use only the Fioreetal.(2009). reflection limiti for assessing the munber of Comptou-thick sources.," Throughout the paper, we use only the \citet{Fiore2009} reflection limit for assessing the number of Compton-thick sources."
 We see that most (ten out of eleven) of the sources classified as Comptou-tlick. on the basis of N-ray spectroscopy lie below the Coupton-thick lines., We see that most (ten out of eleven) of the sources classified as Compton-thick on the basis of X-ray spectroscopy lie below the Compton-thick lines.
 However. we see that many more sources (twelve) which are not Comptou-thick. according to the N-xav. spectroscopic diagnostics. would be classified as caucidate Comptou-thick AGN on the basis of the low 6 ratio.," However, we see that many more sources (twelve) which are not Compton-thick, according to the X-ray spectroscopic diagnostics, would be classified as candidate Compton-thick AGN on the basis of the low 6 ratio."
 All these are presented in Table l., All these are presented in Table \ref{brightman}.
 It is uot/Lo unlikely that some of the low 6 sources are indeed Comptou-thick sources. but tailed to identify them owing to the limited photon statistics.," It is not unlikely that some of the low 6 sources are indeed Compton-thick sources, but failed to identify them owing to the limited photon statistics."
 At least two of them διςI21.. Gwasawactal..2001:Burlonot20101). and UCCs058 (AGK231). Braitoetal.(2001)... are Comptou-thick ACN according toSWIFT audBeppoSaN spectroscopy.," At least two of them NGC424, \citep{Iwasawa2001,Burlon2011} and UGC8058 (Mrk231), \citet{Braito2004}, are Compton-thick AGN according to and spectroscopy."
 We exclude all sources with anΣΕ huuünositv of Lx«107 units. as many of these could be associated with normal ealaxies.," We exclude all sources with an X-ray luminosity of $\rm L_X<10^{42}$ , as many of these could be associated with normal galaxies."
 282 sourceshave bhpuumosities above this level., 283 sourceshave luminosities above this level.
 The 6 diagram ix shown in Fig. 3.., The 6 diagram is shown in Fig. \ref{cdfs-lxl6}.
 There are iue sources/Leo which would be characterised as Comptou-hick ou the beds of the low 6 ratio (see Table 2))., There are nine sources which would be characterised as Compton-thick on the basis of the low 6 ratio (see Table \ref{cdfs}) ).
 In Table 2. we preseut the spectral Bt results., In Table \ref{cdfs} we present the spectral fit results.
 Throughout the paper. we use the Luoctal.(2008) identification nunbers.," Throughout the paper, we use the \citet{Luo2008} identification numbers."
 Four 6 sources could be classified as Couptou-thick: CDFES-2309. CDFS-95. CDFS-197. and CDES-398.," Four 6 sources could be classified as Compton-thick: CDFS-309, CDFS-95, CDFS-197, and CDFS-398."
 CDFS-309 has beeu discussed in detail iu Feruelioctal.(2011) acl was found to present a Comptou-thick source on the basis of au Felklko line with a high EW., CDFS-309 has been discussed in detail in \citet{Feruglio2011} and was found to present a Compton-thick source on the basis of an $\alpha$ line with a high EW.
 Due to the limuted photon-statistics we cannot tell whether the source is fat Dom0.6 or obseured with Ny~3«1075 ., Due to the limited photon-statistics we cannot tell whether the source is flat $\Gamma\approx0.6$ or obscured with $\rm N_H\sim 3\times 10^{23}$ .
. However. the addition of an Fela line miproves the statistic bv ACzm20 yielding an EW of 3.641.8 keV. CDFS-95and CDES-197 have colin densities somewhat below the Coiptou-thick limit (~9«102? ," However, the addition of an $\alpha$ line improves the statistic by $\Delta C \approx20$ yielding an EW of $\pm 1.8$ keV. CDFS-95and CDFS-197 have column densities somewhat below the Compton-thick limit $\sim 9 \times 10^{23}$ "
about 0.4’.,about 0.4 $\!''$.
 The observations were made in “track-sharing”™. to observe LkCa 15 and the Herbig star MWC 480 with a common calibration curve.," The observations were made in “track-sharing”, to observe LkCa 15 and the Herbig star MWC 480 with a common calibration curve."
 Any morphological difference between the two targets is thus genuine., Any morphological difference between the two targets is thus genuine.
 They found a clear inner hole around LkCa 15 while MWC 480 is centrally peaked., They found a clear inner hole around LkCa 15 while MWC 480 is centrally peaked.
 This is confirmed by the modeling of the data. with a derived inner hole of ~46 AU (re. the size of our own Solar System).," This is confirmed by the modeling of the data, with a derived inner hole of $\sim 46$ AU (i.e. the size of our own Solar System)."
 A multi isotopes analysis of CO rotational lines led ? to derive the physical properties of the outer circumstellar disk surrounding the star., A multi isotopes analysis of CO rotational lines led \cite{2007A&A...467..163P} to derive the physical properties of the outer circumstellar disk surrounding the star.
 Contrary to the other system studied in a similar way. no clear vertical temperature gradient was found in the disk structure. possibly due to the peculiar geometry of LkCa 15 disk.," Contrary to the other system studied in a similar way, no clear vertical temperature gradient was found in the disk structure, possibly due to the peculiar geometry of LkCa 15 disk."
 ? discussed the possible mechanisms that could explain the inner hole structure (notcompletelyempty.assomeIRex-cessIspresentabovethestellarblackbody.See ?)..," \cite{2006A&A...460L..43P} discussed the possible mechanisms that could explain the inner hole structure \citep[not completely empty, as
 some IR excess is present above the stellar black
 body. See][]{2004ApJ...614L.133B}."
 Planetary formation or the presence of a low-mass companion seems to be plausible since the LkCal5 disk is quite massive., Planetary formation or the presence of a low-mass companion seems to be plausible since the LkCa15 disk is quite massive.
" In fact ? estimated. from the kinematics. a total mass for the system of Mq,=1.004O.LM.."," In fact \cite{2000ApJ...545.1034S}
 estimated, from the kinematics, a total mass for the system of $M_{\rm{dyn}}=1.00\pm 0.1 M_{\odot}$."
 This set the maximal mass of the putative companion to 0.2M. since the mass of LkCa 15 could hardly be less than 0.8M... due to its spectral type.," This set the maximal mass of the putative companion to $0.2~M_{\odot}$, since the mass of LkCa 15 could hardly be less than $0.8~M_{\odot}$, due to its spectral type."
 This value is in agreeement with a recent estimate of the mass of the star obtained by ?.., This value is in agreeement with a recent estimate of the mass of the star obtained by \cite{2007A&A...473L..21B}.
 Using a new distance estimate by ?.. they confirmed LkCa 15 as a kinematic member of the Taurus association seen through a hole in the molecular cloud. retrieving a mass value of 1.12+0.08M...," Using a new distance estimate by \cite{2006A&A...460..499B}, they confirmed LkCa 15 as a kinematic member of the Taurus association seen through a hole in the molecular cloud, retrieving a mass value of $1.12\pm 0.08 M_{\odot}$."
 Additional results obtained by ? also suggest that other mechanisms such as photo-evaporation are unlikely to. be effective due to the large disk density. which ts significantly higher than the one expected if the photo-evaporation starts propagating beyond 20-30 AU.," Additional results obtained by \cite{2006MNRAS.369..229A} also suggest that other mechanisms such as photo-evaporation are unlikely to be effective due to the large disk density, which is significantly higher than the one expected if the photo-evaporation starts propagating beyond 20-30 AU."
 ? suggest that a ~5—lOMjup planet orbiting at 30 AU would be sufficient to evacuate the inner 50 AU of the LkCa 15 disk.," \cite{2006A&A...460L..43P} suggest that a $\sim 5-10
M_{\rm{Jup}}$ planet orbiting at 30 AU would be sufficient to evacuate the inner 50 AU of the LkCa 15 disk."
 Similar conclusions were reached by ? with Spitzer observations and the modeling of the SED of LkCa 15., Similar conclusions were reached by \cite{2007ApJ...670L.135E} with Spitzer observations and the modeling of the SED of LkCa 15.
 Their analysis suggests that a gap is present in the disk of LkCa 15. with an inner disk going from AU to 4-5 AU. and an outer disk of inner radius 46 AU. in perfect agreement with the findings of ?..," Their analysis suggests that a gap is present in the disk of LkCa 15, with an inner disk going from 0.12-0.15 AU to 4-5 AU, and an outer disk of inner radius 46 AU, in perfect agreement with the findings of \cite{2006A&A...460L..43P}."
 They also concluded that planetary formation or the presence of a close-in stellar or sub-stellar companion are the most probable explanations for the circumstellar material shape around LkCal5 The observations were performed on December 26th. 2007 at ESO/Paranal. using NACO. the AO-assisted near-IR camera NAOS-CONICA (?) mounted on one of the Nasmyth focus of the UT4 8m-telescope.," They also concluded that planetary formation or the presence of a close-in stellar or sub-stellar companion are the most probable explanations for the circumstellar material shape around LkCa15 The observations were performed on December 26th, 2007 at ESO/Paranal, using NACO, the AO-assisted near-IR camera NAOS-CONICA \citep{2003SPIE.4839..140R} mounted on one of the Nasmyth focus of the UT4 8m-telescope."
 Among the numerous NACO observing modes (?).. the classical and coronagraphic imaging modes were used.," Among the numerous NACO observing modes \citep{2003SPIE.4841..944L}, the classical and coronagraphic imaging modes were used."
" Coronagraphic observations were performed with the four-quadrant phase mask (4QPM) optimized for K, band observations.", Coronagraphic observations were performed with the four-quadrant phase mask (4QPM) optimized for $K_s$ band observations.
" The S13 objective (FoV of 14""x and plate-scale of 13.25 mas/pixel) was chosen for a more precise centering with the +QPM and a better sampling of the PSF.", The S13 objective (FoV of $14~\!'' \times 14~\!''$ and plate-scale of 13.25 mas/pixel) was chosen for a more precise centering with the 4QPM and a better sampling of the PSF.
 The 4QPM splits the focal plane into four equal areas. two of which are phase-shifted by π.," The 4QPM splits the focal plane into four equal areas, two of which are phase-shifted by $\pi$."
 As a consequence. a destructive interference occurs in the relayed pupil and the on-axis starlight rejected on the edge of the geometric pupil is filtered with a Lyot stop. being a circular hole of the pupil size.," As a consequence, a destructive interference occurs in the relayed pupil and the on-axis starlight rejected on the edge of the geometric pupil is filtered with a Lyot stop, being a circular hole of the pupil size."
 The advantage over the classical Lyot mask 1s the possibility to access inner angular separations lower than 0.35” (the smallest NACO occulting mask) at relatively large contrast 2008 ).., The advantage over the classical Lyot mask is the possibility to access inner angular separations lower than $0.35~\!''$ (the smallest NACO occulting mask) at relatively large contrast \citep[see][2008]{2004PASP..116.1061B}.
 However. a significant part of the starlight is left in the focal plane due to uncorrected aberrations composed of a dynamical halo averaging over time plus a quasi-static halo corresponding to optical aberrations along the optical train (from telescope to detector).," However, a significant part of the starlight is left in the focal plane due to uncorrected aberrations composed of a dynamical halo averaging over time plus a quasi-static halo corresponding to optical aberrations along the optical train (from telescope to detector)."
 To mitigate this problem. a coronagraphie image of a reference star was taken just after our science target observations. with the same instrumental settings. to serve for speckle calibration in an image-subtraction process.," To mitigate this problem, a coronagraphic image of a reference star was taken just after our science target observations, with the same instrumental settings, to serve for speckle calibration in an image-subtraction process."
 This reference star (BD + 22 729. V=Τ.Σ. K= 7.9) has similar visible and NIR magnitudes to ensure similar AO correction and signal-to-noise at the detector.," This reference star (BD + 22 729, $V
=11.5$, $K= 7.9$ ) has similar visible and NIR magnitudes to ensure similar AO correction and signal-to-noise at the detector."
 Moreover. observations were performed to match the parallactic angle with LkCaló5. 1.9 the instrumental pupil configuration. in order to optimize the overlap of the speckle pattern and diffraction spikes position in the final subtracted image of LkCal5.," Moreover, observations were performed to match the parallactic angle with LkCa15, i.e the instrumental pupil configuration, in order to optimize the overlap of the speckle pattern and diffraction spikes position in the final subtracted image of LkCa15."
 The coronagraphie observations were preceded by a classical imaging sequence that provides a photometric reference., The coronagraphic observations were preceded by a classical imaging sequence that provides a photometric reference.
" A neutral density transmission) and a narrow band filters at 2.17 yam were used to serve as an image quality check and as photometric reference for the coronagraphic observations,", A neutral density transmission) and a narrow band filters at 2.17 $\mu$ m were used to serve as an image quality check and as photometric reference for the coronagraphic observations.
 Table | summarizes the observing parameters., Table \ref{tab:obs1} summarizes the observing parameters.
 During the coronagraphie observing sequence. the precise centering of the science target behind the focal plane mask was critical to maximize the central star attenuation.," During the coronagraphic observing sequence, the precise centering of the science target behind the focal plane mask was critical to maximize the central star attenuation."
 The coronagraphic, The coronagraphic
"We start by exploring the initial distribution in the form (33)) with zo=5, c=0.7 but now we take Tj9= so that the disk mass Mg=350Mp>>Mo.","We start by exploring the initial distribution in the form \ref{eq:prof}) ) with $x_0=5$, $\sigma=0.7$ but now we take $\tau_{\parallel,0}=280$ so that the disk mass $M_d=350M_0\gg M_0$."
 Resulting evolution of rj is shown in Figure 3.., Resulting evolution of $\tau_\parallel$ is shown in Figure \ref{fig:gauss5}.
" Apart from the theoretically expected steady decrease of rj with time in the bulk of the disk where rj>>1 one immediately notices two interesting features of the solutions: (1) a sharp drop of the disk surface density at the outer edge of the disk, and (2) an optically thin (rj<1) tail of debris extending from the optically thick region all the way to Rin, which remains stable throughout the evolution of the disk."," Apart from the theoretically expected steady decrease of $\tau_\parallel$ with time in the bulk of the disk where $\tau_\parallel\gg 1$ one immediately notices two interesting features of the solutions: (1) a sharp drop of the disk surface density at the outer edge of the disk, and (2) an optically thin $\tau_\parallel\lesssim 1$ ) tail of debris extending from the optically thick region all the way to $R_{in}$, which remains stable throughout the evolution of the disk."
 We now discuss the nature of these features in more detail., We now discuss the nature of these features in more detail.
" Sharp outer edge appears because the evolution timescale of a particular disk region ttnickcxΤι, see equation (40))."," Sharp outer edge appears because the evolution timescale of a particular disk region $t_{thick}\propto\tpar$, see equation \ref{eq:T_thick}) )."
 Initial density distribution in the form (33)) has increasinginwards in the outer part of the ring., Initial density distribution in the form \ref{eq:prof}) ) has $\tpar$ increasing in the outer part of the ring.
 This τιcauses the outermost extremity of the ring is optically thin and evolves on timescale t5;) to (whichmigrate towards the WD faster than the parts of the ring closer to the WD do., This causes the outermost extremity of the ring (which is optically thin and evolves on timescale $t_{thin}$ ) to migrate towards the WD faster than the parts of the ring closer to the WD do.
" As a result, the outermost disk material catches up with the debris at smaller x having higher leading to a debris pileup there and formation of a sharp7j, outer edge."," As a result, the outermost disk material catches up with the debris at smaller $x$ having higher $\tpar$, leading to a debris pileup there and formation of a sharp outer edge."
" This edge appears on a timescale ~μή Since this is the characteristic time, on which the material from the optically thin outer disk region reaches x£x9 where most of the ring mass is concentrated."," This edge appears on a timescale $\sim t_{thin}$ since this is the characteristic time, on which the material from the optically thin outer disk region reaches $x\approx x_0$ where most of the ring mass is concentrated."
 We provide a more rigorous justification for the outer edge appearance in Appendix A.., We provide a more rigorous justification for the outer edge appearance in Appendix \ref{app:1}.
 Origin of the quasi-stationary tail with τιS;1 at small r can be understood as follows., Origin of the quasi-stationary tail with $\tau_\parallel\lesssim 1$ at small $r$ can be understood as follows.
" Even if initially the disk has τι>>1 everywhere such tail would rapidly form since, according to equation (38)), at any distance from the WD 7 should switch into the optically thin regime within finite time."," Even if initially the disk has $\tpar\gg 1$ everywhere such tail would rapidly form since, according to equation \ref{eq:thick_sol}) ), at any distance from the WD $\tpar$ should switch into the optically thin regime within finite time."
 This happens fastest at the inner edge of the disk., This happens fastest at the inner edge of the disk.
 After the tail has formed it evolves on atimescale tinin given by equation (26))., After the tail has formed it evolves on atimescale $t_{thin}$ given by equation \ref{eq:T_thin}) ).
 Inner parts of the tail evolve more rapidly than the outer ones and a quasi-steady state is attained in the tail with M almost independent of r., Inner parts of the tail evolve more rapidly than the outer ones and a quasi-steady state is attained in the tail with $\dot M$ almost independent of $r$.
" This point is illustrated in Figure 4,, where we show the run of M(x)/M. for a calculation displayed in Figure 3 at different moments of time."," This point is illustrated in Figure \ref{fig:y}, where we show the run of $\dot M(x)/\dot M_\infty$ for a calculation displayed in Figure \ref{fig:gauss5}
 at different moments of time."
 One can easily see that after some evolution has taken place M(x) indeed becomes independent of x in the tail., One can easily see that after some evolution has taken place $\dot M(x)$ indeed becomes independent of $x$ in the tail.
" Thus, the behavior of rj(z, in the tail can be described by the equation (cf."," Thus, the behavior of $\tpar(x,T)$ in the tail can be described by the equation (cf."
" equationT) (19))) (T)-- where Mtaii(t) is the mass accretion rate through the tail yia; is the dimensionless analog of this quantity), (andwhich in general is a function of time."," equation \ref{eq:y}) )) (T)= where $\dot M_{tail}(t)$ is the mass accretion rate through the tail (and $y_{tail}$ is the dimensionless analog of this quantity), which in general is a function of time."
 Equation (41)) describes the behavior of 7 in Figure 3 as function of x quite well., Equation \ref{eq:tail}) ) describes the behavior of $\tpar$ in Figure \ref{fig:gauss5} asa function of $x$ quite well.
" At asome radius z;,21=r;-1/R;, the tail becomes optically thick, i.e. rj(z,—1)= 1."," At some radius $x_{\tau=1}=r_{\tau=1}/R_{in}$ the tail becomes optically thick, i.e. $\tpar(x_{\tau=1})=1$ ."
" A this point the tail merges with the optically thick part of the disk and the evolution timescale, which is now given by ttnice, rapidly increases, see equation (40))."," A this point the tail merges with the optically thick part of the disk and the evolution timescale, which is now given by $t_{thick}$ , rapidly increases, see equation \ref{eq:T_thick}) )."
" Accretion rate through the tail Misü(T) is set almostexclusively by the value of z,—; since by definition 1—exp[-ry(x,—,,T)]~ independent of T, so that"," Accretion rate through the tail $\dot M_{tail}(T)$ is set almostexclusively by the value of $x_{\tau=1}$ since by definition $1-\exp\left[-\tau_\parallel(x_{\tau=1},T)\right]\sim 1$ independent of $T$ , so that"
‘These 7shells. bubbles. ares” could be associated mainly o explosions of SNe and the final phase of the ealacticwind. ie. the blowout phase of galactic bubbles (Norman Ikeuchi. 1989).,"These “shells, bubbles, arcs"" could be associated mainly to explosions of SNe and the final phase of the galactic--wind, i.e., the blow–out phase of galactic bubbles (Norman Ikeuchi 1989)."
 Multiple explosion of SNe. from massive o»ogenitors. are the main ealactic objects capable to generate this blowout phase (Norman Lkeuchi 1989).," Multiple explosion of SNe, from massive progenitors, are the main galactic objects capable to generate this blow–out phase (Norman Ikeuchi 1989)."
 Alassive YSC's ancl associations of giant regions are the aces. where multiple SN explosions are expected.," Massive YSCs and associations of giant regions are the places, where multiple SN explosions are expected."
 We have started a. multiwavelength studies of the voung stellar. cluster candidates (¥SCe) in NGC 3256. using XCS/UV. NEDPC2/Optical. NICALOS/near-LR images.," We have started a multi–wavelength studies of the young stellar cluster candidates (YSCc) in NGC 3256, using ACS/UV, WFPC2/Optical, NICMOS/near-IR images."
 From these cata we study the colours. ages. masses. subgroups. spatial distribution and. metallicities of these YSCc.," From these data we study the colours, ages, masses, subgroups, spatial distribution and metallicities of these YSCc."
 Here we report new results from this study., Here we report new results from this study.
 First. the colour. distribution. the colourmagnitude diagram. and the spatial distribution of YSCc (of. NGC 3256) were analysed.," First, the colour distribution, the colour–magnitude diagram and the spatial distribution of YSCc (of NGC 3256) were analysed."
 Using mainly the Catalogue of YSCc published by ζορ ct al. (, Using mainly the Catalogue of YSCc published by Zepf et al. (
1999: their Table. 1).,1999: their Table 1).
 We note hat from this catalogue. the positions of YSC'c in. pixels (ancl also the corrected. values of RA and. DEC) were used: since the published data of RA and DEC have small errors. which were generated in the transformation [rom jxxels to RADEC.," We note that from this catalogue, the positions of YSCc in pixels (and also the corrected values of RA and DEC) were used; since the published data of RA and DEC have small errors, which were generated in the transformation from pixels to RA–DEC."
 An electronic corrected. version of this Catalogue of YSCe is available at zepfepa.msu.edu. and iparittoac.uncor.edu., An electronic corrected version of this Catalogue of YSCc is available at zepfpa.msu.edu and liparioac.uncor.edu.
" From this study. the following results were found: These results are in agreemento with a ""composite multiple merger model proposed. for Νας 3256 (Lipari et al."," From this study, the following results were found: These results are in agreement with a “composite"" multiple merger model proposed for NGC 3256 (Lipari et al."
 2000)., 2000).
 In this model the main merger could be between two gas rich massive spiral galaxies (Sc). and then with a satellite galaxy (of one of these Se systems).," In this model the main merger could be between two gas rich massive spiral galaxies (Sc), and then with a satellite galaxy (of one of these Sc systems)."
 In particular. we proposed that the kinematics ancl spectroscopic properties of the core of region 2 could be associated to a 3rd.," In particular, we proposed that the kinematics and spectroscopic properties of the core of region 2 could be associated to a 3rd."
 nucleus. arisen in a satellite of the two original spiral galaxies that collided.," nucleus, arisen in a satellite of the two original spiral galaxies that collided."
 In. order to explore these type of scenarios. we are working in detailed numerical simulations. which are in progress (Lipari ct al.," In order to explore these type of scenarios we are working in detailed numerical simulations, which are in progress (Lipari et al."
 2000. 2004. in preparation).," 2000, 2004, in preparation)."
 The results of these kind of simulations suggest that the two, The results of these kind of simulations suggest that the two
already decayed before the grain fictionally couple to the gas (tp>τι) for eddy sizes /«li.,already decayed before the grain fictionally couple to the gas $\tau_f>\tau_t$ ) for eddy sizes $l<l_{\rm cri}$.
 The erains equilibrate with the gas. hence {μον move together with the turbulent. eclelies only if zy<7 lor which />/44 follows.," The grains equilibrate with the gas, hence they move together with the turbulent eddies only if $\tau_f<\tau_t$ for which $l>l_{\rm cri}$ follows."
 Figure LL demonstrates that the behaviour is rather diverse in the Brown Dwarf case for our choice of /=lem. pointing to a multi-scale problem that only. very approximately can be described bv one pre-chosen scale (see also Telling 2005).," Figure \ref{fig:v_turb_ind} demonstrates that the behaviour is rather diverse in the Brown Dwarf case for our choice of $l=1$ cm, pointing to a multi-scale problem that only very approximately can be described by one pre-chosen scale (see also Helling 2005)."
 In the planetary. alinosphere case. our choice of /=lem suggests (hat all erains are [rictionallv coupled.," In the planetary atmosphere case, our choice of $l=1$ cm suggests that all grains are frictionally coupled."
 If 7 is treated in more detail. the impact of turbulent motions on dust-dust collisions will be larger in the planetary. atmosphere because /4;<Lem.," If $l$ is treated in more detail, the impact of turbulent motions on dust-dust collisions will be larger in the planetary atmosphere because $l_{\rm cri}<1$ cm."
 Note. however. that grains will decouple from the gas due to gravity on (he largest scales in the almosphere.," Note, however, that grains will decouple from the gas due to gravity on the largest scales in the atmosphere."
 Fieure 11 (bottom panels) shows the turbulence-induced relative (drift) velocity Avingi for collisions of erains of different. sizes (dashed line. Eq. 14))," Figure \ref{fig:v_turb_ind} (bottom panels) shows the turbulence-induced relative (drift) velocity ${\Delta \rm v}_{\rm
ind,t}$ for collisions of grains of different sizes (dashed line, Eq. \ref{eq:vindt}) )"
 and for grains of the same size (solid ancl dotted line. Eq. 15))," and for grains of the same size (solid and dotted line, Eq. \ref{eq:vindt2}) )"
 over the whole aüimospheric pressure range., over the whole atmospheric pressure range.
 The turbulence driving velocity. μα). is plotted in comparison.," The turbulence driving velocity, $u_{\rm max}(p_{\rm gas})$, is plotted in comparison."
" There is a maximum ol ναι in both atmosphere cases. brown dwarf and gas planet. whereas Avia, clearly decreases towards higher pressures."," There is a maximum of ${\Delta \rm v}_{\rm
ind,t}$ in both atmosphere cases, brown dwarf and gas planet, whereas ${\Delta \rm v}_{\rm ind,t}$ clearly decreases towards higher pressures."
 In these inner atmospheric regions. 7; decreases because the σας density increases. ancl therefore the dust grains can adjust more quickly to the gas motion.," In these inner atmospheric regions, $\tau_f$ decreases because the gas density increases, and therefore the dust grains can adjust more quickly to the gas motion."
" The collisional energies according to Eqs. 2.. 5.. Ον,"," The collisional energies according to Eqs. \ref{eq:Ecol}, \ref{eq:vsed1}, \ref{eq:vsed2},"
 aud Eqs. 7..," and Eqs. \ref{eq:
vrela1a2},"
 and Eq., and Eq.
 15. are shown in Fig. 12..," \ref{eq:vindt2} are shown in Fig. \ref{fig:col_energy},"
 together with the ionisation potential (work Dunction) for a mix of solid materials as il is expected for a grain surface based on our dust formation model., together with the ionisation potential (work function) for a mix of solid materials as it is expected for a grain surface based on our dust formation model.
 We explore the, We explore the
The value >=1 under the above conditions (ιοπιαπας A = const.),The value $\beta = 1$ under the above conditions (including $\Delta$ = const.)
 is only a theoretical lower linit: the XD relation could be steeper due to the these effects. as can be seen from Equation (26).," is only a theoretical lower limit; the $\Sigma-D$ relation could be steeper due to the these effects, as can be seen from Equation (26)."
 The relation given here is derived under the assumption of spherical svuuucetric expansion., The relation given here is derived under the assumption of spherical symmetric expansion.
 [owever. we know that the large number of PNe show a high asvunuetry aud can have a very complex morphology.," However, we know that the large number of PNe show a high asymmetry and can have a very complex morphology."
 This may explain why the theoretical XD relation is not necessary connected to the empirical relations that we discuss in Section Ll., This may explain why the theoretical $\Sigma-D$ relation is not necessary connected to the empirical relations that we discuss in Section 4.
 Even if we ad a spherically svuuuetric PNe with ©«D1. the coustaut of proportionality would ο different for ditterci phases of evolution. as can be inferred roni Section 2.," Even if we had a spherically symmetric PNe with $\Sigma \propto D^{-1}$, the constant of proportionality would be different for different phases of evolution, as can be inferred from Section 2."
 Moreover. xuanmeters eoveruiug the evolution (V. v. AT. ii) ave not he sae for all PNe. so 10 unique enmipirical relation for PNe is feasible.," Moreover, parameters governing the evolution $V$, $v$, $\dot{M}$, $\dot{m}$ ) are not the same for all PNe, so no unique empirical relation for PNe is feasible."
 Iu the final stage of evolution. after the fast wind las ceased. the isothermal shock wave continuously rturbs ACB wind.," In the final stage of evolution, after the fast wind has ceased, the isothermal shock wave continuously perturbs AGB wind."
 Nonetheless. the surface briglituess steeply declines. as demonstrated in Paper L because the nechanical cuerey input (acceleration) from the ceutral star no longer exists.," Nonetheless, the surface brightness steeply declines, as demonstrated in Paper I, because the mechanical energy input (acceleration) from the central star no longer exists."
". Using. p,;5;w5Πτι the YD relation in the final stage of PN evolution can be expressed an: The most mniportaut prerequisite for defining a proper clupiticalS—D relationis the extraction of a valid sample of PNe."," Using $\rho_{\scriptscriptstyle AGB}\propto R_s^{-2}$, the $\Sigma-D$ relation in the final stage of PN evolution can be expressed as: The most important prerequisite for defining a proper empirical $\Sigma-D$ relation is the extraction of a valid sample of PNe."
 This sample should cousist of calibrators having well determined distances., This sample should consist of calibrators having well determined distances.
 The distances to the calibrators have to be determined by reliable methods., The distances to the calibrators have to be determined by reliable methods.
 These may inclide trigonometric or spectroscopic parallaxes of PN ceutral stars and inethods based on nebulae expansion., These may include trigonometric or spectroscopic parallaxes of PN central stars and methods based on nebulae expansion.
 Although limited by current technology. trigonometric parallax is the most direct method of measuring distances.," Although limited by current technology, trigonometric parallax is the most direct method of measuring distances."
 All extracted PN sauples suffer from severe sclectiou effects caused bv Tinitations d survey seusitfivitv aud resolutiou., All extracted PN samples suffer from severe selection effects caused by limitations in survey sensitivity and resolution.
 For Galactic samples. Maliiquist may be the most severe effect: this is simular to the situation for Galactic SNR samples (Urosevicctal.2005) ).," For Galactic samples, Malmquist may be the most severe effect; this is similar to the situation for Galactic SNR samples \cite{Uro05}) )."
 For the literature extracted PN samples in this paper. we used fux deusities at 5 Πε.," For the literature extracted PN samples in this paper, we used flux densities at 5 GHz."
 This is because most PNe are expected to be optically thin at this frequency (seo Section 5)., This is because most PNe are expected to be optically thin at this frequency (see Section 5).
 Radio diameters would be preferable. however for most PNe used here. only optical diameters are available.," Radio diameters would be preferable, however for most PNe used here, only optical diameters are available."
" Correspouding X, aud L,—D diagrams are shown in Fie.", Corresponding $\Sigma_\nu-D$ and $L_\nu-D$ diagrams are shown in Fig.
 1., 1.
 Alahuquist bias tends to increase the slope of the YD velation for SNRs (Uroseviéetal. 2005))., Malmquist bias tends to increase the slope of the $\Sigma-D$ relation for SNRs \cite{Uro05}) ).
 Similar iulicatious should be valid for PNe as well., Similar implications should be valid for PNe as well.
 Slopes from Table | Sugeest that some of our sample sets may sitter roni a Maluiquist bias., Slopes from Table 1 suggest that some of our sample sets may suffer from a Malmquist bias.
 Iu spite of the possibility of Maliiquist. bias. the uajoritv of the XD slopes listed Tables 1 aud 2. are very close to the so-called trivial X/—D relation with )=2 (Arbutinaetal. 20011).," In spite of the possibility of Malmquist bias, the majority of the $\Sigma-D$ slopes listed Tables 1 and 2, are very close to the so-called trivial $\Sigma-D$ relation with $\beta=2$ \cite{Arbutina04}) )."
 This can be uuderstood nathematically frou: where 5 is the flux density. O is the solid angle occupied by the source aud d dis the source distance.," This can be understood mathematically from: where $S$ is the flux density, $\Omega$ is the solid angle occupied by the source and $d$ is the source distance."
 The usefulness, The usefulness
For collisionswe between equal-sized. particles.. with: Ty. the expressionBE: reduces. to A time-dependent numerical solution of à collisional particle system must take collisions into account when choosing the time-step.,"For collisions between equal-sized particles, with $\tau_{\rm
f}^{(j)}=\tau_{\rm f}^{(k)}$ , the expression reduces to A time-dependent numerical solution of a collisional particle system must take collisions into account when choosing the time-step."
" The time-step criterion of the Monte Carlo collision scheme originates in the requirement that two particles can collide at most once during a time-step. 1.89. the collision probability P=ór/r, between any two particles in the same grid cell must be much smaller the unity."," The time-step criterion of the Monte Carlo collision scheme originates in the requirement that two particles can collide at most once during a time-step, i.e. the collision probability $P=\delta t/\tau_{\rm c}$ between any two particles in the same grid cell must be much smaller the unity."
 This step is Independent of the maximum density in a grid cell. since particles in dense grid cells have many collision partners and hence can suffer more collisions in the same time-step.," This time-step is independent of the maximum density in a grid cell, since particles in dense grid cells have many collision partners and hence can suffer more collisions in the same time-step."
" In the streaming instability simulations presented in and we observe a typical particle rms speed ὃν~ 0.025c,."," In the streaming instability simulations presented in and we observe a typical particle rms speed $\delta
v\sim0.025 c_{\rm s}$ ."
 The mass density represented by a single superparticle is Pp50.2190. and the friction time is Qr;=0.3 (we ormalise here by the Keplerian frequency 2 which we define in 4)., The mass density represented by a single superparticle is $\hat{\rho}_{\rm p} \approx 0.219 \rho_{\rm g}$ and the friction time is $\varOmega \tau_{\rm f}=0.3$ (we normalise here by the Keplerian frequency $\varOmega$ which we define in ).
" This gives Or,=18 from(A5).."," This gives $\varOmega
\tau_{\rm c} \approx 18$ from."
 The Courant eriterion for the hydrodynamical part of the streaming instability gives the time-step Ofjyure=0.000625-! for 647 and Offydre=0.000312527! for 128? simulations.," The Courant criterion for the hydrodynamical part of the streaming instability gives the time-step $\delta t_{\rm hydro}=0.000625\varOmega^{-1}$ for $64^3$ and $\delta t_{\rm
hydro}=0.0003125\varOmega^{-1}$ for $128^3$ simulations."
" Therefore we can Ignore the collision time-scale in the simulations when ""Setermining the numerical time-step.", Therefore we can ignore the collision time-scale in the simulations when determining the numerical time-step.
 defines the collisional time-scale between particles of two sizes., defines the collisional time-scale between particles of two sizes.
" For two superparticles of equal internal particle number (A) we have D=τι, because the cross section ση and relative speed ovj; are symmetric in (7.4&)."," For two superparticles of equal internal particle number $\hat{n}$ ) we have $\tau_{\rm c}^{(k)}=\tau_{\rm c}^{(j)}$, because the cross section $\sigma_{jk}$ and relative speed $\delta v_{jk}$ are symmetric in $(j,k)$."
 However. equal particle number per superparticle is numerically expensive. as the mass of a superparticle in that case scales as A. requiring many more superparticles to represent an equal mass of smaller particles.," However, equal particle number per superparticle is numerically expensive, as the mass of a superparticle in that case scales as $R^3$, requiring many more superparticles to represent an equal mass of smaller particles."
 The second complication is that the collision scale becomes very short for smaller particles., The second complication is that the collision time-scale becomes very short for smaller particles.
 A more common approach is to have equal mass per superparticle., A more common approach is to have equal mass per superparticle.
 In that case we can define a collision time-scale as the time for all mass in particle / to interact with all mass in particle &., In that case we can define a collision time-scale as the time for all mass in particle $j$ to interact with all mass in particle $k$.
" This time-scale is shared between the two particle species and is given by To illustrate this. take small particles of friction time ""- and large particles of friction time D=100."," This time-scale is shared between the two particle species and is given by To illustrate this, take small particles of friction time $\tau_{\rm
f}^{(j)}=1$ and large particles of friction time $\tau_{\rm f}^{(k)}=100$."
 The collision time-scale for the large particles is 100 times shorter than for the small particles. because the superparticle with small physical particles contains 100 times more particles in the swarm.," The collision time-scale for the large particles is 100 times shorter than for the small particles, because the superparticle with small physical particles contains 100 times more particles in the swarm."
 However. the time-scale for collision between a large and a small particle does not imply that all small particles have collided during that time.," However, the time-scale for collision between a large and a small particle does not imply that all small particles have collided during that time."
 The correct time-scale is the time for small particles to collide with large particles., The correct time-scale is the time for small particles to collide with large particles.
 When an average small particle has experienced a collision. then all small particles have collided with a large particle. and all the mass in the two superparticles have interacted.," When an average small particle has experienced a collision, then all small particles have collided with a large particle, and all the mass in the two superparticles have interacted."
 After waiting the common collision time τς. the collision outcome can be solved as if the two colliding particles had equal mass. since effectively a large particle collides with my/nij small particles during this tme.," After waiting the common collision time $\tau_{\rm c}$, the collision outcome can be solved as if the two colliding particles had equal mass, since effectively a large particle collides with $m_k/m_j$ small particles during this time."
 This approach is slightly inconsistent since discrete collisions with N particles of mass a; is not equal to a single collision with a particle of mass Λι., This approach is slightly inconsistent since discrete collisions with $N$ particles of mass $m_j$ is not equal to a single collision with a particle of mass $N m_j$.
 A collision between a particle of velocity v; and a stationary particle results in the new velocity After N such collisions the velocity of particle & is In the limit where v;—=Avi«vy. this equation describes à velocity damping with characteristic time-scale tyΞτος4)/(211;).," A collision between a particle of velocity $v_k$ and a stationary particle results in the new velocity After $N$ such collisions the velocity of particle $k$ is In the limit where $v_k - v_k' = \Delta v_k \ll v_k$, this equation describes a velocity damping with characteristic time-scale $\tau_{\rm d}=\tau_{\rm c} (m_k+m_j)/(2 m_j)$."
 Completely braking down the large particle requires infinite time. whereas a single diserete collision with an equal-mass particle would remove all the momentum from particle j in one collision time.," Completely braking down the large particle requires infinite time, whereas a single discrete collision with an equal-mass particle would remove all the momentum from particle $j$ in one collision time."
 To really get the collisional energy equipartition right between particles of different sizes one would have to allow for collisions between a large particle and individual smaller particles., To really get the collisional energy equipartition right between particles of different sizes one would have to allow for collisions between a large particle and individual smaller particles.
 This could either be done by letting superparticles not represent the same mass. but rather the same number of particles.," This could either be done by letting superparticles not represent the same mass, but rather the same number of particles."
 However. this approach would become unpractical to model a large span in particle sizes. since a huge number of superparticles would be required to represent the low-mass particle bins.," However, this approach would become unpractical to model a large span in particle sizes, since a huge number of superparticles would be required to represent the low-mass particle bins."
 Alternatively the collision between a swarm of large and small particles could be modelled on the collistor time-scale of individual collisions. distributing afterwards the energy and momentum of the particle that suffered the collisior among the entire swarm of small particles or among all particle swarms within its mean free path.," Alternatively the collision between a swarm of large and small particles could be modelled on the collision time-scale of individual collisions, distributing afterwards the energy and momentum of the particle that suffered the collision among the entire swarm of small particles or among all particle swarms within its mean free path."
 However. for a large spar in particle sizes. this would still require a very small time-step and is therefore unpractical.," However, for a large span in particle sizes, this would still require a very small time-step and is therefore unpractical."
" We simply note here that while collisions between unequal-sized particles can be modellec with the right conservation properties. actual equipartition of particle energies would require an adaptation of the collistor algorithm,"," We simply note here that while collisions between unequal-sized particles can be modelled with the right conservation properties, actual equipartition of particle energies would require an adaptation of the collision algorithm."
 During the gravitational contraction of particle clumps the number of particles in a grid cell can become very large. on the order of 1000s of 10000s.," During the gravitational contraction of particle clumps the number of particles in a grid cell can become very large, on the order of 1000s of 10000s."
 Tracking (1/2)NCN—1) possible collisions per grid cell then becomes very computationally expensive., Tracking $(1/2)N(N-1)$ possible collisions per grid cell then becomes very computationally expensive.
 However. particles do not collide with all possible partners during a single time-step.," However, particles do not collide with all possible partners during a single time-step."
 One can limit the number of collision partners. while maintaining the overall collision rate. by sampling only a subset Muay of the possible collisions.," One can limit the number of collision partners, while maintaining the overall collision rate, by sampling only a subset $N_{\rm max}$ of the possible collisions."
9 close to 907. while the model predicts dominance of the linearly polarized 7 components for 9 near 90°.,"$\theta$ close to $90^\circ$, while the model predicts dominance of the linearly polarized $\pi$ –components for $\theta$ near $90^\circ$."
 Our identification of the linearly polarized spots as σ components was based on the presence of some circularly. polarized emission in all spots except spot E. This seenis to be a very firm argument in favor of σ components unless an explanation lor the circular polarization contribution to {ποπ components is found., Our identification of the linearly polarized spots as $\sigma$ –components was based on the presence of some circularly polarized emission in all spots except spot E. This seems to be a very firm argument in favor of $\sigma$ –components unless an explanation for the circular polarization contribution to the $\pi$ –components is found.
 There is an alternative interpretation of the observed linear polarization as a result of (he saturated maser amplification with a weak magnetic field. when the Zeeman splitting is less than the bandwidth.," There is an alternative interpretation of the observed linear polarization as a result of the saturated maser amplification with a weak magnetic field, when the Zeeman splitting is less than the bandwidth."
 For sin?9€1/3. or €.<35.5° there will be a single LOO per cent linearly polarized emission line with E-vector parallel to the magnetic field direction (Case 2a in Goldreich.IxXeeley.andKwan (1973))).," For $\sin^2\theta\le1/3$, or $\theta\le35.5^\circ$ there will be a single 100 per cent linearly polarized emission line with E-vector parallel to the magnetic field direction (Case 2a in \citet*{goldreich73}) )."
 For 6>35.57 the linear polarization will be lower. and the circular polarization appears in the line wings (see Section 3.7).," For $\theta>35.5^\circ$ the linear polarization will be lower, and the circular polarization appears in the line wings (see Section 3.7)."
 This model could explain observed linearly polarized emission in W75N but the observed admixture of circular polarization remains unaccounted. (the same as for the x. components in (he Zeeman splitting.," This model could explain observed linearly polarized emission in W75N but the observed admixture of circular polarization remains unaccounted, the same as for the $\pi$ –components in the Zeeman splitting."
 Also this model limits the magnetic field to about 0.5 milligauss which is much less than was deduced for spots A. D. F. Fl and M. O in Section 3.6.," Also this model limits the magnetic field to about 0.5 milligauss which is much less than was deduced for spots A, B, F, F1 and M, O in Section 3.6."
 More likely this explanation could be applied to the observed linear polarization of WoO and methanol masers., More likely this explanation could be applied to the observed linear polarization of $_2$ O and methanol masers.
 And. finally. the detection of the Zeeman counterpart Fl the spot F. if confirmed. is a definite proof οἱ the interpretation of the linearly polarized. emission features as σ components of Zeeman pattern.," And, finally, the detection of the Zeeman counterpart F1 the spot F, if confirmed, is a definite proof of the interpretation of the linearly polarized emission features as $\sigma$ –components of Zeeman pattern."
 We have mapped 14 maser spots in W75N (Fig., We have mapped 14 maser spots in W75N (Fig.
 5). wilh two or possibly three coinciding Zeeman components.," 5), with two or possibly three coinciding Zeeman components."
 Every maser spot has its own radial velocity. different from the radial velocity of other spots. and a narrow line width of about 0.2 0.3 |.," Every maser spot has its own radial velocity, different from the radial velocity of other spots, and a narrow line width of about 0.2 – 0.3 $^{-1}$."
 1667-MlIZ spots do 10 coincide in position with 1665-MllIz spots., 1667-MHz spots do not coincide in position with 1665-MHz spots.
 Almost every spot is 100 per cent ellipticallv polarized., Almost every spot is 100 per cent elliptically polarized.
 We interpret them as aq components of the Zeeman pattern., We interpret them as $\sigma$ – components of the Zeeman pattern.
 The magnetic field in the Zeeman pairs was determined to be from 5 (to 11 imilligauss. the same values as found in other OII masers (Garcia-Darretoοἱal.1933).," The magnetic field in the Zeeman pairs was determined to be from 5 to 11 milligauss, the same values as found in other OH masers \citep{garcia88}."
. One can assume that in other spots where the second σ component is not visible the magnetic field strength is of the same order., One can assume that in other spots where the second $\sigma$ –component is not visible the magnetic field strength is of the same order.
 If it is ac component. (hen the linear polarization vector must be perpendicular to the direction ol the magnetic field. and (he magnetic field direction seems to be roughly the same for all spots.," If it is a $\sigma$ –component, then the linear polarization vector must be perpendicular to the direction of the magnetic field, and the magnetic field direction seems to be roughly the same for all spots."
 This can be seen on the histogram in Fie 6., This can be seen on the histogram in Fig 6.
 Unfortunately the direction of (he magnetic field can not be determined reliably [rom (he linear polarization measurements because of ihe uncertain amount of Faraday rotation in the Galaxy., Unfortunately the direction of the magnetic field can not be determined reliably from the linear polarization measurements because of the uncertain amount of Faraday rotation in the Galaxy.
 The OIL masers are located in the Galactie disk where the Faraday rotation is al maximum., The OH masers are located in the Galactic disk where the Faraday rotation is at maximum.
 Total Faraday rotation in the, Total Faraday rotation in the
"emission all the way to the extremities of our field of view (24x 24), but suspect that there is still diffuse emission beyond.","emission all the way to the extremities of our field of view $24\times24'$ ), but suspect that there is still diffuse emission beyond."
 Other pointings and/or a larger field of view would be necessary to study the total extent of ionized gas in NGC 300., Other pointings and/or a larger field of view would be necessary to study the total extent of ionized gas in NGC 300.
" In the first section, we briefly discuss the results concerning this galaxy."," In the first section, we briefly discuss the results concerning this galaxy."
" We then address in greater detail the case of NGC 247, where we were able to detect Ho emission on the entirety of the disc."," We then address in greater detail the case of NGC 247, where we were able to detect $\alpha$ emission on the entirety of the disc."
" NGC 300 is a late-type spiral galaxy with a very rich distribution of bright regions(?),, and a very extended stellar disc reaching r—24' that shows no sign of truncation(?)."," NGC 300 is a late-type spiral galaxy with a very rich distribution of bright regions, and a very extended stellar disc reaching $r=24'$ that shows no sign of truncation."
". The faintest level of the stellar disc reaches a surface brightness of 32 mag arcsec” (in V and B band), and the stellar population was found to be mostly old(?)."," The faintest level of the stellar disc reaches a surface brightness of 32 mag $^{-2}$ (in $V$ and $B$ band), and the stellar population was found to be mostly old."
". The disc extends to r~20' (?),, and seems to be flat beyond r—15' (8 kpc).", The disc extends to $r\sim20'$ and seems to be flat beyond $r=15'$ (8 kpc).
 We have obtained very deep Ha observations of this galaxy., We have obtained very deep $\alpha$ observations of this galaxy.
" We reach EMs on the order of 0.1 cm? pc, and our rotation curve extends up to r—13.5'."," We reach EMs on the order of $0.1$ $^{-6}$ pc, and our rotation curve extends up to $r=13.5'$."
 These are to our knowledge the deepest Ho observations of this galaxy., These are to our knowledge the deepest $\alpha$ observations of this galaxy.
 Other Ho studies include that of and?., Other $\alpha$ studies include that of and.
". The first obtained a rotation curve extending only to r—10', and the second reached EMs between 6 and 500 cm pc."," The first obtained a rotation curve extending only to $r=10'$, and the second reached EMs between 6 and 500 $^{-6}$ pc."
 We detect significantly more diffuse emission., We detect significantly more diffuse emission.
 Fig., Fig.
 4 Ó(bottom-left panel) shows that the bright regions are embedded in a large diffuse background., \ref{a2f4} (bottom-left panel) shows that the bright regions are embedded in a large diffuse background.
 estimated that the diffuse fraction of Ha gas after scattering correction lies between 44 to 54 per cent of the total Ha emission., estimated that the diffuse fraction of $\alpha$ gas after scattering correction lies between 44 to 54 per cent of the total $\alpha$ emission.
" Our study, which finds more diffuse gas, suggests that this fraction could be significantly larger."," Our study, which finds more diffuse gas, suggests that this fraction could be significantly larger."
" Since we were limited by our field of view,we suspect that there is more diffuse gas beyond our maximum radius (r— 13.5)."," Since we were limited by our field of view,we suspect that there is more diffuse gas beyond our maximum radius $r=13.5'$ )."
2010).,.
. We investigate this cascade emission effect in two nearby starburst galaxies. M82 and NGC 253.," We investigate this cascade emission effect in two nearby starburst galaxies, M82 and NGC 253."
 Very recently two nearby bright starburst galaxies. M82 and NGC 253. have been detected with >-ray observations byFermi. H.E.S.S.. and VERITAS (Abdoetal.2010a:Acero 2009).," Very recently two nearby bright starburst galaxies, M82 and NGC 253, have been detected with $\gamma$ -ray observations by, H.E.S.S., and VERITAS \citep{abd10_sb,ace09,ver09}."
. These -- ray detected starburst galaxies would not suffer from >-ray absorption with the intergalactic radiation field. since they are very close to us.," These $\gamma$ -ray detected starburst galaxies would not suffer from $\gamma$ -ray absorption with the intergalactic radiation field, since they are very close to us."
 The distance to M82 and NGC 253 is 3.630.3 Mpe (Freedmanetal.1994) and 3.9c0.4 Mpe (Karachentsevetal.2003).. respectively.," The distance to M82 and NGC 253 is $3.6\pm0.3$ Mpc \citep{fre94} and $3.9\pm0.4$ Mpc \citep{kar03}, respectively."
 The starburst region is located in the inner part of the galaxy with a size of 500 pe and 280 pe for M82 and NGC 253. respectively (Mayyaetal.2006;Ulvestad 2000).," The starburst region is located in the inner part of the galaxy with a size of 500 pc and 280 pc for M82 and NGC 253, respectively \citep{may06,ulv00}."
. The central starburst region 1n these galaxies where SNRs are produced is theoretically expected to emit most of the 5-ray photons 2005. hereinafter DSTOS: 2009. hereinafter dCdPO09).," The central starburst region in these galaxies where SNRs are produced is theoretically expected to emit most of the $\gamma$ -ray photons 2005, hereinafter DST05; 2009, hereinafter dCdP09)."
 In the case of the Large Magellanic Cloud. the association between starburst region and 5-ray emission region Is observationally confirmed by (Abdo 2010c).," In the case of the Large Magellanic Cloud, the association between starburst region and $\gamma$ -ray emission region is observationally confirmed by \citep{abd10_lmc}."
. It is also known that there are plenty of interstellar radiation field photons from star forming activities., It is also known that there are plenty of interstellar radiation field photons from star forming activities.
 Therefore. we need to take into account the >-ray absorption effect in the central star formation region.," Therefore, we need to take into account the $\gamma$ -ray absorption effect in the central star formation region."
 Siebenmorgen&Kriigel(2007.hereinafterSK07) developed optical-infrared spectral energy distribution (SED) models for the nuclei of starburst galaxies and ultra luminous infrared galaxies by computing a radiative transfer SED of spherical. dusty galactic nuclei.," \citet[][hereinafter SK07]{sie07} developed optical-infrared spectral energy distribution (SED) models for the nuclei of starburst galaxies and ultra luminous infrared galaxies by computing a radiative transfer SED of spherical, dusty galactic nuclei."
 They include silicates. amorphous carbon. graphite grains. and polyeyelic aromatic hydrocarbons 1n the form of dust.," They include silicates, amorphous carbon, graphite grains, and polycyclic aromatic hydrocarbons in the form of dust."
 Fig., Fig.
 1. shows the optical-infrared spectral energy distribution (SED) of M82 and NGC 253., \ref{fig:SED} shows the optical-infrared spectral energy distribution (SED) of M82 and NGC 253.
 We adopt their spectrum model for M82 and NGC 253 shown in Fig., We adopt their spectrum model for M82 and NGC 253 shown in Fig.
 2 and 3 in their paper. which i5 consistent with the data.," 2 and 3 in their paper, which is consistent with the data."
 In the case of M82. we add a blackbody component (T=2500 K) to fit the data below 5 jm as in Siebenmorgen (2007).. which was not taken into account in previous works.," In the case of M82, we add a blackbody component $T=2500$ K) to fit the data below 5 $\mu {\rm m}$ as in \citet{sie07}, which was not taken into account in previous works."
 For comparison. we also show the interstellar SED model in dCdP09 and DSTOS for M82 and NGC 253. respectively.," For comparison, we also show the interstellar SED model in dCdP09 and DST05 for M82 and NGC 253, respectively."
 Since SKO7 models well reproduce the observed data over a wide range of wavelengths for both galaxies. hereafter we use the SKO7 optical-infrared interstellar SED models as our baseline nuclei SED models.," Since SK07 models well reproduce the observed data over a wide range of wavelengths for both galaxies, hereafter we use the SK07 optical-infrared interstellar SED models as our baseline nuclei SED models."
 Fig., Fig.
 2. shows the expected optical depth. 7... of M82 for SK07 and dCdP09 models and NGC 253 for SKO7 and DSTOS models.," \ref{fig:tau} shows the expected optical depth, $\tau_{\gamma\gamma}$, of M82 for SK07 and dCdP09 models and NGC 253 for SK07 and DST05 models."
 We do not take into account the interaction of a 5-ray with a nuclei. since this would not absorb +- significantly (Domingo-Santam," We do not take into account the interaction of a $\gamma$ -ray with a nuclei, since this would not absorb $\gamma$ -rays significantly \citep{ds05}."
aría&Torres2005).. The ~—~ absorption becomes significant (Le. 7..= 1) above ~5 TeV and ~9 TeV for M82 and NGC 253. respectively.," The $\gamma-\gamma$ absorption becomes significant (i.e. $\tau_{\gamma\gamma}=1$ ) above $\sim 5$ TeV and $\sim9$ TeV for M82 and NGC 253, respectively."
 The >-ray emissions from starburst galaxies are modeled by many papers (Voelketal.1989;Akyuz199];CeadelPozoetal.2009:Rephaeli 2010).," The $\gamma$ -ray emissions from starburst galaxies are modeled by many papers \citep{voe89,aky91,pag96,tor04,ds05,tho07,per08,dcdp09,rep10}."
. In this paper. we substitute the zo-meson decay spectrum for the intrinsic 2-ray spectrum. dN;/dE-. using code provided by Karlsson&Kamae(2008) (seealsoKamaeetal.2005.2006).," In this paper, we substitute the $\pi_0$ -meson decay spectrum for the intrinsic $\gamma$ -ray spectrum, $dN_i/dE_\gamma$, using code provided by \citet{kar08} \citep[see also][]{kam05,kam06}."
. For the total inclusive inelastic cross section. they include the non-diffractive (with Feymann scaling violation) and diffractive components. plus the AC 1232) and Res(1600) resonance excitation contributions.," For the total inclusive inelastic cross section, they include the non-diffractive (with Feymann scaling violation) and diffractive components, plus the $\Delta$ (1232) and Res(1600) resonance excitation contributions."
 The normalization of the total x-ray flux including absorption and cascade effects is adjusted to the observed photon flux data in the 0.1—5 GeV band., The normalization of the total $\gamma$ -ray flux including absorption and cascade effects is adjusted to the observed photon flux data in the 0.1--5 GeV band.
" The intrinsic. 5-ray spectrum is given by the highest proton kinetic energy. 7,,,4,. in units of [TeV] and the spectral index of the intrinsic proton spectrum. TL. where we set dN,/d,xPE dN,/d, the proton flux. >, the Lorentz factor of protons."," The intrinsic $\gamma$ -ray spectrum is given by the highest proton kinetic energy, $T_{p, \rm max}$, in units of [TeV] and the spectral index of the intrinsic proton spectrum, $\Gamma$, where we set $dN_p/d\gamma_p\propto \gamma_p^{-\Gamma}$, $dN_p/d\gamma_p$ the proton flux, $\gamma_p$ the Lorentz factor of protons."
 We do not take into account inverse Compton scattering and bremsstrahlung radiation. since their contribution is expected to be less than 10 above 0.1 GeV (DSTOS. dCdP09).," We do not take into account inverse Compton scattering and bremsstrahlung radiation, since their contribution is expected to be less than 10 above 0.1 GeV (DST05, dCdP09)."
 Although we need to solve the cosmic-ray propagation in starburst galaxies for detailed calculations. we adopt a simple analytical model to see the absorption and cascade effects.," Although we need to solve the cosmic-ray propagation in starburst galaxies for detailed calculations, we adopt a simple analytical model to see the absorption and cascade effects."
 For the calculation. of the cascade emission. we ignore magnetic fields. since the >-ray emissions from starburst galaxies are isotropic and continuous.," For the calculation of the cascade emission, we ignore magnetic fields, since the $\gamma$ -ray emissions from starburst galaxies are isotropic and continuous."
 We calculate the resulting cascade emissions as follows: (Blumenthal&Gould1970:Aharonianetal.1994;WangTotani2009:Venters 2010).. where fjc is the inverse Compton (IC) cooling time scale and the electron injection spectrum is and the scattered photon spectrum per unit time by IC scattering Is: with f(y)22vInQyyt+a41—ον. (Qcx« D) and y=EB.| ," We calculate the resulting cascade emissions as follows: \citep{blu70,aha94,wan01,dai02,raz04,and04,mur07,kne08,ino09,ven10}, where $t_{IC}$ is the inverse Compton (IC) cooling time scale and the electron injection spectrum is and the scattered photon spectrum per unit time by IC scattering is: with $f(x)=2x \ln(x)+x+1-2x^2$, $0<x<1$ ) and $x = E_{\gamma,i}/4\gamma_e^2\epsilon$ ."
"Here. E.;=2μμ is the energy of intrinsic photons. dere.4N;/dE. is the intrinsic x-ray spectrum. 7, the classical electron. radius. d/de(e) is the photon density at the starburst region."," Here, $E_{\gamma,i}=2\gamma_e m_e c^2$ is the energy of intrinsic photons, ${dN_i}/{dE_{\gamma,i}}$ is the intrinsic $\gamma$ -ray spectrum, $r_0$ the classical electron radius, $dn/d\epsilon(\epsilon)$ is the photon density at the starburst region."
" The integration region over the Lorentz factor. το. is ..min*5eπολ. 5eInax-E.INAS2 and ecanin=ΠΙαΧΠΟ/2e.(CE.fey -/2). where E.4, is the maximum energy of 5-ray photons."," The integration region over the Lorentz factor, $\gamma_e$, is $\gamma_{e,{\rm min}}<\gamma_e<\gamma_{e,{\rm max}}$, $\gamma_{e,{\rm max}}=E_{\gamma,{\rm max}}/2$ and $\gamma_{e,{\rm min}}={\rm max}[m_ec^2/2\epsilon,(E_\gamma/\epsilon)^{1/2}/2]$ , where $E_{\gamma,{\rm max}}$ is the maximum energy of $\gamma$ -ray photons."
 We iteratively calculate Eq., We iteratively calculate Eq.
 1. by substituting dN./dE-(1—-c )inorder to include IC scattering due to generated pairs from reabsorbed secondary photons.," \ref{eq:cascade} by substituting $dN_c/dE_\gamma (1-e^{-\tau_{\gamma,\gamma}})$ in order to include IC scattering due to generated pairs from reabsorbed secondary photons."
 Fig., Fig.
" 3. shows the >-ray spectra of M82 and NGC 253 in units of |MeV/em""/s].", \ref{fig:f_sb} shows the $\gamma$ -ray spectra of M82 and NGC 253 in units of $^2$ /s].
 We show the intrinsic (the spectrum without taking into account absorption). absorbed. and cascade components as well as total (absorbed + cascade) spectra.," We show the intrinsic (the spectrum without taking into account absorption), absorbed, and cascade components as well as total (absorbed + cascade) spectra."
 For the intrinsic spectra. we adopt 1.Των)= and (2.3.512TeV) for M82 and NGC 253 as our standard models. respectively.," For the intrinsic spectra, we adopt $\Gamma, T_{p, \rm max}) = (2.4, 512\ {\rm TeV})$ and $(2.3, 512 \ {\rm TeV})$ for M82 and NGC 253 as our standard models, respectively."
 Inthe case of M82. our >- spectrum nicely fits to the observed data in the range of the uncertainty of the data by taking into account the —5 absorption and the cascade emissions.," Inthe case of M82, our $\gamma$ -ray spectrum nicely fits to the observed data in the range of the uncertainty of the data by taking into account the $\gamma-\gamma$ absorption and the cascade emissions."
 However. in the case," However, in the case"
"range, one might become suspicious about the reality of the effect.","range, one might become suspicious about the reality of the effect."
" A direct test is plotting the quasar fraction as function of optical (R-band in this case) magnitude, instead of radio flux density as in Fig. 7.."," A direct test is plotting the quasar fraction as function of optical (R-band in this case) magnitude, instead of radio flux density as in Fig. \ref{morphRF}."
" This has been done in Fig. 8,,"," This has been done in Fig. \ref{morphMAG},"
 with again our variable sample in green and the control sample in red., with again our variable sample in green and the control sample in red.
 It is reassuring to see that most of the excess quasar identifications in the variable sample are made for magnitudes brighter than R~21., It is reassuring to see that most of the excess quasar identifications in the variable sample are made for magnitudes brighter than $R \sim 21$.
 Stoughton et al. (, Stoughton et al. (
2002) estimate a confidence level for classifications up to this magnitude.,2002) estimate a confidence level for classifications up to this magnitude.
 It becomes much less certain for fainter sources., It becomes much less certain for fainter sources.
" Without doubt some of the current classifications are wrong, but there is no indication that our consistently higher quasar fraction in the variable sample, both as function of radio flux density and optical magnitude, is significantly affected by this."," Without doubt some of the current classifications are wrong, but there is no indication that our consistently higher quasar fraction in the variable sample, both as function of radio flux density and optical magnitude, is significantly affected by this."
 We have established that there is a significant difference between the variable and non-variable radio samples in terms of their optical counterparts., We have established that there is a significant difference between the variable and non-variable radio samples in terms of their optical counterparts.
" The former have a higher matching rate with optical identifications (given a fixed detection limit), which is at least in part ascribed to the higher fraction of quasars."," The former have a higher matching rate with optical identifications (given a fixed detection limit), which is at least in part ascribed to the higher fraction of quasars."
" We have already seen that the magnitude distributions are not significantly different (Fig. 6)),"," We have already seen that the magnitude distributions are not significantly different (Fig. \ref{apmids}) ),"
" however, this does not rule out that significant optical color differences may be present between the samples."," however, this does not rule out that significant optical color differences may be present between the samples."
 The clearest difference between galaxies and quasars (at least in the restframe) is provided by the (u—r) vs (r—z) color plot (see Fig 13 in Stoughton et al., The clearest difference between galaxies and quasars (at least in the restframe) is provided by the $-$ r) vs $-$ z) color plot (see Fig 13 in Stoughton et al.
 2002 for intermediate color-color plots)., 2002 for intermediate color-color plots).
 This allows for sampling of the bbreak in galaxies up to a redshift of about unity., This allows for sampling of the break in galaxies up to a redshift of about unity.
" This break will offset the galaxies from the quasars, which usually exhibit fairly flat, powerlaw continua (see Fig."," This break will offset the galaxies from the quasars, which usually exhibit fairly flat, powerlaw continua (see Fig."
 3 in de Vries et al., 3 in de Vries et al.
 2003)., 2003).
 Our data are presented in Fig. 9.., Our data are presented in Fig. \ref{sdssCol}.
 The sample means and their lo errors (indicated by the crosses in Fig. 9)), The sample means and their $\sigma$ errors (indicated by the crosses in Fig. \ref{sdssCol}) )
 are presented in Table 6.., are presented in Table \ref{sdsscoltable}.
 Both the variable quasar and galaxy distributions in the color-color plane are (statistically) distinct from the non-variable control sample., Both the variable quasar and galaxy distributions in the color-color plane are (statistically) distinct from the non-variable control sample.
" This trend for both the galaxies and quasars to become bluer when they vary has been seen before (e.g., Giveon et al."," This trend for both the galaxies and quasars to become bluer when they vary has been seen before (e.g., Giveon et al."
 1999; Trevese Vagnetti 2002; Vagnetti et al., 1999; Trevese Vagnetti 2002; Vagnetti et al.
 2003; de Vries et al., 2003; de Vries et al.
 2003)., 2003).
" Still, the color spread among the sources in the sample is much larger than theéndividual color change as a source goes through an outburst."," Still, the color spread among the sources in the sample is much larger than the color change as a source goes through an outburst."
" Indeed, the studies of Giveon, Trevese, Vagnetti and their respective collaborators all rely on following individual objects through their outbursts, instead of using a statistical sample."," Indeed, the studies of Giveon, Trevese, Vagnetti and their respective collaborators all rely on following individual objects through their outbursts, instead of using a statistical sample."
" 'The presence of the mean optical color offsets in our radio variable sample implies two things: 1) on average, radio variable sources also exhibit optical (color) variations, and 2) the optical variation time-scales are shorter than our 7 year base-line, since the SDSS data were not taken at the epoch of radio outburst."," The presence of the mean optical color offsets in our radio variable sample implies two things: 1) on average, radio variable sources also exhibit optical (color) variations, and 2) the optical variation time-scales are shorter than our 7 year base-line, since the SDSS data were not taken at the epoch of radio outburst."
" However, optical colors for an individual source, while useful in differentiating between quasars and galaxies, do not provide an effective selection mechanism for (radio) variability."," However, optical colors for an individual source, while useful in differentiating between quasars and galaxies, do not provide an effective selection mechanism for (radio) variability."
" Nonetheless, as we have seen in the previous sections, the presence of radio"," Nonetheless, as we have seen in the previous sections, the presence of radio"
The disk of the Milkv. Was is warped like an integral sien. rising above the plane on one side axl falling below the plane on the other.,"The disk of the Milky Way is warped like an integral sign, rising above the plane on one side and falling below the plane on the other."
 This warp is seen both in maps of neutral hvdrogen (e.g..Diplas&Savage1991). and in the stellar distribution (Reed 1996: Drimumnel. Smart. Lattanzi 2000: Lóppez-Corredoira et al.," This warp is seen both in maps of neutral hydrogen \citep[e.g.,][]{diplas and savage91} and in the stellar distribution (Reed 1996; Drimmel, Smart, Lattanzi 2000; Lóppez-Corredoira et al."
 2002b)., 2002b).
 The Sun lies along the line of nodes of tje warp. where Ullec outer rings cross the inner plane (see Figure 1)).," The Sun lies along the line of nodes of the warp, where tilted outer rings cross the inner plane (see Figure \ref{schematic}) )."
 Despite (heir tendency to disperse when isolated (IIunter&Toomre1969).. warps are common 1 external galaxies (Dosma 1981: Driggs 1990: Christodoulou. Tohline. Cameron 1993: Reshetnikov Combes 1993).," Despite their tendency to disperse when isolated \citep{hunter and toomre69}, warps are common in external galaxies (Bosma 1981; Briggs 1990; Christodoulou, Tohline, Steiman-Cameron 1993; Reshetnikov Combes 1998)."
 This has driven many authors to search [or, This has driven many authors to search for
"the hot shell and, therefore, absolute values of the generalized spectral functions at frequencies of 90 and 217 GHz should be higher for the internal region in the corresponding temperature range (see Figs. [I]. D], ","the hot shell and, therefore, absolute values of the generalized spectral functions at frequencies of 90 and 217 GHz should be higher for the internal region in the corresponding temperature range (see Figs. \ref{Fig90}, , \ref{Fig217}, ,"
and[7))., and \ref{Tsim}) ).
 We find that the pressure values are higher inside the internal region in comparison with the thermal hot shell pressure., We find that the pressure values are higher inside the internal region in comparison with the thermal hot shell pressure.
" Therefore, the SZ effect in the cocoon internal region should be more significant than that produced by electrons located inside the hot shell."," Therefore, the SZ effect in the cocoon internal region should be more significant than that produced by electrons located inside the hot shell."
 The pressure inside the simulation cocoon is two orders of magnitude higher than that found in the diffuse ISM., The pressure inside the simulation cocoon is two orders of magnitude higher than that found in the diffuse ISM.
" Such a pressure-jump corresponds to the Mach number of a shock equaled M~V4P2/5P,10 (Landau Lifshitz 1959).", Such a pressure-jump corresponds to the Mach number of a shock equaled $M\approx\sqrt{4 P_2/5 P_1}\sim 10$ (Landau Lifshitz 1959).
 The plasma inside the AGN cocoon is largely over-pressured relative to the ambient ISM., The plasma inside the AGN cocoon is largely over-pressured relative to the ambient ISM.
 The radio lobes of the nearby radio galaxy Centaurus A is expanding into the ISM at the Mach number of 8.5 (Kraft et al., The radio lobes of the nearby radio galaxy Centaurus A is expanding into the ISM at the Mach number of 8.5 (Kraft et al.
 2003) is representative of such over-pressured plasmas., 2003) is representative of such over-pressured plasmas.
 To produce the 3D pressure and temperature maps which are necessary to derive the intensity map of the SZ effect we rotate the 2D pressure and temperature maps along the jet axis., To produce the 3D pressure and temperature maps which are necessary to derive the intensity map of the SZ effect we rotate the 2D pressure and temperature maps along the jet axis.
 We calculated the SZ effect using the values of the generalized spectral functions found in Sect., We calculated the SZ effect using the values of the generalized spectral functions found in Sect.
 2., 2.
" Since the gas temperatures in a cocoon are high, such a cocoon should be a source of the SZ effect at a frequency 217 GHz."," Since the gas temperatures in a cocoon are high, such a cocoon should be a source of the SZ effect at a frequency 217 GHz."
 The intensity map of the SZ effect at a frequency 217 GHz derived from the simulation maps of the gas pressure and temperature is plotted in Fig. 9}. , The intensity map of the SZ effect at a frequency 217 GHz derived from the simulation maps of the gas pressure and temperature is plotted in Fig. \ref{sz-sim}. .
A comparison of Figs., A comparison of Figs.
 [7] and [8] with Fig., \ref{Tsim} and \ref{Psim} with Fig.
 B] strongly suggests that the dominant contribution to the SZ effect at a frequency 217 GHz originates from the cocoon internal region., \ref{sz-sim} strongly suggests that the dominant contribution to the SZ effect at a frequency 217 GHz originates from the cocoon internal region.
" Therefore, measurements of the SZ effect at a frequency 217 GHz are a powerful tool for potentially revealing the dynamically dominant component inside AGN cocoons (see Fig. [8))."," Therefore, measurements of the SZ effect at a frequency 217 GHz are a powerful tool for potentially revealing the dynamically dominant component inside AGN cocoons (see Fig. \ref{Psim}) )."
" The proposed method based on the generalized spectral function is a way of maximizing the contribution from the gas with a temperature in the range 10? K «T,10? K (see Fig. D)).", The proposed method based on the generalized spectral function is a way of maximizing the contribution from the gas with a temperature in the range $10^9$ K $< T_{\mathrm{e}}< 10^{10}$ K (see Fig. \ref{Fig217}) ).
 The study of the SZ effect in a plasma at such temperatures is important since numerical simulations predict that such a plasma represents the still invisible dynamically dominant component inside AGN cocoons., The study of the SZ effect in a plasma at such temperatures is important since numerical simulations predict that such a plasma represents the still invisible dynamically dominant component inside AGN cocoons.
" Our simulations are characterized by ISM density and temperature at 1 cm? and 10’ K (ISM pressure at ~10? erg/cm?), respectively, typical of the ISM in the central parts of an elliptical at high redshift (z~1)."," Our simulations are characterized by ISM density and temperature at 1 $^{-3}$ and $10^7$ K (ISM pressure at $\sim 10^{-9}$ $^3$ ), respectively, typical of the ISM in the central parts of an elliptical at high redshift $z\approx1$ )."
 We found that the plasma inside the AGN cocoon is largely over-pressured relative to the ambient ISM., We found that the plasma inside the AGN cocoon is largely over-pressured relative to the ambient ISM.
 The derived intensity AZ of the SZ effect in this case is three orders of magnitude higher than that of cocoons in galaxy clusters., The derived intensity $\Delta I$ of the SZ effect in this case is three orders of magnitude higher than that of cocoons in galaxy clusters.
" Though the linear scales of cocoons are close in both cases, the angular diameter of the cocoon within the distant elliptical galaxy equaled to 6,2;=Lpox/Da~2.5"" is much smaller than that of the cocoon with the same linear size inside the Perseus cluster 6,-0,0179=1.8’, the angular distance is denoted by Da."," Though the linear scales of cocoons are close in both cases, the angular diameter of the cocoon within the distant elliptical galaxy equaled to $\theta_{\mathrm{z}=1}=L_{\mathrm{box}}/D_\mathrm{A}\sim
2.5^{\prime\prime}$ is much smaller than that of the cocoon with the same linear size inside the Perseus cluster $\theta_{\mathrm{z}=0.0179}=1.8^{\prime}$, the angular distance is denoted by $D_\mathrm{A}$ ."
 Then the SZ fluxes from such cocoons which are proportional to A7/D should be of the same order (see Figs., Then the SZ fluxes from such cocoons which are proportional to $\Delta I/ D^2_\mathrm{A}$ should be of the same order (see Figs.
5] and[9))., \ref{sz217} and \ref{sz-sim}) ).
 The main difference in the intensity of the signal is then a result of the higher internal pressure in the simulated cocoon., The main difference in the intensity of the signal is then a result of the higher internal pressure in the simulated cocoon.
" The pressure of high temperature gas in the cavity center, in the framework of the analytical model in Sect."," The pressure of high temperature gas in the cavity center, in the framework of the analytical model in Sect."
" 3, can not be much higher than the ambient pressure of a galaxy cluster 3x10:11 erg/cm? because of the moderate Mach number of the shock of M=2 and a decrease of the pressure from the shock front to the cavity center (e.g., Sedov 1959)."," 3, can not be much higher than the ambient pressure of a galaxy cluster $\sim3\times10^{-11}$ $^3$ because of the moderate Mach number of the shock of $M=2$ and a decrease of the pressure from the shock front to the cavity center (e.g., Sedov 1959)."
 As was shown in Sect., As was shown in Sect.
" 2.2 using the combined generalized spectral functions provides us with an alternative to a measurement of the SZ effect at a frequency 217 GHz for studying a population of electrons with energies higher than 100 keV. In this section we derive the intensity map of the SZ effect by means of the combined function C(2.26,6.51,T.) and consider how the contribution from the kinematic Sunyaev-Zel'dovich (SZ) effect can be eliminated (for review of the kinematical SZ effect, see Birkinshaw 1999)."," 2.2 using the combined generalized spectral functions provides us with an alternative to a measurement of the SZ effect at a frequency 217 GHz for studying a population of electrons with energies higher than 100 keV. In this section we derive the intensity map of the SZ effect by means of the combined function $C(2.26, 6.51, T_{\mathrm{e}})$ and consider how the contribution from the kinematic Sunyaev-Zel'dovich (SZ) effect can be eliminated (for review of the kinematical SZ effect, see Birkinshaw 1999)."
 The generalized spectral function at a frequency 217 GHz belongs to the family of spectral functions defined by Eq. (7))., The generalized spectral function at a frequency 217 GHz belongs to the family of spectral functions defined by Eq. \ref{C}) ).
" Both G(x,Τα) and C(2.26,6.51,T.)functionshave absolutemaximal values at temperatures above 100keV and, therefore,"," Both $G(x, T_{\mathrm{e}})$ and $C(2.26, 6.51, T_{\mathrm{e}})$functionshave absolutemaximal values at temperatures above 100keV and, therefore,"
obtained to an accuracy of roughly one spectral channel. or 0.1 kkniss,"obtained to an accuracy of roughly one spectral channel, or $\sim 0.1$ $^{-1}$."
 We have spectral line data for a total of 20 sources — 14 Class 0 protostars and 6 Class | protostars — in transitions of 52) and ο 16193. 22)., We have spectral line data for a total of 20 sources – 14 Class 0 protostars and 6 Class I protostars – in transitions of $^+$ $\rightarrow$ 2) and $^{13}$ $\rightarrow$ 2).
 Ht ds necessary o obtain data in the optically thin isotope line to check hat the two peaks beingὃν observed in the main Isotope line are not due to two cdillerent clouds along the same line of sight., It is necessary to obtain data in the optically thin isotope line to check that the two peaks being observed in the main isotope line are not due to two different clouds along the same line of sight.
 We could confirm that the rare isotope line was »eaked in cach case., We could confirm that the rare isotope line was single-peaked in each case.
 The data are listed in Table 1., The data are listed in Table 1.
 All 20 sources have been well studied: previously. and have known continuum spectra and. bolometric temperatures.," All 20 sources have been well studied previously, and have known continuum spectra and bolometric temperatures."
 Table 1 ists the source namo. its bolometric temperature and Class.," Table 1 lists the source name, its bolometric temperature and Class."
 Table 1 also lists the measured peak separation. of the clouble-peakec spectra and the source of the literature data., Table 1 also lists the measured peak separation of the double-peaked spectra and the source of the literature data.
 Figure 3 shows a plot of peak separation versus Ti for the sources listed in Table 1., Figure 3 shows a plot of peak separation versus $_{{\rm Bol}}$ for the sources listed in Table 1.
 As shown above. peak separation scales with turbulence. level.," As shown above, peak separation scales with turbulence level."
 Consequently those sources showing high levels of turbulence appear in the upper parts of the graph. while those with Lower levels of turbulence appear in the lower parts of the graph.," Consequently those sources showing high levels of turbulence appear in the upper parts of the graph, while those with lower levels of turbulence appear in the lower parts of the graph."
 The error-bars correspond to ESIX and dO. +., The error-bars correspond to $\pm$ 5K and $\pm$ $^{-1}$.
 ys is used simply as à wav of separating the Class 0 protostars from the Class 1 protostars., $_{{\rm Bol}}$ is used simply as a way of separating the Class 0 protostars from the Class I protostars.
 The definition of a Class 0 protostar corresponds to it having Ty <TOWK. The vertical dashed line is the dividing line between Class 0 and { protostars., The definition of a Class 0 protostar corresponds to it having $_{{\rm Bol}}<$ 70K. The vertical dashed line is the dividing line between Class 0 and I protostars.
 In Figure 3 we see Class 0 protostars with a wide variety of levels of turbulence in their envelopes. although. we see that there are predominantly low levels of turbulence in Class | sources.," In Figure 3 we see Class 0 protostars with a wide variety of levels of turbulence in their envelopes, although we see that there are predominantly low levels of turbulence in Class I sources."
 Figure 3 is consistent with the evolutionary scenario proposed by ANBOS. in which Class 0 sources can form in regions of both high and low turbulence. before evolving into Class | protostars.," Figure 3 is consistent with the evolutionary scenario proposed by AWB93, in which Class 0 sources can form in regions of both high and low turbulence, before evolving into Class I protostars."
 One interesting [facet of the data is that there do not appear to be any. Class E protostars in Figure 3 with high turbulence levels., One interesting facet of the data is that there do not appear to be any Class I protostars in Figure 3 with high turbulence levels.
 This may not be statistically significant. due to the smaller number of Class L sources.," This may not be statistically significant, due to the smaller number of Class I sources."
 However. if future cata prove that this is also significant. then this appears to be saving that any initial turbulence. in the environment of a forming protostar dissipates before it has evolved to the Class 1 stage.," However, if future data prove that this is also significant, then this appears to be saying that any initial turbulence in the environment of a forming protostar dissipates before it has evolved to the Class I stage."
 The more massive envelopes in the earlier Class 0 protostars can sustain relatively high initial levels of turbulence when they occur., The more massive envelopes in the earlier Class 0 protostars can sustain relatively high initial levels of turbulence when they occur.
 However. by the Class E stage when more than half of the envelope has accreted onto the central protostar. a large percentage of any initial turbulence has been cissipatect.," However, by the Class I stage when more than half of the envelope has accreted onto the central protostar, a large percentage of any initial turbulence has been dissipated."
 We independently. estimated. the line optical depths for the data in Table | (using the isotope line data). and conlirmed that they lie in the range ~1.53.5. consistent with the regime of parameters that we are mocdelling (see discussion in section 2 above).," We independently estimated the line optical depths for the data in Table 1 (using the isotope line data), and confirmed that they lie in the range $\sim$ 1.5–3.5, consistent with the regime of parameters that we are modelling (see discussion in section 2 above)."
 We plotted. optical depth versus peak separation and found. no correlation., We plotted optical depth versus peak separation and found no correlation.
 This appears to confirm our assertion that the peak separation is not dominated. by optical depth effects. but rather by urbulence.," This appears to confirm our assertion that the peak separation is not dominated by optical depth effects, but rather by turbulence."
 We note finally that an alternative evolutionary scenario for the Class 0 LI protostellar stages has recently en proposed. (Javawardhana. Hartmann Calvet. 2001 hereafter JHICO in which Class 0. protostars only [orm in ugh density. environments whilst Class E. protostars form in low density environments. and the two stages are in act parallel evolutionary tracks.," We note finally that an alternative evolutionary scenario for the Class 0 I protostellar stages has recently been proposed (Jayawardhana, Hartmann Calvet, 2001 – hereafter JHC) in which Class 0 protostars only form in high density environments whilst Class I protostars form in low density environments, and the two stages are in fact parallel evolutionary tracks."
 In this scenario the high density environmoents required to produce Class O provostars are produced. by invoking colliding turbulent Lows (sec, In this scenario the high density environments required to produce Class 0 protostars are produced by invoking colliding turbulent flows (see
wavelengths (mean rest-frame wavelength of the filters is ~1560À for so that the k-correction between the different samples is negligible.,wavelengths (mean rest-frame wavelength of the filters is $\sim1560\AA$ for ) so that the $k$ -correction between the different samples is negligible.
 The distance modulus between the mean redshifts of the u- and the g-dropouts and between the mean redshifts of the g- and the ;-dropouts is ~ 0.39mag and ~ O.65mag. respectively (again for sin).," The distance modulus between the mean redshifts of the $u$ - and the $g$ -dropouts and between the mean redshifts of the $g$ - and the $r$ -dropouts is $\sim0.39$ mag and $\sim0.65$ mag, respectively (again for )."
 We estimate the angular correlation function w(#) of the different flux-limited subsamples by applying the estimator of ?.., We estimate the angular correlation function $\omega(\theta)$ of the different flux-limited subsamples by applying the estimator of \cite{1993ApJ...412...64L}.
 Errors are estimated from the field-to-field variance of the four fields., Errors are estimated from the field-to-field variance of the four fields.
 We choose a common angular binning for all samples but we check the influence of this choice (see below)., We choose a common angular binning for all samples but we check the influence of this choice (see below).
 The same masks as for the object detection are used for creating the random catalogue. which is necessary to estimate the correlation function.," The same masks as for the object detection are used for creating the random catalogue, which is necessary to estimate the correlation function."
 See Fig., See Fig.
 6 for the correlation functions of the different magnitude-limited LBG samples., \ref{fig:omega} for the correlation functions of the different magnitude-limited LBG samples.
" We fit a power-law. «X0)=A6"". to the angular correlation functions in the angular region 01«610’."," We fit a power-law, $\omega(\theta)=A\,\theta^{-\delta}$, to the angular correlation functions in the angular region $0\farcm1<\theta<10'$."
 A de-projection with Limber's (?) involving the redshift distributions shown above (see of Fig. 5)), A de-projection with Limber's \citep{1953ApJ...117..134L} involving the redshift distributions shown above (see of Fig. \ref{fig:z_dist}) )
 yields the real-space correlation funetions approximated by a power-law. &()(2). with 75 being the comoving correlation length and y=6+| being the slope of the correlation function.," yields the real-space correlation functions approximated by a power-law, $\xi(r)=\left(\frac{r}{r_0}\right)^{-\gamma}$, with $r_0$ being the comoving correlation length and $\gamma=\delta+1$ being the slope of the correlation function."
 We correct iteratively for the integral constraint bias with the method explained in ?.., We correct iteratively for the integral constraint bias with the method explained in \cite{2005ApJ...619..697A}.
 In this way we also estimate values for the galaxy bias factor. br The galaxy fluctuations in spheres of radius 8/7Mpe can be estimated from the power law fit to the correlation function in the following way (2): whereas the corresponding DM fluctuations at redshift <. σκίς). are calculated from theory.," In this way we also estimate values for the galaxy bias factor, $b_{\rm
 gal}$: The galaxy fluctuations in spheres of radius $8h^{-1}{\rm Mpc}$ can be estimated from the power law fit to the correlation function in the following way \citep{1980lssu.book.....P}: whereas the corresponding DM fluctuations at redshift $z$, $\sigma_8(z)$, are calculated from theory."
 The results are summarised in Tables 2--5 and the dependence of different measured quantities on redshift. magnitude. and the assumed redshift distribution is visualised in Fig. 7..," The results are summarised in Tables \ref{tab:r_0_z_mag_BPZ}- \ref{tab:r_0_z_mag_Masim} and the dependence of different measured quantities on redshift, magnitude, and the assumed redshift distribution is visualised in Fig. \ref{fig:r_0}."
 The errors for ro and 5644 in Fig., The errors for $r_0$ and $b_{\rm gal}$ in Fig.
 7 are estimated in Monte Carlo realisatiols from the errors of the power-law fit., \ref{fig:r_0} are estimated in Monte Carlo realisations from the errors of the power-law fit.
 Therein. we assume ulcorrelated Gaussian errors of the amplitude and the slope of the power-law.," Therein, we assume uncorrelated Gaussian errors of the amplitude and the slope of the power-law."
 In the lower S/N domain the error introduced by the binning of the data becomes dominant over, In the lower S/N domain the error introduced by the binning of the data becomes dominant over
inllowing gas (c.g.Governatoctal.2007).,inflowing gas \citep[e.g.][]{Governato07}.
. Thereafter it will sustain itself by erabbing coronal eas., Thereafter it will sustain itself by grabbing coronal gas.
 When a more massive galaxy experiences a major merger. the disrupted: forming disc is less likely to re-form.," When a more massive galaxy experiences a major merger, the disrupted star-forming disc is less likely to re-form."
 LP it does not form from cold inllowing gas. it will not form from coronal gas. because in the absence of a star-forming clisc. catastrophically. cooling coronal gas will feed the central black hole and reheat the corona.," If it does not form from cold inflowing gas, it will not form from coronal gas, because in the absence of a star-forming disc, catastrophically cooling coronal gas will feed the central black hole and reheat the corona."
 Fhus star-formation permanently ceases if there is insullicient. cool gas to re-Form a star-forming disc after a Merecr., Thus star-formation permanently ceases if there is insufficient cool gas to re-form a star-forming disc after a merger.
 In this picture it is likely that cooling at the centre of he corona of a disc galaxy leads to episodic reheating of he central corona. just as in a classical cooling tow.," In this picture it is likely that cooling at the centre of the corona of a disc galaxy leads to episodic reheating of the central corona, just as in a classical cooling flow."
 Tha is. we sugeest that the corona of a star-forming disc galaxy accretes onto central black hole onto the disc.," That is, we suggest that the corona of a star-forming disc galaxy accretes onto central black hole onto the disc."
 For his to be a viable proposal. the central reheating associatec with aceretion by the black hole must not undermine the ability of the star-forming cise to grab coronal eas from the xut of the corona that [ies above it.," For this to be a viable proposal, the central reheating associated with accretion by the black hole must not undermine the ability of the star-forming disc to grab coronal gas from the part of the corona that lies above it."
 This condition coulc »e satisfied if the AGN outburst were sullicientlv small aux sullicientIy directed perpendicular to the galactic plane., This condition could be satisfied if the AGN outburst were sufficiently small and sufficiently directed perpendicular to the galactic plane.
 Fhis is à topic for a later paper. however.," This is a topic for a later paper, however."
 As a cloud moves through the ambient coronal eas. turbulence in the boundary [aver at the interface of the hot ancl cold fluids must cause gas to be stripped from the cloud at some rate.," As a cloud moves through the ambient coronal gas, turbulence in the boundary layer at the interface of the hot and cold fluids must cause gas to be stripped from the cloud at some rate."
 To calculate this rate from ab-initio physics is a daunting task because one would have to consider plasma instabilities in addition to hydrodynamical ones such as the Ixelvin-LHelmholtz instability. and turbulent energy. will be cascading to very small scales.," To calculate this rate from ab-initio physics is a daunting task because one would have to consider plasma instabilities in addition to hydrodynamical ones such as the Kelvin-Helmholtz instability, and turbulent energy will be cascading to very small scales."
" In view of these dillieulties. the obvious way forward is to parametrise the probleni by hypothesising that mass is stripped from the a cloud of mass AM, αἲ à rate aM,"," In view of these difficulties, the obvious way forward is to parametrise the problem by hypothesising that mass is stripped from the a cloud of mass $M_\c$ at a rate $\alpha M_\c$."
 AX turbulent wake of stripped gas will run back through the corona from the moving cloud., A turbulent wake of stripped gas will run back through the corona from the moving cloud.
 In this wake turbulence will mix the stripped gas with ambient gas., In this wake turbulence will mix the stripped gas with ambient gas.
 A key quantity is the cooling time of the plasma that results from this mixing., A key quantity is the cooling time of the plasma that results from this mixing.
 We estimate this under the assumption that mixing occurs so quickly that radiative losses during mixing can be neglected., We estimate this under the assumption that mixing occurs so quickly that radiative losses during mixing can be neglected.
 Let s be a short length of the wake. which has cross- area chy and originally contained a mass clespp of coronal eas.," Let $s$ be a short length of the wake, which has cross-sectional area $A_\w$ and originally contained a mass $A_\w s\rho_\h$ of coronal gas."
" In the time s/r that it took the cloud to pass through the length. a mass aM,s/e of gas was stripped from it."," In the time $s/v$ that it took the cloud to pass through the length, a mass $\alpha M_\c s/v$ of gas was stripped from it."
 After the cold gas has mixed and come into thermal equilibrium with the coronal gas. the temperature of the resulting uid is where Zi and d; are the coronal anc cloud temperatures. respectively.," After the cold gas has mixed and come into thermal equilibrium with the coronal gas, the temperature of the resulting fluid is where $T_\h$ and $T_\c$ are the coronal and cloud temperatures, respectively."
 The natural unit for the cross-sectional area of the wake is the characteristic cross-section. (Mc/pc) of the cloud., The natural unit for the cross-sectional area of the wake is the characteristic cross-section $(M_\c/\rho_\c)^{2/3}$ of the cloud.
" We write cle=SCM,p,7. where 31."," We write $A_\w= \beta(M_\c/\rho_\c)^{2/3}$, where $\beta\sim1$."
" Thus We simplify this expression under the assumption that the cloud is in approximate pressure equilibrium: with the Corona. so pyly&pedi. and have Since CM,p.""afr is the fraction of the clouds mass that is stripped in the time taken for the eloud to travel its own length. and 3~I. the second term in the numerator of equation (3)) must be small compared to unity and may be neglected."," Thus We simplify this expression under the assumption that the cloud is in approximate pressure equilibrium with the corona, so $\rho_\h T_\h\simeq\rho_\c T_\c$, and have Since $(M_\c/\rho_\c)^{1/3}\alpha/v$ is the fraction of the cloud's mass that is stripped in the time taken for the cloud to travel its own length, and $\beta\sim1$, the second term in the numerator of equation \ref{eq:Trat}) ) must be small compared to unity and may be neglected."
" The second term in the denominator is larger by a [actor Zn/1.z200. so Zi, may be significantly lower than Tu."," The second term in the denominator is larger by a factor $T_\h/T_\c\ga200$, so $T_\m$ may be significantly lower than $T_\h$."
 On account of the steepness of the cooling-time curve plotted inLl. this result implies that the cooling time in the wake may fall below the ambient cooling time bv a actor of several.," On account of the steepness of the cooling-time curve plotted in, this result implies that the cooling time in the wake may fall below the ambient cooling time by a factor of several."
 From this back-of-envelope. caleulation we draw the ollowing conclusions, From this back-of-envelope calculation we draw the following conclusions
Recent experiments in terrestrial nuclear laboratories have narrowed down significantly the range of the EOS of neutron-rich nuclear matter although there are still many remaining challenges and uncertainties.,Recent experiments in terrestrial nuclear laboratories have narrowed down significantly the range of the EOS of neutron-rich nuclear matter although there are still many remaining challenges and uncertainties.
" In particular, the EOS for symmetric nuclear matter was constrained up to about five times the normal nuclear matter density by collective flow and particle production in relativistic heavy-ion reactions."," In particular, the EOS for symmetric nuclear matter was constrained up to about five times the normal nuclear matter density by collective flow and particle production in relativistic heavy-ion reactions."
 The density dependence of the symmetry energy was constrained at subsaturation densities by isospin diffusion and isoscaling in heavy-ion reactions at intermediate energies., The density dependence of the symmetry energy was constrained at subsaturation densities by isospin diffusion and isoscaling in heavy-ion reactions at intermediate energies.
 Applying the EOSs constrained by the reaction data we have studied the neutron star momenta of inertia of both slowly and, Applying the EOSs constrained by the heavy-ion reaction data we have studied the neutron star momenta of inertia of both slowly and
to prevent clilfusion and thus lens broadening.,to prevent diffusion and thus lens broadening.
 Images are always scattered perpendicular to the field line direction., Images are always scattered perpendicular to the field line direction.
 ‘Two regimes can be searched to confirm this model of these convergent lenses., Two regimes can be searched to confirm this model of these convergent lenses.
 Phe most direct is the observation of an ESE with VLBL, The most direct is the observation of an ESE with VLBI.
 It. had been done once by 2.. but. with insullicient resolution to discern the multiple images.," It had been done once by \cite{2000ApJ...534..706L}, but with insufficient resolution to discern the multiple images."
 They note a slight broadening of the source at the lensing peak. which might be due to an unresolved mixing of two images.," They note a slight broadening of the source at the lensing peak, which might be due to an unresolved mixing of two images."
 Lensing conserves surface brightness. so one expects a larger image to be brighter.," Lensing conserves surface brightness, so one expects a larger image to be brighter."
 The authors used this argument. to rule out a refractive divergent lens., The authors used this argument to rule out a refractive divergent lens.
 We note. however. that the expected. lens image separation is only mocdestIy. sav a [actor of 2 larger than the source.," We note, however, that the expected lens image separation is only modestly, say a factor of 2 larger than the source."
 This results in two images. each of which is demagnified and smaller. but. next to each other to make them appear more extended.," This results in two images, each of which is demagnified and smaller, but next to each other to make them appear more extended."
 Their synthesized beam is much larger than the source size or the separation. so these two ellects are not separable.," Their synthesized beam is much larger than the source size or the separation, so these two effects are not separable."
 Unfortunately. the observational monitoring progranune of bright compact sources has been discontinued.," Unfortunately, the observational monitoring programme of bright compact sources has been discontinued."
 Pulsars also undergo extreme scattering events. (?).., Pulsars also undergo extreme scattering events \citep{1993Natur.366..320C}.
 In such cases. the three images will coherently interfere with each other. resulting in a distinctive signature in the secondary spectrum. shown in Figure 4..," In such cases, the three images will coherently interfere with each other, resulting in a distinctive signature in the secondary spectrum, shown in Figure \ref{fig:delay}."
 This figure plots the raw delavs which can be obtained. from. interstellar holography(?).., This figure plots the raw delays which can be obtained from interstellar \citep{2008MNRAS.388.1214W}.
 So far. the discussion has focussed. on the quantitative properties of the convergent lens.," So far, the discussion has focussed on the quantitative properties of the convergent lens."
 Some comments on its relation to the physical properties of the lenses are in order., Some comments on its relation to the physical properties of the lenses are in order.
" Typical electron densities. as measured by pulsar dispersion measure are n,0.03.", Typical electron densities as measured by pulsar dispersion measure are $n_e \sim 0.03$.
 This gives a phase velocity about 10.I7 larger than e at frequencies of ~ CllIz., This gives a phase velocity about $10^{-12}$ larger than $c$ at frequencies of $\sim$ GHz.
 For an uncerdense lens. the maximal dellection angle from Snell's law is therefore 10.“fa mas. where a is the angle between the lens surface and the line of sight.," For an underdense lens, the maximal deflection angle from Snell's law is therefore $10^{-4}/\alpha$ mas, where $\alpha$ is the angle between the lens surface and the line of sight."
 Grazing incidence with a10! js needed to explain the observed refractive angles of  mas.," Grazing incidence with $\alpha \sim
10^{-4}$ is needed to explain the observed refractive angles of $\sim$ mas."
 We speculate that. reconnection current sheets cou satisfy the properties required by the phenomenoloey., We speculate that reconnection current sheets could satisfy the properties required by the phenomenology.
 Unfortunately. the quantitative process of reconnection is very poorly understood.," Unfortunately, the quantitative process of reconnection is very poorly understood."
 The thickness of a current. shee is probably related to resistivity., The thickness of a current sheet is probably related to resistivity.
 Phe primary quantitative calculable ohmic restivity leads to very long reconnection time scales. leacing to problems with the generation of magnetic fields.," The primary quantitative calculable ohmic restivity leads to very long reconnection time scales, leading to problems with the generation of magnetic fields."
 Alternative reconnection processes. or Las reconnection. has been proposed.," Alternative reconnection processes, or fast reconnection, has been proposed."
 The reconnection rate arn sheet geometry ους by many orders of magnitude between proposed scenarios (?).., The reconnection rate and sheet geometry differs by many orders of magnitude between proposed scenarios \citep{2010PhPl...17j2302P}.
 Without a known reconnection rate. the incidence rate of current sheets is not predictable. nor is the lens lickness.," Without a known reconnection rate, the incidence rate of current sheets is not predictable, nor is the lens thickness."
 Similarly. the aspect ratio of the sheets is not. known.," Similarly, the aspect ratio of the sheets is not known."
 One qualitatively expects the longest axis to be related to the curvature of field. lines. which determines the amount of area that approaches sullicienthy closely to undergo reconnection.," One qualitatively expects the longest axis to be related to the curvature of field lines, which determines the amount of area that approaches sufficiently closely to undergo reconnection."
 The shortest axis is given by the unknown resistive scale. which could be as short as the electron. gvromagnetic radius.," The shortest axis is given by the unknown resistive scale, which could be as short as the electron gyromagnetic radius."
 “Lo explain the phenomenology of ESE. we require the ratio of the intermediate axis to the short axis to be 710.," To explain the phenomenology of ESE, we require the ratio of the intermediate axis to the short axis to be $\gtrsim 10^4$."
 Unfortunately. the statistical properties. of SEs is. not predictable: within our understanding of plasma physics.," Unfortunately, the statistical properties of ESEs is not predictable within our understanding of plasma physics."
 Insteacl. we can use the properties of the LSE to learn about reconnection.," Instead, we can use the properties of the ESE to learn about reconnection."
 We have introduced triaxial convergent plasma lenses to explain three separate. but related. refractive plasma lensing ellects.," We have introduced triaxial convergent plasma lenses to explain three separate, but related, refractive plasma lensing effects."
 Reconnection events generically lead το highly Hattenec current sheets., Reconnection events generically lead to highly flattened current sheets.
 We have found that generic physical conditions can account for all the known phenomena. which are otherwise challenging to explain.," We have found that generic physical conditions can account for all the known phenomena, which are otherwise challenging to explain."
 Fhese include the apparent distance localization of scatters. the modest frequeney cdepencenee of ESE. overpressure problem.," These include the apparent distance localization of scatters, the modest frequency dependence of ESE, overpressure problem."
 Lt also sugeests that the apparent dilfractive phenomena in pulsar may be due to interference of refractive images. rather than stochastic Lresnel scale lensing.," It also suggests that the apparent diffractive phenomena in pulsar may be due to interference of refractive images, rather than stochastic fresnel scale lensing."
 U-LP thanks NSERC and the KLCC for support and hosting the visit which lead to this paper. and LJI thanks the Roval Society.," U-LP thanks NSERC and the KICC for support and hosting the visit which lead to this paper, and LJK thanks the Royal Society."
 We would like to thank Peter Goldreich. Barney lückett. Ethan Vishniae and Chris Thompson for helpful discussions.," We would like to thank Peter Goldreich, Barney Rickett, Ethan Vishniac and Chris Thompson for helpful discussions."
"the fgQSOs is the same as that of the cosmic average, which are clearly offset from the observations.","the fgQSOs is the same as that of the cosmic average, which are clearly offset from the observations."
" Note that the errors in the estimates of the fgQSO systemic redshift are also considered in our simulation here by adopting the method same as that used in i.e., the error is randomly chosen over the range of (2007),,0—1500kms~!."," Note that the errors in the estimates of the fgQSO systemic redshift are also considered in our simulation here by adopting the method same as that used in , i.e., the error is randomly chosen over the range of $0-1500\kms$."
" The decrease of the simulated optical depth with decreasing distance to the QSOs appears more significant than the observation results, as the effect of the density enhancement in the proximity regions of these QSOs is ignored."," The decrease of the simulated optical depth with decreasing distance to the QSOs appears more significant than the observation results, as the effect of the density enhancement in the proximity regions of these QSOs is ignored."
" Assuming that the density in the QSO proximity regions is enhanced by a factor of (A(R))~1+Co(R/Cg)?exp|-(R/Cg)3], the observations can be reproduced by simulations (blue solid circles) if Co=0.9, Cr=11 Mpc, p=0.7, and q=0.5.3 For this case, we set the number of spectra in each realization to be 50 and we simulate 100 realizations."," Assuming that the density in the QSO proximity regions is enhanced by a factor of $\left<\Delta(R)\right>\sim 1+ C_0(R/C_R)^{-p} \exp[-(R/C_R)^q]$, the observations can be reproduced by simulations (blue solid circles) if $C_0=0.9$, $C_R=11$ Mpc, $p=0.7$, and $q=0.5$ For this case, we set the number of spectra in each realization to be 50 and we simulate 100 realizations."
 Thus the mean median optical depth and its standard deviation can be estimated from these realizations., Thus the mean median optical depth and its standard deviation can be estimated from these realizations.
" As seen from Figure 8,, the observations can be well matched by the simulation results (blue solid points)."," As seen from Figure \ref{fig:f8}, the observations can be well matched by the simulation results (blue solid points)."
 also suggest that the density enhancement in (2007)the near zones of QSOs is required in order to explain their measurements on the LOSPE., also suggest that the density enhancement in the near zones of QSOs is required in order to explain their measurements on the LOSPE.
" Compared with their estimates on the density enhancement in the near zones of those high luminosity QSOs, our estimates are smaller by a factor of 1.5—5 at a distance of RyS10 Mpc-1 Mpc but similar at a distance of RjZ15 Mpc."," Compared with their estimates on the density enhancement in the near zones of those high luminosity QSOs, our estimates are smaller by a factor of $1.5-5$ at a distance of $R_{||}\la 10$ $1$ Mpc but similar at a distance of $R_{||}\ga 15$ Mpc."
" The main reason for this difference is that the Poisson variance of absorption lines in different sight lines to the near zones of QSOs, which is included in the Monte-Carlo simulations here, leads to an averaged absorption that is larger than a simple estimation obtained without considering of the Poisson variance2007)."," The main reason for this difference is that the Poisson variance of absorption lines in different sight lines to the near zones of QSOs, which is included in the Monte-Carlo simulations here, leads to an averaged absorption that is larger than a simple estimation obtained without considering of the Poisson variance."
". Therefore, the density enhancement required here is less significant compared with that obtained by (see also discussions on the Poisson variance in 2008))."," Therefore, the density enhancement required here is less significant compared with that obtained by (see also discussions on the Poisson variance in )."
" Although so far there is no observational measurement on the TPE of QSOs (with luminosity l0?lergs! similar to the sample in(2007),, we Hz~')illustrate here the effects of different QSO properties on the TPE and demonstrate that both 7, and Og can be simultaneously constrained by the TPE if the density enhancement in the fgQSO proximity region has been constrained by the LOSPE."," Although so far there is no observational measurement on the TPE of QSOs (with luminosity $\sim 8\times 10^{31}\ergshz$ ) similar to the sample in, we illustrate here the effects of different QSO properties on the TPE and demonstrate that both $\tau_{\rm lt}$ and $\Theta_0$ can be simultaneously constrained by the TPE if the density enhancement in the fgQSO proximity region has been constrained by the LOSPE."
 We first assume that the effect of the density enhancement in the proximity regions of those QSOs is the same as that indicated by the LOSPE obtained by (see Figure 8))., We first assume that the effect of the density enhancement in the proximity regions of those QSOs is the same as that indicated by the LOSPE obtained by (see Figure \ref{fig:f8}) ).
" Under this assumption, we synthesize(2007) a large number of Lya forest spectra of bgQSOs with different choices of the QSO lifetime and the half opening angle of the associated torus (see parameter settings in Section ??))."," Under this assumption, we synthesize a large number of $\alpha$ forest spectra of bgQSOs with different choices of the QSO lifetime and the half opening angle of the associated torus (see parameter settings in Section \ref{sec:Geom}) )."
 We then extract the median optical depth distribution in the fgQSO proximity regions from these spectra and show the results in Figure 9.., We then extract the median optical depth distribution in the fgQSO proximity regions from these spectra and show the results in Figure \ref{fig:f9}. .
" As seen from Figure 9aa, given the half opening angle of the torus associated with the fgQSO, e.g., Og=60°, generally the optical depth in the proximity region to the QSO (e.g., |Rj|S10 Mpc) decrease with increasing the QSO lifetime; and thedecrease is significant initially at the frontside to the QSO, and becomes significant at its backside (e.g., Rj—4 Mpc) when the QSO lifetime τι is long enough so that the region can be reached by the radiation from the fgQSO."," As seen from Figure \ref{fig:f9}a a, given the half opening angle of the torus associated with the fgQSO, e.g., $\Theta_0=60\degr $, generally the optical depth in the proximity region to the QSO (e.g., $|R_{\parallel}|\la 10$ Mpc) decrease with increasing the QSO lifetime; and thedecrease is significant initially at the frontside to the QSO, and becomes significant at its backside (e.g., $R_{\parallel}\sim-4$ Mpc) when the QSO lifetime $\tau_{\rm lt}$ is long enough so that the region can be reached by the radiation from the fgQSO."
" That tendency is manifest especially for large Oo, as (1) for the case with short QSO lifetime (e.g., τι~10°° yr; see Figure 9bb), the insignificant decrease of the optical depth is not sensitive to different choices of Oo; and (2) as shown in Figure 9cc and d, for large τις, the decrease of the optical depth at RyS;0 Mpc increases with increasing large Oo."," That tendency is manifest especially for large $\Theta_0$, as (1) for the case with short QSO lifetime (e.g., $\tau_{\rm
lt}\sim 10^{6.3}$ yr; see Figure \ref{fig:f9}b b), the insignificant decrease of the optical depth is not sensitive to different choices of $\Theta_0$; and (2) as shown in Figure \ref{fig:f9}c c and d, for large $\tau_{\rm lt}$, the decrease of the optical depth at $R_{\parallel}\la 0$ Mpc increases with increasing large $\Theta_0$."
" As shown in the figure, the QSO lifetime and the half opening angle of the torus affect the optical depth curves in different ways, so that they can be simultaneously constrained by the LOSPE and TPE of QSOs if τι and Θρ are not too small n10° yr and O9= "," As shown in the figure, the QSO lifetime and the half opening angle of the torus affect the optical depth curves in different ways, so that they can be simultaneously constrained by the LOSPE and TPE of QSOs if $\tau_{\rm lt}$ and $\Theta_0$ are not too small (e.g., $\tau_{\rm lt} \ga
10^6$ yr and $\Theta_0\ga 30\degr $ )."
"Note here that the standard(e.g., 2deviation, on the order 30°).similar to that shown in Figure 9,, is small enough so that the constraints on the QSO lifetime and the torus half opening angle may be accurately extracted through the PEs of a sample with a few hundreds of high luminosity QSOs."," Note here that the standard deviation, on the order similar to that shown in Figure \ref{fig:f9}, is small enough so that the constraints on the QSO lifetime and the torus half opening angle may be accurately extracted through the PEs of a sample with a few hundreds of high luminosity QSOs."
" Considering the possibility of inaccuracy in the errors of the QSO systemic redshifts used above, we also test that our results obtained above remain the same qualitatively, if the systemic redshift errors is set to be the same as that for the low-luminosity low-redshift sample inSection ??.."," Considering the possibility of inaccuracy in the errors of the QSO systemic redshifts used above, we also test that our results obtained above remain the same qualitatively, if the systemic redshift errors is set to be the same as that for the low-luminosity low-redshift sample inSection \ref{sec:tpe1}. ."
" The overall shape of the optical depth distribution around Ry)~0 Mpc does not change significantly, though it may shift slightly toward the"," The overall shape of the optical depth distribution around $R_{\parallel}\sim0$ Mpc does not change significantly, though it may shift slightly toward the"
We also analyze the cluster heat flux from our stored time slices at the conclusion of a run.,We also analyze the cluster heat flux from our stored time slices at the conclusion of a run.
" We will focus on the radial component of the heat fluxes in a spherical shell of width 100 kpc, centered on a radius of 450 kpc from the cluster center."," We will focus on the radial component of the heat fluxes in a spherical shell of width 100 kpc, centered on a radius of 450 kpc from the cluster center."
" We begin by defining fiducial conductivity to be the Spitzer conductivity in athis shell if the conduction were isotropic at the Spitzer value, namely: This value is the same as for anisotropic thermal conduction with purely radial field lines."," We begin by defining a fiducial conductivity to be the Spitzer conductivity in this shell if the conduction were isotropic at the Spitzer value, namely: This value is the same as for anisotropic thermal conduction with purely radial field lines."
" In Figure 6,, we compare the heat flux due to conduction to this Spitzer value."," In Figure \ref{fig:A2-heatflux}, we compare the heat flux due to conduction to this Spitzer value."
" However, there remains some ambiguity as the temperature profile evolves over the course of the simulation."," However, there remains some ambiguity as the temperature profile evolves over the course of the simulation."
" Thus, we normalize the conductive heat flux to both the Spitzer heat flux and the Spitzer heat flux."," Thus, we normalize the conductive heat flux to both the Spitzer heat flux and the Spitzer heat flux."
" We define the Spitzer fraction, the oft-quoted paramater as In Table 4,, we utilize the instantaneous fiducial heat flux at the time of maximum #|)for normalization of this quantity."," We define the Spitzer fraction, the oft-quoted paramater as In Table \ref{tab:satprop}, we utilize the instantaneous fiducial heat flux at the time of maximum for normalization of this quantity."
" Figure 6 shows that as the ICM temperature gradient is depressed, the conductive heat flux normalized to the instantaneous heat flux increases substantially, until is larger than 0.4."," Figure \ref{fig:A2-heatflux} shows that as the ICM temperature gradient is depressed, the conductive heat flux normalized to the instantaneous heat flux increases substantially, until $f_\textrm{Spitz}$ is larger than 0.4."
" Meanwhile, calculatingfgpitz the heat flux due to the transport of enthalpy is worth a careful treatment."," Meanwhile, calculating the heat flux due to the transport of enthalpy is worth a careful treatment."
" The advection of enthalpy can be written as where we neglect the transport of kinetic and magnetic energy since the flows are both low mach number and high-8, respectively."," The advection of enthalpy can be written as where we neglect the transport of kinetic and magnetic energy since the flows are both low mach number and $\beta$, respectively."
 We can split this heat flux into an advective term proportional to a mass flux (cxnv) and a term that is a purely convective heat flux., We can split this heat flux into an advective term proportional to a mass flux $\propto n\boldsymbol{v}$ ) and a term that is a purely convective heat flux.
" Writing the pressure in Eqn.[36]] as nkpgT', we define several quantities: where, for example (n) is the spatial average of the number density in the shell and ὅτι is the local deviation of a quantity from its average."," Writing the pressure in \ref{eqn:qenthalpy}] ] as $nk_B T$ we define several quantities: where, for example $\langle n\rangle$ is the spatial average of the number density in the shell and $\delta n$ is the local deviation of a quantity from its average."
" In this formulation the new total heat flux becomes In future equations, we drop the r subscript."," In this formulation the new total heat flux becomes In future equations, we drop the $r$ subscript."
" In simplifying this expression, terms that are proportional to the average of a delta quantity, like (6n) vanish by definition."," In simplifying this expression, terms that are proportional to the average of a delta quantity, like $\langle \delta n\rangle$ vanish by definition."
" We can then separate it into two quantities: one due to a mass flux, identified as terms that scale as nv,and the remaining terms into a convective flux."," We can then separate it into two quantities: one due to a mass flux, identified as terms that scale as $nv$ ,and the remaining terms into a convective flux."
both exacerbates Che lithium problem but also opens the possibility of using Li abundances in different cosmic environments to probe both the early universe ancl cosmic-ray acceleration in structure formation.,both exacerbates the lithium problem but also opens the possibility of using Li abundances in different cosmic environments to probe both the early universe and cosmic-ray acceleration in structure formation.
 Lithium also plavs a role in the primordial chemistry leading to the lormation of the first stars., Lithium also plays a role in the primordial chemistry leading to the formation of the first stars.
 Namely. the lithium hydride (Lill) molecule (Bougleux&Galli1997:Palla1993). contributes to the cooling of the dust-Iree primordial gas.," Namely, the lithium hydride (LiH) molecule \citep{bg97,gp98}
 contributes to the cooling of the dust-free primordial gas."
 Current caleulations suggest however that the impact of Lill is not dominant due to the small primordial Li/ll abundance. but the chemistry of primordial Lill is complex aud a subject of ongoing interest (Dovinoetal.2011).," Current calculations suggest however that the impact of LiH is not dominant due to the small primordial Li/H abundance, but the chemistry of primordial LiH is complex and a subject of ongoing interest \citep{bovino}."
. Lithium (hus takes a uniquely important role in the cosmologv of the early universe., Lithium thus takes a uniquely important role in the cosmology of the early universe.
 Unfortunately. lithium is dillieult to measure due to its tinv abundance: in the solar svstem. Li/Il ?. while in metal-poor halo stars. Li/IL o x 10...," Unfortunately, lithium is difficult to measure due to its tiny abundance: in the solar system, Li/H $\sim$ $^{-9}$, while in metal-poor halo stars, Li/H $\sim$ $\times$ $^{-10}$."
 To date. optical spectroscopy of Galactic stellar atmospheres has been the only probe of Li in (hie early Galaxy. and stubborn svstematic uncertainties could account for the Li problem.," To date, optical spectroscopy of Galactic stellar atmospheres has been the only probe of Li in the early Galaxy, and stubborn systematic uncertainties could account for the Li problem."
 Any alternative means of observing Li in any source would be extraordinarily useful for many reasons. not least of which would be the different svstematics.," Any alternative means of observing Li in any source would be extraordinarily useful for many reasons, not least of which would be the different systematics."
 Furthermore. to date there are no stellar observations of extragalactic Li. and so anv extragalactic information would shed qualitatively new light onto this problem.," Furthermore, to date there are no stellar observations of extragalactic Li, and so any extragalactic information would shed qualitatively new light onto this problem."
 since the atomic lines of Li have been so challenging to observe. we have chosen (o search for the simplest Li containing molecule. lithium hydride (Lill).," Since the atomic lines of Li have been so challenging to observe, we have chosen to search for the simplest Li containing molecule, lithium hydride (LiH)."
 Lill has been searched for previously by Combes&Wiklind(1998) towardDO2182-351.. where (he authors report a tentative detection in a svstem that is 5 km/s ollset from the svstemic velocity but within the CO velocity prolile.," LiH has been searched for previously by \citet{combes98} toward, where the authors report a tentative detection in a system that is 5 km/s offset from the systemic velocity but within the CO velocity profile."
 The Lill molecular isolopomers have a 0:0. J—0-1 rotational transition al 443.95293 Gllz (Lill) and 453.16028 GIIz (LilD (Bellinietal.1994:Plummer.cia LO84).," The LiH molecular isotopomers have a =0, =0-1 rotational transition at 443.95293 GHz $^7$ LiH) and 453.16028 GHz $^6$ LiH) \citep{bellini94,plummer84}."
. This region of the mm-wave spectrum is hiehlv opaque through the alinosphere., This region of the mm-wave spectrum is highly opaque through the atmosphere.
 Tlowever these lines are accessible in the 1 3 mm bands at the Combined Array for Research in \lillimeter-wave Astronomy. (CABRMA) for extragalactic sources in redshift ranges of 2—0.64-1.1 and 2.85-4.2. respectively.," However these lines are accessible in the 1 3 mm bands at the Combined Array for Research in Millimeter-wave Astronomy (CARMA) for extragalactic sources in redshift ranges of =0.64-1.1 and 2.85-4.2, respectively."
 Our strategy [or detecting. Lill in these redshift ranges is (ο search. [or interstellar Lill absorption towarel lensed quasars with high millimeter continuum [flux densitv., Our strategy for detecting LiH in these redshift ranges is to search for interstellar LiH absorption toward lensed quasars with high millimeter continuum flux density.
 The absorplion originates [from molecular clouds in the lensing galaxy along the line of sight., The absorption originates from molecular clouds in the lensing galaxy along the line of sight.
At the white-clhwarl surface. the temperature is which is also finite.,"At the white-dwarf surface, the temperature is which is also finite."
" 1n the limit of m,©0. we have στ)= land Ga/Or=O."," In the limit of $\tau_{\rm m} \rightarrow 0$, we have $\sigma(\tau) \rightarrow 1$ and $\partial \sigma/\partial \tau \rightarrow 0$ ."
 Moreover. €(0)»x. 60)=0. and ze(0)=1. tthe cold “stationary-wall” boundary condition is recovered.," Moreover, $\zeta(0) \rightarrow \infty$, $\theta(0) \rightarrow 0$, and $\varpi(0) = 1$, the cold “stationary-wall” boundary condition is recovered."
 Not only are the hyerocdvnamic variables of the conventional and our formulations are cillerent at the lower boundary. but they also have very dilferent. eracicnts.," Not only are the hydrodynamic variables of the conventional and our formulations are different at the lower boundary, but they also have very different gradients."
 In the conventional formulation. the density gradient is given by As Or/OZx777 [or small 7 (Leon equation (2) in Wu 1994). the density eracient is proportional to 7Αι," In the conventional formulation, the density gradient is given by As $\partial \tau/ \partial \xi \propto \tau^{-3/2}$ for small $\tau$ (from equation (2) in Wu 1994), the density gradient is proportional to $\tau^{-7/2}$."
 =0.7=0: hence the density gracient is infinite.," At $\xi = 0$, $\tau = 0$; hence the density gradient is infinite."
" The pressure eradient. which is and the temperature eradient. which is are also infinito at £=0: For à ""leakv accretion column allowing an energy Lux emerging from below. the velocity gradient at the base is The velocity gradient is zero at £=0. provided that X(0)=Au(O)."," The pressure gradient, which is and the temperature gradient, which is are also infinite at $\xi = 0$; For a “leaky” accretion column allowing an energy flux emerging from below, the velocity gradient at the base is The velocity gradient is zero at $\xi = 0$, provided that $\Lambda_{\rm c}(0) = \Lambda_{\rm h}(0)$."
 This is in [act the conditionof radiative equilibrium in a hyelrostatic white-cdwarl atmosphere., This is in fact the conditionof radiative equilibrium in a hydrostatic white-dwarf atmosphere.
" Dilferentiating equation (S) with respect to £ vields the density. gradient: ὃν substituting the expression of στ) given in equation (15) into the equation above. we obtain where The power series C, converges. with a value of the order of unity when τ=0 (Appendix A.1)."," Differentiating equation (8) with respect to $\xi$ yields the density gradient: By substituting the expression of $\sigma(\tau)$ given in equation (15) into the equation above, we obtain where The power series $C_n$ converges, with a value of the order of unity when $\tau \rightarrow 0$ (Appendix A.1)."
 The densitygracient is therefore finite and equals zero at £= 0., The densitygradient is therefore finite and equals zero at $\xi = 0$ .
 The pressure gradient is, The pressure gradient is
so 1 AL. is ejected from the disk.,so $\sim$ 1 $\Msun$ is ejected from the disk.
 Most of this will not be neutron-rich enough or expand fast enough to make the r-process. but consider the implications if only of the mass that is lost comes [rom the inner disk at times when the accretion rate is over 0.1 AL.see1 and Y.«04.," Most of this will not be neutron-rich enough or expand fast enough to make the $r$ -process, but consider the implications if only of the mass that is lost comes from the inner disk at times when the accretion rate is over 0.1 $\Msunsec$ and $Y_e
< 0.4$."
 50 close to the hole. signilicant entropy will be added by any. acceleration process (hat produces a strong outflow.," So close to the hole, significant entropy will be added by any acceleration process that produces a strong outflow."
" A erude estimate is s/A~GMmj(Ohl,μα.)o4005,(2MeV/T,,,,) (Qian&Woosley19000). where {ως is the temperature where (he energy deposition is the ereatest."," A crude estimate is $s/k_b \sim G M m_n/(r k T_{max}) \sim 400 \, r_6^{-1} (2 \ {\rm
MeV}/T_{max})$ \citep{qia96}, where $T_{max}$ is the temperature where the energy deposition is the greatest."
" In fact the conditions in the inner disk lor accretion rates 0.1. AL.sec.*1 are not so different - in terms of neutrino luminosity. neutron excess. temperature. and gravitational potential - from the neutrino-driven wind in an ordinary supernova, aud one might expect a similar r-process (Wooslevetal.1994)."," In fact the conditions in the inner disk for accretion rates $\sim$ 0.1 $\Msunsec$ are not so different - in terms of neutrino luminosity, neutron excess, temperature, and gravitational potential - from the neutrino-driven wind in an ordinary supernova, and one might expect a similar $r$ -process \citep{woo94}."
. The gravitational potential though. can in principle. be ereater ancl a larger entropy and shorter time scale might be realized.," The gravitational potential though, can in principle, be greater and a larger entropy and shorter time scale might be realized."
 Assume that the ejecta will be very rich in a-particles and contain of order r-process by mass., Assume that the ejecta will be very rich in $\alpha$ -particles and contain of order $r$ -process by mass.
 Putting the numbers together then gives an r-process svnthesis equivalent to LO? M. per supernova. making this possibility well worth further investigation.," Putting the numbers together then gives an $r$ -process synthesis equivalent to $^{-5}$ $\Msun$ per supernova, making this possibility well worth further investigation."
 The authors acknowledge helpful correspondence with Weiqua Zhang regarding the conditions in eollapsar jets., The authors acknowledge helpful correspondence with Weiqun Zhang regarding the conditions in collapsar jets.
 This research has been supported bv NASA (NAGS5-8128. and MIT-292701) and the DOE Program for Scientific Discovery through Advanced Computing (SciDAC:; DE-FCO2-01EBRAT1176).," This research has been supported by NASA (NAG5-8128, NAG5-12036, and MIT-292701) and the DOE Program for Scientific Discovery through Advanced Computing (SciDAC; DE-FC02-01ER41176)."
 A portion of this work was performed under the auspices of the U.S. Department of Energy by University of California Lawrence Livermore Laboratory under contract W-7405-ENG-43., A portion of this work was performed under the auspices of the U.S. Department of Energy by University of California Lawrence Livermore Laboratory under contract W-7405-ENG-48.
(TF) relation and the galaxy stellar mass function (sec.ee. 22).,"(TF) relation and the galaxy stellar mass function \citep[see, e.g.,][]{Croton_etal06,Guo_etal09}."
 According to the barvonic TF relation of ?.. the mass of the bhuninous component of galaxies with rotation speeds of order 120kins+ is AL;~13ς LOMAL...," According to the baryonic TF relation of \citet{Noordermeer_Verheijen07}, the mass of the luminous component of galaxies with rotation speeds of order $120 \kms$ is $M_d \sim 1.3
\times 10^{10} M_{\odot}$ ."
 We aodel these galaxies as exponential disks with three different values of the radial scale leneth. Ry= 28.8.0. and 22.3 kpc. (," We model these galaxies as exponential disks with three different values of the radial scale length, $R_d=2.8, 8.0$ , and $22.3$ kpc. ("
For the disk vertical structure we asstune an isothermal sheet with scale height z;=0.2 Ry. consistent with observations of external galaxies.),"For the disk vertical structure we assume an isothermal sheet with scale height $z_d=0.2\,R_d$ , consistent with observations of external galaxies.)"
 These values span the range of observed scale leneths for galaxies of this mass. frou: the very exteuded low surface brightuess (LSB) disks to intermediate (ISB) to concentrated hieh surface brightuess GISB) “normal” spirals.," These values span the range of observed scale lengths for galaxies of this mass, from the very extended low surface brightness (LSB) disks to intermediate (ISB) to concentrated high surface brightness (HSB) “normal” spirals."
 Within the disk radius the circular speed of the | halo system is roughly the same for all models. respective ofthe surface brightuess/deusity of the disk.," Within the disk radius the circular speed of the $+$ halo system is roughly the same for all models, irrespective of the surface brightness/density of the disk."
 This is iu aerecment with observations. which show that he zero poiut ofthe TF relation is roughly independent of surface brightuess (?)..," This is in agreement with observations, which show that the zero point of the TF relation is roughly independent of surface brightness \citep{Zwaan_etal95}."
 Iu addition to surface brightness. we have also explored he response of the halo to varving the timescale of disk assembly.," In addition to surface brightness, we have also explored the response of the halo to varying the timescale of disk assembly."
 This is accomplished by allowingthe disk nass to erow linearly over three different timescales. Τζ Cr. d Corr. and 10 Cr.," This is accomplished by allowingthe disk mass to grow linearly over three different timescales, $\tau_d$ $0$ Gyr, $1$ Gyr, and $10$ Gyr."
 The third choice eusures hat the response of the halo is adiabatic. whereas the first. although unrealistic. provides a useful check of the sensitivity of our results to z;," The third choice ensures that the response of the halo is adiabatic, whereas the first, although unrealistic, provides a useful check of the sensitivity of our results to $\tau_d$."
 We have also varied the uode of disk assembly by simulating cases where the mass of the disk grows in three discrete steps over 1 Cyr. instead of linearly.," We have also varied the mode of disk assembly by simulating cases where the mass of the disk grows in three discrete steps over $1$ Gyr, instead of linearly."
" None of the results we obtain depend 1 either the choices for 7, or the mode of disk assembly (see also. 2)). so we shall mainly concentrate on the ""adiabatic (7;=10 Cr) experiments where the iiss of the disk grows linearly with time."," None of the results we obtain depend on either the choices for $\tau_d$ or the mode of disk assembly (see also \citealt{Berentzen_Shlosman06}) ), so we shall mainly concentrate on the “adiabatic” $\tau_d=10$ Gyr) experiments where the mass of the disk grows linearly with time."
 Because the precise aliguiment between the disk aneular imonenutuni axis and the host halo principal axes ds still a matter of debate (0.9.777). the final xuwanmeter we investigated is the oricutation of the disk ane relative to the principal axes of the triaxial halo.," Because the precise alignment between the disk angular momentum axis and the host halo principal axes is still a matter of debate \citep[e.g.,][]{Faltenbacher_etal05,Zentner_etal05,Bailin_etal05}, the final parameter we investigated is the orientation of the disk plane relative to the principal axes of the triaxial halo."
 For cach halo. we placed the disk plane perpendicular ο its mmajor as well as its minor axis.," For each halo, we placed the disk plane perpendicular to its major as well as its minor axis."
 In total. we verformed 21 simulations of the erowth of a ceutral disk ealaxv in each of the triaxial halo models A aud D. ‘or a total of [8 experiments.," In total, we performed $24$ simulations of the growth of a central disk galaxy in each of the triaxial halo models A and B, for a total of $48$ experiments."
 After the disk growth is complete. all simulations are coutiuued for at least a further 1.5 Gar in order to allow for equilibrimu fc ve reached.," After the disk growth is complete, all simulations are continued for at least a further $1.5$ Gyr in order to allow for equilibrium to be reached."
 Iu the experiments where the disk grows im liree distinet events. one-third of the disk mass jucreascs instantaneously at 0.1. 0.6. and 1.1 Cir.," In the experiments where the disk grows in three distinct events, one-third of the disk mass increases instantaneously at $0.1$, $0.6$, and $1.1$ Gyr."
 We ouly consider he final equilibriunu coufguratious of the halos when ucasurue the shape of the poteutial., We only consider the final equilibrium configurations of the halos when measuring the shape of the potential.
 All ummevical simulations were carried out with the xwallel N-body code PIDGRAV. (?).., All numerical simulations were carried out with the parallel $N$ -body code PKDGRAV \citep{Stadel01}. .
 The softening was set to e=0.5 kpc., The softening was set to $\epsilon=0.5$ kpc.
 Each disk was modeled with N=10° static particles iu order to ensure a smooth represcutation of the disk potential., Each disk was modeled with $N=10^5$ static particles in order to ensure a smooth representation of the disk potential.
 (Balbus2001).., \citep{b01}.
restored with Algorithm 2 among the 1000 sources of the simulation.,restored with Algorithm \ref{alg4} among the $1000$ sources of the simulation.
" On noisy data, many fluctuations due to Poisson noise are detected as sources by SExtractor, which leads to a big number of false detections (more than 2000 in the case of Fermi data)."," On noisy data, many fluctuations due to Poisson noise are detected as sources by SExtractor, which leads to a big number of false detections (more than 2000 in the case of Fermi data)."
" As it is difficult to model the background precisely, it is important to study the sensitivity of the method to model errors."," As it is difficult to model the background precisely, it is important to study the sensitivity of the method to model errors."
" We add a stationary Gaussian noise to the background model, we compute the MS-VSTS + IUWT with threshold 3e; on the simulated Fermi Poisson data with extraction of the noisy background, and we study the percent of true and false detections with respect to the total number of sources of the simulation and the signal-noise ratio (SNR(dB)=20log(dsignat/Onocise)) versus the standard deviation of the Gaussian perturbation."," We add a stationary Gaussian noise to the background model, we compute the MS-VSTS + IUWT with threshold $3\sigma_j$ on the simulated Fermi Poisson data with extraction of the noisy background, and we study the percent of true and false detections with respect to the total number of sources of the simulation and the signal-noise ratio $\text{SNR} (dB) = 20 \log (\sigma_{signal} / \sigma_{noise})$ ) versus the standard deviation of the Gaussian perturbation."
" Table 3 shows that, when the standard deviation of the noise on the background model becomes of the same range the mean of the Poisson intensity distribution (Amean= 68.764), the number of false detections increases, the number of true detections decreases and the signal noise ratio decreases."," Table \ref{table1} shows that, when the standard deviation of the noise on the background model becomes of the same range the mean of the Poisson intensity distribution $\lambda_{\text{mean}} = 68.764$ ), the number of false detections increases, the number of true detections decreases and the signal noise ratio decreases."
case ICS-NTVG is favored for larger radii of curvature Ry>0.05. while CR-NTVG requires lower values of Ry«0.1.,"case ICS-NTVG is favored for larger radii of curvature ${\cal R}_6>0.05$, while CR-NTVG requires lower values of ${\cal R}_6<0.1$."
 In J86 case. only ICS-NTVG corresponding to Ry<0.1 can develop.," In J86 case, only ICS-NTVG corresponding to ${\cal R}_6<0.1$ can develop."
 If the actual cohesive energies correspond to some intermediate case between ASO! and J86 cases. they will also be associated with ICS-NTWVG.," If the actual cohesive energies correspond to some intermediate case between AS91 and J86 cases, they will also be associated with ICS-NTVG."
 Therefore. we can generally conclude that ICS-NTVG model is apparently [avored in pulsars with drifting subpulses.," Therefore, we can generally conclude that ICS-NTVG model is apparently favored in pulsars with drifting subpulses."
" CR-NTVG is also possible. although it requires a relatively low cohesive energies. as well as an extremely strong (D.,/D, 1) and/or curved (Re~0.01) surface magnetic field at the polar cap."," CR-NTVG is also possible, although it requires a relatively low cohesive energies, as well as an extremely strong $B_s/B_q\sim 1$ ) and/or curved ${\cal R}_6\sim 0.01$ ) surface magnetic field at the polar cap."
 There is a growing evidence (hat the radio emission of pulsars with svstematically drifting subpulses (erazing cuts of the lme-of-sight) or periodic intensity modulations (central cuts of the line-ol-sight) is based on the inner vacuum gap developed just above the polar cap (Deshpande&Rankin1999.2001:VivekanandJoshi1999:GilSendyk2000).," There is a growing evidence that the radio emission of pulsars with systematically drifting subpulses (grazing cuts of the line-of-sight) or periodic intensity modulations (central cuts of the line-of-sight) is based on the inner vacuum gap developed just above the polar cap \citep{dr99,dr01,vj99,gs00}."
.. To overcome (he binding energy problem: Xuetal.(1999.2001). put forward an attractive bul exolic conjecture (hat pulsars showing the drifGne subpulses represent bare polar cap strange sts (BPCSS) rather (han neutron stars.," To overcome the binding energy problem \citet{xqz99,xzq01} put forward an attractive but exotic conjecture that pulsars showing the drifting subpulses represent bare polar cap strange stars (BPCSS) rather than neutron stars."
 However. as demonstrated in (his paper. invoking ihe BPCSS conjecture is not necessary to explain the ciilting subpulse phenomenon.," However, as demonstrated in this paper, invoking the BPCSS conjecture is not necessary to explain the drifting subpulse phenomenon."
" The quasi steady vacuum gap. wilh either curvature radiation or inverse Compton scattering seed photons. ean form im pulsars with £2-DB«0. provided that the actual surface magnetic field ab the polar cap is extremely strong DJ,~LOM G and curved R.<LO em. irrespective of the value of dipolar component measured [rom the pulsar spindown rate."," The quasi steady vacuum gap, with either curvature radiation or inverse Compton scattering seed photons, can form in pulsars with ${\bf\Omega}\cdot{\bf B}<0$, provided that the actual surface magnetic field at the polar cap is extremely strong $B_s\sim 10^{13}$ G and curved ${\cal
R}<10^6$ cm, irrespective of the value of dipolar component measured from the pulsar spindown rate."
 We have used (wo sets of the cohesive (bounding) energies of the surface iron ions: higher values obtained by A891 and about five times lower values obtained by J86., We have used two sets of the cohesive (bounding) energies of the surface iron ions: higher values obtained by AS91 and about five times lower values obtained by J86.
 If the actual cohesive energies are close to J86 values. then only [CS controlled VG can form and CHR. controlled VG. never forms even under the most extreme conditions (see Fig.," If the actual cohesive energies are close to J86 values, then only ICS controlled VG can form and CR controlled VG never forms even under the most extreme conditions (see Fig."
 1)., 1).
" It is worth noting that IC5 discussed in this paper is the so-called ""resonant scattering."," It is worth noting that ICS discussed in this paper is the so-called ""resonant"" scattering."
 The choice of this mode of ICS is justified by the superstrong magnetic fields relevant [or 1 near-(hreshokl regine., The choice of this mode of ICS is justified by the superstrong magnetic fields relevant for the near-threshold regime.
 However. il is clear that the thermal photons also contribute sone pair-producing ganuna-ravs (e.g.Zhangelal.1997).," However, it is clear that the thermal photons also contribute some pair-producing gamma-rays \citep[e.g.][]{zetal97}."
. This process cannot dominate 1 breakdown of the gap. since it requires higher surface temperatures. which may inhibit 1e formation of VG* in the first place.," This process cannot dominate the breakdown of the gap, since it requires higher surface temperatures, which may inhibit the formation of VG in the first place."
 One can only speculate that the dominating thermal ICS mode is associated with the pulse nulling phenomenon., One can only speculate that the dominating thermal ICS mode is associated with the pulse nulling phenomenon.
 These issues require further mjvestigalion., These issues require further investigation.
 The pulsars with drifting subpulses and/or periodic intensity modulation do not seem, The pulsars with drifting subpulses and/or periodic intensity modulation do not seem
The original SOC model could not vield lognormal flux distribution because of its iutrinsic additive processes through avalauche aud diffusion.,The original SOC model could not yield lognormal flux distribution because of its intrinsic additive processes through avalanche and diffusion.
 For cach avalanche or diffusiou process. the energy released is where Ai; is the total amount of mass transported frou the ving/ to /| Lina time step by avalanche aud diffusiou.," For each avalanche or diffusion process, the energy released is where $\Delta m_i$ is the total amount of mass transported from the ring$i$ to $i+1$ in a time step by avalanche and diffusion."
 The enerey produced frou an avalanche iu a cell is stinply added to those of the other cells in the sale time step. and assumed that the energy release is totally radiated locally.," The energy produced from an avalanche in a cell is simply added to those of the other cells in the same time step, and assumed that the energy release is totally radiated locally."
 It is reasonable to think that uot all of the cuerey produced are radiated locally., It is reasonable to think that not all of the energy produced are radiated locally.
 Di other words. some fraction of euergv remains in the disk aud increases the energv content of the disk.," In other words, some fraction of energy remains in the disk and increases the energy content of the disk."
 Such mechanisin has been considered in ADAF model (c.gB.Abramowiczctal. L988)., Such mechanism has been considered in ADAF model \citep[e.g.][]{Abra88}.
. However it was not applied in the framework of SOC inodel., However it was not applied in the framework of SOC model.
 Tu the other avalanche. the part of the stored cuerey will also be racdated along with the currently produced energv.," In the other avalanche, the part of the stored energy will also be radiated along with the currently produced energy."
 This will create a imultiplicative effect needed to vield a lognormal fix distribution., This will create a multiplicative effect needed to yield a lognormal flux distribution.
 So. when a succeeding avalanche occurred in the same tine step. the cuerey cuiitted is larger because the energy couteut of the disk has increased due to the absorption of some energv from the previous avalanche.," So, when a succeeding avalanche occurred in the same time step, the energy emitted is larger because the energy content of the disk has increased due to the absorption of some energy from the previous avalanche."
" Asstune that the enerey produced from an avalanche or diffusion is Let the fraction of the euergv cuutted from au avalanche or diffusion is ας then the emitted energy is The rest. 5; will be stored iu the dix. Tu the next avalanche. still iu the same time step. the energv produced is £o. a part of it will be ciitted. also a part of the stored cherey from the previous avalanche will be emitted: The energy stored in the disk will be: This euergvOo, will influence the amount of energvOo, cuutted iu the uext avalauche."," Assume that the energy produced from an avalanche or diffusion is Let the fraction of the energy emitted from an avalanche or diffusion is $a$, then the emitted energy is The rest, $S_j$ will be stored in the disk, In the next avalanche, still in the same time step, the energy produced is $E_2$, a part of it will be emitted, also a part of the stored energy from the previous avalanche will be emitted: The energy stored in the disk will be: This energy will influence the amount of energy emitted in the next avalanche."
 Then. iu general the energv cuuitted after the wth avalanche. will be: Aud the stored cnerey will be: A black hole has a very stroug eravity. even light cau rot escape from its eravitational pull. so that eventually here will be a certain amount of euerey produced iu the avalanches which will be swallowed by the black hole.," Then, in general, the energy emitted after the $n$ -th avalanche, will be: And the stored energy will be: A black hole has a very strong gravity, even light can not escape from its gravitational pull, so that eventually there will be a certain amount of energy produced in the avalanches which will be swallowed by the black hole."
 Iu this model it is asstimed that the amount of energv hat falls iuto the black hole is equal o the amount of euergy stored in the disk after the last avalauche in he same time step., In this model it is assumed that the amount of energy that falls into the black hole is equal to the amount of energy stored in the disk after the last avalanche in the same time step.
 Such iiechauismi occur in the slim disk model. iu which the energy released by avalanche. ciffision or viscous flow is stored as radiatiou eutropy and transported inward with accretion (Ilxatoetal.2008).," Such mechanism occur in the slim disk model, in which the energy released by avalanche, diffusion or viscous flow is stored as radiation entropy and transported inward with accretion \citep{Kato08}."
. The balancing between the remiainiug energy in the disk aud the euergv absorption by black hole may sound artificial. but it is necessary to keep energy Dalance in the disk so that the disk always in a stable critical state.," The balancing between the remaining energy in the disk and the energy absorption by black hole may sound artificial, but it is necessary to keep energy balance in the disk so that the disk always in a stable critical state."
 We did a nmuuerical simulation based on cellular automaton nechanismn which was previously done by Iunjava.Mineshige&Ivezió(2005)., We did a numerical simulation based on cellular automaton mechanism which was previously done by \cite{Kunjaya05}.
.. The main aiu of the simulation is to test whether it is possible to generate lognormal flux distribution., The main aim of the simulation is to test whether it is possible to generate lognormal flux distribution.
 Therefore we chose a set of fixed values of the previous parameters. that is in. the amount of each avalanche i. the amount of diffusion aud i”. the amount of raudoi mass input in the outermost ring of the accretion disk.," Therefore we chose a set of fixed values of the previous parameters, that is $m$, the amount of each avalanche $m'$, the amount of diffusion and $m''$, the amount of random mass input in the outermost ring of the accretion disk."
 Each time the avalanche or diffusion occurs. the amount of enerey produced are calculated using equation (5).," Each time the avalanche or diffusion occurs, the amount of energy produced are calculated using equation (5)."
 The accummlated energv to be enmütted after a series of n avalanche in a time step is calculated usiug equation (11) aud the accunulated energy stored iu the disk is calculated using equation (12)., The accumulated energy to be emitted after a series of n avalanche in a time step is calculated using equation (11) and the accumulated energy stored in the disk is calculated using equation (12).
 Iu the simulation. we tried several values of a. 0.01<< 0.8.," In the simulation, we tried several values of $a$, $0.01 \leq a \leq 0.8$ ."
 The system with ο<0.5 means most of the dissipated energy from cach avalanche or diffusion are rot directly ciuitted but is stored iu the disk aud to )o emitted part by part in the avalanches or diffusious afterwards at smaller radii., The system with $a<0.5$ means most of the dissipated energy from each avalanche or diffusion are not directly emitted but is stored in the disk and to be emitted part by part in the avalanches or diffusions afterwards at smaller radii.
 This settiug is simular ) hose of advection cominated accretion flow model (Ikxatoetal. 2008).., This setting is similar to those of advection dominated accretion flow model \citep{Kato08}. .
 We ran the snuulatiou for 10° iterations., We ran the simulation for $10^5$ iterations.
 The first ew thousauds caleulatious were not used in the analysis, The first few thousands calculations were not used in the analysis
"are often referred to as ""disk lockine™ models (Choi&Herbst 1996).",are often referred to as “disk locking” models \citep{choiherbst96}.
. The predicted equilibrium spin rate for both types of iiodols can be expressed as where C is a dineusionless constant that takes into account various assuniptionus iu the models. aud W is the sale as defined in equation (13)).," The predicted equilibrium spin rate for both types of models can be expressed as where $C$ is a dimensionless constant that takes into account various assumptions in the models, and $\Psi$ is the same as defined in equation \ref{eq_psi}) )."
 The various disk-locking models in the literature differ in their assumptions. which is reflected in a different value of C.," The various disk-locking models in the literature differ in their assumptions, which is reflected in a different value of $C$ ."
 The star-disk torque model of Matt&Pudritz(2005b) (used iu the present paper) is a GL-tvpo model hat includes the effect of magnetic field openiug (via the waralucter 5.) aud magnetic coupling streneth (via the xumeter 3)., The star-disk torque model of \citet{mattpudritz05} (used in the present paper) is a GL-type model that includes the effect of magnetic field opening (via the parameter $\gamma_c$ ) and magnetic coupling strength (via the parameter $\beta$ ).
 That work showed that the € is a functiou of oulv 5. and 2. and different choices of hese can result in C values that span the range of inodels presented in the literature.," That work showed that the $C$ is a function of only $\gamma_c$ and $\beta$, and different choices of these can result in $C$ values that span the range of GL-type models presented in the literature."
" Uzdeuskyctal.(2002) demonstrated that a value 5,=1 (adoped in the xeseut work) is the most appropriate to properly take into account the effect of the opening of maguetic Ποιά ines due to differential twisting between the star and disk.", \citet{uzdensky3ea02} demonstrated that a value $\gamma_c = 1$ (adoped in the present work) is the most appropriate to properly take into account the effect of the opening of magnetic field lines due to differential twisting between the star and disk.
" The dashed lines in Figure 5. show the prediction of GL-type models for >,=1 aud two different values of P (0.0L and 1.0. correspouding to Cz10.1 and Czx 2.19. respectivolv)."," The dashed lines in Figure \ref{fig_feqs} show the prediction of GL-type models for $\gamma_c=1$ and two different values of $\beta$ (0.01 and 1.0, corresponding to $C \approx 10.1$ and $C \approx
2.49$ , respectively)."
 The upperauost line corresponds to he CGL-tvpe model used to compute the maguetic star-disk torques iu the spiu-caleulations of the preseut paper (i.c.. for 9= 0.01).," The upper-most line corresponds to the GL-type model used to compute the magnetic star-disk torques in the spin-calculations of the present paper (i.e., for $\beta = 0.01$ )."
 It is clear that this GL-type model xedicets an equilibrium spin rate that is faster than the APSW iodel. for the eutire plotted rauge of V. aud as ong as X is greater than a few thousauths.," It is clear that this GL-type model predicts an equilibrium spin rate that is faster than the APSW model, for the entire plotted range of $\Psi$, and as long as $\chi$ is greater than a few thousanths."
 This is just another way of demonstrating that. when the mass loss is veh cnough. APSWs are much more effective at spinning down stars than a GL-type star-disk interaction alone.," This is just another way of demonstrating that, when the mass loss is high enough, APSWs are much more effective at spinning down stars than a GL-type star-disk interaction alone."
 The effective value of .} in real systems is not well constrained., The effective value of $\beta$ in real systems is not well constrained.
" In the GL-tvpoe models with 5,.=1. Matt&Pudvritz(2005b) showed that the value of JJ=1 results iu the 1ininmia value for C and thus represeuts the case where the CL-tvpe model torques would be the most effective at spinning the star down."," In the GL-type models with $\gamma_c=1$, \citet{mattpudritz05} showed that the value of $\beta=1$ results in the minimum value for $C$ and thus represents the case where the GL-type model torques would be the most effective at spinning the star down."
 The bottom-most dashed line in Figure Ὁ shows the prediction for this special case., The bottom-most dashed line in Figure \ref{fig_feqs} shows the prediction for this special case.
 It is clear that the APSW inodoel eenerallv predicts a slower spin rate for y~0.1 than the GL-tvpe models with auy of 3., It is clear that the APSW model generally predicts a slower spin rate for $\chi \sim 0.1$ than the GL-type models with any of $\beta$.
 For y~0.01. the APSW model predictions are of a similar magnitude as those of a GL-type model with 3=0.1.," For $\chi \sim 0.01$, the APSW model predictions are of a similar magnitude as those of a GL-type model with $\beta=0.1$."
" While the effective value of ο in real svsteius is πουταπα, it secs likely that the inner disks of voung stars exist in a state of ligh magnetic Reyvuolds umber (2« 1). aud thus that the CGL-tvpoe spin-down torques are relatively weak."," While the effective value of $\beta$ in real systems is uncertain, it seems likely that the inner disks of young stars exist in a state of high magnetic Reynolds number $\beta \ll 1$ ), and thus that the GL-type spin-down torques are relatively weak."
" Tn the N-wind model the aueular momen ds asstuned not to accrete onto the star. but rather is intercepted at the inner edee of the disk (tle ""N-poiut') aud then carried out of the svstem by a wind (the wind”) from the N-point."," In the X-wind model, the angular momenum is assumed not to accrete onto the star, but rather is intercepted at the inner edge of the disk (the “X-point”) and then carried out of the system by a wind (the ``X-wind'') from the X-point."
 The fiducial model of Sliuetal.la) and Ostriker&Shu(1995) supposes that the wind outflow rate from the N-point is ~1/3 times the accretion rate. which results in a predicted equilibrimu spin rate that is given by equation (11)). with a value of Cz(0.971.," The fiducial model of \citet{shuea94}
 and \citet{ostrikershu95} supposes that the wind outflow rate from the X-point is $\sim 1/3$ times the accretion rate, which results in a predicted equilibrium spin rate that is given by equation \ref{eq_feqdl}) ), with a value of $C\approx0.974$."
 The dotted line in Figure 5 shows this xedieted spin rate of the N-wind model., The dotted line in Figure \ref{fig_feqs} shows this predicted spin rate of the X-wind model.
 It is clear that lis ος] also predicts slower spin rates than the CGL-ype models., It is clear that this model also predicts slower spin rates than the GL-type models.
 Both the APSW aud N-wind models eiiploy. a wind orelove augular momentum fom the system., Both the APSW and X-wind models employ a wind to remove angular momentum from the system.
 Thus. it is perhaps not surprising that the predicted spin rate of he APSW iuodel with the same mass outflow rate as he X-wind. \=1/3 (bottom solid line in isquantitatioclysiinilartothepredictionse flex," Thus, it is perhaps not surprising that the predicted spin rate of the APSW model with the same mass outflow rate as the X-wind, $\chi=1/3$ (bottom solid line in \\ref{fig_feqs}) ), is quantitatively similar to the predictions of the X-wind."
" Furthermore, toanbeshouwnthatthe factorC inthe windimnodeldependsupontheinasslossratetothepowero f- 3/Towhichisecryclosctothcdependenecco fthe APSVmodelpredictio: 0.122. for the adopted value of 1zz0.223)."," Furthermore, it can be shown that the factor $C$ in the X-wind model depends upon the mass loss rate to the power of $-3/7$ , which is very close to the dependence of the APSW model prediction on the mass loss rate (in ], $f\propto\chi^{-0.42}$ , for the adopted value of $m\approx0.223$ )."
 Although the predicted equilibriiuni spin rates are quantitatively similar. the physical picture of APSW and N-wind iodels have some important differences.," Although the predicted equilibrium spin rates are quantitatively similar, the physical picture of APSW and X-wind models have some important differences."
 Tn the X-wiud. the wind is magnetically connected to. and removes aneular momentum from. the N-point (not from the star. as in a stellar wind).," In the X-wind, the wind is magnetically connected to, and removes angular momentum from, the X-point (not from the star, as in a stellar wind)."
 The X-wiud nodel is developed under the asstuuption that there is uo net torque on the star aud that the N-point is ocated at the corotation radius., The X-wind model is developed under the assumption that there is no net torque on the star and that the X-point is located at the corotation radius.
 Thus. the balancing of aneular monmoentun is evaluated at the X-poiut (rather han on the star) and this leads to a requirement Or a certain amount of stellar maguetie flux to be xuched. or “trapped.” within the N-point.," Thus, the balancing of angular momentum is evaluated at the X-point (rather than on the star), and this leads to a requirement for a certain amount of stellar magnetic flux to be pinched, or “trapped,” within the X-point."
IIaviug the Hux is trapped at the N-point avoids the large-scale opening of magnetic field (πο is a problem for the CGL-tvpe iodels) aud also produces a geometry du which magueto-ceutzüifusal forces may provide a natural explanation for the acceleration of the wind.,Having the flux is trapped at the X-point avoids the large-scale opening of magnetic field (which is a problem for the GL-type models) and also produces a geometry in which magneto-centrifugal forces may provide a natural explanation for the acceleration of the wind.
 However. he biggest problem for the N-wind model is that he amount of fux trapping required at the N-poiut raises questions of clvnamical stability (e.e..Uzdeuskv 2001). and such a iaeucetic field configuration has rot vet been demonstrated to exist in a dynamical uodel.," However, the biggest problem for the X-wind model is that the amount of flux trapping required at the X-point raises questions of dynamical stability \citep[e.g.,][]{uzdensky04}, and such a magnetic field configuration has not yet been demonstrated to exist in a dynamical model."
 Furthermore. since the model assunues no not orque on the star (ie. if assumes the star to be iu spin equilibrium). the model provides little information about how angular moneutun may flow from accreting uaterial outward to the N-point. and accross the N-point. iu order to be removed iu the N-awind aud/or via transport processes in the disk (Ferreiractal.2000)..," Furthermore, since the model assumes no net torque on the star (i.e., it assumes the star to be in spin equilibrium), the model provides little information about how angular momentum may flow from accreting material outward to the X-point, and accross the X-point, in order to be removed in the X-wind and/or via transport processes in the disk \citep[][]{Ferreira:2000p1706}."
" The assunptioun of equilibrium also means that the model cannot address how these svstenis evolve (e.g... one cannot make plots such as panels [c] iu refiie,0 1). oorehetherorhowasgstemwillrcachorapproaimeate "," The assumption of equilibrium also means that the model cannot address how these systems evolve (e.g., one cannot make plots such as panels [c] in \\ref{fig_b0}- \ref{fig_01b2000}) ), nor whether or how a system will reach or approximate the equilibrium state and magnetic field configuration described."
By contrast. the APSW model calculates the spiu-up and spin-down torques acting on the star from various processes (1.0.. at differcut ecoegraphicallocations ou the star).," By contrast, the APSW model calculates the spin-up and spin-down torques acting on the star from various processes (i.e., at different geographicallocations on the star)."
 The theoretical equilibrium state correspouds to when the net torque is zero. but the uodel cau be used to compute the net torque in any state.," The theoretical equilibrium state corresponds to when the net torque is zero, but the model can be used to compute the net torque in any state."
 The magnetic field configuration is that of the stellar magnetic field. which has a significant amount of flux that isopened by various dynanical processes (e.@..seediscussioninMatt 2005b).," The magnetic field configuration is that of the stellar magnetic field, which has a significant amount of flux that isopened by various dynamical processes \citep[e.g., see discussion in][]{mattpudritz05}."
. Numerical. ανασα simulations," Numerical, dynamical simulations"
compoucut is sufficient to explain the spectral variation.,component is sufficient to explain the spectral variation.
 Then. the addition of a second conrponeut is avoided.," Then, the addition of a second component is avoided."
 At the heeranining of the nieht the brightuess of this quasar was my= 0.01., At the beginning of the night the brightness of this quasar was $m_V=17.55 \pm 0.01$ .
 But asthe frst A data sot was lost (sce PIT). then the second group was used to fit the model," But asthe first $R$ data set was lost (see PII), then the second group was used to fit the model."
" A single non-thermal component fairly fits the data. with ay,=3.0640.19 (Figure 7-9a))."," A single non-thermal component fairly fits the data, with $\alpha_{t_0} = -3.06 \pm 0.49$ (Figure \ref{fig89a}) )."
 This value was used for the initial data of B and V hands o predict the R value., This value was used for the initial data of $B$ and $V$ bands to predict the $R$ value.
 It is possible to explain a wicrovariation event as observed with chanees iu lis non-thermal component (Figure 7-9c)}., It is possible to explain a microvariation event as observed with changes in this non-thermal component (Figure \ref{fig89c}) ).
" The spectral iudex would have iucreased in 9.8<10 3, vetween the first aud the third data set. while ο would have staved constant."," The spectral index would have increased in $9.8 \times 10^{-4}$ , between the first and the third data set, while $n_n$ would have stayed constant."
" The sinultaucity criterion indicates that observations were talen at he same level of brightuess (Sp=0.09. 93. 135, and Sp= 0.02)."," The simultaneity criterion indicates that observations were taken at the same level of brightness $S_B=0.09$, $S_V=0.35$ , and $S_R=0.02$ )."
 The brightuess was at the samc evel as for August 17. n=17.56c0.03.," The brightness was at the same level as for August 17, $m_V=17.56 \pm 0.03$."
 The first set of observations at the beeiuniug of the session cau be fitted by a noun-thenual component with αγ=2.51£0.35 (Figure 7-10a))., The first set of observations at the beginning of the session can be fitted by a non-thermal component with $\alpha_{t_0} = -2.51 \pm 0.35$ (Figure \ref{fig810a}) ).
 The variation can be described by changes in this conrponeut., The variation can be described by changes in this component.
" Iu such a case. the spectral index would have decreased by a factor of «10.? aud the auplitude would have increased by a factor of 1 (ay,~ 1.01) (Figure 7-10c))."," In such a case, the spectral index would have decreased by a factor of $\times10^{-3}$ and the amplitude would have increased by a factor of $1\%$ $n_{n_{t_1}} \sim 1.01$ ) (Figure \ref{fig810c}) )."
 The siunultancitv criterion indicates that the spectral variation is reliable (Sp=0.27. Sy=0.39. aud Sp=0.15).," The simultaneity criterion indicates that the spectral variation is reliable $S_B=0.27$, $S_V=0.39$, and $S_R=0.15$ )."
 However. although the model predicts that the variation could be observed in the £2 baud. it also predicts that a variatious should be observed also in the FR baud. which is not the case.," However, although the model predicts that the variation could be observed in the $B$ band, it also predicts that a variations should be observed also in the $R$ band, which is not the case."
 The brightuess diving this night. my=17.66 0:01. was a tenth of a magnitude weaker with respect to the values measured in 1999 August.," The brightness during this night, $m_V = 17.66 \pm 0.01$ , was a tenth of a magnitude weaker with respect to the values measured in 1999 August."
" Initial data can be described by a power law. with ay,ιδ] (Figure 7-Llaj})."," Initial data can be described by a power law, with $\alpha_{t_0} = -1.81 \pm 0.04$ (Figure \ref{fig811a}) )."
 The mucrovariability event can be mareinally explained with a variation in the spectral iudex of 0.001 while the amplitude stavs coustant (Figure 7- llc)., The microvariability event can be marginally explained with a variation in the spectral index of $\sim 0.001$ while the amplitude stays constant (Figure \ref{fig811c}) ).
 The simaultaneity criterion establishes that observations have been taken without a clhauee iu brightness (55=0.17. 54= 0.50. and Sp= 0.37): therefore. thecolor change is considered real.," The simultaneity criterion establishes that observations have been taken without a change in brightness $S_B=0.17$, $S_V=0.50$ , and $S_R=0.37$ ); therefore, thecolor change is considered real."
 Oc Although the σωμα emission cau be fitted in all cases by a unique nou-thermal component. iniuauv of them the broad baud spectral variation," 0.5cm Although the continuum emission can be fitted in all cases by a unique non-thermal component, inmany of them the broad band spectral variation"
"model, Wgοςa-19/169-9/1611/2. and this explains the approximately straight inclination of the thick gray scale curve of Fig.","model, $\dot{W_B} \propto \alpha^{-19/16} \beta^{-9/16} M^{1/2}$, and this explains the approximately straight inclination of the thick gray scale curve of Fig."
 3., 3.
" This result is consistent with the recently found empirical correlation between the radio and the (hard) X-ray luminosity, mentioned in Sect."," This result is consistent with the recently found empirical correlation between the radio and the (hard) X-ray luminosity, mentioned in Sect."
" 1, for sources spanning from magnetically active stars to some BHXRBs and radio quiet AGNs)(Laor Behar 2008; see also Soker Vrtilek 2009)."," 1, for sources spanning from magnetically active stars to some BHXRBs and radio quiet AGNs)(Laor Behar 2008; see also Soker Vrtilek 2009)."
" As remarked before, this correlation clearly suggests that the radio emission in some BHXRBs and low-luminous AGNs comes mainly from magnetic activity in the coronae above the accretion disk and is therefore nearly independent of the intrinsic physics of the central source and the accretion disk, just like in the model above (see also de Gouveia Dal Pino 2006; de Gouveia Dal Pino et al."," As remarked before, this correlation clearly suggests that the radio emission in some BHXRBs and low-luminous AGNs comes mainly from magnetic activity in the coronae above the accretion disk and is therefore nearly independent of the intrinsic physics of the central source and the accretion disk, just like in the model above (see also de Gouveia Dal Pino 2006; de Gouveia Dal Pino et al."
 2009)., 2009).
 The correlation found in Fig., The correlation found in Fig.
" 3 does not hold for radio-loud AGNs, possibly because their surroundings are much denser and then ""mask"" the emission due to coronal magnetic activity."," 3 does not hold for radio-loud AGNs, possibly because their surroundings are much denser and then ""mask"" the emission due to coronal magnetic activity."
" In this case, particle re-acceleration behind shocks will probably prevail further out in the jet-launching region and will be the main responsible for the radio emission."," In this case, particle re-acceleration behind shocks will probably prevail further out in the jet-launching region and will be the main responsible for the radio emission."
 Young stellar objects differ in many aspects from microquasars., Young stellar objects differ in many aspects from microquasars.
 For instance they produce thermal rather than relativistic jets and exhibit emission lines from which their physical properties (such as density and temperature) are inferred., For instance they produce thermal rather than relativistic jets and exhibit emission lines from which their physical properties (such as density and temperature) are inferred.
 But they may also exhibit an intense magnetic activity that results in a strong and variable x-ray emission (e.g. Bouvier et al., But they may also exhibit an intense magnetic activity that results in a strong and variable x-ray emission (e.g. Bouvier et al.
" 2007a, b)."," 2007a, b)."
 Observed flares in x-rays are often attributed to magnetic activity in the stellar corona (Feigelson Montmerle 1999)., Observed flares in x-rays are often attributed to magnetic activity in the stellar corona (Feigelson Montmerle 1999).
" However, some COUP (Chandra Orion Ultra-deep Project) sources have revealed strong flares that were related to peculiar gigantic magnetic loops linking the magnetosphere of the central star with the inner region of the accretion disk."," However, some COUP (Chandra Orion Ultra-deep Project) sources have revealed strong flares that were related to peculiar gigantic magnetic loops linking the magnetosphere of the central star with the inner region of the accretion disk."
" It has been argued that this x-ray emission could be due to magnetic reconnection in these gigantic loops (Favata et al.,"," It has been argued that this x-ray emission could be due to magnetic reconnection in these gigantic loops (Favata et al.,"
 2005)., 2005).
" In this section we examine this in more detail, investigating the role of magnetic reconnection events in the inner disc/corona of YSOs to explain the X-ray flares and check if magnetic reconnection may be related to the thermal jets of these sources."," In this section we examine this in more detail, investigating the role of magnetic reconnection events in the inner disc/corona of YSOs to explain the X-ray flares and check if magnetic reconnection may be related to the thermal jets of these sources."
" We note that previous numerical studies of star-disk interactions have shown that differential rotation between the star and the inner disk regions where the stellar magnetosphere is anchored may lead to field lines opening and reconnection, which eventually restores the initial magnetospheric configuration (e.g., Goodson Winglee 1999; Matt, Goodson, Winglee, et al."," We note that previous numerical studies of star-disk interactions have shown that differential rotation between the star and the inner disk regions where the stellar magnetosphere is anchored may lead to field lines opening and reconnection, which eventually restores the initial magnetospheric configuration (e.g., Goodson Winglee 1999; Matt, Goodson, Winglee, et al."
" 2002; Uzdensky, Kónnigl Litwin 2002; Romanova, Ustyugova, Koldoba, et al."," 2002; Uzdensky, Könnigl Litwin 2002; Romanova, Ustyugova, Koldoba, et al."
 2004; von Rekowski Brandenburg 2004; Zanni 2009; see also Alencar 2007 for a review)., 2004; von Rekowski Brandenburg 2004; Zanni 2009; see also Alencar 2007 for a review).
" Magnetospheric reconnection cycles are expected by most numerical models to develop in a few Keplerian periods at the inner disk, and be accompanied by time dependent accretion onto the star and by episodic outflow events as reconnection takes place."," Magnetospheric reconnection cycles are expected by most numerical models to develop in a few Keplerian periods at the inner disk, and be accompanied by time dependent accretion onto the star and by episodic outflow events as reconnection takes place."
" Also, recently several observational results have indicated that the surface magnetic fields of young low-mass stars may be very complex and have high-order multipoles (Valenti Johns-Krull 2004; Jardine, Cameron, Donati, et al."," Also, recently several observational results have indicated that the surface magnetic fields of young low-mass stars may be very complex and have high-order multipoles (Valenti Johns-Krull 2004; Jardine, Cameron, Donati, et al."
" 2006; Daou, Johns-Krull Valenti 2006; Donati, Jardine, Petit, et al."," 2006; Daou, Johns-Krull Valenti 2006; Donati, Jardine, Petit, et al."
 2007)., 2007).
 Accretion may then eventually proceed through multipole field lines if the interaction region between star and disk is not very far away from the stellar surface for only the dipole component to be important., Accretion may then eventually proceed through multipole field lines if the interaction region between star and disk is not very far away from the stellar surface for only the dipole component to be important.
" This is more likely to happen in stars that present high accretion rates (see below), because them the inner gas disk will extend closer to the star than in the low mass accretion rate systems."," This is more likely to happen in stars that present high accretion rates (see below), because them the inner gas disk will extend closer to the star than in the low mass accretion rate systems."
 Multipolar axisymmetric fields were recently modeled by von Rekowski Brandenburg (2006)., Multipolar axisymmetric fields were recently modeled by von Rekowski Brandenburg (2006).
 Their time-dependent simulations include a dynamo generated field in the star and the disk and results in a very complex and variable field configuration., Their time-dependent simulations include a dynamo generated field in the star and the disk and results in a very complex and variable field configuration.
 In their simulations accretion tends to be irregular and episodic., In their simulations accretion tends to be irregular and episodic.
 To analytically estimate the amount of magnetic power that may be released by violent reconnection in the inner disk-star region we will consider here a most simple geometrical configuration to the problem as in Fig., To analytically estimate the amount of magnetic power that may be released by violent reconnection in the inner disk-star region we will consider here a most simple geometrical configuration to the problem as in Fig.
 1., 1.
" However, while the inner disk regions of microquasars are dominated by radiation pressure, the inner disk region of YSOs are dominated by gas pressure."," However, while the inner disk regions of microquasars are dominated by radiation pressure, the inner disk region of YSOs are dominated by gas pressure."
" Also, instead of modeling the disk/corona of these systems, we used mean values of the coronal density and"," Also, instead of modeling the disk/corona of these systems, we used mean values of the coronal density and"
201 on all eleneuts of he original matrix.,depend on all elements of the original matrix.
 We have implemented a serial as well as a parallel version the aleorithiu., We have implemented a serial as well as a parallel version of the algorithm.
 In the parallel version the preconditioner is applied separately in cach subdomain., In the parallel version the preconditioner is applied separately in each subdomain.
" FDIPS ds an accurate (down to jo1"" relative error) solution on a 180"" erid iu less than 1000 iterations.", FDIPS finds an accurate (down to $10^{-10}$ relative error) solution on a $180^3$ grid in less than 1000 iterations.
 Even mune serially. tjs fakes ess than au hour ou todays computers.," Even running serially, this takes less than an hour on today's computers."
 ~)uce the solution is found in terms of the discrete potential 0;4. we apply the original boundary xditious mclucdi1ο (23) ) audi calculae the magnetic field with equatiou(21)).," Once the solution is found in terms of the discrete potential $\Phi_{i,j,k}$, we apply the original boundary conditions including \ref{eq:innerbcnum}) ) and calculate the magnetic field with \ref{eq:gradient}) )."
 The divergence of the magnetic field will he zc voto the accuracy of t16 Poisson solver., The divergence of the magnetic field will be zero to the accuracy of the Poisson solver.
 The curl of the magnetic field will be zero iu a finite differLUCE sense SImply because it is constructed as the discrete eradieut of the potential., The curl of the magnetic field will be zero in a finite difference sense simply because it is constructed as the discrete gradient of the potential.
" The boundary COxdiion at the 1ner boundary is also satisfied exactly: D,4/5]=Me"," The boundary condition at the inner boundary is also satisfied exactly: $B_{r,i=1/2,j,k}=M'_{j,k}$."
" Averaging (rD,, aud rBy to the rypie=RB position at the outer botudary also eives exactlv zero tanecutial fields due to the equatious (25)) and (21)).", Averaging $r B_{\phi}$ and $r B_{\theta}$ to the $r_{N_r+1/2} = R$ position at the outer boundary also gives exactly zero tangential fields due to the equations \ref{eq:outerbcnum}) ) and \ref{eq:gradient}) ).
 DeOTLling on the applicaion. we max interpolate the poteutial maguetic field onto a collocated eyid. or use it on he original stagecred eril.," Depending on the application, we may interpolate the potential magnetic field onto a collocated grid, or use it on the original staggered grid."
 Figure 6 shows the solution of lie magnetic field obtained with the fuite differcuce solver FDIPS ou a 150«360 eyid., Figure \ref{fig:fdips} shows the solution of the magnetic field obtained with the finite difference solver FDIPS on a $150\times 360\times 180$ grid.
 Since we use tje samme uniforni-cosÓ erid as the magnuetogram. the obtained radial field is identical with the maguectoeramOoC» at kr= 1.," Since we use the same $\cos\theta$ grid as the magnetogram, the obtained radial field is identical with the magnetogram at $r=1$ ."
 The taugeutialC» conipoueuts agreeC» well with the remeshed spherical hiiinonics solution shown in Figure 5., The tangential components agree well with the remeshed spherical harmonics solution shown in Figure \ref{fig:remeshed_fft}.
" It took 1166 iterations to eet a solution with a relative accuracy of 101"".", It took 1166 iterations to get a solution with a relative accuracy of $10^{-10}$.
 The run time was amost exactly one hour on a 2.66 GIIz Iutel CPU., The run time was almost exactly one hour on a 2.66 GHz Intel CPU.
 We have discussed various wavs ο obtain t16 potential field solution based ou solar 10agnueogramis., We have discussed various ways to obtain the potential field solution based on solar magnetograms.
 While spherical harmouics provide an efficieut and elegant method. there are some subtle restrictious that require attention.," While spherical harmonics provide an efficient and elegant method, there are some subtle restrictions that require attention."
 If one wants to use many sprerical harmonics (the same order as the umber ot mnaenetograni jxels in the colatitude direction). je Inaegnetoeram data on the Ngο exid has to be POMCshed onto a uniform exid with AS«NS poiuts. NT unist be an odd umber. aud the new erid must mcelude both voles.," If one wants to use many spherical harmonics (the same order as the number of magnetogram pixels in the colatitude direction), the magnetogram data on the $N_{\theta}\times N_{\phi}$ grid has to be remeshed onto a $\theta$ grid with $N'_{\theta}\times N_{\phi}$ points, $N'_{\theta}$ must be an odd number, and the new grid must include both poles."
 After the reΠΟΣΟ the maximaun deeree of harmonies Nis only limited by the aul-ali:ws lint to πω3./3. ONEBj.," After the remeshing the maximum degree of harmonics $N$ is only limited by the anti-alias limit to $\min(2N'_{\theta}/3$, $N_{\phi}/3)$."
 We used a simple linear iuerpolation for the remeshine., We used a simple linear interpolation for the remeshing.
 T1ο remeshing can be avoked »* he use of a 3D finite differeuce scheme., The remeshing can be avoided by the use of a 3D finite difference scheme.
 ο)1C Call 1se the original naguctoeram erid. aud the ouly freeοι is iu Choosing the radial discretization.," One can use the original magnetogram grid, and the only freedom is in choosing the radial discretization."
 The finite difference scheme oovides a solution that is fully coupatil ewith the boundary conditions. aud the soluion has zero cliverecuce aud curl in the fiuite difference seuse.," The finite difference scheme provides a solution that is fully compatible with the boundary conditions, and the solution has zero divergence and curl in the finite difference sense."
" Figure 7 compares the sohtions obtained with the three methods along the raclal cürection or a fixed latitude 0=76.5"" and longitude o=307.", Figure \ref{fig:radial} compares the solutions obtained with the three methods along the radial direction for a fixed latitude $\theta=76.5\degr$ and longitude $\phi=30\degr$.
 The spjerical haruionics series were trilica edat N=120 for both the naive aud remeshed metlocs., The spherical harmonics series were truncated at $N=120$ for both the naive and remeshed methods.
 The naive spherical harmonics aleorithin eives dacorrect resuls close to the soar surface where the high order harmonics dominate., The naive spherical harmonics algorithm gives incorrect results close to the solar surface where the high order harmonics dominate.
 This is most ovios for the racial οςn1üponoent. which is eiveu at r=lL by the uaenetoeran. and it is exactly reproduce by he finite differeuc' sehen.," This is most obvious for the radial component, which is given at $r=1$ by the magnetogram, and it is exactly reproduced by the finite difference scheme."
 The latitudinal comiponent at r=1 is also very different from the values eiveu by the remeshed iumionies and the finite differences., The latitudinal component at $r=1$ is also very different from the values given by the remeshed harmonics and the finite differences.
 The latter two methods agree reasonably well wit leac1r other., The latter two methods agree reasonably well with each other.
 For radia listaces above r—1.05 all three iiethods agree quite weLB, For radial distances above $r=1.05$ all three methods agree quite well.
 So far we restricted our cxaiple to à CONC magnuctogram taken at t1C SOzar nmnuunnmuu., So far we restricted our example to a GONG magnetogram taken at the solar minimum.
 If o10 SCs all MDI inaenetoeraim during solar maxiuuu. the largest magnetic fields are ich stronger (order of CC) and the resolution of the maenetogram is mach fuer (order of 1000 pixels).," If one uses an MDI magnetogram during solar maximum, the largest magnetic fields are much stronger (order of G) and the resolution of the magnetogram is much finer (order of 1000 pixels)."
 The finer maeuctoeramresolution, The finer magnetogramresolution
radius ry>à will have an infall velocity of The orbital velocity of the companion is ve»=(GM./a)*?.,radius $r_0>a$ will have an infall velocity of The orbital velocity of the companion is $v_2=(G M_\ast/{a})^{1/2}$.
" The relative velocity of the inflowing. gas and the companion. is. given. by (ez,»=e$+0jth>cort2CAL,/a. where we assume that ry>1.5e."," The relative velocity of the inflowing gas and the companion is given by $v_{\rm rel}^2=v_2^2+v_{\rm in}^2 \simeq {2G M_\ast}/{a}$, where we assume that $r_0 >1.5 a$."
 The accretion radius of the companion is the Dondi-Hovle-Lyttletoi one where M» is the mass of the companion., The accretion radius of the companion is the Bondi-Hoyle-Lyttleton one where $M_2$ is the mass of the companion.
 The mass with an impact parameter of b<Ry is accreted by the companion., The mass with an impact parameter of $b \la R_{a2}$ is accreted by the companion.
 Mass at larger distances applies drag force on the secondary. with a magnitude of (Alexander et al.," Mass at larger distances applies drag force on the secondary, with a magnitude of (Alexander et al."
 1976) The same force is applied by the star on the gas., 1976) The same force is applied by the star on the gas.
 The azimuthal component of this force is fy=Feo/t.y., The azimuthal component of this force is $F_{\theta} = Fv_2/v_{\rm rel}$.
" We also take (he maximum radius of influence to be Rua,ca. which by equation (13)) eives Ες{μις{το}cIn(M,/Ma)1. for Ma2OAL,0.181..."," We also take the maximum radius of influence to be $R_{\rm max}\simeq a$ , which by equation \ref{racc1}) ) gives $\ln (R_{\rm max}/R_{a2}) \simeq \ln(M_\ast/{M_2}) \simeq 1$, for $M_2 \simeq 0.1 M_\ast \simeq 0.1 M_\odot$."
" The angular momentum imparted to the inflowing gas by (his gravitational interaction ls The total inflow rate through radius à is Mi,=4ampe;,,.", The angular momentum imparted to the inflowing gas by this gravitational interaction is The total inflow rate through radius $a$ is $\dot M_{\rm in} = 4 \pi a^2 \rho v_{in}$.
 The specific angular momentum of the inflowing gas is j=(dJ/dl)/Mi., The specific angular momentum of the inflowing gas is $j=({dJ}/{dt})/\dot M_{\rm in}$.
" Using our approximations. the value of j is To examine the possibility of disk formation around the central star. jj, should be compared with (he specific angular momentum of a Weplerian circular orbit on the star equator 4,=(GAM,R,)7."," Using our approximations, the value of $j$ is To examine the possibility of disk formation around the central star, $j_{\rm in}$ should be compared with the specific angular momentum of a Keplerian circular orbit on the star equator $j_\ast=(G M_\ast R_\ast)^{1/2}$."
 We lind. alter using the expression [or re. where in the second equality we have substituted values appropriate for a central star ofa PN in ourstudied case of a wide low mass companion.," We find, after using the expression for $v_2$, where in the second equality we have substituted values appropriate for a central star of a PN in ourstudied case of a wide low mass companion."
The original population svnthesis models of W94 cover an aec-metallicity parameter space of 1.5 <rx 17 Gye and (0.5 x Fe/l] 0.5.,The original population synthesis models of W94 cover an age-metallicity parameter space of 1.5 $\leq \tau \leq$ 17 Gyr and –0.5 $\leq$ [Fe/H] $\leq$ 0.5.
 ‘Lo encompass old. metal-poor populations. these models were subsequently extended: to bracket 2.0 € Μο < 0.5 for ages S <τς 1T Gyr. calibrated. using Galactic GCs.," To encompass old, metal-poor populations, these models were subsequently extended to bracket –2.0 $\leq$ [Fe/H] $\leq$ –0.5 for ages 8 $\leq \tau \leq$ 17 Gyr, calibrated using Galactic GCs."
 The W94 models do not cover the voung. metal-poor characteristics of the majority of the LMC clusters. making them of value only for the old GCs in this present. study.," The W94 models do not cover the young, metal-poor characteristics of the majority of the LMC clusters, making them of value only for the old GCs in this present study."
 However. since the W94 moclels predict. line-strength indices for the full range of Lick/IDS fitting functions. they are invaluable for checking the calibration of these indices.," However, since the W94 models predict line-strength indices for the full range of Lick/IDS fitting functions, they are invaluable for checking the calibration of these indices."
 Prior to obtaining age and metallicity estimates for the LAIC clusters using the SSP models. it is important to ensure that we have been able to correct our data onto the Lick/1DS system.," Prior to obtaining age and metallicity estimates for the LMC clusters using the SSP models, it is important to ensure that we have been able to correct our data onto the Lick/IDS system."
 To this end. we compare different Lick/LDS indices of the elusters which measure similar chemical species (or are at least influenced by these same elements 1995).," To this end, we compare different Lick/IDS indices of the clusters which measure similar chemical species (or are at least influenced by these same elements 1995)."
 δν plotting these indices against cach other onto SSP erids. which are effectively. degenerate in age ancl metallicity. we can look for evidence of any systematic olfsets in these data 1998).," By plotting these indices against each other onto SSP grids, which are effectively degenerate in age and metallicity, we can look for evidence of any systematic offsets in these data 1998)."
" In Figure ο we show index-index plots of the Mgbó. Meg, and Meg» indices of the clusters. compared. to the WO94 models."," In Figure \ref{fig:mg.indices} we show index-index plots of the Mg, $_1$ and $_2$ indices of the clusters, compared to the W94 models."
 We find that. despite the apparent small range of metallicity covered. by the clusters due to the ‘squashing of the model erids. the agreement between the models and these cata is good.," We find that, despite the apparent small range of metallicity covered by the clusters due to the 'squashing' of the model grids, the agreement between the models and these data is good."
 The absence of any significant ollsets from the W94 mocdels indicates that our resolution corrections. and the corrections derived from the Lick standard stars ave accurate.," The absence of any significant offsets from the W94 models indicates that our resolution corrections, and the corrections derived from the Lick standard stars are accurate."
 The agreement between the iron indices (Figure 6)) is also generally good., The agreement between the iron indices (Figure \ref{fig:fe.indices}) ) is also generally good.
 The scatter in the figure is somewhat, The scatter in the figure is somewhat
Globular cluster are often the only objects that can be detected in the halos of other ealaxies.,Globular cluster are often the only objects that can be detected in the halos of other galaxies.
 Since most elobular cluster are very old. stuclving (hem provides vital information about the formation and chemical history of these galaxies (Grattonοἱal.2004).," Since most globular cluster are very old, studying them provides vital information about the formation and chemical history of these galaxies \citep{gs04}."
. The similar ages of the constituent stars and low velocity dispersion makes (hese objects relatively easy to study in integrated light., The similar ages of the constituent stars and low velocity dispersion makes these objects relatively easy to study in integrated light.
 The most reliable wav to obtain information about age and abundance patterns of a stellar population is star bv star analvsis via a color-magnitude diagram (CMD)., The most reliable way to obtain information about age and abundance patterns of a stellar population is star by star analysis via a color-magnitude diagram (CMD).
 However (his technique is only feasible for nearby objects. due to the limits of current instrumentation.," However this technique is only feasible for nearby objects, due to the limits of current instrumentation."
 Therefore. we need to develop reliable methods for analvzing (he composite light from all the stars in (hese svstems.," Therefore, we need to develop reliable methods for analyzing the composite light from all the stars in these systems."
 This is no easy feat. since many factors can complicate the analvsis.," This is no easy feat, since many factors can complicate the analysis."
In the right panel of Fig.,In the right panel of Fig.
 4 wevary (ηΠλ, \ref{simovidistvary} wevary $(n/\bar{n})_{\rmn{cut}}$.
 The onset of the correlation moves to larger optical depth., The onset of the correlation moves to larger optical depth.
 Unfortunately. this dependence is weak.," Unfortunately, this dependence is weak."
 The left. panels. of Fig., The left panels of Fig.
 5 show the elfect of changing the simulated noise., \ref{simnoisecontinvary} show the effect of changing the simulated noise.
 As expected. the errors. increase significanthy with increasing noise level.," As expected, the errors increase significantly with increasing noise level."
 The οσοι of an error in the continuum Lit is shown in the right panels of Fig. r5..," The effect of an error in the continuum fit is shown in the right panels of Fig. \ref{simnoisecontinvary},"
 where we have raised. and Iowered the continuum as a whole by in each svnthetic spectrum., where we have raised and lowered the continuum as a whole by in each synthetic spectrum.
 When we raise the continuum level the apparent optical depth increases because the optical depth of spuriously coincident absorption increases., When we raise the continuum level the apparent optical depth increases because the optical depth of spuriously coincident absorption increases.
 Note that if the continuum is placed too low the correlation. between and. optical depth misleadingly appears to extend to smaller optical depth., Note that if the continuum is placed too low the correlation between and optical depth misleadingly appears to extend to smaller optical depth.
 ere we perform a detailed. comparison between four observed. QSO spectra and sets of synthetic spectra., Here we perform a detailed comparison between four observed QSO spectra and sets of synthetic spectra.
 We use the same S/N. Dov. Yap and wavelength. range as," We use the same S/N, $\bar D_{\rmn{O\textsc{VI}}}$ , $\bar D_{\rmn{H\textsc{I}}}$ and wavelength range as"
If one considers high resolution data. there appears to be a quite steep metallicity gradient up to a Galactocentric distance Κως10-II kpc. and then a flattening at larger distances.,"If one considers high resolution data, there appears to be a quite steep metallicity gradient up to a Galactocentric distance $R_{\rm{gc}}\lesssim10-11$ kpc, and then a flattening at larger distances."
" More in detail. using a weighted linear fit. we find a slope of —0.17+0.02 dex kpc! (with a correlation coefficient 20.83) in the disk region with A,.<11 kpe. while in the outer disk (δω>11 kpe) the slope is. within the errors. consistent with zero."," More in detail, using a weighted linear fit, we find a slope of $-0.17\pm0.02$ dex $^{-1}$ (with a correlation coefficient $-0.83$ ) in the disk region with $R_{\rm{gc}}\leq11$ kpc, while in the outer disk $R_{\rm{gc}}>11$ kpc) the slope is, within the errors, consistent with zero."
 The sample by Friel et al. (, The sample by Friel et al. (
2002) 15 instead limited to Ree<16 kpe. with only one cluster in the distance range R«~14—16 kpe: these authors found a negative slope of —0.063x0.010 dex kpe7! between ~7 and 16 kpe. and they did not discuss the possibility of a flattening of the gradient outside of a given radius.,"2002) is instead limited to $R_{\rm{gc}}\lesssim16$ kpc, with only one cluster in the distance range $R_{\rm{gc}}\sim14-16$ kpc: these authors found a negative slope of $-0.063\pm0.010$ dex $^{-1}$ between $\sim$ 7 and 16 kpc, and they did not discuss the possibility of a flattening of the gradient outside of a given radius."
 Considering only the clusters at δω<[1 kpe. the [Fe/H] gradient from low resolution data has a slope of —0.09x0.02 dex kpe7!.," Considering only the clusters at $R_{\rm{gc}}\leq11$ kpc, the [Fe/H] gradient from low resolution data has a slope of $-0.09\pm0.02$ dex $^{-1}$."
 The difference between this value of the slope in the inner disk and that found from high resolution data within the same radius (Ry.€11 κρο) is important. since even small differences in the slope might lead to changes in the infall rate and star formation efficiency adopted in the models.," The difference between this value of the slope in the inner disk and that found from high resolution data within the same radius $R_{\rm{gc}}\leq11$ kpc) is important, since even small differences in the slope might lead to changes in the infall rate and star formation efficiency adopted in the models."
 This point will be discussed in more detail in a future paper of the series. where we plan to compare our empirical results with updated Galactic chemical evolution models.," This point will be discussed in more detail in a future paper of the series, where we plan to compare our empirical results with updated Galactic chemical evolution models."
 On the other hand. considering only high resolution data. our sample and the one from the literature give similar results; in. particular. the possibility of a plateau in the Fe gradient for clusters in the outer disk was already evidenced by Yong et al. (2005))," On the other hand, considering only high resolution data, our sample and the one from the literature give similar results; in particular, the possibility of a plateau in the Fe gradient for clusters in the outer disk was already evidenced by Yong et al. \cite{yong05}) )"
 who studied Be 20 and Be 29 (also in our sample). and two clusters at Ry.~12 kpc. Be 31 and NGC 2141. as well as by Carraro et al. (2004.. 2007)).," who studied Be 20 and Be 29 (also in our sample), and two clusters at $R_{\rm{gc}}\sim12$ kpc, Be 31 and NGC 2141, as well as by Carraro et al. \cite{carraro04}, \cite{carraro07}) )."
 Our clusters significantly improve the statistics. confirming on a more solid basis the change in the slope of the gradient. and in particular of a flattening for radii larger than - 11 kpe: the average metallicity outside of this Galactocentrie distance is —0.27+0.13.," Our clusters significantly improve the statistics, confirming on a more solid basis the change in the slope of the gradient, and in particular of a flattening for radii larger than $\sim$ 11 kpc; the average metallicity outside of this Galactocentric distance is $-0.27\pm0.13$."
" However. we have to remember that the ""inner"" slope is based on a short baseline (about 4 kpe). since the OC closest to the Galactic centre has Rec slightly less than 7 kpe. and that the number or ""outer"" OCs is still scarce."," However, we have to remember that the ""inner"" slope is based on a short baseline (about 4 kpc), since the OC closest to the Galactic centre has $R_{\rm
GC}$ slightly less than 7 kpc, and that the number or ""outer"" OCs is still scarce."
 To improve the situation. we should concentrate on deriving the metallicity of clusters both nearer to the Galactic centre and in the outer part of the disk. and in the transition region between the two slopes.," To improve the situation, we should concentrate on deriving the metallicity of clusters both nearer to the Galactic centre and in the outer part of the disk, and in the transition region between the two slopes."
 The growing number of photometric studies dedicated to overlooked OCs (e.g.. Carraro et al. 2006.. 2007..," The growing number of photometric studies dedicated to overlooked OCs (e.g., Carraro et al. \cite{carraro06}, \cite{carraro07},"
 Ortolani et al. 2005)), Ortolani et al. \cite{ortolani}) )
 and the new catalogues of candidate clusters (e.g.. Froebrich et al. 2007..," and the new catalogues of candidate clusters (e.g., Froebrich et al. \cite{froebrich},"
 Kronberger et al. 2006)), Kronberger et al. \cite{kronberger}) )
 will provide a good database for future abundance studies., will provide a good database for future abundance studies.
 Figure 12. shows [X/Fe] ratios vs. [Fe/H] for the open clusters in our sample and for field disk stars., Figure \ref{confronto_disco} shows [X/Fe] ratios vs. [Fe/H] for the open clusters in our sample and for field disk stars.
 We show the a-elements Si. Ca. Ti and the Fe-peak elements Cr and Ni.," We show the $\alpha$ -elements Si, Ca, Ti and the Fe-peak elements Cr and Ni."
 We considered the disk star catalogue by Soubiran Girard (2005)). which is a compilation and rehomogenization of several recent literature studies. but we also included the data by Bensby et al. (2005)).," We considered the disk star catalogue by Soubiran Girard \cite{soubiran}) ), which is a compilation and rehomogenization of several recent literature studies, but we also included the data by Bensby et al. \cite{bensby}) ),"
 since the latter authors derived also Cr abundances. not included in the other catalogue.," since the latter authors derived also Cr abundances, not included in the other catalogue."
 It is evident that the open clusters follow the same distribution of chemical abundances with [Fe/H] of disk stars. with average values close to (but slightly larger than) solar.," It is evident that the open clusters follow the same distribution of chemical abundances with [Fe/H] of disk stars, with average values close to (but slightly larger than) solar."
 Note that the a-elements have a certam amount of scatter. in particular Ti.," Note that the $\alpha$ -elements have a certain amount of scatter, in particular Ti."
 Figure 13. shows the radial gradients for c-elements. Fe-peak elements and the s-process element Ba. based on our sample and the other high resolution data from the literature.," Figure \ref{grad_altri} shows the radial gradients for $\alpha$ -elements, Fe-peak elements and the s-process element Ba, based on our sample and the other high resolution data from the literature."
 The plots indicate some amount of dispersion at any given Galactocentric distance (this ts especially true for Ba). with no obvious trends. with the only exception of Ca: in this case there seems to be a gradient with a positive slope (0.0820.02. with correlation coefficient0.64) in the inner disk.," The plots indicate some amount of dispersion at any given Galactocentric distance (this is especially true for Ba), with no obvious trends, with the only exception of Ca: in this case there seems to be a gradient with a positive slope $\pm$ 0.02, with correlation coefficient 0.64) in the inner disk."
 The interpretation of the distributions of these elements is rather complex and it is strictly related to the role of SNela and in the Galactic chemical erichment (see. e.g. Yong et al.," The interpretation of the distributions of these elements is rather complex and it is strictly related to the role of a and in the Galactic chemical enrichment (see, e.g. Yong et al."
 2005)., 2005).
 The comparison between our data and Galactic chemical evolution models. planned for a future paper of the series. will be aimed to distinguish the main contribution in element production by different nucleosynthesis process.," The comparison between our data and Galactic chemical evolution models, planned for a future paper of the series, will be aimed to distinguish the main contribution in element production by different nucleosynthesis process."
 Whereas we find nearly solar [X/Fe] (a and Fe-peak elements) in all the clusters. Yong et al. (," Whereas we find nearly solar [X/Fe] $\alpha$ and Fe-peak elements) in all the clusters, Yong et al. ("
2005) and Carraro et al. (,2005) and Carraro et al. (
2004) claim that the outer disk clusters Sau 1. Be 20 and Be 29 (the two latter in common with the present study) have an enhancement in the a-elements.,"2004) claim that the outer disk clusters Sau 1, Be 20 and Be 29 (the two latter in common with the present study) have an enhancement in the $\alpha$ -elements."
 Note. however. that Yong et al. (," Note, however, that Yong et al. ("
2005) derive [Si+Ca/Fe]~+0.1. te. values similar to those we find for some stars. and which are not much higher than solar. and have been also found in field stars of similar metallicity; on the other hand. they derive |Mg+Ti/Fe]~+0.2- 0.3.,"2005) derive $\sim$ +0.1, i.e. values similar to those we find for some stars, and which are not much higher than solar, and have been also found in field stars of similar metallicity; on the other hand, they derive $\sim$ +0.2--0.3."
 Also the values for the a-elements found by Carraro et al. (, Also the values for the $\alpha$ -elements found by Carraro et al. (
2004) are actually only slightly oversolar (average [a/Fe] ratios of less than 0.1 for Be 29 and about 0.15 for Sau 1): anyway. in the most recent analysis of five outer disk clusters Carraro et al. (2007)),"2004) are actually only slightly oversolar (average $\alpha$ /Fe] ratios of less than 0.1 for Be 29 and about 0.15 for Sau 1); anyway, in the most recent analysis of five outer disk clusters Carraro et al. \cite{carraro07}) )"
 do seem to converge on solar-scaled [c /Fe] ratios. similar to open clusters in the solar vicinity and thin disk stars.," do seem to converge on solar-scaled $\alpha$ /Fe] ratios, similar to open clusters in the solar vicinity and thin disk stars."
 With the present data. we do not think there ts a strong argument for a significantly increased [a/Fe] at large Galactocentric distances.," With the present data, we do not think there is a strong argument for a significantly increased $\alpha$ /Fe] at large Galactocentric distances."
 The only s-process element analyzed by us is Ba. which ts clearly enhanced for our open clusters.," The only s-process element analyzed by us is Ba, which is clearly enhanced for our open clusters."
 At the Fe abundances of our clusters. the Ba production is dominated by the s-processes in low mass stars (see e.g. Travaglio et al.," At the Fe abundances of our clusters, the Ba production is dominated by the s-processes in low mass stars (see e.g. Travaglio et al."
 2001. and Cescutti et al., \cite{travaglio01} and Cescutti et al.
 2006 for a complete discussion about this element)., \cite{cescutti06} for a complete discussion about this element).
 The Ba enhancement. mainly in the younger clusters. might be explained by its high production in low mass stars. mainly from | to 2 M.. that were formed from ~ 7 to 2 Gyr ago (e.g. Charbonnel et al. 1996)).," The Ba enhancement, mainly in the younger clusters, might be explained by its high production in low mass stars, mainly from 1 to 2 $_{\odot}$, that were formed from $\sim$ 7 to 2 Gyr ago (e.g. Charbonnel et al. \cite{charbonnel96}) )."
 These stars contribute to the Ba enrichment exactly in the metallicity range covered by the examined clusters. while at lower metallicities the main contribution is from r-processes which are less effective (see also Fig.," These stars contribute to the Ba enrichment exactly in the metallicity range covered by the examined clusters, while at lower metallicities the main contribution is from r-processes which are less effective (see also Fig."
 4 by Cescutti et al. 2006))., 4 by Cescutti et al. \cite{cescutti06}) ).
 However. Ba abundances should be treated with caution. since the hyperfine structure could be important. as a preliminary analysis on available data seems to show (D'Orazi and Randich. in preparation). contrary to what found e.g.. by Mashonkina and Gehren (2001): this would decrease the abundances.," However, Ba abundances should be treated with caution, since the hyperfine structure could be important, as a preliminary analysis on available data seems to show (D'Orazi and Randich, in preparation), contrary to what found e.g., by Mashonkina and Gehren (2001); this would decrease the abundances."
 Furthermore. Ba abundances (ours and literature ones) show a scatter much higher than for any other element here considered. suggesting that we," Furthermore, Ba abundances (ours and literature ones) show a scatter much higher than for any other element here considered, suggesting that we"
Nitrogeuarich: WolfRavet stars (WN) produce very. fast winds (e21000 kins +. which produce strong N-rav cluitting shocks.,"Nitrogen-rich Wolf-Rayet stars (WN) produce very fast winds $v\ga1000$ km $^{-1}$, which produce strong X-ray emitting shocks."
 About of WN stars lave Lx>510 (Wessolowski1996).. so the detectionof the SerA-A WN6 star ds not surprising (Table 3)).," About of WN stars have $L_{\rm X} > 5\times10^{31}$ \citep{wes96}, so the detection of the SgrA-A WN6 star is not surprising (Table \ref{tab:spec}) )."
 The two most hIuninous X-ray sources. aand I2. arethe brightest infrared sources in Table 1..," The two most luminous X-ray sources, and H2, arethe brightest infrared sources in Table \ref{tab:prop}."
 We compared their spectra (Fig. 3)), We compared their spectra (Fig. \ref{fig:spex}) )
 to those of the Pistol star (Figeretal.1998). and to the atlases of Morriset.al.(1996) and Hanson. Conti. Rieke (1996).," to those of the Pistol star \citep{fig98} and to the atlases of \citet{mor96} and Hanson, Conti, Rieke \nocite{han96}."
.. Our low-resolution spectra cannot be used to classify stars with the detail achievable with highiresultion optical spectra., Our low-resolution spectra cannot be used to classify stars with the detail achievable with high-resultion optical spectra.
 However. the line emission from Di-5. He I. and Me TW in Figure 3 resenibles that from the Pistol star aud other Bic] supergiauts. Of stars. and LDVs.," However, the line emission from $\gamma$, He I, and Mg II in Figure \ref{fig:spex}
 resembles that from the Pistol star and other B[e] supergiants, Of stars, and LBVs."
 All of these spectral types represent imassive stars that have evolved off of the main sequence. but not vet reached the WoltBavoet phase. so the similarities are wuderstandable.," All of these spectral types represent massive stars that have evolved off of the main sequence, but not yet reached the Wolf-Rayet phase, so the similarities are understandable."
 For IT2. our results match the conchisions of Fieger(1995) and Coteraetal.(1999).," For H2, our results match the conclusions of \citet{fig95} and \citet{cot99}."
 Our suggestion that these stars are D|e|. Of. or LBV stars suggests that they are unusuallv huuinous.," Our suggestion that these stars are B[e], Of, or LBV stars suggests that they are unusually luminous."
 To coufiim this. we estimate the intrinsic colors aud 1uimosities of aand II2 as follows.," To confirm this, we estimate the intrinsic colors and luminosities of and H2 as follows."
 First. we assume that they de rear the Calactic center (D=s kpc). because (1) the interstellar absorption colununus that we measure frou heir Naav spectra (Table 3)) are quite simular tothe value towardA. Ny=ον107 7. and (2) if we ignore the heavil-absorbed half of the Calaxy nore distant than 10 kpe. ~80% of the Galactic stellar uass along this line-ofsight is located within +200 pc of the Calactic center (Lauuhardtctal.2002).," First, we assume that they lie near the Galactic center $D$$=$ 8 kpc), because (1) the interstellar absorption columns that we measure from their X-ray spectra (Table \ref{tab:spec}) ) are quite similar tothe value toward, $N_{\rm H} = 6\times10^{22}$ $^{-2}$, and (2) if we ignore the heavily-absorbed half of the Galaxy more distant than 10 kpc, $\sim80$ of the Galactic stellar mass along this line-of-sight is located within $\pm$ 200 pc of the Galactic center \citep{lzm02}."
. The distance modulus is therefore11.5., The distance modulus is therefore14.5.
 Second. we estimate he reddening using the relationship Nyyfly=1.79«107!cin7o from Predehl&Sclunitt(1995). and the extinction law from from Moneti (2001)..," Second, we estimate the reddening using the relationship $N_{\rm H}/A_V = 
1.79 \times 10^{21}~{\rm cm}^{-2}$ from \citet{ps95}, and the extinction law from from \citet{mon01}. ."
 For 1516.1.. aking Ng=(LdΕυ)«10°? 7. we find Ay 0.3. and Ay=2.940.2.," For , taking $N_{\rm H} = (4.4\pm0.3)\times10^{22}$ $^{-2}$, we find $A_H = 4.6\pm0.3$ , and $A_K = 2.9 \pm 0.2$."
 Its iutriusie color is (77Koy= OLAOM. and absolute Aj=9.540.2.," Its intrinsic color is $(H - K_s)_0 = -0.4 \pm 0.1$ , and absolute $M_K = -9.5\pm0.2$."
" For II2. Ny=(7.21.0&1022 2, we find Ay=7.54 LL. and “Ay=La+0.9."," For H2, $N_{\rm H} = (7.7\pm1.4)\times10^{22}$ $^{-2}$, we find $A_H = 7.5\pm1.4$ , and $A_K = 4.8 \pm 0.9$."
 Its intrinsic color is 11Wey=0.6 30.5. and Wy=10.1c0.9.," Its intrinsic color is $(H - K_s)_0 = -0.6 \pm 0.5$ , and $M_K = -10.1\pm0.9$."
" In both cases. the colors are within lo of (/T.Ay020,1 expected for O supergiant stars (Woeener199D)."," In both cases, the colors are within $\sigma$ of $(H-K)_0 = -0.3 \pm 0.1$ expected for O supergiant stars \citep{weg94}."
. Towever. the colors are inconsistent with the (/7.—Λου20.8 iufrared excesses seen from stars in the Magellanic clouds 1986).. Ble}," However, the colors are inconsistent with the $(H-K)_0 \ga 0.8$ infrared excesses seen from B[e] stars in the Magellanic clouds \citep{zic86}."
Therefore. we tentatively classify aud I2 aseither Of or candidate LBV stars.," Therefore, we tentatively classify and H2 aseither Of or candidate LBV stars."
 We estimate the bolometric luminosities of these stars by assunue that they have temperatures comparable to those of the Pistol star aud other hot. massive stars iu transition. Tzz1.5«10! K (Morrisefal.1996:Fiecrctal.1998:Clarket 2003).," We estimate the bolometric luminosities of these stars by assuming that they have temperatures comparable to those of the Pistol star and other hot, massive stars in transition, $T \ga 1.5\times10^{4}$ K \citep{mor96,fig98,cla03}."
". The bolometric corrections to the absolute A, magnitudes will then be eiven by BC=rllogf|28.5 (Blum.DePox.&Sellgren 1995).. so that BC: 0.9."," The bolometric corrections to the absolute $K_s$ magnitudes will then be given by $BC_K = -7.1\log T + 28.8$ \citep{blu95}, , so that $BC_K \la -0.9$ ."
 Therefore. we fud upper limits outhe bolometric=magnitudes for oof Mp;S 10.2. aud for II2 of Mp 10.1.," Therefore, we find upper limits onthe bolometricmagnitudes for of $M_{\rm Bol} \la -10.2$ , and for H2 of $M_{\rm Bol} \la 10.1$ ."
 We converted these to dower lits on the luminosities (listed in Table 19). audfind for tthat Lp> LOPE... aud for H2that Lp ου. ," We converted these to limits on the luminosities (listed in Table \ref{tab:lxlbol}) ), andfind for that $L_{\rm Bol} > 10^{6.1} L_\odot$ , and for H2that $L_{\rm Bol} > 10^{5.9} L_\odot$ ."
Their true Iuninositiesmay be significautly higher., Their true luminositiesmay be significantly higher.
 For, For
4.3 (Freemanetal.2001;Fruscione2006).,"4.3 \citep{Freeman01, Fruscione06}."
". We implement Nelder-Mead simplex optimization (Nelder&Mead1965) using Cstat statistics, a modification of the statistics of Cash(1979)."," We implement Nelder-Mead simplex optimization \citep{Nelder65} using Cstat statistics, a modification of the statistics of \citet{Cash79}."
". For sources within the cluster half-light radius that have greater than 200 net counts in the EPIC cameras (X1, X2, X3, X4, and X5) we do model determination between five different XSPEC models: thermal bremsstrahlung, APEC thermal plasma etal. blackbody, neutron star atmosphere (Smith(Pavlovet2001),al.1991;Zavlin1996),, and power law."," For sources within the cluster half-light radius that have greater than 200 net counts in the EPIC cameras (X1, X2, X3, X4, and X5) we do model determination between five different XSPEC models: thermal bremsstrahlung, APEC thermal plasma \citep{Smith01}, , blackbody, neutron star atmosphere \citep{Pavlov91, Zavlin96}, and power law."
" All source models also take into account absorption, including an absorption componentusing the"," All source models also take into account absorption, including an absorption componentusing the"
cofthemostprobabletouchdown flic,of the most probable touchdown flux.
" Dithesiesourceswithonlgoncpitirof PREburstscach, thewilthoftheunderlgingdistributi ."," In the six sources with only one pair of PRE bursts each, the width of the underlying distribution is consistent with being at a similar level."
 ," When the latter group of sources is taken as a representative sample, the most likely fractional width of their touchdown fluxes is $\simeq 11$ ."
The only clear exception is Aql N-1. where the systematic uncertainties exceed ~20%," The only clear exception is Aql X-1, where the systematic uncertainties exceed $\sim 20\%$."
 As we explored in Section 5. the distribuion of the touchdown fluxes is expected to have a fiute width for à uunboer observaional and physical reasons.," As we explored in Section 5, the distribution of the touchdown fluxes is expected to have a finite width for a number of observational and physical reasons."
 For a nuuberoft jese effects; we were able to estimate that they imtroduce a5! fsystematicuncertaintginthe fliwes," For a number of these effects, we were able to estimate that they introduce a $-$ level of systematic uncertainty in the fluxes."
 Thetwounkiownsthatpotentiallgiutrodueclur JCrsuystemati, The two unknowns that potentially introduce larger systematic uncertainties are the asymmetry of the PRE event and the composition of the material at the photosphere.
cuncertaintiesurcthe," Our results show, however, that even these unknowns do not introduce uncertainties larger than."
 The PRE bursts allow us to measure he Eddington luit at the πιrace of the neutroji star for cach source., The PRE bursts allow us to measure the Eddington limit at the surface of the neutron star for each source.
 Doeterumiie the Eddington luit requires a1 absolute fliXx 1ueasuremoenut. which is affected by tli! overall fiux calibration of the Na:w detector used.," Determining the Eddington limit requires an absolute flux measurement, which is affected by the overall flux calibration of the X-ray detector used."
 Such calibratiojs are notorkxlv difficult to achieve aud are usualyv based ou a particular set of asst111tious regarding the spectra and variabiity of the Crab 1cula (Jahoda et 22006: see also Toor Seward 19Tl Ikisch et 22005: WeisskoXf et 220)0., Such calibrations are notoriously difficult to achieve and are usually based on a particular set of assumptions regarding the spectrum and variability of the Crab nebula (Jahoda et 2006; see also Toor Seward 1974; Kirsch et 2005; Weisskopf et 2010).
 Anv bias in fιο absolute fiux calixation cannot mierease the spread « touchdown fluxes that we iler here for eacji source., Any bias in the absolute flux calibration cannot increase the spread of touchdown fluxes that we infer here for each source.
 However. it can affect the mean ouchdown fiux. which. in tur1. enters iuto the measurement of neutron star masses and raci.," However, it can affect the mean touchdown flux, which, in turn, enters into the measurement of neutron star masses and radii."
 We will quantity he potential svstematic ncertainties introduced by the absolute fi1x calibration o PCÀ in Paper TI of this series., We will quantify the potential systematic uncertainties introduced by the absolute flux calibration of PCA in Paper III of this series.
 It is also iuportaut to eniplasize here tjat our results are based on a statistical analysis of the cutive sample of PRE bursts CL source aid do not preclude the possibility that anv one 1idividual burst παν show a rather differeut touchdown flix., It is also important to emphasize here that our results are based on a statistical analysis of the entire sample of PRE bursts per source and do not preclude the possibility that any one individual burst may show a rather different touchdown flux.
 Indeed. here is at least o1ο burst observe| from LU 536 (ID#116)4.. for which the touchdown fux Was uualler compar’ed to the averaec| value bv a facor of L.7 Callavay et ((2006).," Indeed, there is at least one burst observed from 4U $-$ 536 (ID, for which the touchdown flux was smaller compared to the average value by a factor of 1.7 Galloway et (2006)."
 Tn this particular case. a variation 1i the hydroeen mass fractioji from X0 oN 0.7 beween the bursts has been cousidered as a natural expl:tation of fιο difference in touchdown fluxes (Sueinoto et 11981: Galloway ct 22006).," In this particular case, a variation in the hydrogen mass fraction from $X=0$ to $X=0.7$ between the bursts has been considered as a natural explanation of the difference in touchdown fluxes (Sugimoto et 1984; Galloway et 2006)."
 The fact trat such outhers may aid do exist makes it esseutial tha proper statistical tools are used in all iufereuces based ou measurements of the touchdown fluxes of PRE bursts., The fact that such outliers may and do exist makes it essential that proper statistical tools are used in all inferences based on measurements of the touchdown fluxes of PRE bursts.
 Iu conclusion. our results iudicate that the systematic uucertanties in the measurements of touchdown fluxes in radius-expausion bursts youn low-mass N-ray binaries are within ~105. for nearly the eutire source sample.," In conclusion, our results indicate that the systematic uncertainties in the measurements of touchdown fluxes in radius-expansion bursts from low-mass X-ray binaries are within $\simeq 10$, for nearly the entire source sample."
 Such svstenmiatie uncertainties do not preclude. in and of themselves. neutron star mass-radius measurements with high enough precision to distiretish between differcut equations of state of neutron-star matter.," Such systematic uncertainties do not preclude, in and of themselves, neutron star mass-radius measurements with high enough precision to distinguish between different equations of state of neutron-star matter."
with respect to the central line wavelength.,with respect to the central line wavelength.
 This was a clear evidence for a one-armed spiral structure in the Z Tau disk., This was a clear evidence for a one-armed spiral structure in the $\zeta$ Tau disk.
" On the other hand, despite these symmetrical differential phases, our visibilities across these 3 spectral lines are asymmetrical, especially for Bry and lines, with a blue wing of the visibility systematically smaller than the red wing."," On the other hand, despite these symmetrical differential phases, our visibilities across these 3 spectral lines are asymmetrical, especially for $\gamma$ and lines, with a blue wing of the visibility systematically smaller than the red wing."
" The line profiles are also clearly asymmetrical, with Ha and Bry line profiles showing a 4i, lwhereastheS /Nandspectralre solution for iarenothighenovstitdeshPpeestinaths Thus, 6 Sco’s circumstellar disk is globally rotating, with a larger (in size) and brighter emitting region in the blue part of the line, i.e. coming towards us, and a smaller (more compact), fainter or more absorbed region in the red part of the line, ie. rotating away from us."," The line profiles are also clearly asymmetrical, with $\alpha$ and $\gamma$ line profiles showing a $> 1$ whereas the S/N and spectral resolution for are not high enough to exhibit any line Thus, $\delta$ Sco's circumstellar disk is globally rotating, with a larger (in size) and brighter emitting region in the blue part of the line, i.e. coming towards us, and a smaller (more compact), fainter or more absorbed region in the red part of the line, i.e. rotating away from us."
 These emitting regions are also responsible for the Ho and Bry line profiles with 41., These emitting regions are also responsible for the $\alpha$ and $\gamma$ line profiles with $> 1$.
" Moreover, sincethediff erentialphaseissymmetrical, itmeanseZEAEp kpStntekaWledpeirtivagtia impport fymm?iNSWitimas pecttothec ar"," Moreover, since the differential phase is symmetrical, it means that the photocenter of the emitting regions are symmetric with respect to the central star (or the rotational axis)."
medoscillationscenariowoulddif f icultlymatchtheinhomogeneiNbNae9edfedihó Sco's disk since it would have produced non-symmetrical spectrally resolved visibilitiesand phases as measured by Carciofi et al. (," Therefore, the one-armed oscillation scenario would difficultly match the inhomogeneities detected in $\delta$ Sco's disk since it would have produced non-symmetrical spectrally resolved visibilities phases as measured by Carciofi et al. ("
2009) for Z Tau.,2009) for $\zeta$ Tau.
" In our case, the binarity of the system should also play an important role in the shaping and the kinematics of the circumstellar environment."," In our case, the binarity of the system should also play an important role in the shaping and the kinematics of the circumstellar environment."
" For example, one could think of a tidally-warped disk from previous periastron passages of the companion star (like what is described in Moreno et al."," For example, one could think of a tidally-warped disk from previous periastron passages of the companion star (like what is described in Moreno et al."
 2011)., 2011).
 This definitely needs a dedicated modeling effort to really constrain the detected disk inhomogeneities., This definitely needs a dedicated modeling effort to really constrain the detected disk inhomogeneities.
" Using all available binary separation measurments including our VLTI/AMBER ones, we check the consistency of previous estimates of the ó Sco orbital elements."," Using all available binary separation measurments including our VLTI/AMBER ones, we check the consistency of previous estimates of the $\delta$ Sco orbital elements."
 We found parameters that totally agrees with the latest work from Tycner et al. (, We found parameters that totally agrees with the latest work from Tycner et al. (
2011).,2011).
" Thanks to the combination of high spectral and spatial resolution we managed to constrain for the first time, not only the geometry but also the kinematics of the ó Sco circumstellar environment."," Thanks to the combination of high spectral and spatial resolution we managed to constrain for the first time, not only the geometry but also the kinematics of the $\delta$ Sco circumstellar environment."
 As in the case of Be stars with stable or nearly stable disks such as « Ara (Meilland et al., As in the case of Be stars with stable or nearly stable disks such as $\alpha$ Ara (Meilland et al.
" 2007a), κ CMa (Meilland et al."," 2007a), $\kappa$ CMa (Meilland et al."
" 2007b), V Per, and 48 Per (Delaa et al."," 2007b), $\Psi$ Per, and 48 Per (Delaa et al."
" 2011), the line emission originates from a dense equatorial disk dominated by rotation."," 2011), the line emission originates from a dense equatorial disk dominated by rotation."
" As for most of these objects, the rotation appears to be Keplerian, with an inner boundary (photosphere/disk interface) rotating at the critical velocity (Vc)."," As for most of these objects, the rotation appears to be Keplerian, with an inner boundary (photosphere/disk interface) rotating at the critical velocity $\mathrm{c}$ )."
" The expansion is negligible and considering an outburst scenario, should be on the order of ss~!."," The expansion is negligible and considering an outburst scenario, should be on the order of $^{-1}$."
" Considering the measured vsini, i.e. ss! and the measured inclination angle, i.e. 30.2+ 0.7, the star rotates, at about Finally, as the measured extension in the Ha line, +1.5 mas, is on the order of the binary separation at the periastron, i.e. 6.14x0.07 mas (i.e., according to Tycner 2011), we expect strong interaction between the previously ejected matterand the companion."," Considering the measured vsini, i.e. $^{-1}$ , and the measured inclination angle, i.e. $\pm$ $^o$, the star rotates at about Finally, as the measured extension in the $\alpha$ line, $\pm$ 1.5 mas, is on the order of the binary separation at the periastron, i.e. $6.14 \pm 0.07$ mas (i.e., according to Tycner 2011), we expect strong interaction between the previously ejected matterand the companion."
in type 1 objects. as observed.,"in type 1 objects, as observed."
 However. to date the idea that the Fe vij] emission lines are associated with the torus has not been thoroughly tested using detailed emission line ratios that measure physical conditions accurately.," However, to date the idea that the [Fe ] emission lines are associated with the torus has not been thoroughly tested using detailed emission line ratios that measure physical conditions accurately."
 Although it seems. plausible that at least. some of the FULL emission arises in the torus. the causes of the diversity in both the relative strength and. kinematies of the ΕΠΗ: remain uncertain.," Although it seems plausible that at least some of the FHIL emission arises in the torus, the causes of the diversity in both the relative strength and kinematics of the FHILs remain uncertain."
 One possibility is that the unusually strong ΕΠΗ observed. in some objects are. symptomatic ofa recent energetic occurrence which is able to illuminate the region that produces these lines (Ixomossactal.2009): suggestions for such events. include the rapid. accretion of material from the ISM. the tidal disruption of a stars by the supermassive black hole (SMDBLLD. gamma-ray bursts ancl supernovae (Ixomossactal.2009).," One possibility is that the unusually strong FHILs observed in some objects are symptomatic of a recent energetic occurrence which is able to illuminate the region that produces these lines \citep{komossa2}; suggestions for such events include the rapid accretion of material from the ISM, the tidal disruption of a stars by the supermassive black hole (SMBH), gamma-ray bursts and supernovae \citep{komossa2}."
 The unusual strength of the FLULs in objects such as LLL Zw 77. Tololo 0109-383 and ESO 138 Cl provides us with a rare opportunity to study the nature of the region of the AGN which emits them. and thus helps us to better understand the structure of AGN in general.," The unusual strength of the FHILs in objects such as III Zw 77, Tololo 0109-383 and ESO 138 G1 provides us with a rare opportunity to study the nature of the region of the AGN which emits them, and thus helps us to better understand the structure of AGN in general."
 This paper reports an investigation of the object QI131]16. which has the richest spectrum of Ες vet reported. for an AGN.," This paper reports an investigation of the object Q1131+16, which has the richest spectrum of FHILs yet reported for an AGN."
 QUIS1|16 was discovered. in. the 2ALASS survey and was classified as a Seyfert 2 galaxy witha redshift of 0.174., Q1131+16 was discovered in the 2MASS survey and was classified as a Seyfert 2 galaxy with a redshift of 0.174.
 It belongs to sample of quasar-like objects identified on the basis of their red near-Hi colours (J-15 72.0) in the PALASS survey (Cutrietal.1995)., It belongs to sample of quasar-like objects identified on the basis of their red near-IR colours $>$ 2.0) in the 2MASS survey \citep{cutri}.
.. In this paper we present. deep optical and infrared spectra of QLIS1|16. and use these to deduce the physical conditions ancl kinematics of the FHL emission lines. with the aim of investigating their origin.," In this paper we present deep optical and infrared spectra of Q1131+16, and use these to deduce the physical conditions and kinematics of the FHIL emission lines, with the aim of investigating their origin."
 “Phe cosmological parameters used. throughout this paper are adopted from WNLAP: ff. = Tl km | Qj = 0.27 and O4 = 0.73 (Spergelοἱal.2003) resulting in a spatial scale of 3.46 προ aresec," The cosmological parameters used throughout this paper are adopted from WMAP: $H_\circ$ = 71 km $^{-1}$ , $\Omega_M$ = 0.27 and $\Omega_\Lambda$ = 0.73 \citep{spergel}, , resulting in a spatial scale of 3.46 kpc $^{-1}$."
 Low resolution optical spectroscopic observations of QII31I]16 were taken on 9th. February 2007 with the ISIS clual-arm spectrograph on the 4.2-m William Llerschel Telescope (IIT) on La Palma as part of a spectroscopic survey of a representative sample red. quasar-like objects.," Low resolution optical spectroscopic observations of Q1131+16 were taken on 9th February 2007 with the ISIS dual-arm spectrograph on the 4.2-m William Herschel Telescope (WHT) on La Palma as part of a spectroscopic survey of a representative sample red, quasar-like objects."
 The full sample for this survey comprised a complete ItA-limited sub-sample of 24 objects with (I-I572) and 2<0.2 selected from the list of Hutehingsetal.J)2003.. which is itself representative of the population of red. 24LASS-selected quasars.," The full sample for this survey comprised a complete RA-limited sub-sample of 24 objects with $>$ 2) and $z<0.2$ selected from the list of \citealt{hutchings03}, which is itself representative of the population of red, 2MASS-selected quasars."
 The unusually strong ΕΤΗ in QI131116. were discovered serendipitously in the course of this survey., The unusually strong FHIL in Q1131+16 were discovered serendipitously in the course of this survey.
 In the red. the RI5sh. erating was used with the RIEDPLUS CCD. and in the blue. the R300B erating was used. with the EEVI2 CCD.," In the red, the R158R grating was used with the REDPLUS CCD, and in the blue, the R300B grating was used with the EEV12 CCD."
 A dichroic at was cmploved to obtain spectra with useful wavelength ranges ~3250-5250A in the blue. and ~5200-9500A in the red.," A dichroic at was employed to obtain spectra with useful wavelength ranges $\sim$ in the blue, and $\sim$ in the red."
 To reduce the elfects of differential refraction. all— exposures were taken when Q1131|16 was at low airmass (sec z « 1.1) and with the slit aligned close to the parallactic angle.," To reduce the effects of differential refraction, all exposures were taken when Q1131+16 was at low airmass (sec $z$ $<$ 1.1) and with the slit aligned close to the parallactic angle."
 The secing for the night of the observations varied. from 0.8 to 1.3 aresecondsLIADHT., The seeing for the night of the observations varied from 0.8 to 1.3 arcseconds.
. Sets of three 6008 exposures were taken on both armis simultaneously. giving a total exposure time of 1800s.," Sets of three 600s exposures were taken on both arms simultaneously, giving a total exposure time of 1800s."
 The data were taken with a 1.5 aresecond. slit along a position angle P315., The data were taken with a 1.5 arcsecond slit along a position angle PA315.
 To eliminate contamination from second order emission. a G@G495 blocking filter was introduced into the ISIS red arm.," To eliminate contamination from second order emission, a GG495 blocking filter was introduced into the ISIS red arm."
 The data were reduced in the standard wav (bias subtraction. [lat [ielding. cosmic rav removal. wavelength calibration. (lux calibration) using packages inIRALFT.," The data were reduced in the standard way (bias subtraction, flat fielding, cosmic ray removal, wavelength calibration, flux calibration) using packages in."
. The two-dimensional spectra were also corrected. for. spatial distortions of the CCD., The two-dimensional spectra were also corrected for spatial distortions of the CCD.
 To reduce wavelength. calibration errors due to flexure of the telescope and. instrument. arc spectra were taken at the position of the object on the sky: the estimated wavelength calibration accuracy is iin both the blue and red respectively (however. this may be an unclerestimate at the extreme edges of the spectra).," To reduce wavelength calibration errors due to flexure of the telescope and instrument, arc spectra were taken at the position of the object on the sky; the estimated wavelength calibration accuracy is in both the blue and red respectively (however, this may be an underestimate at the extreme edges of the spectra)."
 The atmospheric absorption features were removed. by. dividing by the spectrum of a telluric standard (BD|17 2352). taken close in time and airmass to the observations of QI131|16.," The atmospheric absorption features were removed by dividing by the spectrum of a telluric standard (BD+17 2352), taken close in time and airmass to the observations of Q1131+16."
 The spectral resolution. calculated: using the widths of the night sky emission lines. was 6.11 + in the blue and 10.96 + in the red. (measured in the observed. frame of the spectrum).," The spectral resolution, calculated using the widths of the night sky emission lines, was 6.11 $\pm$ in the blue and 10.96 $\pm$ in the red (measured in the observed frame of the spectrum)."
 The spatial pixel scales of the 2D spectra ave 0.4 arceseconds in the blue and O44 areseconds in the rec. and the relative Dux calibration uncertainty — based on 1l observations of δ [lux standard. stars taken throughout the run. ds estimated to be ρα...," The spatial pixel scales of the 2D spectra are 0.4 arcseconds in the blue and 0.44 arcseconds in the red, and the relative flux calibration uncertainty – based on 11 observations of 8 flux standard stars taken throughout the run – is estimated to be $\pm$."
 The spectra were extracted. ancl analysed using theSTARLINK packages andDIP'SO., The spectra were extracted and analysed using the packages and.
 In order to test the possibility that the emission lines of Q113110 are variable. a follow-up spectrum. was taken using GALOS on the Gemini south telescope at. Cerro Pachon.," In order to test the possibility that the emission lines of Q1131+16 are variable, a follow-up spectrum was taken using GMOS on the Gemini south telescope at Cerro Pachon."
 Spectroscopic and imaging observations using GALOS were taken on the 21st February 2010 during Gemini queue observation time as part of the program CiS-2009D-CO-NT., Spectroscopic and imaging observations using GMOS were taken on the 21st February 2010 during Gemini queue observation time as part of the program GS-2009B-Q-87.
 The spectral observations comprised. three GOO second exposures using a central wavelength aat an airmass of 1.46. and a 1.5 arceseconds slit. aligned along position angle. D'A163 close to the parallactic angle for the centre of the observations.," The spectral observations comprised three 600 second exposures using a central wavelength at an airmass of 1.46, and a 1.5 arcseconds slit aligned along position angle $^{\circ}$ — close to the parallactic angle for the centre of the observations."
 This resulted. in a spectrum with a useful wavelength range 23700-6450A., This resulted in a spectrum with a useful wavelength range $\sim$.
.. ‘The observations of the [ux standard star. Feige56. had the same set-up but with an exposure time of LO seconds. taken at an airmass of 1.61 and with a 5 arcsecond. slit.," The observations of the flux standard star, Feige56, had the same set-up but with an exposure time of 10 seconds, taken at an airmass of 1.61 and with a 5 arcsecond slit."
 Deep CGMOS images were also taken using ane’ filter at an airmass of 1.56., Deep GMOS images were also taken using an r' filter at an airmass of 1.56.
 Each exposure of the imaging observations had an exposure time of250 seconds. and the data were taken in à 4 point dither pattern. resulting ina total exposure time of," Each exposure of the imaging observations had an exposure time of250 seconds, and the data were taken in a 4 point dither pattern, resulting ina total exposure time of"
more evolved aud fainter. brown dwarfs produced by these birth rates.,"more evolved and fainter, brown dwarfs produced by these birth rates."
" The differcuces are most pronounced for the halo age distribution. which predicts a substantial deficiency of Ty,=1200.2000 I& brown dwarts. conrprised primarily of uuid- aud late-tvpe L and carly T dwarts."," The differences are most pronounced for the halo age distribution, which predicts a substantial deficiency of $_{eff} \approx 1200-2000$ K brown dwarfs, comprised primarily of mid- and late-type L and early T dwarfs."
 Note that this deficiency is likely to be more pronounced ina veal halo population. as the reduced imetallicity typical for halo dwarts (Cüzis1997) inuplv more transparent atimospleres. culauced buninosities. aud hence more rapid cooliug 2001).," Note that this deficiency is likely to be more pronounced in a real halo population, as the reduced metallicity typical for halo dwarfs \citep{giz97}
 imply more transparent atmospheres, enhanced luminosities, and hence more rapid cooling \citep{bur01}."
". It is interesting to note that all of the distributions are generally consistent between 18=M,z21 (500 = ToppmLOOO WS). which euconipasses mid- aud late-type T dwarfs."," It is interesting to note that all of the distributions are generally consistent between $18 \lesssim M_{bol} \lesssim 21$ (500 $\lesssim$ $_{eff} \lesssim 1000$ K), which encompasses mid- and late-type T dwarfs."
 These cousisteucies suggestoOo that while the field T dwarf population may be highly scusitive to the underlving ME. (5 3.1). it is generally insensitive to the birth rate.," These consistencies suggest that while the field T dwarf population may be highly sensitive to the underlying MF $\S$ 3.1), it is generally insensitive to the birth rate."
 Iu contrast. the L chart field population is somewhat less scusitive to the ALF but may be an excellent probe of extreme Cadactic birth rates.," In contrast, the L dwarf field population is somewhat less sensitive to the MF but may be an excellent probe of extreme Galactic birth rates."
 These treuds are also ποσα in the steeper ATFs., These trends are also seen in the steeper MFs.
" Figure 11 compares (ΑΓ) aud (T,,,) for minimum ages of 1 to 100 Myr for the α = 0.5 aud 1.5 AIEs.", Figure 11 compares $\Phi(M_{bol})$ and $\Phi({\rm T}_{eff})$ for minimum ages of 1 to 100 Myr for the $\alpha$ = 0.5 and 1.5 MFs.
 No siguificant differences are seen between these LEs. primarily because very voung (f100 My) brown dwarfs coutribute munimally (~1' (fora flat birthrate) to the 10 Cyr field LF over the mass rauge exanuned.," No significant differences are seen between these LFs, primarily because very young $t < 100$ Myr) brown dwarfs contribute minimally $\sim$ for a flat birthrate) to the 10 Gyr field LF over the mass range examined."
 Hence. the minima age of the substellar field population oulv iufiueuces the observed LF if it is of order 1 Gyr or later. as seen with the halo birthrate discussed above.," Hence, the minimum age of the substellar field population only influences the observed LF if it is of order 1 Gyr or later, as seen with the halo birthrate discussed above."
 One of the key parameters for low-1uass star formation theories is the minimum formation mass. which depends not only on the thermodvuamical conditions of the initial gas reservoir (through the Jeans mass. Jeans 1902). but also on the eficiency and history of accretioncarly iu the formation process.," One of the key parameters for low-mass star formation theories is the minimum formation mass, which depends not only on the thermodynamical conditions of the initial gas reservoir (through the Jean's mass, Jeans 1902), but also on the efficiency and history of accretionearly in the formation process."
" Figure 12 compares 9615,) aud 9T.44,) for πάπα formation masses M,,;; — 0.001. 0.010. and 0.015 AL. aud the a = 0.5 and 1.5 MEs."," Figure 12 compares $\Phi(M_{bol})$ and $\Phi({\rm T}_{eff})$ for minimum formation masses $_{min}$ = 0.001, 0.010, and 0.015 $_{\sun}$ and the $\alpha$ = 0.5 and 1.5 MFs."
" As expected. reducing M,,;; results in many more intrinsically faint objects. and the low-telpcrature turnover in @(T,47) (Figure 5) is essentially abseut for M,,;, = 0.001 AL..."," As expected, reducing $_{min}$ results in many more intrinsically faint objects, and the low-temperature turnover in $\Phi({\rm T}_{eff})$ (Figure 5) is essentially absent for $_{min}$ = 0.001 $_{\sun}$."
" Dowever. the differences between these distributions are uceligible for Ap,m=20 and Tyει2500 I for both power-law ATFs."," However, the differences between these distributions are negligible for $M_{bol} \lesssim 20$ and $_{eff} \gtrsim 500$ K for both power-law MFs."
" This is consisteut with the iiass breakdown of 9T.,;) in Figure 6. which shows that the lowest mass brown dwarfs contribute siguificautlv only to the lowest teiiperature/Iunuinositv bius."," This is consistent with the mass breakdown of $\Phi({\rm T}_{eff})$ in Figure 6, which shows that the lowest mass brown dwarfs contribute significantly only to the lowest temperature/luminosity bins."
" Thus. the signature of a init brown dwarf formation mass. unless it is larecr than 0.015 ML... cannot be detected in the ctnrently known sauple of field brown chvarfs. which exteud only to Thy,~TOO IS (Coliniowsld 2001)..."," Thus, the signature of a minimum brown dwarf formation mass, unless it is larger than 0.015 $_{\sun}$, cannot be detected in the currently known sample of field brown dwarfs, which extend only to $_{eff} \sim 700$ K \citep{gol04,vrb04}."
" Determining M,,;, from field measurements will require the discovery of substantiallv cooler brown dwiurfs.", Determining $_{min}$ from field measurements will require the discovery of substantially cooler brown dwarfs.
 Any observed sample of stars or brown dwarfs may include somepercentage of unresolved. iuultiple systems., Any observed sample of stars or brown dwarfs may include somepercentage of unresolved multiple systems.
 Indeed. stellar iiultiples are more frequent amonest Solaw-mass stars than suele svstems (Abt . and this frequeney may be even higher during the carly T-Taun phase (Chez.Neugebauer.&Mattews 1993)..," Indeed, stellar multiples are more frequent amongst Solar-mass stars than single systems \citep[$\sim 60$\%]{abt76,duq91}, , and this frequency may be even higher during the early T-Tauri phase \citep{ghe93}. ."
 Recent ligh-resolution imagine studies of low-mass, Recent high-resolution imaging studies of low-mass
"The authors in (Alietal.2009) proposed a GUP which is consistent with DSR theory. String theory aud black hole plivsics aud which savs and where | EN Here 7, is the Plank length (+ Pn) aud it is geucrally assumed that αρ=1.","The authors in \cite{z6} proposed a GUP which is consistent with DSR theory, String theory and black hole physics and which says and where $ l=\frac{a_0 l_{p}}{\hbar} $ Here $ l_{p} $ is the Plank length $ \approx 10^{-35} m $ ) and it is generally assumed that $ a_0 = 1 $."
" It is evident that this new physical leneth scale which is of the order of ayl, cannot exceed the olectroseak leugth scale ~Ley, which surely dunplies ayx 108) ", It is evident that this new physical length scale which is of the order of $a_0 l_{p} $ cannot exceed the electroweak length scale $ \sim 10^{17} l_{p}$ which surely implies $a_0 \leq 10^{17}$ .
The above equations are approximately covariaut nuder DSR transformations but not Loreutz covariant (Cortesetal. 2005)., The above equations are approximately covariant under DSR transformations but not Lorentz covariant \cite{z4}. .
". These equations also imply and where AZ, is the Plauk mass aud e is the velocity of heht in vacua.", These equations also imply and where $ M_p $ is the Plank mass and $c$ is the velocity of light in vacuum.
 The effect of this proposed GUP is well studied recently for some well known plivsical svstcuis πι (Alietal.2009).. (Dasetal.20093... (Mauuuder 2011).," The effect of this proposed GUP is well studied recently for some well known physical systems in \cite{z6}, \cite{z7}, \cite{bar1}."
. Iu this article we will reinvestieate the results of (Caietal.2008). (Zhuetal.2009) uxiug this ιοην proposed GUP (ALietal.2009).," In this article we will reinvestigate the results of \cite{a13}, \cite{a} using this newly proposed GUP \cite{z6}."
.. We will apply his GUP for the calculation of the modified cutropy., We will apply this GUP for the calculation of the modified entropy.
 Uere we will make the assumption that the apparcut jorizon las the eutropy which is the caleulated modified cutropy., Here we will make the assumption that the apparent horizon has the entropy which is the calculated modified entropy.
 Then we will apply the Clausius relation of hermocdvuaiics to the apparent horizon to get the rewly modified Fricdimuanun equations which governs the dynamical evolution of the universe., Then we will apply the Clausius relation of thermodynamics to the apparent horizon to get the newly modified Friedmann equations which governs the dynamical evolution of the universe.
 If we now sce equation (2)) we can easily notice that the uucertaiutv relation has a linear term in the Plank leneth which is abseut in the earlier version of (UP. (Usecmpt.MaugauoandMaun 1995)..., If we now see equation \ref{e2}) ) we can easily notice that the uncertainty relation has a linear term in the Plank length which is absent in the earlier version of GUP \cite{a16}. .
 Tere we write the GUP of equation (2)) iu a amore ecueral form as Throughout the whole process we will consider 7— where G is the Newton coustaut. c he velocity of light in vactuun anc Ay the Boltzimanu constant.," Here we write the GUP of equation \ref{e2}) ) in a more general form as Throughout the whole process we will consider $\hbar=G=c=k_B=1$ , where $G$ is the Newton constant, $c$ the velocity of light in vacuum and $k_B$ the Boltzmann constant."
" Tore 7, is the Plauk length.", Here $l_p$ is the Plank length.
 J is the coefficieut of ἐν aud a? ijs the coefficient. of p., $\beta$ is the coefficient of $l_p$ and $\alpha^2$ is the coefficient of $l_p^2$.
 > essentially üehlights the effect of the linear terii in Plank length of the nuucertaimty relation., $\beta$ essentially highlights the effect of the linear term in Plank length of the uncertainty relation.
 We keep the freedom im our rand to choose the parameters o and 2., We keep the freedom in our hand to choose the parameters $\alpha$ and $\beta$.
 For example >=j0 gives back the earlier version of CUP as discussed in (Ikenipf.MauganuoaudMaun1995)., For example $\beta=0$ gives back the earlier version of GUP as discussed in \cite{a16}.
. We can further uno 3}j—.2ay mud eD=deg so that we oget back (3))., We can further tune $\beta=-2a_0$ and $\alpha^2=4a_0^2$ so that we get back \ref{ee1}) ).
 From equation (6)) we can write the form of the Ποιοτα uncertainty as where Fepta?) measures the inuouut of departure frou our usual IHeiseuberg uncertainty principle., From equation \ref{n1}) ) we can write the form of the momentum uncertainty as where $f_{GUP}(\delta x^2)$ measures the amount of departure from our usual Heisenberg uncertainty principle.
 We consider the negative sien for ferp(ér7) as the positive sigu has no physical iicaniug., We consider the negative sign for $f_{GUP}(\delta x^2)$ as the positive sign has no physical meaning.
 We now consider a (6|1» dimensional FRW universe with line element where dO74 is the line clement of a {ων 1)dimensional uit radius sphere. e is the scale factor aud k defines the curvature of the spatial section.," We now consider a $n+1$ )-dimensional FRW universe with line element where $\ud \Omega_{n-1}^2$ is the line element of a $n-1$ )-dimensional unit radius sphere, $a$ is the scale factor and $k$ defines the curvature of the spatial section."
 Without the whole evolution history of the universe we cannot sav anything about the cosinological eveut horizon., Without the whole evolution history of the universe we cannot say anything about the cosmological event horizon.
But a dvnaucal apparent horizonalwaysexist in the FRW universe because it is a local quantity of spacetime.,But a dynamical apparent horizonalwaysexist in the FRW universe because it is a local quantity of spacetime.
 The apparent horizon is a mareinally trapped surface with vanishing expansion., The apparent horizon is a marginally trapped surface with vanishing expansion.
 The location of the apparent horizon is given by where ry is due to the re-definition P4= ar., The location of the apparent horizon is given by where $\tilde{r}_A$ is due to the re-definition $\tilde{r}_A = ar$ .
 Here J7(—7) is the IDibble parzineter., Here $H(=\frac{\dot{a}}{a})$ is the Hubble parameter.
 The appareut, The apparent
Taking. advantage of the extremely stable PSF. of IIST/WFETPC?2. we performed a PSF subtraction ou the six images of DENIS-P. J020529.0-115925 using 9 differeut unresolved objects from programs CGOnS5720 (P.LBrauduer.Douyvetal.2003) aud GOshst (P.LReid.Cagisetal.2003).. aud 1 svuthetic PSF obtained with (IKkxist&Took2003) as reference PSF stars.,"Taking advantage of the extremely stable PSF of HST/WFPC2, we performed a PSF subtraction on the six images of DENIS-P J020529.0-115925 using 9 different unresolved objects from programs GO8720 \citep[P.I. Brandner,][]{2003AJ....126.1526B} and GO8581 \citep[P.I. Reid,][]{2003AJ....125.3302G}, and 1 synthetic PSF obtained with \citep{TINY_TIM} as reference PSF stars."
 Figure 7 shows that on { of he 6 images.° the ]primary is extremely well subtracted. while stroug residuals appearsp clearly afier subtraction of the secoudary. indicating he presence of a third component.," Figure \ref{psf_subtraction} shows that on 4 of the 6 images, the primary is extremely well subtracted, while strong residuals appears clearly after subtraction of the secondary, indicating the presence of a third component."
 The results witi the 9 other reference PSF stars are simular‘ o that presented in Figure 7.., The results with the 9 other reference PSF stars are similar to that presented in Figure \ref{psf_subtraction}.
 Using this technique. we tried to define a criterion that would allow to determine quantitatively if the secondary is a uniltiple svsteni or not.," Using this technique, we tried to define a criterion that would allow to determine quantitatively if the secondary is a multiple system or not."
 We performed a detailed PSF analysis of all the WFPC2 images of programs COshsl (PI.ReidCagisetal. 2003).., We performed a detailed PSF analysis of all the WFPC2 images of programs GO8581 \citep[P.I. Reid][]{2003AJ....125.3302G}.
 Tus large sample has the advantage to be homogeneous (all targets were observed with. the FSIIW ∖↽⇁⋅filter of the ↴⋝⋯⋜∐⋅↕↸∖↴∖↴∪↕↑∐↸∖↴∖↴⋜∐⊔↸∖↴∖↴⋜⋯∏⋈∖∙↴⋝∏↑↖↽↸∖↥⋅⋅↖↽↴∖↴⋯∐⋜∐⋅ ↴∖↴↑⋜∐⋅↴∖↴⋜⋯≺∏⋝↥⋅∪↖↖↽∐≼↧↖↖↽⋜∐⋅↸∖↴↖↖↽↕↑↕↴∖↴↻↸∖↸⊳⊓⋅⋜↧↕↻↥⋅∪⋉∖↥⋅↑↕↸∖↴∖↴ (aud therefore. PSF). ⋅⋅ similartoDENTS-Pe J020529.0-," This large sample has the advantage to be homogeneous (all targets were observed with the F814W filter of the WFPC2/PC), and all targets are very low mass stars and brown dwarfs with spectral properties (and therefore PSF) similar to DENIS-P J020529.0-115925."
"very8 iutensity of the residuals.shows defined inte""i of affer PSF nsubtraction. questedivided by the the integratedresidualsiuteusita of the objectJ ‘as expressed1 in equationN 1..appe:appears too |be a ]powerfulf criteriontori to ddidentifvüt binaryt candidates."," The relative intensity of the residuals, defined as the integrated intensity of the residuals after PSF subtraction, divided by the integrated intensity of the object as expressed in equation \ref{eq:eq1},appears to be a powerful criterion to identify binary candidates."
"‘üd"", where: os d8 the ΠΡΟ of pixels. Z5) is the flux -3in pixelEUN / joafterftpry PSEDS subtraction.:syeti au. Fo(i) is the flux in pixel / before PSF subtraction."," where: $n$ is the number of pixels, $\mathcal{F}_{R}(i)$ is the flux in pixel $i$ after PSF subtraction, and $\mathcal{F}_{O}(i)$ is the flux in pixel $i$ before PSF subtraction."
 Figure δ shows that πι is very low aud very stable for unresolved objects. while it is always düeher than ~3-0 above the median value in the case of gauultüple systems.," Figure \ref{residuals} shows that $\mathcal{R.I}$ is very low and very stable for unresolved objects, while it is always higher than $\sim$ $\sigma$ above the median value in the case of multiple systems."
 Thus technique found casily aud automatically all the multiple svstems prescut in the sauple (seeBouvctal2003:Cüzisal. 2003).. except one (2MASSW Jü856179|223518).," This technique found easily and automatically all the multiple systems present in the sample \citep[see ][]{2003AJ....126.1526B,2003AJ....125.3302G}, except one (2MASSW J0856479+223518)."
 This latter uultiple system is indeed a good example to illustrate” the liutatious° of tdis technique., This latter multiple system is indeed a good example to illustrate the limitations of this technique.
 2\TASSW JO856179|223512 is a very close binary (60-0711) with a relatively aree clifference of magnitude (AMag=2.8mae.seeDouvetal.2003)..," 2MASSW J0856479+223518 is a very close binary $\delta$ 1) with a relatively large difference of magnitude \citep[$\Delta$Mag=2.8~mag, see ][]{2003AJ....126.1526B}."
 The relative intensity of the residuals after PSF subtraction of the ΠΑΝ. corresponding roughly to the ratio of he iuteusitv of the secoudary over that of the o»mmary. is therefore of the order of a few perceut only (AMae=2.8 mae is equivalent to a flux ratio of fpífa~0.07).," The relative intensity of the residuals after PSF subtraction of the primary, corresponding roughly to the ratio of the intensity of the secondary over that of the primary, is therefore of the order of a few percent only $\Delta$ Mag=2.8 mag is equivalent to a flux ratio of $f_{B}/f_{A}\sim$ 0.07)."
 The relative⋅ iuteusitv⋅ of. the residuals after PSF subtraction is therefore a good ∐∐∖↑∐∪≼↧↑∪∐∐≼↧∐⋯↕↑∏≻↕↸∖↴∖↴⋅↖⇁↴∖↴↸∖∐↓↴∖↴↖↖↽↕↑∐⋯∪≼∐∖↥⋅⋜↧↑↸∖ . . differences of magnitude., The relative intensity of the residuals after PSF subtraction is therefore a good method to find multiple systems with moderate differences of magnitude.
 Fortunately it is the case of the possible third component of P J020529.0-115925. as shown below.," Fortunately, it is the case of the possible third component of DENIS-P J020529.0-115925, as shown below."
" Using this property, we compare the relative lutensity of the residuals after PSF subtraction on the secondary of DENIS-P. 1020529.0-115925. using the same library of PSF stars."," Using this property, we compare the relative intensity of the residuals after PSF subtraction on the secondary of DENIS-P J020529.0-115925, using the same library of PSF stars."
 Figure ὃ shows the restIts., Figure \ref{residuals} shows the results.
 The resknals at six cliffereut epochs are all ligher than that of unresolved objects. and unresolved companions. of. kuown ∖↖⇪↕≼−≽↕≼⋟∙⋜⋯≼↧⋜↧∐↑⋜∐⋅∶↴∙⊾↸∖↴∖↴⋜∐⋅↸∖↖↽↸∖↥⋅⋅↖↽↕∪↖↖↽⋯⋜↧↴∖∷∖↴⋅⋅ ⋅⋅ ⇁⋟⊲⊀∣⋟⊲ ↑∪↑↕⋯↑∪↕≯∪↑∐↸∖↥⋅∐," The residuals at six different epochs are all higher than that of unresolved objects and unresolved companions of known binaries of the same sample, but very similar to that of other multiple systems."
⋯↕↑↕≻↕↸∖↴∖↴⋅↖↽↴∖↴↑↸∖⋯↴∖↴∙⊺∐↕↴∖↴∐∶↴∙⊾⋃⋅↸∖ shows clearly that DENIS-P.⇁≼ 7020529.0-115925D↽↽ likely to ⋅∙⊽be a binary syste itself., This figure shows clearly that DENIS-P J020529.0-115925B is very likely to be a binary system itself.
 Figure 7115925 .is also that even for the two images where The relative the residuals do mot appear clearly by eves (see Figure 7)). the R.Z is slightly hielier tha in any unresolved object.," Figure \ref{residuals} shows also that even for the two images where the residuals do not appear clearly by eyes (see Figure \ref{psf_subtraction}) ), the $\mathcal{R.I}$ is slightly higher than in any unresolved object."
 One- should note that PRZ strouegly depends ou the signal-to-noise⋅ ratio⋅ (S/N)., One should note that $\mathcal{R.I}$ strongly depends on the signal-to-noise ratio (S/N).
"⊲∣⇁ At low S/N.A value of .Z .must indeed increase. due‘ to the , noise in the target."," At low S/N, the value of $\mathcal{R.I}$ must indeed increase due to the increased noise in the target."
 Figure & shows RI-HLA all our images were obtained at sufficieutlv hieh S/N (always greater than 100). so that 'R.Z provides a robust statistical test of multiplicity.," Figure \ref{residuals} shows that all our images were obtained at sufficiently high S/N (always greater than 100), so that $\mathcal{R.I}$ provides a robust statistical test of multiplicity."
 In order to confirm7 that these residuals. are consistent with the presence of a third component. we a performeddual-PSF on fitting.the secondary oulv. using a custom-made program described iu Douvetal.(2003).," In order to confirm that these residuals are consistent with the presence of a third component, we performed a dual-PSF fitting, on the secondary only, using a custom-made program described in \citet{2003AJ....126.1526B}."
 Briefly. aud as presented in Figue 9. the PSF fitting routine builds a model binary using each of he ten cdiffercut PSF stars mentioned above. and then performs a nou- PSF at of the observed image to determine the best-fit values for the tree free paraiueters:," Briefly, and as presented in Figure \ref{companion}, the PSF fitting routine builds a model binary using each of the ten different PSF stars mentioned above, and then performs a non-linear PSF fit of the observed image to determine the best-fit values for the three free parameters:"
the one-dimensional simulations to accurately represent the BL between the star aud disk.,the one-dimensional simulations to accurately represent the BL between the star and disk.
 In (his case. the exact temperature. size ancl rotation rate of the BL are given in Popham and Godonetal.(1995).," In this case, the exact temperature, size and rotation rate of the BL are given in \citet{pop95} and \citet{god95}."
. The temperature is in the range x150.0001 (Godonetal.1995) to a few 100.000 (Popham&Naravan1995).," The temperature is in the range $\approx 150,000$ K \citep{god95} to a few 100,000K \citep{pop95}."
.. A modeling of the extreme UV (Chandra) spectra of some CVs (e.g. OY Car. $5 Cvg WZ See) in this regime of M leads to such temperatures 2004)...," A modeling of the extreme UV (Chandra) spectra of some CVs (e.g. OY Car, SS Cyg WZ Sge) in this regime of $\dot{M}$ leads to such temperatures \citep{mau00,mau04}."
 Here too. the BL emits in the πο X-ray regime with relatively more contribution to (he extreme UV and lar UV as the temperature decreases.," Here too, the BL emits in the soft X-ray regime with relatively more contribution to the extreme UV and far UV as the temperature decreases."
 The WD and (he inner disk emit in the UV. the outer disk emits in the optical.," The WD and the inner disk emit in the UV, the outer disk emits in the optical."
 The modeling we carry out in the presentwork is the modeling of such a DL., The modeling we carry out in the presentwork is the modeling of such a BL.
" For M«3x10WAL, /vr. the BL is optically thin and the only UV emission is coming from the inner disk edge irradiated bv the optically thin BL."," For $\dot{M} < 3 \times 10^{-10} M_{\odot}$ /yr, the BL is optically thin and the only UV emission is coming from the inner disk edge irradiated by the optically thin BL."
" That inner edge forms a rine al a radius r~1.248,—1.6H.. rotating al IXeplerian speed. wilh an ellective temperature T,gyc100. 000IX—50. 000Ix. (Naravan&Popham1993:1999). [or an accretion rate of 3xLOMAL. /vr and 3x10.MAL. /vi. respectively."," That inner edge forms a ring at a radius $r \sim 1.2R_* -1.6 R_*$, rotating at Keplerian speed, with an effective temperature $T_{eff} \sim 
100,000$ $-50,000$ K \citep{pop93b,pop99} for an accretion rate of $ 3 \times 10^{-10} M_{\odot}$ /yr and $ 3 \times 10^{-11} M_{\odot}$ /yr, respectively."
 As the mass accretion rate decreases. (he ring radius increases and its elleclive temperature decreases.," As the mass accretion rate decreases, the ring radius increases and its effective temperature decreases."
 The DL emits in the had X-ray range and the clisk’s inner edge ring emits in the UV if its temperature is high enough., The BL emits in the hard X-ray range and the disk's inner edge ring emits in the UV if its temperature is high enough.
 such UV emission from the inner edge of the disk/BL was identified in VW IIvi and im other dwarf novae in quiescence (Pandeletal.2005)., Such UV emission from the inner edge of the disk/BL was identified in VW Hyi \citep{pan03} and in other dwarf novae in quiescence \citep{pan05}.
.. This UV contribution from the DL has to be taken into account when modeling the UV spectra of accreting WDs in quiescence (Godon&Sion.2005)., This UV contribution from the BL has to be taken into account when modeling the UV spectra of accreting WDs in quiescence \citep{god05}.
. The WD emits in the UV. as does the very inner disk.," The WD emits in the UV, as does the very inner disk."
 The rest of the disk emits in the optical., The rest of the disk emits in the optical.
 As (he mass accretion rate decreases further. the disk eventually ceases to emit in the UV.," As the mass accretion rate decreases further, the disk eventually ceases to emit in the UV."
" As a consequence. the standard. disk model is roughly correct. for (sav) r22, and larger. al both high aud low mass accretion rates."," As a consequence, the standard disk model is roughly correct for (say) $r \approx 2 R_*$ and larger, at both high and low mass accretion rates."
 The exact value of r for which the standard disk model breaks down depends on the mass accretion rate and mass of the WD., The exact value of $r$ for which the standard disk model breaks down depends on the mass accretion rate and mass of the WD.
 At smaller radi the temperature in the disk-DL region increases rapidly compared to the standard model prediction., At smaller radii the temperature in the disk-BL region increases rapidly compared to the standard model prediction.
 AIV Lyr was observed withFUSE on May. 6. 2003. (at 201:071:238. UT). just after it climbed into a high state that lasted for about 4vr.," MV Lyr was observed with on May 6, 2003 (at 20h:07m:23s UT), just after it climbed into a high state that lasted for about 4yr."
 The FUSE exposure (D9050901) consists ol 1637s of good exposure time., The exposure (D9050901) consists of 7637s of good exposure time.
 The data were processed with the latest and final version, The data were processed with the latest and final version
for the types of molecular accretion flows that are observed around HII regions.,for the types of molecular accretion flows that are observed around HII regions.
 Rather they assume the accretion flow is in the form of an infinitely thin disk., Rather they assume the accretion flow is in the form of an infinitely thin disk.
 In the absence of a more complete calculation we may conceptually combine the flow morphology from Keto(2007) with the velocities and densities from the PED models to interpret the the observations of W51e2., In the absence of a more complete calculation we may conceptually combine the flow morphology from \citet{Keto2007} with the velocities and densities from the PED models to interpret the the observations of W51e2.
 The spin-up of the flow seen in the molecular line observations suggests that the large scale molecular accretion flow becomes progressively more flattened at smaller radii., The spin-up of the flow seen in the molecular line observations suggests that the large scale molecular accretion flow becomes progressively more flattened at smaller radii.
" If the flow is sufficiently dense at its mid-plane, it is able to resist ionization and extend into the HII region as described in Keto(2007)."," If the flow is sufficiently dense at its mid-plane, it is able to resist ionization and extend into the HII region as described in \citet{Keto2007}."
. Across the top and bottom boundaries of the flattened molecular flow there is an ionization front that continuously supplies gas to the HII region as described in the PED models., Across the top and bottom boundaries of the flattened molecular flow there is an ionization front that continuously supplies gas to the HII region as described in the PED models.
" Since there is no obvious force to immediately stop the rotation of the gas as it crosses the ionization front, the ionized gas is also rotating."," Since there is no obvious force to immediately stop the rotation of the gas as it crosses the ionization front, the ionized gas is also rotating."
 Observations of high frequency RRL such as the H53a are dominated by the densest gas that is near the ionization front., Observations of high frequency RRL such as the $\alpha$ are dominated by the densest gas that is near the ionization front.
 The velocities of the H53a line thus show the rotation in the ionized gas coming off the rotating molecular accretion flow., The velocities of the $\alpha$ line thus show the rotation in the ionized gas coming off the rotating molecular accretion flow.
 'The density of the ionized gas decreases rapidly as it accelerates off the ionization front around the flattened molecular accretion flow., The density of the ionized gas decreases rapidly as it accelerates off the ionization front around the flattened molecular accretion flow.
 The exact rate of the density decrease depends on the geometry of the flow., The exact rate of the density decrease depends on the geometry of the flow.
" The Parker wind model assumes a purely spherical flow while the PED models describe the flow off a flat, rotating disk."," The Parker wind model assumes a purely spherical flow while the PED models describe the flow off a flat, rotating disk."
" In both cases, the density distribution is approximately exponential (hydrostatic) where the flow velocities are subsonic and closer to a power law past the transonic point."," In both cases, the density distribution is approximately exponential (hydrostatic) where the flow velocities are subsonic and closer to a power law past the transonic point."
 Observations in the radio continuum should thus measure a rising spectral energy distribution (SED) consistent with the steep density gradient in the outflowing ionized wind., Observations in the radio continuum should thus measure a rising spectral energy distribution (SED) consistent with the steep density gradient in the outflowing ionized wind.
" In the case of W51e2, the observed SED may be reproduced by a model with a density gradient scaling with radius as the —2.5 power 2008)."," In the case of W51e2, the observed SED may be reproduced by a model with a density gradient scaling with radius as the $-2.5$ power \citep{KetoZhangKurtz2008}."
". The PED models show that the ionized gas flows along the rotation axis, and the structures in Keto(2007) show that this flow can be further channeled by the density structure of the molecular accretion flow."," The PED models show that the ionized gas flows along the rotation axis, and the structures in \citet{Keto2007} show that this flow can be further channeled by the density structure of the molecular accretion flow."
" This outflowing ionized gas could drive a molecular BPO, but this BPO would be quite different in collimation and velocity structure from those BPOs associated with the accretion flows of lower mass stars."," This outflowing ionized gas could drive a molecular BPO, but this BPO would be quite different in collimation and velocity structure from those BPOs associated with the accretion flows of lower mass stars."
" At the moment, the driver of the BPO that we see in CO in not known."," At the moment, the driver of the BPO that we see in CO in not known."
" It may be the original BPO, but now flowing through the HII region, or it may be a remnant of a now terminated flow or it might be driven by the ionized outflow."," It may be the original BPO, but now flowing through the HII region, or it may be a remnant of a now terminated flow \citep{Klaassen2006}
 or it might be driven by the ionized outflow."
probe the οσοι of secondary. infall onto groups.,probe the effect of secondary infall onto groups.
 Members and interlopers were redefined after the identification. of gaps in the redshift distribution according to the technique described. by Lopes et al. (, Members and interlopers were redefined after the identification of gaps in the redshift distribution according to the technique described by Lopes et al. (
2009).,2009).
 Aefore selecting group members and rejecting interlopers we first. refine the spectroscopic redshift of cach group and identify its velocity imits., Before selecting group members and rejecting interlopers we first refine the spectroscopic redshift of each group and identify its velocity limits.
 For this purpose. we emplov the egap-technique described in WKatecrt et al. (," For this purpose, we employ the gap-technique described in Katgert et al. ("
1996) ancl Olsen ct al. (,1996) and Olsen et al. (
2005) o identify gaps in the redshift. distribution.,2005) to identify gaps in the redshift distribution.
 X variable eap. calledgap (Aclami et al.," A variable gap, called (Adami et al."
 1998). is considered.," 1998), is considered."
 To determine the group redshift. only galaxies within 0.50 + Alpe are considered.," To determine the group redshift, only galaxies within 0.50 $^{-1}$ Mpc are considered."
 Details about this procedure are ound in Lopes ct al. (, Details about this procedure are found in Lopes et al. (
2009): see also Ribeiro ct al. (,2009); see also Ribeiro et al. (
2009) or applications of this technique to 241 galaxy groups.,2009) for applications of this technique to 2dF galaxy groups.
 With the new redshift’ ancl velocity limits. we apply an algorithm for interloper rejection to define the final list of eroup members.," With the new redshift and velocity limits, we apply an algorithm for interloper rejection to define the final list of group members."
" We use the ""shifting gapper” technique (Fackela et al.", We use the “shifting gapper” technique (Fadda et al.
 1996). which consists of the application of the gap-technique to racial bins from the group center.," 1996), which consists of the application of the gap-technique to radial bins from the group center."
 We consider a bin size of 0.42 + Mpe (0.60 Mpe for h = 0.7) or larger to ensure that at least 15 galaxies are selected., We consider a bin size of 0.42 $^{-1}$ Mpc (0.60 Mpc for h = 0.7) or larger to ensure that at least 15 galaxies are selected.
 Galaxies not associated with the main body of the group are discarded., Galaxies not associated with the main body of the group are discarded.
 This procedure is repeated. until the number of eroup members is stable and no further galaxics are eliminated as intrucers., This procedure is repeated until the number of group members is stable and no further galaxies are eliminated as intruders.
 La the present. work. we have sampled galaxies out to LO times the hamonic mean radius of the systems. including galaxies whose distances to the centers can reach ~S Alpe.," In the present work, we have sampled galaxies out to 10 times the hamonic mean radius of the systems, including galaxies whose distances to the centers can reach $\sim$ 8 Mpc."
 To avoid contamination of nearby structures. we select galaxies within the maximunt racius Rij.= 4.0 Alpe (see La Barbera ct al.," To avoid contamination of nearby structures, we select galaxies within the maximum radius $R_{max}=$ 4.0 Mpc (see La Barbera et al."
 2010)., 2010).
 After applying the shifting eapper procedure we have a list. of eroup members anc we call fy the aperture equivalent to the radial olfset of the most distant member (normally close to Rive)., After applying the shifting gapper procedure we have a list of group members and we call $R_A$ the aperture equivalent to the radial offset of the most distant member (normally close to $R_{max}$ ).
 We estimate the velocity dispersion. (0) within Ry and then the physical radius (Reo) of each group., We estimate the velocity dispersion $\sigma$ ) within $R_A$ and then the physical radius $_{200}$ ) of each group.
 Finally. a virial analysis is perfomed for. mass estimation (Aloo).," Finally, a virial analysis is perfomed for mass estimation $_{200}$ )."
 Further details regarding the interloper removal and estimation of global properties (o. physical radius and mass) are found in Lopes et al. (," Further details regarding the interloper removal and estimation of global properties $\sigma$, physical radius and mass) are found in Lopes et al. ("
2009).,2009).
 The first step in our analysis is to apply the AD test (see Hou et al., The first step in our analysis is to apply the AD test (see Hou et al.
 2009 for a good description of the test) to the velocity distributions of galaxies in groups., 2009 for a good description of the test) to the velocity distributions of galaxies in groups.
 This is done for cillerent distances. producing the following ratios of non-Gaussian eroups: (I: ΕΝ x 24004). and. (2:3459s and £2: 44505).," This is done for different distances, producing the following ratios of non-Gaussian groups: $R \leq 1R_{200}$ ), $R\leq 2R_{200}$ ), and $R\leq 3R_{200}$ and $R\leq 4R_{200}$ )."
 Approximately οἱ all galaxies in our sample have distances :Reoy., Approximately of all galaxies in our sample have distances $\leq 4R_{200}$.
 Vhis is the natural cutolf in space we have made in this work., This is the natural cutoff in space we have made in this work.
 Some properties of galaxy groups are presented in Table 1., Some properties of galaxy groups are presented in Table 1.
 We have classified eroups according to the AD test (at 0.05 significance level) done at A?x4?soo. encompassing all groups with evidence for normality deviations.," We have classified groups according to the AD test (at 0.05 significance level) done at $R\leq 4R_{200}$, encompassing all groups with evidence for normality deviations."
 Properties of non-Ciaussian. (NC) eroups in Table 1 were computed twice. with anc without a correction based on iterative removal of galaxies whose absence in the sample cause the groups become Caussian. following a procedure similar to Perea. del Olmo Moles (1990).," Properties of non-Gaussian (NG) groups in Table 1 were computed twice, with and without a correction based on iterative removal of galaxies whose absence in the sample cause the groups become Gaussian, following a procedure similar to Perea, del Olmo Moles (1990)."
 The corrected properties are just those the system would have i£ it was mace only with galaxies consistent with the normal velocity distribution., The corrected properties are just those the system would have if it was made only with galaxies consistent with the normal velocity distribution.
 Fhis correction allows one to honestly compare typical properties of G and NG groups., This correction allows one to honestly compare typical properties of G and NG groups.
 Not doing that. NC. groups could. have their properties overestimated by a factor of ~1.5. taking all members within 4ogy.," Not doing that, NG groups could have their properties overestimated by a factor of $\sim 1.5$, taking all members within $4R_{200}$."
 After this procedure. we see in Table 1 that G and corrected NG groups have similar properties.," After this procedure, we see in Table 1 that G and corrected NG groups have similar properties."
 A suitable way to investigate galaxics in multiple galaxy systems is to combine them in stacked objects (Biviano et al., A suitable way to investigate galaxies in multiple galaxy systems is to combine them in stacked objects (Biviano et al.
 1992)., 1992).
 Thus. we built two composite groups. Gaussian.C (composed of 48 systems) and non-GaussianNCG (composed of 9 systems).," Thus, we built two composite groups, Gaussian–G (composed of 48 systems) and non-Gaussian–NG (composed of 9 systems)."
 Galaxies in theses composite groups have distances to group centers normalized. by σου ancl their velocities refers to the group median velocities and are scaled by the group velocity dispersions where / and j are. respectively. the galaxy and the group indices.," Galaxies in theses composite groups have distances to group centers normalized by $R_{200}$ and their velocities refers to the group median velocities and are scaled by the group velocity dispersions where $i$ and $j$ are, respectively, the galaxy and the group indices."
 Velocity dispersions of the composite groups refer to the dimensionless quantity v;., Velocity dispersions of the composite groups refer to the dimensionless quantity $u_i$.
 Absolute magnitudes. Alp. are obtained [rom Super-COSALIOS ]t band. a 24) photometric information.," Absolute magnitudes, $M_R$ , are obtained from Super-COSMOS R band, a 2dF photometric information."
" Cosmology. is defined. by Q,,, = 03. Ὃν = (7. and Jy=100hkmsAlpe! Distance-cdependent quantities are calculated: using fo= 0.7."," Cosmology is defined by $\Omega_m$ = 0.3, $\Omega_\lambda$ = 0.7, and $H_0 = 100~h~{\rm km~s^{-1}Mpc^{-1}}$ Distance-dependent quantities are calculated using $h=0.7$ ."
 ALL figures presented. in the next section. correspond to cumulative data in 2/4599 or Alp., All figures presented in the next section correspond to cumulative data in $R/R_{200}$ or $M_R$.
 Error-bars in. our analysis are obtained [rom a bootstrap technique with 1000 resamplinegs., Error-bars in our analysis are obtained from a bootstrap technique with 1000 resamplings.
 Segregation analysis is a powerful tool to evaluate galaxy evolution in galaxy svstems (e.g. Goto. 2005).," Segregation analysis is a powerful tool to evaluate galaxy evolution in galaxy systems (e.g. Goto, 2005)."
 We probe scerceation phenomena out to 4/?»oo. looking for dilferences in galaxies with respect to the dynamical state of the groups.," We probe segregation phenomena out to $4R_{200}$ , looking for differences in galaxies with respect to the dynamical state of the groups."
 First. we test the presence. of luminosity segregation in the velocity. space by computing the normalized. velocity dispersion. o. of the stacked G and NG groups.," First, we test the presence of luminosity segregation in the velocity space by computing the normalized velocity dispersion, $\sigma_u$, of the stacked G and NG groups."
" In Figure l. we plot a, of the composite groups as a function. of the absolute magnitude in the Ro band."," In Figure 1, we plot $\sigma_u$ of the composite groups as a function of the absolute magnitude in the R band."
 We clearly see that. at Aly:21.5. the velocity clispersions decreases towards brighter absolute magnitudes.," We clearly see that, at $M_R \leq -21.5$, the velocity dispersions decreases towards brighter absolute magnitudes."
 On the other hand. for fainter absolute magnitudes. the velocity dispersions are approximately constant.," On the other hand, for fainter absolute magnitudes, the velocity dispersions are approximately constant."
 More interestingly. although the result is similar for both stacked. &roups. for the NC group we see a steeper correlation in the bright end than that we see for the G group.," More interestingly, although the result is similar for both stacked groups, for the NG group we see a steeper correlation in the bright end than that we see for the G group."
" HE one assumes a constant galaxy niass-to-light ratio. enerey equipartition implies e,x107-4 (e.g. Adami. Biviano Mazure. 1998)."," If one assumes a constant galaxy mass-to-light ratio, energy equipartition implies $\sigma_u \propto 10^{0.2M_R}$ (e.g. Adami, Biviano Mazure, 1998)."
" The regression lines between loge, and Aly have slopes QO.184:0.05. and 0.380.038. for €i and. NG groups. respectively."," The regression lines between $\log{\sigma_u}$ and $M_R$ have slopes $\pm$ 0.05 and $\pm$ 0.03, for G and NG groups, respectively."
 Phat is. the brightest galaxies are moving more slowlythan other group galaxies.," That is, the brightest galaxies are moving more slowlythan other group galaxies."
 Such a segregation in the velocity space maybe, Such a segregation in the velocity space maybe
The magnetic field evolution is given by Iq. (1)).,"The magnetic field evolution is given by Eq. \ref{bz}) ),"
" where ος denotes the velocity of the charged. particles (the same for electrons ancl protons). e, is the velocity of the neutrons. and oymeDy ds the ambipolar illusion The first term on the right-hand. side of this equation. causes an advection of the magnetic [ux with a velocity. 0. while the second term describes the ohmie diffusion of the field. where a is the electrical conductivity."," where $v_c$ denotes the velocity of the charged particles (the same for electrons and protons), $v_n$ is the velocity of the neutrons, and $v_A\equiv v_c-v_n$ is the ambipolar diffusion The first term on the right-hand side of this equation causes an advection of the magnetic flux with a velocity $v_c$, while the second term describes the ohmic diffusion of the field, where $\sigma_0$ is the electrical conductivity."
 In neutron star core conditions. the electrical conductivity is my1078.+. thus. the evolution of the large-scale magnetic field through ohmic diffusion proceeds very. slowly. with a time-scale vr. longer than the age of the universe (Davi.Pethick.&Pines 1969).," In neutron star core conditions, the electrical conductivity is $\sigma_0
\sim 10^{28}s^{-1}$, thus, the evolution of the large-scale magnetic field through ohmic diffusion proceeds very slowly, with a time-scale $t_{ohmic}\sim 10^{11}~ \mathrm{yr}$ , longer than the age of the universe \citep{Bay-69}."
. Qualitatively similar conditions hold in essentially all astrophysical settings., Qualitatively similar conditions hold in essentially all astrophysical settings.
 Pherefore. in the rest of this paper. we neglect the ohmic term in the magnetic field evolution equation and focus only on the magnetic evolution due to the advective terni.," Therefore, in the rest of this paper, we neglect the ohmic term in the magnetic field evolution equation and focus only on the magnetic evolution due to the advective term."
 However. the time-scale for ohimic diffusion scales with the square of the characteristic length of the magnetic field. variations (see Paper D). thus. the influence of ohmic dissipation can be important in regions with strong spatial magnetic field variations.," However, the time-scale for ohmic diffusion scales with the square of the characteristic length of the magnetic field variations (see Paper I), thus, the influence of ohmic dissipation can be important in regions with strong spatial magnetic field variations."
 The evolution of the particle number density perturbations. for the different species. is given by Eqs. (2))," The evolution of the particle number density perturbations, for the different species, is given by Eqs. \ref{nB}) )"
" and (3)). where δη,=an,dn. is the perturbation of the charged. particle number density keeping charge neutrality and Ong=On,|On. is the perturbation of the barvon number density."," and \ref{nc}) ), where $\delta n_e=\delta n_p \equiv \delta n_c$ is the perturbation of the charged particle number density keeping charge neutrality and $\delta n_B=\delta n_n+\delta n_c$ is the perturbation of the baryon number density."
" Weak interactions cause beta decays (conversion of charged: particles into neutrons and the opposite) that tend. to. reduce. deviations from. the state between charged: particles and neutrons. with a net rate coefficient A and 05654.—Oye,|Of6."," Weak interactions cause beta decays (conversion of charged particles into neutrons and the opposite) that tend to reduce deviations from the state between charged particles and neutrons, with a net rate coefficient $\lambda$ and $\delta \mu_c\equiv \delta
\mu_e+\delta \mu_p$."
 The chemical equilibrium is achieved when 9/5.=955., The chemical equilibrium is achieved when $\delta \mu_c=\delta \mu_n$.
 Continuing our description. Iq. (5))," Continuing our description, Eq. \ref{vA}) )"
" shows that the ambipolar diffusion velocity ey is controlled by the collision rate between the charged. particles and. neutrons. which is oportional to the parametre «.,."," shows that the ambipolar diffusion velocity $v_A$ is controlled by the collision rate between the charged particles and neutrons, which is proportional to the parametre $\gamma_{cn}$."
 This velocity is driven » the Lorentz force but. choked by the charged: particle oessure gradient it. produces. (Pethick1992:Goldreich&teiscnegecr 1992).," This velocity is driven by the Lorentz force but choked by the charged particle pressure gradient it produces \citep{P-92,GR-92}."
 Reearcding the neutrons. we see from the right-hand side of Eq. (4))," Regarding the neutrons, we see from the right-hand side of Eq. \ref{vn}) )"
 that they move with a velocity ον às ong as the system is not in a state ofequilibrium., that they move with a velocity $v_n$ as long as the system is not in a state of.
 In this state. the Lorentz force is balanced. by he pressure eracients of all the particles (charged. particles and neutrons).," In this state, the Lorentz force is balanced by the pressure gradients of all the particles (charged particles and neutrons)."
 We also see that the neutron velocity is controlled by the parametre à. whose meaning we explain in the next. paragraph.," We also see that the neutron velocity is controlled by the parametre $\alpha$, whose meaning we explain in the next paragraph."
 Since we are interested in a numerical solution of our equations. we have to take into account that mocdelling the true dvnamical. time-scales (sound or. Alfyvéenn time-scales of milliseconds to seconds) would require a time step many orders of magnitude shorter than that. required. to simulate the long-term. evolution. making the simulation computationally unfeasible.," Since we are interested in a numerical solution of our equations, we have to take into account that modelling the true dynamical time-scales (sound or Alfvénn time-scales of milliseconds to seconds) would require a time step many orders of magnitude shorter than that required to simulate the long-term evolution, making the simulation computationally unfeasible."
 For. this reason. in Paper 1 we introduced a approrimalion. neglecting the acceleration terms and instead introducing a small. artificial. friction-like force acting on the neutrons (the most abundant species) of the form. —noí00s. where the parametre a is chosen in such a wav that the time to reach magneto-hvelrostatic equilibrium is long enough for the numerical code to be able to deal. with it (and. therefore. much longer than the real cvnamical times). but shorter than the long time-scales over which the magnetic field evolves.," For this reason, in Paper I we introduced a , neglecting the acceleration terms and instead introducing a small, artificial, friction-like force acting on the neutrons (the most abundant species) of the form $-n_{0n}\alpha
v_n$, where the parametre $\alpha$ is chosen in such a way that the time to reach magneto-hydrostatic equilibrium is long enough for the numerical code to be able to deal with it (and therefore, much longer than the real dynamical times), but shorter than the long time-scales over which the magnetic field evolves."
 We showed in Paper E that the latter time-scales. which are the astrophysically interesting ones. are unallected by the choice ob a.," We showed in Paper I that the latter time-scales, which are the astrophysically interesting ones, are unaffected by the choice of $\alpha$."
 We require the conservation of both the magnetic Lux p=hBede and the barvon number perturbation ὁδρ=A9ndr during the evolution of our system. which spans the segment. O<uexd.," We require the conservation of both the magnetic flux $\Phi=\int_0^{d}B~dx$ and the baryon number perturbation $\delta
N_B=\int_0^{d}\delta n_B~dx$ during the evolution of our system, which spans the segment $0\leq x\leq d$."
 We ensure this through the boundary. conclitions The system of qs. (1--3)), We ensure this through the boundary conditions The system of Eqs. \ref{bz}- \ref{nc}) )
 describes the evolution of three coupled. variables: the magnetic field. the charged. particle density perturbation. anc the barvon density perturbation.," describes the evolution of three coupled variables: the magnetic field, the charged particle density perturbation, and the baryon density perturbation."
 In. Paper I. we estimated the three associated. characteristic cvolutionary time-scales of this set of equations corresponding to exponentially decaving cigenmodes in. the linear approximation and. showed that they characterize the approach to three successive equilibrium states.," In Paper I, we estimated the three associated characteristic evolutionary time-scales of this set of equations corresponding to exponentially decaying eigenmodes in the linear approximation and showed that they characterize the approach to three successive equilibrium states."
 Here we brielly describe the evolutionary stages and summarize the relevant time-scales with the main &oal of establishing the basic ideas that will be used in our subsequent analysis (see Paper L for more details about their derivation)., Here we briefly describe the evolutionary stages and summarize the relevant time-scales with the main goal of establishing the basic ideas that will be used in our subsequent analysis (see Paper I for more details about their derivation).
 The shortest time corresponds to the approach of the magneto-hyelrostatic equilibrium. controlled. by à in our model. as already explained.," The shortest time corresponds to the approach of the magneto-hydrostatic equilibrium, controlled by $\alpha$ in our model, as already explained."
 In a longer time. two alternative processes compete: Ifthe relative motion of charged and neutral particles proceedsmuch faster than the conversion from one into another. a diffusive equilibrium is achieved in the system. characterized by the balance equations and," In a longer time, two alternative processes compete: If the relative motion of charged and neutral particles proceedsmuch faster than the conversion from one into another, a diffusive equilibrium is achieved in the system, characterized by the balance equations and"
the core with an apparent velocity that ranges between 15e and 20e. and all the jet components are well aligned along the same position angle. suggesting that no jet. precession is taking place on a time lag longer than a decade in our frame.,"the core with an apparent velocity that ranges between $c$ and $c$, and all the jet components are well aligned along the same position angle, suggesting that no jet precession is taking place on a time lag longer than a decade in our frame."
 In Fig. 9.," In Fig. \ref{1510_arrow},"
 we report the time of zero separation of all the components cjectec between 1994. ancl 2010. either. analysed in this paper. or from the literature.," we report the time of zero separation of all the components ejected between 1994 and 2010, either analysed in this paper, or from the literature."
 The ejection of a new component likely occurs roug= once per vear., The ejection of a new component likely occurs roughly once per year.
 “Phe gaps between 2000 and 2004 is likely due to sparse observations available. precluding to reliably follow the evolution of the source components in this period.," The gaps between 2000 and 2004 is likely due to sparse observations available, precluding to reliably follow the evolution of the source components in this period."
 Since 1995. in addition to the typical variability interchanging low-activitv with high-activity states. the source shows on average a slight increase of the total [Lux density (Fig. 9)).," Since 1995, in addition to the typical variability interchanging low-activity with high-activity states, the source shows on average a slight increase of the total flux density (Fig. \ref{1510_arrow}) )."
 However. the sparse distribution of the data points does not allow us to unambieuoush relate the ejection date of the superluminal components. represented bv the arrows in Fig. 9..," However, the sparse distribution of the data points does not allow us to unambiguously relate the ejection date of the superluminal components, represented by the arrows in Fig. \ref{1510_arrow},"
 with the various states of the source From the analysis of the parsec-scale. resolution data we could. construct the lighteurves for the core ane jet components separately. and to study a connection between he total intensity and the polarized. emission.," with the various states of the source From the analysis of the parsec-scale resolution data we could construct the lightcurves for the core and jet components separately, and to study a connection between the total intensity and the polarized emission."
 The core has high degree of variability. both in. the otal intensity Εαν density and in polarization properties. out. the data. points are too sparse (o compare changes in the total intensity ancl polarized. emission. throughout he entire time interval.," The core has high degree of variability, both in the total intensity flux density and in polarization properties, but the data points are too sparse to compare changes in the total intensity and polarized emission throughout the entire time interval."
 An interesting case Is represented ov the simultaneous enhancement. of the total intensity lux density and the decrease of the fractional polarization observed in the second. half of 1999. and the rotation of EVPA around. Alav-July 2001. with a swing of 85 and 130 at SA and 15 Cllz respectively (Fig. 6)).," An interesting case is represented by the simultaneous enhancement of the total intensity flux density and the decrease of the fractional polarization observed in the second half of 1999, and the rotation of EVPA around May-July 2001, with a swing of $^{\circ}$ and $^{\circ}$ at 8.4 and 15 GHz respectively (Fig. \ref{plot_mojave99}) )."
 Possible explanations mav be related either. to variation in the opacity in a newly ejected self-absorbed component. or to a highly ordered magnetic field produced by the compression of tangled magnetic fields by shocks.," Possible explanations may be related either to variation in the opacity in a newly ejected self-absorbed component, or to a highly ordered magnetic field produced by the compression of tangled magnetic fields by shocks."
 In the former scenario the evolution of the jet knot implies a change between the Opacity regimes., In the former scenario the evolution of the jet knot implies a change between the opacity regimes.
 As à consequence both the total intensity and the polarized IHux densities should decrease during the transition. whereas the magnetic field. should jump by 90 degrees.," As a consequence both the total intensity and the polarized flux densities should decrease during the transition, whereas the magnetic field should jump by 90 degrees."
 As the opacity decreases. the total flux. density should increase in the optically-thick part of the spectrum. while the polarization percentage decreases.," As the opacity decreases, the total flux density should increase in the optically-thick part of the spectrum, while the polarization percentage decreases."
 As the emitting region expands becoming optically-thin. the radio emission decreases. while the fractional polarization increases and the magnetic field flips of 90 degrees. going back to its original value.," As the emitting region expands becoming optically-thin, the radio emission decreases, while the fractional polarization increases and the magnetic field flips of 90 degrees, going back to its original value."
 Ht must be noted that changes of the IZVPAX of 90° are expected from. opacity effects. during the transition between the two regimes., It must be noted that changes of the EVPA of $^{\circ}$ are expected from opacity effects during the transition between the two regimes.
 Since the process is quite fast. à dense time sampling of the polarization information (as the one obtained between 2007 ancl 2010. Fig. 10))," Since the process is quite fast, a dense time sampling of the polarization information (as the one obtained between 2007 and 2010, Fig. \ref{multi_chi}) )"
 is crucial to detect. the transition between the regimes., is crucial to detect the transition between the regimes.
 (X. sparser sampling. like the one between 1999 and 2001 (Fig. 6)).," A sparser sampling, like the one between 1999 and 2001 (Fig. \ref{plot_mojave99}) ),"
 may provide EVDPA changes larecr/smaller than the expected 90., may provide EVPA changes larger/smaller than the expected $^{\circ}$.
 Phe transition between the optically-thieck ancl -thin regimes occurs in the presence of high opacity values (e.g. Toc 5 - JO. see for example Aller (1970) ancl Pacholezvk (1970) for a detailed discussion).," The transition between the optically-thick and -thin regimes occurs in the presence of high opacity values (e.g. $\tau
 \sim$ 5 - 10, see for example Aller (1970) and Pacholczyk (1970) for a detailed discussion)."
 Such opacity would. cause a dramatic drop of the total intensity [lux density. during the transition., Such opacity would cause a dramatic drop of the total intensity flux density during the transition.
 The lack of observations during the Lirst half of 1999. Le. when the new jet component should have originated. do not allow us to unambieuously confirm the opacity scenario.," The lack of observations during the first half of 1999, i.e. when the new jet component should have originated, do not allow us to unambiguously confirm the opacity scenario."
 However. the increase. of the total [ux density observed. between 1999. and 2000. is. clillicult to explain in the presence of such high. values of the opacity. making this scenario The alternative scenario assumes the formation of a shock that causes the compression of the plasma along the propagation axis.," However, the increase of the total flux density observed between 1999 and 2000 is difficult to explain in the presence of such high values of the opacity, making this scenario The alternative scenario assumes the formation of a shock that causes the compression of the plasma along the propagation axis."
 This generates the amplification of the perpendicular component of the magnetic field. with respect to the parallel one. anc an enhancement. of the luminosity.," This generates the amplification of the perpendicular component of the magnetic field with respect to the parallel one, and an enhancement of the luminosity."
 This scenario would precict an increase of, This scenario would predict an increase of
much of the atmosphere.,much of the atmosphere.
 Since this extra source of heating in the deeper regions of the atmosphere is spatially cuform (Fig., Since this extra source of heating in the deeper regions of the atmosphere is spatially non-uniform (Fig.
 D). it should lead to some form of Joule-driven circulation in the “inert” layers located well below the radiatively-forced “weather” layers.," 1), it should lead to some form of Joule-driven circulation in the “inert” layers located well below the radiatively-forced “weather” layers."
 This conclusion is largely independent of the specific stellar insolation model adopted here since the existence of radiatively inactive layers at levels deeper than a few bars is a rather generic property of irradiated hot Jupiter atmospheres (e.g.. Hansen 2008).," This conclusion is largely independent of the specific stellar insolation model adopted here since the existence of radiatively inactive layers at levels deeper than a few bars is a rather generic property of irradiated hot Jupiter atmospheres (e.g., Hansen 2008)."
 However. it remains to be seen what the nature of such Joule-driven circulation might be given the uneven heating pattern shown in Fig.," However, it remains to be seen what the nature of such Joule-driven circulation might be given the uneven heating pattern shown in Fig."
 1. the non-uniformity of stellar insolation from the day-side to the night-side and the presence of an additional. spatially uniform net flux emerging from the deep planetary interior.," 1, the non-uniformity of stellar insolation from the day-side to the night-side and the presence of an additional, spatially uniform net flux emerging from the deep planetary interior."
 Finally. it is worth remembering that. since z;xB. a stronger field could produce a substantially more dominant ohmic dissipation than estimated above for B=3G*.," Finally, it is worth remembering that, since $\tau_J\propto B^{-2}$, a stronger field could produce a substantially more dominant ohmic dissipation than estimated above for $B=3$."
. Another potentially important consequence of ohmic dissipation ts the possibility that 1t contributes to the tendency for hot Jupiters to have inflated radii (Batygin Stevenson 2010)., Another potentially important consequence of ohmic dissipation is the possibility that it contributes to the tendency for hot Jupiters to have inflated radii (Batygin Stevenson 2010).
" To evaluate the magnitude of this effect on the basis of our three-dimensional atmospheric models. we compute the cumulative ohmic power dissipated in successive atmospheric shells. from the top level at a pressure p, of 1 mbar. down to a pressure pcr) within the atmosphere: this is readily obtained by integration of Eq.(6)). Fig."," To evaluate the magnitude of this effect on the basis of our three-dimensional atmospheric models, we compute the cumulative ohmic power dissipated in successive atmospheric shells, from the top level at a pressure $p_t$ of 1 mbar, down to a pressure $p(r)$ within the atmosphere; this is readily obtained by integration of \ref{eq:QJ}) ), Fig."
 2 displays this ohmic power as a function of pressure for our fiducial. drag-free atmospheric model and for the model with strongest drag described in PMRIO. with zonal wind speeds typically reduced by about 30—5056.," 2 displays this ohmic power as a function of pressure for our fiducial, drag-free atmospheric model and for the model with strongest drag described in PMR10, with zonal wind speeds typically reduced by about $30-50\%$."
 For each model. we display ohmic power for two different values of the magnetic field strength adopted in the ohmie dissipation calculation. B=3 G. and B=10 G. This allows us to separately explore the effects of drag on the zonal winds and of the magnetic field strength on the resulting magnitude of ohmic dissipation.," For each model, we display ohmic power for two different values of the magnetic field strength adopted in the ohmic dissipation calculation, $B=3$ G, and $B=10$ G. This allows us to separately explore the effects of drag on the zonal winds and of the magnetic field strength on the resulting magnitude of ohmic dissipation."
 This is a useful exercise until models with a self-consistent treatment of magnetic drag and ohmie dissipation become available., This is a useful exercise until models with a self-consistent treatment of magnetic drag and ohmic dissipation become available.
 Fig., Fig.
 2 shows that much of the dissipated ohmic power builds up at pressure levels from a few bars to several tens of bars., 2 shows that much of the dissipated ohmic power builds up at pressure levels from a few bars to several tens of bars.
 The power scales as B and. relative to the drag-free model. it is reduced by a factor ~3 1n the circulation model with strongly dragged winds. at fixed magnetic field strength.," The power scales as $B^2$ and, relative to the drag-free model, it is reduced by a factor $\sim 3$ in the circulation model with strongly dragged winds, at fixed magnetic field strength."
 Since the stellar insolation power received by HD 209458b is ~3«107 W. Fig.," Since the stellar insolation power received by HD 209458b is $\sim 3 \times 10^{22}$ W, Fig."
 2 shows that the total ohmic power dissipated in these model atmospheres typically approaches or exceeds of the stellar insolation power., 2 shows that the total ohmic power dissipated in these model atmospheres typically approaches or exceeds of the stellar insolation power.
 Notice that. while the Ohmic power in Fig.2 is integrated down to the lowest grid points (~200 bar) available in the model of Rauscher Menou (2010). it is not yet seen to saturate at those levels.," Notice that, while the Ohmic power in Fig.2 is integrated down to the lowest grid points $\sim 200$ bar) available in the model of Rauscher Menou (2010), it is not yet seen to saturate at those levels."
 This is because. while the zonal velocities decrease with depth. they are still non-zero in the deepest. ‘inert’ model layers.," This is because, while the zonal velocities decrease with depth, they are still non-zero in the deepest, 'inert' model layers."
 It is in fact not a priori clear where exactly the transition to the convection zone should be located in hot Jupiters and this constitutes an uncertainty for Joule heating models of the deepest atmospheric layers., It is in fact not a priori clear where exactly the transition to the convection zone should be located in hot Jupiters and this constitutes an uncertainty for Joule heating models of the deepest atmospheric layers.
 Overall. our results on the dissipated ohmic power are consistent with the estimates made by Batygin Stevenson (2010) on the basis of a more idealized. one-dimensional atmosphere model. using a parametrized zonal wind profile.," Overall, our results on the dissipated ohmic power are consistent with the estimates made by Batygin Stevenson (2010) on the basis of a more idealized, one-dimensional atmosphere model, using a parametrized zonal wind profile."
 Like them. we find that a large fraction of the ohmic power ts dissipated in the deeper atmospheric levels. below ~10 bar.," Like them, we find that a large fraction of the ohmic power is dissipated in the deeper atmospheric levels, below $\sim 10$ bar."
 Contrary to Batygin Stevenson (2010). however. we do not solve for radial currents 1n our models. only meridional ones (Eq. 5)).," Contrary to Batygin Stevenson (2010), however, we do not solve for radial currents in our models, only meridional ones (Eq. \ref{eq:j}) ),"
 and our results are thus largely independent of the presence or absence of an insulating layer high in the atmosphere., and our results are thus largely independent of the presence or absence of an insulating layer high in the atmosphere.
 In fact. we find that an insulating layer is present on the nightside in our three-dimensional models. but not necessarily on the much hotter. less resistive dayside.," In fact, we find that an insulating layer is present on the nightside in our three-dimensional models, but not necessarily on the much hotter, less resistive dayside."
 This raises the possibility that. in more detailed. three-dimensional MHD models. current loops could actually close high in the atmosphere. rather than deep in the planetary interior as conjectured by Batygin Stevenson (2010).," This raises the possibility that, in more detailed three-dimensional MHD models, current loops could actually close high in the atmosphere, rather than deep in the planetary interior as conjectured by Batygin Stevenson (2010)."
 While our results. relying on leading-order meridional currents in the atmospheric region alone. are presumably not strongly sensitive to these conditions affecting radial currents. it remains to be seen how currents would flow in more realistic," While our results, relying on leading-order meridional currents in the atmospheric region alone, are presumably not strongly sensitive to these conditions affecting radial currents, it remains to be seen how currents would flow in more realistic"
The Sun accelerates as il travels in the galaxy.,The Sun accelerates as it travels in the galaxy.
 The acceleration is small (approximately. Gmms lj and it takes millions of vears to change the Sun's velocity vector significantly.," The acceleration is small (approximately, $6$ mm $^{-1}$ $^{-1}$ ), and it takes millions of years to change the Sun's velocity vector significantly."
 The instantaneous velocity of the Sun in a non-rotating galactic reference frame with the origin at the center-of-mass of the \lilky Way causes all observed directions to both galactic and extragalactic sources (ο deflect from their “true” directions in a svstematic wav.," The instantaneous velocity of the Sun in a non-rotating galactic reference frame with the origin at the center-of-mass of the Milky Way causes all observed directions to both galactic and extragalactic sources to deflect from their “true"" directions in a systematic way."
 This phenomenon has (the same special-relativistic nature as (he annual stellar aberration induced bv the orbital motion of the Earth around the Sun., This phenomenon has the same special-relativistic nature as the annual stellar aberration induced by the orbital motion of the Earth around the Sun.
 Due to the aberration the observed position of a light source is displaced (toward (he cirection of the instantaneous velocity of the observer wilh respect to an inertial reference frame al rest by an amount proportional to the velocity magnitude (Novalevsky&Seicelmann2OO4)., Due to the aberration the observed position of a light source is displaced toward the direction of the instantaneous velocity of the observer with respect to an inertial reference frame at rest by an amount proportional to the velocity magnitude \citep{ks}.
. For an observer located at (he barveenter of the solar svstem. the instantaneous ellect ol the relativistic aberration due to the galactic motion of the solar svstem is not directly observable at anv fixed instant of Gime because the velocits-induced aberration pattern is constant.," For an observer located at the barycenter of the solar system, the instantaneous effect of the relativistic aberration due to the galactic motion of the solar system is not directly observable at any fixed instant of time because the velocity-induced aberration pattern is constant."
 But since the motion of the solar barvcenter wilh respect to a fixed. non-rotatine ealactic reference Drame is non-rectilinear. (he euasi-stationary (secular) change of its velocity vector causes (he aberration pattern to echange on (the skv gradually as (ime progresses.," But since the motion of the solar barycenter with respect to a fixed, non-rotating galactic reference frame is non-rectilinear, the quasi-stationary (secular) change of its velocity vector causes the aberration pattern to change on the sky gradually as time progresses."
 Given astrometric measurements al a mieroaresecond. (jas)) level of accuracy conducted over a lime span of al least several vears. (he changing aberration pattern can be observed in (wo wavs.," Given astrometric measurements at a microarcsecond ) level of accuracy conducted over a time span of at least several years, the changing aberration pattern can be observed in two ways."
 An astrometric [acilitv. such as Verv Long Baseline Interferometry. capable of measuring lone arcs between light sources al (wo separate epochs. could directly detect ancl measure the distortion of the erid of a global astrometric frame based on quasars taken as relerence objects (Fomalont&Reid2004).," An astrometric facility, such as Very Long Baseline Interferometry, capable of measuring long arcs between light sources at two separate epochs, could directly detect and measure the distortion of the grid of a global astrometric frame based on quasars taken as reference objects \citep{freid}."
. This distortion would be a direct consequence of the directional difference in Sun's galactic velocity between the two epochs. as well as," This distortion would be a direct consequence of the directional difference in Sun's galactic velocity between the two epochs, as well as"
"expansion: I5, 53.5 log eres.",expansion: $_{Weav}$ = 53.5 log ergs.
 Energies of this magnitucle are equivalent to that provided. by stellar winds. or SNe ooduced. from. 100 O-twpe stars (Chevalier1974).," Energies of this magnitude are equivalent to that provided by stellar winds, or SNe produced from $\sim$ 100 O-type stars \citep{chevalier}."
.. 1n any other larger svsteni an association of a few huncdre stars would be unremarkable.," In any other larger system, an association of a few hundred stars would be unremarkable."
 However. OB associations numbering more than 10 are uncommon in the Magellanic Bridge (Biea&Schmitt1995.:Dica.priv.comm.2001). aux is not clear how such a relatively well populated association mav come to be at the observed location.," However, OB associations numbering more than $\sim$ 10 are uncommon in the Magellanic Bridge \citep[][;Bica, priv. comm. 2001]{bica} and is not clear how such a relatively well populated association may come to be at the observed location."
 For this reason. ormation of the loop by stellar winds (or SNe) is considere implausible.," For this reason, formation of the loop by stellar winds (or SNe) is considered implausible."
 Works by Tenorio-Fagle&Bodenheimer(1988)— anc 'Tenorio-Taeleetal(1986). have shown through a variety of two dimensional numerical simulations. how an cecloucl infalling into a stratified σας laver can generate an expanding shell-like structure.," Works by \cite{tb88} and \cite{tt86} have shown through a variety of two dimensional numerical simulations, how an cloud infalling into a stratified gas layer can generate an expanding shell-like structure."
" From. Tenorio-Iagle.et (1986).. we can relate the expansion energy of the hole to the kinetic energy of a hypothetical infalling cloud. via the cloucs density ariel raclius: nem""]29.78. pEu Using the kinetic energy estimated previously. £i, (erg). we can probe the ranges of LVC properties that are capable of generating a hole having the same observed parameters of the Bridge loop."," From \citet{tt86}, we can relate the expansion energy of the hole to the kinetic energy of a hypothetical infalling cloud, via the cloud's density and radius: $^{-3}$ $\times$ $^{-43}\frac{E_{kin}}{R_c^3V_c^2}$ Using the kinetic energy estimated previously, $E_{kin}$ (erg), we can probe the ranges of HVC properties that are capable of generating a hole having the same observed parameters of the Bridge loop."
" We find that after limiting the density of the candidate infalling ΗΝ to 0.2<p 75 em ὃς according to the range of LIVC densitiesΟ estimated. by Jpüns(2003).. a 0.6. 10"". AL. cloud. need only move at a velocity of ~-350 iin the LS1t frame to attain the estimated kinetic energy to create the observed loop."," We find that after limiting the density of the candidate infalling HVC to $< \rho >$ 5 cm $^{-3}$, according to the range of HVC densities estimated by \cite{brunsphd}, a 0.6 $\times$ $^6$ ${\rm M}_{\odot}$ cloud need only move at a velocity of $\sim$ -350 in the LSR frame to attain the estimated kinetic energy to create the observed loop."
 Such a velocity is not unusual for many of the known UVC population (c.g.Wakker.Oster-loo&Putman2002:ctal. 2002).," Such a velocity is not unusual for many of the known HVC population \citep[e.g.][]{wop2002,putman2002}."
. Recently Bekki&Chiba(2006) have found that massive sub-halos can create kpe-scale giant. HIE holes such as those seen in the Western Alaecllanic Bridge., Recently \cite{bekki06} have found that massive sub-halos can create kpc-scale giant HI holes such as those seen in the Western Magellanic Bridge.
 These results imply that if the ALB interacted with low-mass subhalos. which are predicted by a hierarchical clustering scenario to be ubiquitous in the Galactic halo region. giant LIE holes can be formed.," These results imply that if the MB interacted with low-mass subhalos, which are predicted by a hierarchical clustering scenario to be ubiquitous in the Galactic halo region, giant HI holes can be formed."
 We plan to investigate this possibility in our forthcoming papers., We plan to investigate this possibility in our forthcoming papers.
Cosmological gamaua-ray bursts are the most relativistic events lu the Universe. at au observed rate of about two pe ‘day.,"Cosmological gamma-ray bursts are the most relativistic events in the Universe, at an observed rate of about two per day."
 The BATSE catalogue shows a bi-mocal distribution of GRB durations. with short bursts of about a secoucl aud long bursts of about oue ninute (Ixouveliotou .," The BATSE catalogue shows a bi-modal distribution of GRB durations, with short bursts of about a second and long bursts of about one minute \citep{kou93,pac99}."
. These events are probably associated witi the formation of high-aneilar momenttun low-tass black hole-torIs systelus representing liyperi0vae (Woosley1993:Paczy iusar-forming regious (Paczyüskü1995:Bloometal.2000).," These events are probably associated with the formation of high-angular momentum low-mass black hole-torus systems representing hypernovae \citep{woo93,pac97,pac98,bro00} in star-forming regions \citep{pac98,blo00}."
. General arguime1s suggest tliat the rotational energy ina blacx hole can be similarly imporal as accretion. provided there is a charuel in place to tap this euergy.," General arguments suggest that the rotational energy in a black hole can be similarly important as accretion, provided there is a channel in place to tap this energy."
 A maximally rotating black hole of mass Al~TAL. σοιtains about 2:I. iu rotational energy., A maximally rotating black hole of mass $M\sim7M_\odot$ contains about $2M_\odot$ in rotational energy.
 RaleieLs criterion (e.©... (1990))) remiuds s that roating systems lave a teicdeucy to shed of angular momentuui 1o larger distances. ai da rotating black hole is no exception.," Raleigh's criterion (e.g., \cite{kip90}) ) reminds us that rotating systems have a tendency to shed off angular momentum to larger distances, and a rotating black hole is no exception."
 This may be seell from the first law of black hole thermocvilaics., This may be seen from the first law of black hole thermodynamics.
"For a black |ole of mass A. angular momentui Jy aud augular velocity Qy; we have (Barceenetal.1973:Hayvkine1975.1976): 6034=OydA+Ty 657. where Ty, the horiZO temperature a dSy (Sj> 0)te entropy.","For a black hole of mass $M$, angular momentum $J_H$ and angular velocity $\Omega_H$ we have \citep{bar73,haw75,haw76}: $\delta M=\Omega_H\delta J_H+T_H\delta S_H$ , where $T_H$ the horizon temperature and $S_H$ $\delta S_H\ge0)$ the entropy."
" Thus. the specific angula “OMeun ofa radiated paricle exceeds that of the black hole: e,>1/04,=2MMJu/=a. Aneular nmonuentunm ls sicwed αἱ a lower cos in radiation than in the black hoe. ald is emittect with ai elliciency of at most50%."," Thus, the specific angular momentum $a_p=\delta J_H/\delta M$ of a radiated particle exceeds that of the black hole: $a_p\ge1/\Omega_H\ge 2M> M\ge J_H/M=a.$ Angular momentum is stored at a lower cost in radiation than in the black hole, and is emitted with an efficiency of at most."
. In vacuuu. the decay of a Ixerr black hole 1ito Seliwarzschild black hole is safeguardecd by ai angular momenutum barrier. aud spontaneous emissiou of particles is expoientially small (Teukolsky19723:PressaudTeukolsky1973:197 1).," In vacuum, the decay of a Kerr black hole into Schwarzschild black hole is safeguarded by an angular momentum barrier, and spontaneous emission of particles is exponentially small \citep{teu73,pre73,teu74}."
. Yet. black holes formed in hyperuovae will be surrounded by maguetizecl matter. in the form of au," Yet, black holes formed in hypernovae will be surrounded by magnetized matter, in the form of an"
(admittedly. not too strong) support for the optimism comes from the detection of a possible GLMXB near the turnoff of M30 (Luggeretal.2006).,"(admittedly, not too strong) support for the optimism comes from the detection of a possible qLMXB near the turnoff of M30 \citep{lug06}."
. Main-sequence secondaries similar to those in VI7 and V36 are also found in galactic X-ray binaries whose primary components are black hole candidates: e.g. VI033 Sco. GRS 1739-278 or V821 Ara (Ziolkowski. priv.," Main-sequence secondaries similar to those in V17 and V36 are also found in galactic X-ray binaries whose primary components are black hole candidates; e.g. V1033 Sco, GRS 1739–278 or V821 Ara (Ziolkowski, priv."
 comm)., comm).
 The fact that all our blue stragglers except V240 are binary supports the currently leading hypothesis concerning the nature of these objects. namely that they result rather from an eXtensive mass exchange between the component stars than from stellar collisions (Kniggeetal.2009).," The fact that all our blue stragglers except V240 are binary supports the currently leading hypothesis concerning the nature of these objects, namely that they result rather from an extensive mass exchange between the component stars than from stellar collisions \citep{knig09}."
. In all cases the brighter component is the primary which according to the mass-transfer scenario must have acquired significant amounts of hydrogen-rich material from the envelope of the originally more massive secondary., In all cases the brighter component is the primary which according to the mass-transfer scenario must have acquired significant amounts of hydrogen-rich material from the envelope of the originally more massive secondary.
 In fact. our results seem to indicate that blue stragglers are Algol-like systems in which the original mass ratio has been reversed. causing the mass transfer to effectively stop.," In fact, our results seem to indicate that blue stragglers are Algol-like systems in which the original mass ratio has been reversed, causing the mass transfer to effectively stop."
 Our stragglers may be similar to the well-studied star V228 in 47 Tue (Kaluznyetal.2007).. however we cannot exclude the possibility that in some of them the primary is already at the beginning of the subgiant branch and has become large enough to approach its Roche lobe.," Our stragglers may be similar to the well-studied star V228 in 47 Tuc \citep{kal07}, however we cannot exclude the possibility that in some of them the primary is already at the beginning of the subgiant branch and has become large enough to approach its Roche lobe."
 V2]4 and V254 in w Cen have primaries with masses likely exceeding | M4 (maybe even 1.5 M). and their total masses approach 2 M.," V214 and V254 in $\omega$ Cen have primaries with masses likely exceeding 1 $M_\odot$ (maybe even 1.5 $M_\odot$ ), and their total masses approach 2 $M_\odot$."
 Systems that massive must have originated from binaries subject to conservative or nearly conservative mass transfer., Systems that massive must have originated from binaries subject to conservative or nearly conservative mass transfer.
 They also have large amplitudes of light variations (0.12. mag and 0.10 mag. respectively). in our sample rivaled only by that of NV33 (~0.3 mag): however a significant contribution to this effect must originate from relatively large inclinations found for all three systems.," They also have large amplitudes of light variations (0.12, mag and 0.10 mag, respectively), in our sample rivaled only by that of NV334 $\sim$ 0.3 mag); however a significant contribution to this effect must originate from relatively large inclinations found for all three systems."
 The remaining blue stragglers are significantly less massive — theirtotal masses are lower than | M., The remaining blue stragglers are significantly less massive – their masses are lower than 1 $M_\odot$.
 As far as we know. our results are the first to indicate such a broad (and possibly bimodal) mass distribution.," As far as we know, our results are the first to indicate such a broad (and possibly bimodal) mass distribution."
 If this effect is real. it probably reflects differences in the efficiency of mass transfer and/or mass loss.," If this effect is real, it probably reflects differences in the efficiency of mass transfer and/or mass loss."
 We note in passing that some mass may still be flowing between the components of V20. as this system exhibits an asymmetric and variable light. curve. and is a weak X-ray source.," We note in passing that some mass may still be flowing between the components of V20, as this system exhibits an asymmetric and variable light curve, and is a weak X-ray source."
 Another possibility to account for the apparent large width of the mass distribution is related to the peculiarities of the chemical composition of ω Cen. in which strongly He-enriched subpopulations are suggested to exist. with Y reaching up to ~0.4 (e.g.. Norris 2004:; Piottoetal. 2005:; but see also Catelanetal. 2010 for a recent review and some caveats).Our fits are based on the standard value Y=0.24. and for a given B—V color they can overestimate the temperature of enriched stars by up to a few hundred Kelvin (Dotteretal.2008).," Another possibility to account for the apparent large width of the mass distribution is related to the peculiarities of the chemical composition of $\omega$ Cen, in which strongly He-enriched subpopulations are suggested to exist, with $Y$ reaching up to $\sim0.4$ (e.g., \citeauthor{nor04} \citeyear{nor04}; \citeauthor{pio05} \citeyear{pio05}; but see also \citeauthor{mcea10} \citeyear{mcea10} for a recent review and some caveats).Our fits are based on the standard value $Y=0.24$, and for a given $B-V$ color they can overestimate the temperature of enriched stars by up to a few hundred Kelvin \citep{dot08}."
. If this is what happened. then our standard analysis caused the He-rich systems to artificially shrink both in size and mass.," If this is what happened, then our standard analysis caused the He-rich systems to artificially shrink both in size and mass."
 Such hypothesis would also explain why stragglers with signficantly different “apparent” masses are able to form a relatively tight group on the H-R diagram., Such hypothesis would also explain why stragglers with signficantly different “apparent” masses are able to form a relatively tight group on the H-R diagram.
 The subdwarfs NV33 and V27 are very similar to each other in that they have large mass ratios and low-mass primaries with practically the same effective temperatures., The subdwarfs NV334 and V27 are very similar to each other in that they have large mass ratios and low-mass primaries with practically the same effective temperatures.
 Our estimate of the primary's mass in NV334 is consistent with the canonical hot subdwarf mass of 0.47 M. while the estimated mass of the secondary falls near the peak of the mass distribution of the unseen companions to field hot subdwarf stars (Geieral.2009).," Our estimate of the primary's mass in NV334 is consistent with the canonical hot subdwarf mass of 0.47 $M_\odot$, while the estimated mass of the secondary falls near the peak of the mass distribution of the unseen companions to field hot subdwarf stars \citep{gei09}."
. On the other hand. and as pointed out by Montal.(2008).. the close binary fraction among subdwarf stars appears to be much lower in GCs than in the field. which may point to different formation mechanisms for many of the GC subdwarfs.," On the other hand, and as pointed out by \citet{mbea08}, the close binary fraction among subdwarf stars appears to be much lower in GCs than in the field, which may point to different formation mechanisms for many of the GC subdwarfs."
" In V27 the estimated masses are significantly smaller (see Table 12)); note however that they are also less reliable due to problems discussed in Section 3.3.2.. and that the solution with mw,=0.57 M. and m»=0.11 Ma is also acceptable."," In V27 the estimated masses are significantly smaller (see Table \ref{tab: V27}) ); note however that they are also less reliable due to problems discussed in Section \ref{sect: NGC6397}, and that the solution with $m_1 = 0.57$ $M_\odot$ and $m_2 = 0.11$ $M_\odot$ is also acceptable."
 Taken at face value. the starred parameters in Table 12 would indicate that V27 is very similar to HS 223142241. an HW Vir-type system with a companion at or below the very low mass limit for M-dwarfs.," Taken at face value, the starred parameters in Table \ref{tab: V27} would indicate that V27 is very similar to HS 2231+2241, an HW Vir-type system with a companion at or below the very low mass limit for M-dwarfs."
 The ~0.26 M. primary of HS 223]42241] is a non-helium burning. post-RGB star (Ostensenetal.2008).," The $\sim0.26$ $M_\odot$ primary of HS 2231+2241 is a non-helium burning, post-RGB star \citep{ost08}."
. Another example of such a type of star is the 0.24 M.. component of HD 188112 (Heberetal.2003)., Another example of such a type of star is the 0.24 $M_\odot$ component of HD 188112 \citep{heb03}.
. The primaries of NV334 and V27 are much cooler (7~7500 K compared to 7=28400 K and 7=20500 K. respectively. for HS 2231-224] and HD 188112). but this difference may result from their more advanced evolutionary stage.," The primaries of NV334 and V27 are much cooler $T\sim7500$ K compared to $T = 28400$ K and $T = 20500$ K, respectively, for HS 2231+2241 and HD 188112), but this difference may result from their more advanced evolutionary stage."
 Our sample is too small for a meaningful statistical analysis: nevertheless the above-discussed results indicate that further research on ellipsoidal binaries in globular clusters is a wortwhile task. and suggest its most promising direction: 1t seems that future sampling of the photometric variables should be biased towards objects located on or to the right of the main sequence. including those identified as weak X-ray sources.," Our sample is too small for a meaningful statistical analysis; nevertheless the above-discussed results indicate that further research on ellipsoidal binaries in globular clusters is a wortwhile task, and suggest its most promising direction: it seems that future sampling of the photometric variables should be biased towards objects located on or to the right of the main sequence, including those identified as weak X-ray sources."
 Obviously. for purely economical reasons (exposure times!)," Obviously, for purely economical reasons (exposure times!)"
 the targets sould not fall much below the turnoff point., the targets sould not fall much below the turnoff point.
 Such an approach ts tedious. but at the same time it is the only one on which a fair census of degenerate binaries in globular clusters can be based.," Such an approach is tedious, but at the same time it is the only one on which a fair census of degenerate binaries in globular clusters can be based."
"which is calculated by assuming all unmeasured emission to be zeros, is shown in figure](b).","which is calculated by assuming all unmeasured emission to be zeros, is shown in figure (b)."
 The result of RM-CLEAN is shown in figure2 (c)., The result of RM-CLEAN is shown in figure (c).
 Note that RM-CLEAN cannot correctly reconstruct the magnitude or phase of the Faraday dispersion function in this case., Note that RM-CLEAN cannot correctly reconstruct the magnitude or phase of the Faraday dispersion function in this case.
 It has been demonstrated previously that RM-CLEAN has difficulty with the separation of sources below the FWHM(?)., It has been demonstrated previously that RM-CLEAN has difficulty with the separation of sources below the FWHM.
" In figure| (a), from left to right, the distance between the first and the second sources is smaller than the FWHM."," In figure (a), from left to right, the distance between the first and the second sources is smaller than the FWHM."
" In figureEJ (c), the first source and second source are merged to form an unrealistic source."," In figure (c), the first source and second source are merged to form an unrealistic source."
" Results of CS-RM-Thin, CS-RM-Thick and CS-RM-Mix are shown in the figure] (e) and (f), respectively."," Results of CS-RM-Thin, CS-RM-Thick and CS-RM-Mix are shown in the figure (d), (e) and (f), respectively."
" As one might expect, the result of (d),CS-RM-Thick is poor, because the Faraday dispersion function has no Faraday thick sources."," As one might expect, the result of CS-RM-Thick is poor, because the Faraday dispersion function has no Faraday thick sources."
" In figure[5] though CS-RM-Mix does a better job than CS-RM-Thick and RM-CLEAN in terms of magnitude and phase of the reconstructed disperse functions, the result is far from good enough."," In figure, though CS-RM-Mix does a better job than CS-RM-Thick and RM-CLEAN in terms of magnitude and phase of the reconstructed disperse functions, the result is far from good enough."
 The result of CS-RM-Thin is shown in figure[2) (d)., The result of CS-RM-Thin is shown in figure ) (d).
" We can see that CS-RM-Thin gives the best result, reconstructing the Faraday dispersion function without any error."," We can see that CS-RM-Thin gives the best result, reconstructing the Faraday dispersion function without any error."
" This is consistent with CS theory which says that we can reconstruct the sparse signal exactly with “overwhelming probability""(???)."," This is consistent with CS theory which says that we can reconstruct the sparse signal exactly with “overwhelming probability""."
" In this test, we select ¢g=3.6 m? which is around one quarter of FWHM, so N=480."," In this test, we select $\phi_R=3.6$ $^2$ which is around one quarter of FWHM, so $N=480$."
" To carry out a numerical comparison, we use the root mean square (RMS) error to characterise the difference between the reconstructed f and the Faraday dispersion function f: The RMS error is calculated for all the candidate methods, and the results are listed in table["," To carry out a numerical comparison, we use the root mean square (RMS) error to characterise the difference between the reconstructed $\widetilde{\mathbf{f}}$ and the Faraday dispersion function $\mathbf{f}$: The RMS error is calculated for all the candidate methods, and the results are listed in table."
I] Results for this test can be found in the first row of the table., Results for this test can be found in the first row of the table.
 CS-RM-Thin gives the best result., CS-RM-Thin gives the best result.
 We now test our CS-based methods for a Faraday dispersion function with Faraday thick sources., We now test our CS-based methods for a Faraday dispersion function with Faraday thick sources.
 Here we assume that the Faraday dispersion function includes two sources 2-2i Jy n? rad! where -120<à40 and F(¢)=—-6—3i Jy m? rad“! where 30<$70., Here we assume that the Faraday dispersion function includes two sources $F(\phi)=2-2\mathrm{i}$ Jy $^2$ $^{-1}$ where $-120\leq \phi \leq 40$ and $F(\phi)=-6-3\mathrm{i}$ Jy $^2$ $^{-1}$ where $30\leq \phi \leq 70$.
 The simulated Faraday dispersion function is shown in figure[3] (a)., The simulated Faraday dispersion function is shown in figure (a).
 We adopt the previous observing window for this test and the dirty Faraday dispersion function is shown in figure (b)., We adopt the previous observing window for this test and the dirty Faraday dispersion function is shown in figure (b).
 Note that thisonly provides us with the approximate B]shape of the Faraday dispersion function., Note that thisonly provides us with the approximate shape of the Faraday dispersion function.
" As mentioned in?, the magnitude of F(¢) indicates the polarized emission of the region with Faraday depth ¢ and its phase defines the intrinsic position angle."," As mentioned in, the magnitude of $F(\phi)$ indicates the polarized emission of the region with Faraday depth $\phi$ and its phase defines the intrinsic position angle."
" For the study of polarized emission of galaxies, the magnitude of F(Q) is very important, and for the study of orientation of the magnetic field perpendicular to the line of sight, the phase information of F($) is crucial."," For the study of polarized emission of galaxies, the magnitude of $F(\phi)$ is very important, and for the study of orientation of the magnetic field perpendicular to the line of sight, the phase information of $F(\phi)$ is crucial."
" Unfortunately, Brentjens & de Bruyn' method can not reconstruct reliable phase information for F(9)."," Unfortunately, Brentjens & de Bruyn' method can not reconstruct reliable phase information for $F(\phi)$."
 The cleaned version is shown in figureB] (c)., The cleaned version is shown in figure (c).
 RM- also failed to reconstruct the phase information., RM-CLEAN also failed to reconstruct the phase information.
" RM-Thin does not work well for this case, because F(9) is not sparse."," CS-RM-Thin does not work well for this case, because $F(\phi)$ is not sparse."
 The result of CS-RM-Thin is shown in figureB](d)., The result of CS-RM-Thin is shown in figure (d).
 CS-RM-Mix is also used for this test by assuming that there are both Faraday thin sources and thick sources., CS-RM-Mix is also used for this test by assuming that there are both Faraday thin sources and thick sources.
 The result of CS-RM-Mix is shown in figureB] (f)., The result of CS-RM-Mix is shown in figure (f).
 The reconstructed result from CS-RM-Thick is shown in figure[3] (e)., The reconstructed result from CS-RM-Thick is shown in figure (e).
" Even though CS-RM-Mix gives a much better result than CS-RM-Thin and RM-CLEAN, the result is not as good as that of CS-RM-Thick."," Even though CS-RM-Mix gives a much better result than CS-RM-Thin and RM-CLEAN, the result is not as good as that of CS-RM-Thick."
 CS-RM-Thick provides the best approximation to the original Faraday dispersion function in terms of both magnitude and phase., CS-RM-Thick provides the best approximation to the original Faraday dispersion function in terms of both magnitude and phase.
 This is also supported by the numerical comparison in table[, This is also supported by the numerical comparison in table.
"I] From the second row of the table, we can see that CS- gives the smallest RMS error 0.72 which is slightly better than that of CS-RM-Mix 0.77."," From the second row of the table, we can see that CS-RM-Thick gives the smallest RMS error $0.72$ which is slightly better than that of CS-RM-Mix $0.77$."
" So far we have tested our methods for both Faraday thin sources and Faraday thick sources, then we will test our CS-based methods for the mixed circumstances - both Faraday thin sources and thick sources."," So far we have tested our methods for both Faraday thin sources and Faraday thick sources, then we will test our CS-based methods for the mixed circumstances - both Faraday thin sources and thick sources."
 The simulated Faraday dispersion function for this test is shown in the top left corner of figure[, The simulated Faraday dispersion function for this test is shown in the top left corner of figure.
ῇ There are three sources along the line of sight., There are three sources along the line of sight.
" From left to right, these sources are: F(-58)=-4+3i Jy m? rad, F(-30)=10-4i Jy m? rad!, ΕΦ)=2—6i Jy m? rad""! where 41<¢100."," From left to right, these sources are: $F(-58)=-4+3\mathrm{i}$ Jy $^2$ $^{-1}$, $F(-30)=10-4\mathrm{i}$ Jy $^2$ $^{-1}$, $F(\phi)=2-6\mathrm{i}$ Jy $^2$ $^{-1}$ where $41\leq \phi \leq 100$."
 The original Faraday dispersion function is shown in the top left corner of figureD, The original Faraday dispersion function is shown in the top left corner of figure.
] We use the previous observing window for this test., We use the previous observing window for this test.
" Based on the observing window (with 126 observing channels distributed between 0.036 m to 0.5 m), the RMSF is calculated and shown in figureβ (0)."," Based on the observing window (with 126 observing channels distributed between 0.036 m to 0.5 m), the RMSF is calculated and shown in figure (b)."
 the dirty curve is shown in figure[] (c)., the dirty curve is shown in figure (c).
 The cleaned version by RM-CLEAN is shown in figure[4] (d)., The cleaned version by RM-CLEAN is shown in figure (d).
" RM-CLEAN performs badly in this test, because there is a Faraday thick source."," RM-CLEAN performs badly in this test, because there is a Faraday thick source."
" In general, RM-CLEAN can only work well when there are Faraday thin sources along the line of sight."," In general, RM-CLEAN can only work well when there are Faraday thin sources along the line of sight."
" The results of CS-RM-Thin and CS-RM-Thick are shown in figure] (e) and (f), respectively."," The results of CS-RM-Thin and CS-RM-Thick are shown in figure (e) and (f), respectively."
" We can see that CS-RM-Thin reconstructs the two Faraday thin sources nicely, but failed to reconstruct the Faraday thick source."," We can see that CS-RM-Thin reconstructs the two Faraday thin sources nicely, but failed to reconstruct the Faraday thick source."
" On the contrary, CS-RM-Thick can properly reconstruct the Faraday thick source, but expends the two Faraday thin sources by mistake."," On the contrary, CS-RM-Thick can properly reconstruct the Faraday thick source, but expends the two Faraday thin sources by mistake."
" As introduced above, CS-RM-Mix can separate the Faraday thin components fi, and the Faraday thick components fiy; during the reconstruction."," As introduced above, CS-RM-Mix can separate the Faraday thin components $\mathbf{f}_{\rm{thin}}$ and the Faraday thick components $\mathbf{f}_{\rm{thick}}$ during the reconstruction."
" In this test, the soft-threshold t=1 and 6=0.001 in the proposed algorithm for CS-RM-Mix."," In this test, the soft-threshold $\tau=1$ and $\delta=0.001$ in the proposed algorithm for CS-RM-Mix."
 The results of Mix are shown in the last row of figure|, The results of CS-RM-Mix are shown in the last row of figure.
" The separated Faraday components fipin and finick are shown in figure Al (g) and (h), respectively."," The separated Faraday components $\mathbf{f}_{\rm{thin}}$ and $\mathbf{f}_{\rm{thick}}$ are shown in figure (g) and (h), respectively."
 The sum of the separated components is the reconstructed Faraday dispersion function which is shown in figure B] (i)., The sum of the separated components is the reconstructed Faraday dispersion function which is shown in figure (i).
 We can see that the result of CS-RM-Mix takes advantage of the results of both CS-RM-Thin and CS-RM- and it gives the closestapproximation to the original F(@¢).," We can see that the result of CS-RM-Mix takes advantage of the results of both CS-RM-Thin and CS-RM-Thick, and it gives the closestapproximation to the original $F(\phi)$ ."
" From objective evaluation point of view, CS-RM-Mix gives the minimum RMS error 0.80 which can be seen from the third row of table ["," From objective evaluation point of view, CS-RM-Mix gives the minimum RMS error $0.80$ which can be seen from the third row of table ."
" From objective evaluation point of view, CS-RM-Mix gives the minimum RMS error 0.80 which can be seen from the third row of table [I"," From objective evaluation point of view, CS-RM-Mix gives the minimum RMS error $0.80$ which can be seen from the third row of table ."
" From objective evaluation point of view, CS-RM-Mix gives the minimum RMS error 0.80 which can be seen from the third row of table [I]"," From objective evaluation point of view, CS-RM-Mix gives the minimum RMS error $0.80$ which can be seen from the third row of table ."
[function tells us about the tvpical conditions in which disk stars form. namely that over the lifetime of the Galactic disk. most stars formed in clusters with velocity dispersions of order a lew hundred meters per second aud that (hese clusters disrupt on timescales of order a [ew hundred million vears.,"function tells us about the typical conditions in which disk stars form, namely that over the lifetime of the Galactic disk, most stars formed in clusters with velocity dispersions of order a few hundred meters per second and that these clusters disrupt on timescales of order a few hundred million years."
 Figure 10 of Chanamé&Gould(2004) hints that the halo binary breakpoint occurs at al least a somewhat smaller semi-major axis (han Chat of the disk., Figure 10 of \citet{chaname04} hints that the halo binary breakpoint occurs at at least a somewhat smaller semi-major axis than that of the disk.
 If this is confirmed by future observations. i would imply that halo stars were born in environments that were kinematically at least slightly hotter than disk stars. (," If this is confirmed by future observations, it would imply that halo stars were born in environments that were kinematically at least slightly hotter than disk stars. ("
Note that ib is also possible. in principle. to produce a smaller ej; bv selecting a sample of halo binaries with svstematicallv smaller companion masses.,"Note that it is also possible, in principle, to produce a smaller $a_{\rm break}$ by selecting a sample of halo binaries with systematically smaller companion masses."
 However. (he POOL sample actually has the opposite bias: since the halo stars are a [actor ~4 further away. (hev are actually biased toward being more luminous and so more massive — stars.)," However, the \citealt{chaname04} sample actually has the opposite bias: since the halo stars are a factor $\sim 4$ farther away, they are actually biased toward being more luminous – and so more massive – stars.)"
 Detailed verification of the scenario presented here will require simulations (hat simultaneously model the disruption of the binary. and of the cluster in which it initially resides., Detailed verification of the scenario presented here will require simulations that simultaneously model the disruption of the binary and of the cluster in which it initially resides.
 II verified. (he same simulations will permit more precise characterization of the typical clusters in which today’s field binaries were born.," If verified, the same simulations will permit more precise characterization of the typical clusters in which today's field binaries were born."
 We (hank Scott Gaudi for valuable comments on (he manuscript., We thank Scott Gaudi for valuable comments on the manuscript.
 This work was supported by NSF grant AST 042758., This work was supported by NSF grant AST 042758.
 Any opinions. findings. and conclusions or recommendations expressed in (his material are those of the authors and do not necessarily rellect the views of the NSF.," Any opinions, findings, and conclusions or recommendations expressed in this material are those of the authors and do not necessarily reflect the views of the NSF."
We also verified that the non-gray logy corrections are fairly independent of the Ty corrections by redoing our fits at fixed temperatures (Te=Tap).,We also verified that the non-gray $\log g$ corrections are fairly independent of the $T_{\rm eff}$ corrections by redoing our fits at fixed temperatures $T_{\rm eff}=T_{\rm 3D}$ ).
 The corrections we obtain with this method are within ~20% of those found in Fig. 5.., The corrections we obtain with this method are within $\sim20$ of those found in Fig. \ref{fg:f5}.
 Finally. the fact that we have a set of 3D gray models allows us to evaluate the impact on our results of the opacity binning procedure in the non-gray models.," Finally, the fact that we have a set of 3D gray models allows us to evaluate the impact on our results of the opacity binning procedure in the non-gray models."
 The shifts in logy that we obtain with | opacity bin (gray models) are smaller than those with 7 bins (non-gray models). but still correct to a large degree the high-logy problem.," The shifts in $\log g$ that we obtain with 1 opacity bin (gray models) are smaller than those with 7 bins (non-gray models), but still correct to a large degree the $\log g$ problem."
 This shows that further improvements in the binning procedure of the RHD models are not expected to change our differential analysis significantly., This shows that further improvements in the binning procedure of the RHD models are not expected to change our differential analysis significantly.
 We conducted a differential analysis. of synthetic. spectra for cool convective DA white dwarfs computed from model atmospheres with a detalled 3D radiation-hydrodynamical treatment of convective motions. and with the standard 1D mixing-length theory.," We conducted a differential analysis of synthetic spectra for cool convective DA white dwarfs computed from model atmospheres with a detailed 3D radiation-hydrodynamical treatment of convective motions, and with the standard 1D mixing-length theory."
 We introduced the concept of 3D atmospheric parameter corrections. which are defined as the difference in Tey and logg between {νι and 3D spectra when fitted with an independent grid.," We introduced the concept of 3D atmospheric parameter corrections, which are defined as the difference in $T_{\rm eff}$ and $\log g$ between $_{\rm MLT}$ and 3D spectra when fitted with an independent grid."
 These corrections illustrate by how much the atmospheric parameters of DA stars would change if we were to use a grid of 3D spectra to fit the observations., These corrections illustrate by how much the atmospheric parameters of DA stars would change if we were to use a grid of 3D spectra to fit the observations.
 We find that the 3D surface gravity corrections have the correct amplitude (~ 0.15 dex) to solve the high-logy problem for DA white dwarfs in the range 11.300 K <Tey 12.800 K. This indubitably suggests that a weakness with the standard ΜΙΤ theory ts creating the high logg values reported in recent spectroscopic analyses (e.g.. TBIO).," We find that the 3D surface gravity corrections have the correct amplitude $\sim$ $0.15$ dex) to solve the $\log g$ problem for DA white dwarfs in the range 11,300 K $< T_{\rm eff} <$ 12,800 K. This indubitably suggests that a weakness with the standard MLT theory is creating the high $\log g$ values reported in recent spectroscopic analyses (e.g., TB10)."
 Our next goal will be to review the microphysics and the opacity binning procedure in the CO°BOLD code to connect the absolute properties of these hydrodynamical models with those of standard white dwarf models., Our next goal will be to review the microphysics and the opacity binning procedure in the $^5$ BOLD code to connect the absolute properties of these hydrodynamical models with those of standard white dwarf models.
 We will then be able to compute improved 3D model spectra that can actually be compared with observations to determine the atmospheric parameters., We will then be able to compute improved 3D model spectra that can actually be compared with observations to determine the atmospheric parameters.
 We also plan to look at convective white dwarfs cooler than those studied in this work. where the MLT theory predicts an increasingly adiabatic convection. which is less sensitive to the mixing-length parameter.," We also plan to look at convective white dwarfs cooler than those studied in this work, where the MLT theory predicts an increasingly adiabatic convection, which is less sensitive to the mixing-length parameter."
 It remains to be seen how the 3D logg corrections would behave in this regime., It remains to be seen how the 3D $\log g$ corrections would behave in this regime.
he 3D correlation of Eq. (16)).,the 3D correlation of Eq. \ref{eq:GRB-3D_1}) ).
 So. in addition to exanune the quality of the lhunuinositv relations * colparing their intrinsic scatters. we also calculated the correlation coefficients between ιο residual of the fit of 2D correlations aud ie possible hidden parameters we introduced. to xteud correlations from 2D to 3D. Iu the fit of the Iuninositv relations. we used 1ο techniques presented in D’Agostini(2005).. ‘ollowing which the likelihood function for the cocfiicicuts ο and the iutrinsic scatter ds (for re cases of Eq. (16)).," So, in addition to examine the quality of the luminosity relations by comparing their intrinsic scatters, we also calculated the correlation coefficients between the residual of the fit of 2D correlations and the possible hidden parameters we introduced to extend correlations from 2D to 3D. In the fit of the luminosity relations, we used the techniques presented in \citet{Agostini:2005fe}, following which the likelihood function for the coefficients $c$ and the intrinsic scatter is (for the cases of Eq. \ref{eq:GRB-3D_1}) )."
 The other cases of 3D ununuositv relations are simular aud the likelihood πιοΊο for the 2D luminosity relations can be obtained just by setting e9= 0) where ᾱ runs over GRBs with corresponding quantities available., The other cases of 3D luminosity relations are similar and the likelihood function for the 2D luminosity relations can be obtained just by setting $c_2 = 0$ ) where $k$ runs over GRBs with corresponding quantities available.
 In the calculation. Markov chain Monte Carlo techniques are used.," In the calculation, Markov chain Monte Carlo techniques are used."
 For each huuinositv relations. A Markov chain with samples of order 1089 is generated according to the likelihood. function aud then properly burned in and thinned to derive statistics of interested parameters.," For each luminosity relations, A Markov chain with samples of order $10^6$ is generated according to the likelihood function and then properly burned in and thinned to derive statistics of interested parameters."
 For theCRD data. we used the compilation πι Schaefer(2007).. which includes 69 CRBs.," For theGRB data, we used the compilation in \citet{Schaefer:2006pa}, which includes 69 GRBs."
When cousideriug error propagation from a quantity. sav € with error σε. to its loganithnun we set logiétol)5losec.) aud offal):2lowiga.) as the center value and the error of the logaritlin correspondinely.,"When considering error propagation from a quantity, say $\xi$ with error $\sigma_{\xi}$, to its logarithm, we set $
\frac
{
 \log(\xi + \sigma_{\xi}^{+})
 +
 \log(\xi - \sigma_{\xi}^{-})
}
{
 2
}
$ and $
\frac
{
 \log(\xi + \sigma_{\xi}^{+})
 -
 \log(\xi - \sigma_{\xi}^{-})
}
{
 2
}
$ as the center value and the error of the logarithm correspondingly."
" This requies &£7σι, (the quantities we are interested in here ave all positive).", This requires $\xi > \sigma_{\xi}^{-}$ (the quantities we are interested in here are all positive).
" Due to the linütation of the data. for a eiven Iuminositv relation (7.7). uot all the CRBs have all ofthe needed observational quantities available and satisfy ©20, at the same time."," Due to the limitation of the data, for a given luminosity relation $(i, \, j)$, not all the GRBs have all of the needed observational quantities available and satisfy $\xi > \sigma_{\xi}^{-}$ at the same time."
 By set (7.7) we denote the maxim CRB set that can be used in the huminositv relation (7. 7). The wuubers of GRBs of differcut sets are presented im Table 1..," By set $(i, \, j)$ we denote the maximum GRB set that can be used in the luminosity relation $(i, \, j)$ The numbers of GRBs of different sets are presented in Table \ref{tab:fit}. ."
The results for the thermal background/power-law-jet are shown in fie. 2..,The results for the thermal background/power-law-jet are shown in fig. \ref{fig:xspec}.
" From) this fit. we fud that the power law compoucut has a fiux. f(0.3-7keV) = 1.5.10tere 28 banda =0,8+0.3."," From this fit, we find that the power law component has a flux, f(0.3-7keV) = $\times~10^{-14}$ erg $^{-2}$ $^{-1}$ and $\alpha=0.8~\pm~0.3$ ."
 This flux correspouds to 12 not counts in our data., This flux corresponds to 42 net counts in our data.
 Dividing the observed flux according to the ratio of he counts in two small circles centered ou the wo knots provides the intensity values reported im table 1.., Dividing the observed flux according to the ratio of the counts in two small circles centered on the two knots provides the intensity values reported in table \ref{tab:inten}.
 The VLA data are a combination of observations made in three VLA configurations at 1.9 Πε: in the A configuration on 1980 November 9. iu the D coufieuration ou 1981 June 25. aud iu the C configuration on 1981 November 17.," The VLA data are a combination of observations made in three VLA configurations at 4.9 GHz: in the A configuration on 1980 November 9, in the B configuration on 1981 June 25, and in the C configuration on 1981 November 17."
 They were reduced iu the AIPS software package usine standard selfcealibration aud maging methods., They were reduced in the AIPS software package using standard self-calibration and imaging methods.
 The asvinmuetry in brightucss between the northern and southern jets supports the notion that the northern jet is the one coming towards us aud that Doppler favoritisu is operating out to z15” (1.2 kpc)., The asymmetry in brightness between the northern and southern jets supports the notion that the northern jet is the one coming towards us and that Doppler favoritism is operating out to $\approx~15^{\prime\prime}$ (1.2 kpc).
 We do not have radio data of sufficient. spatial resolution at lower frequencies to determine accurate spectral iudices. but the lower resolution data imdicate values αzz 0.65.," We do not have radio data of sufficient spatial resolution at lower frequencies to determine accurate spectral indices, but the lower resolution data indicate values $\alpha~\approx$ 0.65."
 To estimate physical parameters associated with various X-ray cluission mechanisuis. we need to assign volumes for the knots.," To estimate physical parameters associated with various X-ray emission mechanisms, we need to assign volumes for the knots."
" For cach knot we choose a cylindrical volume with leusth of 0.5"" and radii of 0.15"" (N2.5) aud 0.217 (N3.3).", For each knot we choose a cylindrical volume with length of $^{\prime\prime}$ and radii of $^{\prime\prime}$ (N2.5) and $^{\prime\prime}$ (N3.3).
 The overall spectiiun for N3.3 is shown iu figure 3.., The overall spectrum for N3.3 is shown in figure \ref{fig:spec}.
 Iuiuanv instances (0.8. IEuis. Carilli. Perley. 1991) it has been argued that if N-vayv ciission roni racio features were to be frou hot gas rather hana non-thermal mechanism. then the required electron densities together with the (equipartition) nagnetic field strengths would predict departures youn the A? relation of radio polarization position anele. as well as excess depolarization.," In many instances (e.g. Harris, Carilli, Perley, 1994) it has been argued that if X-ray emission from radio features were to be from hot gas rather than a non-thermal mechanism, then the required electron densities together with the (equipartition) magnetic field strengths would predict departures from the $\lambda^2$ relation of radio polarization position angle, as well as excess depolarization."
 Since hese effects have not been fouud. it appears that," Since these effects have not been found, it appears that"
shell is perhaps the forefront of the shock in a relatively low ISM density toward southwest. where the greatest magnetic field compression and particle acceleration occurs.,"shell is perhaps the forefront of the shock in a relatively low ISM density toward southwest, where the greatest magnetic field compression and particle acceleration occurs."
 The optical emission of Braun-101 is the smallest (3 arcsec: 11 pe) compared to other wavelengths and it is encompassed by the X-ray emission (Fig., The optical emission of Braun-101 is the smallest (3 arcsec; 11 pc) compared to other wavelengths and it is encompassed by the X-ray emission (Fig.
 2)., 2).
 The Ho shell roughly traces the inner X-ray ring of the remnant and shows a non-uniform distribution of emission over the shell. with slight enhancement toward the southwest.," The $\alpha$ shell roughly traces the inner X-ray ring of the remnant and shows a non-uniform distribution of emission over the shell, with slight enhancement toward the southwest."
 It also shows a crescent opening and possibly à double-ring structure in the northern region of the optical remnant., It also shows a crescent opening and possibly a double-ring structure in the northern region of the optical remnant.
 In the northeast. a local enhancement of the X-ray emission corresponds to a local optical enhancement.," In the northeast, a local enhancement of the X-ray emission corresponds to a local optical enhancement."
 The optical emission could be produced by shocks driven into cooler and denser material by the high pressure downstream of the reverse shock., The optical emission could be produced by shocks driven into cooler and denser material by the high pressure downstream of the reverse shock.
 It is worth noting that such large discrepancies in SNR size at different wavelengths are not typical (Winkler Long 1997)., It is worth noting that such large discrepancies in SNR size at different wavelengths are not typical (Winkler Long 1997).
 In this case. we may be seeing only the brightest optical knots and not the whole optical shell.," In this case, we may be seeing only the brightest optical knots and not the whole optical shell."
 Braun-101 is only 0.5 aresee from an optically identified planetary nebula candidate (Kong et al., Braun-101 is only 0.5 arcsec from an optically identified planetary nebula candidate (Kong et al.
 2002b: Ciardullo et al., 2002b; Ciardullo et al.
 1959)., 1989).
 Our multiwavelength observations provide strong evidence that the source is actually a SNR., Our multiwavelength observations provide strong evidence that the source is actually a SNR.
 As planetary nebulae are usually identified by their strong [O II] emission (Ciardullo et al., As planetary nebulae are usually identified by their strong [O III] emission (Ciardullo et al.
" 1989), they could resemble the [O III] emission from some SNRs."," 1989), they could resemble the [O III] emission from some SNRs."
" However. the large angular size (from 10 pe in the optical to 40 pe in the radio) and high X-ray luminosity (4.4<10? erg sv!) rule out the possibility of a planetary nebula for which the typical diameter and X-ray luminosity is <| pe (c0.25 aresee) and 10°"" erg s7!. respectively."," However, the large angular size (from 10 pc in the optical to 40 pc in the radio) and high X-ray luminosity $4.4\times10^{35}$ erg $^{-1}$ ) rule out the possibility of a planetary nebula for which the typical diameter and X-ray luminosity is $< 1$ pc $< 0.25$ arcsec) and $10^{30}$ erg $^{-1}$, respectively."
 Like Braun-101. Braun-95 is clearly à shell-like radio remnant (see Fig.," Like Braun-101, Braun-95 is clearly a shell-like radio remnant (see Fig."
 1) and its morphology is also very similar to Braun-101., 1) and its morphology is also very similar to Braun-101.
 The X-ray remnant of Braun-95. however. is patchy and£o is mainly located in the northeastern side of the radio remnant. where the radio emission is brightest.," The X-ray remnant of Braun-95, however, is patchy and is mainly located in the northeastern side of the radio remnant, where the radio emission is brightest."
 Faint X-ray emission is also seen near the radio shell., Faint X-ray emission is also seen near the radio shell.
 The X-ray spectrum likely consists of two components: a Raymond-Smith model and a power-law model., The X-ray spectrum likely consists of two components: a Raymond-Smith model and a power-law model.
 The Raymond-Smith temperature ts similar to Braun- and CXOM31 J004327.74+411829 (Kong et al., The Raymond-Smith temperature is similar to Braun-101 and CXOM31 J004327.7+411829 (Kong et al.
 2002a)., 2002a).
 The 101extra power-law component might hint that Braun-95 contains a pulsar and/or a pulsar nebula., The extra power-law component might hint that Braun-95 contains a pulsar and/or a pulsar nebula.
 Discovery of SNRs near the center of M31 provides some information about the star formation history of the region., Discovery of SNRs near the center of M31 provides some information about the star formation history of the region.
 For instance. type la SNRs suggest an old stellar population. while type Ib/II SNRs are associated with young population.," For instance, type Ia SNRs suggest an old stellar population, while type Ib/II SNRs are associated with young population."
 Unfortunately. we do not have direct evidence about the type for the two SNRs based on the current multi-wavelength data.," Unfortunately, we do not have direct evidence about the type for the two SNRs based on the current multi-wavelength data."
 By comparing with Galactic SNRs. the mismatch in morphology of the two SNRs in different wavelengths might hint the type.," By comparing with Galactic SNRs, the mismatch in morphology of the two SNRs in different wavelengths might hint the type."
 For example. IC 443 (Petre et al.," For example, IC 443 (Petre et al."
 1988: Kawasaki et al., 1988; Kawasaki et al.
 2002) and Ser A East (Maeda et al., 2002) and Sgr A East (Maeda et al.
 2002) are type II SNRs that have different morphology in radio and X-ray. and type Ia SNR like Tycho and SN 1006 generally have a well matched morphology in X-ray and radio.," 2002) are type II SNRs that have different morphology in radio and X-ray, and type Ia SNR like Tycho and SN 1006 generally have a well matched morphology in X-ray and radio."
 On the other hand. the X-ray and radio surface brightness of Tycho does not show correspondence (Hwang Gotthelf 1997). and the X-ray emission of SN 1006 is non-thermal (Long et al.," On the other hand, the X-ray and radio surface brightness of Tycho does not show correspondence (Hwang Gotthelf 1997), and the X-ray emission of SN 1006 is non-thermal (Long et al."
 2003). suggesting à synchrotron origin of X-rays.," 2003), suggesting a synchrotron origin of X-rays."
 In addition. type Ib/II SNRs such as Cas A (compare Hughes. et al.," In addition, type Ib/II SNRs such as Cas A (compare Hughes, et al."
 2000 and Braun. Gull Perley 1987; see also Keohane et al.," 2000 and Braun, Gull Perley 1987; see also Keohane et al."
 1996 and Dickel et al., 1996 and Dickel et al.
 1982) and RCW103 (compare Dickel et al., 1982) and RCW103 (compare Dickel et al.
 1996 and Tuohy Garmire 1980) also have a good match in different wavelengths., 1996 and Tuohy Garmire 1980) also have a good match in different wavelengths.
 Further multi-wavelength observaitons of these 2 SNRs are required to identify their type and hence the assoicated stellar population in the bulge of M31., Further multi-wavelength observaitons of these 2 SNRs are required to identify their type and hence the assoicated stellar population in the bulge of M31.
 The multiwavelength detection of SNRs in M31 can also provide further information about the local ISM density around the SNRs., The multiwavelength detection of SNRs in M31 can also provide further information about the local ISM density around the SNRs.
 The ISM near M3l's SNRs is thought to be low Cz0.1 em“). which is | to 2 orders of magnitude smaller than the measurements by the II observations (Magnier et al.," The ISM near M31's SNRs is thought to be low $< 0.1$ $^{-3}$ ), which is 1 to 2 orders of magnitude smaller than the measurements by the I observations (Magnier et al."
 1997)., 1997).
 Our discovery of X-ray emission from SNRs near the nucleus of M31 indicates that the ISM density in the vicinity of the SNRs could be higher., Our discovery of X-ray emission from SNRs near the nucleus of M31 indicates that the ISM density in the vicinity of the SNRs could be higher.
 We have shown that by combining all the high-resolution instruments across many wavebands. we can study the morphological structure of SNRs in M31 for the first time.," We have shown that by combining all the high-resolution instruments across many wavebands, we can study the morphological structure of SNRs in M31 for the first time."
 This is a first step toward understanding the evolutionary states and shock morphologies of SNRs in M31., This is a first step toward understanding the evolutionary states and shock morphologies of SNRs in M31.
 More specific multiwavelength observations in the future will shed more light on the properties of SNRs in M31. allowing more comprehensive comparative studies with our own Galaxy.," More specific multiwavelength observations in the future will shed more light on the properties of SNRs in M31, allowing more comprehensive comparative studies with our own Galaxy."
— 150A. Ed107 £ « 1077 ~ MC E « 1003 79 ALCENI)o ALL OOFAL.. ,$\sim$ $M_\odot$ $E \sim 1\times 10^{51}$ $E$ $E = (1.0$ $\times$ $10^{51}$ $\sim$ $M_\odot$ $E$ $ \times$ $10^{52}$ $^{56}$ $M(^{56}$$\simeq$ $M_\odot$ $M_\odot$ 
average of the two.,average of the two.
 We should stress that low frequency data are not simultaneous with the X/y-ray data that we selected., We should stress that low frequency data are not simultaneous with the $\gamma$ -ray data that we selected.
 The modeled synchrotron bump has higher flux density than the optical data because we are dealinga with a high state of activity in y-rays as reported by ?.., The modeled synchrotron bump has a higher flux density than the optical data because we are dealing with a high state of activity in $\gamma$ -rays as reported by \citet{2000A&A...354..513C}.
" We only predict a very marginal detection of 2273 by H.E.S.S. in its low energy range, depending on the energy threshold (butseealso9).."," We only predict a very marginal detection of 273 by H.E.S.S. in its low energy range, depending on the energy threshold \citep[but see also][]{2006ApJ...653L...5G}."
 Furthermore it should be recalled that ? report an active state in y-rays and ? a high level of the non-thermal emission at the epoch of the data we are considering., Furthermore it should be recalled that \citet{2000A&A...354..513C} report an active state in $\gamma$ -rays and \citet{2006A&A...451L...1T} a high level of the non-thermal emission at the epoch of the data we are considering.
 So even in a high state we do not expect a strong level of VHE y-ray within our scenario., So even in a high state we do not expect a strong level of VHE $\gamma$ -ray within our scenario.
 A strong detection of 2273 at VHE with the current generation of Ceerenkov arrays would be difficult to explain within our SSC nuclear scenario., A strong detection of 273 at VHE with the current generation of Čeerenkov arrays would be difficult to explain within our SSC nuclear scenario.
 A possibility would be to invoke disparity among the various emitting plasma blobs., A possibility would be to invoke disparity among the various emitting plasma blobs.
 Our model shows the presence of a well peaked inverse Compton bump; different magnetic fields or electron energy distributions among the blobs could result in a tail of the IC bump at VHE that could account for a VHE detection., Our model shows the presence of a well peaked inverse Compton bump; different magnetic fields or electron energy distributions among the blobs could result in a tail of the IC bump at VHE that could account for a VHE detection.
" An alternative would be an extended emission due to external inverse Compton radiation, which is then expected to be not very variable."," An alternative would be an extended emission due to external inverse Compton radiation, which is then expected to be not very variable."
" In all cases, further observations withGLAST GGeV) and H.E.S.S.IIL, which will extend the spectral domain of H.E.S.S.II down to ~20GGeV with a better sensitivity, are required to disentangle the different plausible scenarii."," In all cases, further observations with GeV) and II, which will extend the spectral domain of I down to $\sim$ GeV with a better sensitivity, are required to disentangle the different plausible scenarii."
" AA is the nearest radiogalaxy (z=0.0018,?) and one of the best studied."," A is the nearest radiogalaxy \citep[z=0.0018,][]{1978PASP...90..237G} and one of the best studied."
 The presence of an AGN in AA is evident from the radio-band to the y-rays., The presence of an AGN in A is evident from the radio-band to the $\gamma$ -rays.
 Observations from (?) show a bump that seems to peak αἱ ~200 kkeV as pointed out by ?.., Observations from \citep[][]{1995ApJ...449..105K} show a bump that seems to peak at $\sim$ keV as pointed out by \citet[][]{1998A&A...330...97S}.
" We thus have a real constraint for the parameters of our model, particularly with regard to the description of the population of electrons."," We thus have a real constraint for the parameters of our model, particularly with regard to the description of the population of electrons."
" As in the case of 2273, the IC bump peaks at a rather low energy, leading ? to note that AA could be a misaligned low-energy peaked LLac (LBL) object."," As in the case of 273, the IC bump peaks at a rather low energy, leading \citet{2001MNRAS.324L..33C} to note that A could be a misaligned low-energy peaked Lac (LBL) object."
 The value of the viewing angle of the jet is still a controversial issue., The value of the viewing angle of the jet is still a controversial issue.
" For instance, ? claim 0~ 50?—80? for theparsec-scale jet, whereas ? find 0—15? for the ppc scale jet."," For instance, \citet{1998AJ....115..960T} claim $\theta \sim 50\degr$ $80\degr$ for theparsec-scale jet, whereas \citet{2003ApJ...593..169H} find $\theta \sim 15\degr$ for the pc scale jet."
 We choose here to take an intermediate value of 0~25° (See ? for a discussion about the different values for the viewing angle of AA found in the literature)., We choose here to take an intermediate value of $\theta \sim 25\degr$ (See \citet{2006PASJ...58..211H} for a discussion about the different values for the viewing angle of A found in the literature).
" The SMBH mass inferred from gas kinematical analysis using IIT] line is 1.1x105 (?),, but see also ? who give Mgua~6x10!Mc; using a III] Mo,line."," The SMBH mass inferred from gas kinematical analysis using a III] line is $\sim$ $ \times 10^8 M_{\sun}$ \citep[][]{2006A&A...448..921M}, but see also \citet{2006ApJ...643..226H} who give $M_\mathrm{BH} \sim 6 \times 10^7 M_{\sun}$ using a II] line."
 The data sample chosen here is almost the same as in ?.. , The data sample chosen here is almost the same as in \citet{2001MNRAS.324L..33C}. .
The data were retrieved carefully to take only the nucleus into account., The data were retrieved carefully to take only the nucleus into account.
 The strongly constraining CGRO/COMPTEL y-ray, The strongly constraining /COMPTEL $\gamma$ -ray
We acknowledge the great efforts of all the members of the team.,We acknowledge the great efforts of all the members of the team.
 We are grateful to Prof. Jules Halpern for reduction of optical spectra ancl helpful discussion., We are grateful to Prof. Jules Halpern for reduction of optical spectra and helpful discussion.
 We thank John Moore for his analvses of optical spectra., We thank John Moore for his analyses of optical spectra.
 We also thank an anonvimous releree lor careful reading and helpful comments., We also thank an anonymous referee for careful reading and helpful comments.
 CAT and WAIL gratefully acknowledge support through the GGO 0-11624 erant., CM and KML gratefully acknowledge support through the GO 0-11624 grant.
 ΕΠΗΝΤ is supported by SAO contract SV1-GLOLO to MIT., HLM is supported by SAO contract SV1-61010 to MIT.
ltecent distance measurements of some Type. a supernovae at intermediary and high redshifts indicate that the expansion of the Universe is speeding up. rather than slowing down (Riess et al.,"Recent distance measurements of some Type Ia supernovae at intermediary and high redshifts indicate that the expansion of the Universe is speeding up, rather than slowing down (Riess et al."
 1998: Perlmutter et al., 1998; Perlmutter et al.
 19992)., 1999a).
 Incirectly. these results also mean that the Universe is much older than the one predicted by the standard. CDM Lat model with a critical deceleration. parameter αν= 0.5).," Indirectly, these results also mean that the Universe is much older than the one predicted by the standard CDM flat model with a critical deceleration parameter $q_{o}
=$ 0.5)."
 1t confirmed. from more accurate observations and/or a different class of phenomena. the existence of these data poses a crucial problem. for all CDM. models since. their generic prediction is a decelerating universe (q.2 0). whateverthe sign adopted for the curvature parameter.," If confirmed from more accurate observations and/or a different class of phenomena, the existence of these data poses a crucial problem for all CDM models since their generic prediction is a decelerating universe $q_{o} >
0$ ), whatever the sign adopted for the curvature parameter."
 Another source of dilliculties for the standard. CDA niocel is related to the “age problem or its modern variant. he age of old high redshift objects.," Another source of difficulties for the standard CDM model is related to the “age problem"" or its modern variant, the age of old high redshift objects."
 It should be recallec hat some measurements of the Hubble parameter. vieldec h=Hof100km/sec/Mpe0.7x0.1 (Nevalainen Roos 1998. Friedmann 1998).," It should be recalled that some measurements of the Hubble parameter yielded $h = H_{o}/100
\rm{km/sec/Mpc} = 0.7 \pm 0.1$ (Nevalainen Roos 1998, Friedmann 1998)."
" In. particular. this means hat the expansion age for a FRA Dat matter cdominatec universe(rese (f,—54, ) falls. within: the! intervalyal 6.3Cyrκον»xM0.mep 5Gyr. whileael the Conage inferredinfnrpreed fyfrom globular aroclusters"" ⊔⊳∖↴⋅∖↓⋟⊔↖↧∏⋅∖↓⊔∣∪∣↓⊔↓↖⋯∆⋅≟⊲∣∠∕∢∿⇀⊁ ⋅⋟⋖⋅⋅∖↓∪↓⊔∆⋅≟↓⊔↓ ∖∐∪↓∣∪⊔⊾∖↽∐∪⋏∙≟⋜⋯≤⋗≤⋗⋅↱≻∶↓↴∪⊔↿∢⊾⇂⋜↧↓⊳↓≤⋗≤⋗↖∖⊐⋡"," In particular, this means that the expansion age for a FRW flat matter dominated universe $t_o = {2 \over 3}H_o^{-1}$ ) falls within the interval $8.3 \rm{Gyr} \leq t_o \leq 10.5 \rm{Gyr}$ , while the age inferred from globular clusters lies typically into the range $t_{gc} \sim
13-15\rm{Gyr}$ or higher (Bolton Hogan 1995; Pont et al."
⇀∖⊳∖∖∖⋰, 1998).
⊔⇂∢⊾↓∙∖⇁↳⊔⋯∖⋎⊔⊳ ↓↕⊀↓⊳∖≼⇍∪⊔∐⊲⊔⇍↿↕≻↓↕∢≱↿⋜↧↓⊲↓∖⇁⊲↓⋜⋯⊾∠⇂∐∺↓∖⊽⋖⋅∠⇂⋜↧⋯⋜⊔⋅∢⊾≼⇍∪⊔⊳∖⊲⊔⇂∢⊾↓⋅," As widely known, this conflict is not alliviated if SNe data are considered."
⋖⋅∠⇂⊳↓⊔ he age analysis of Perlmutter at al. (, In the age analysis of Perlmutter at al. (
1998). the favorecl value is f=0.63. while Riess et al. (,"1998), the favored value is $h=0.63$, while Riess et al. ("
1998) found h=0.65+0.02. rom their SNe data.,"1998) found $h=0.65 \pm 0.02$, from their SNe data."
 Age measurements of cxtragalactic objects at high redshifts also provide an alternative route to the “age problem”.," Age measurements of extragalactic objects at high redshifts also provide an alternative route to the “age problem""."
 In this case. the discoveries of a 4.00Civr-old galaxy at z=1.175 (Stockton et αἱ.," In this case, the discoveries of a $4.0-\rm{Gyr}$ -old galaxy at $z=1.175$ (Stockton et al."
 1995). ofa 3.50Givr-old. galaxy at z=1.55 (Dunlop et al.," 1995), of a $3.5-\rm {Gyr}$ -old galaxy at $z=1.55$ (Dunlop et al."
 1996: Spinrard et al., 1996; Spinrard et al.
 1997). and ofa 4.0. Civr-old galaxy at z=1.43 (Dunlop 1998) have been proved to be incompatible with age estimates for a Lat Universe unless the Llubble parameter is very. low.," 1997), and of a $4.0-\rm{Gyr}$ -old galaxy at $z=1.43$ (Dunlop 1998) have been proved to be incompatible with age estimates for a flat Universe unless the Hubble parameter is very low."
 These constraints are even more stringent than the ones from globular cluster age measurements (Dunlop 19060. Krauss 1997. Roos Llarun-or-Raschicl 1998).," These constraints are even more stringent than the ones from globular cluster age measurements (Dunlop 1996, Krauss 1997, Roos Harun-or-Raschid 1998)."
 More. recently. it was shown that the existence of the two OLURG discovered bv Dunlop would be accommodated in a model with no cosmological constant only i£ Q3;50.37 (Meaniz and Lina 1999).," More recently, it was shown that the existence of the two OHRG discovered by Dunlop would be accommodated in a model with no cosmological constant only if $\Omega_M \leq
0.37$ (Alcaniz and Lima 1999)."
 In the last [ew vears. flat models with a relic cosmological constant. CXCDM) have also been considered as à serious candidate for standard cosmology (Ixrauss ‘Turner 1995: Krauss 1997).," In the last few years, flat models with a relic cosmological constant $\Lambda$ CDM) have also been considered as a serious candidate for standard cosmology (Krauss Turner 1995; Krauss 1997)."
 However. although fitting some observations better than other theoretical models (e.g. he first⋅ accoustic. peak of⋅ the relic. radiations angular power spectrum).l} ACDAL cosmologies5 are reasonably restricted by ⇂↥∢⋅⊳∖↿⋜∐↓⊳∖⇂⊔⋱∖∪⇂⋏∙≟↓⋅⋜↧∖⇁↓⋯↿↓∪⊔∥⇂∢⋅⊔⊳∖∢⊾⊳∖↿∖⊾↔≺∣∟∃⋃∖⋯∼↓⋯⊔⋖⋅↓∡≤⋗≤⋈≻⋡ ↓⊲∖," However, although fitting some observations better than other theoretical models (e.g., the first accoustic peak of the relic radiation angular power spectrum), $\Lambda$ CDM cosmologies are reasonably restricted by the statistics of gravitational lenses (SGL) (Kochanek 1996, Falco et al."
⋜↧↓⊓∏⊾↿⋜↧↓⋡↓≤⋗≤⋗↖∖⋡∖∖⊽⋜↧⋏∙≟∥∥," 1998, Waga and Miceli 1999)."
⊔∠↳∖∐⊓⋅∐↓≤⋗≤⋗≤∩⋡∪⊔↿↓⊔⋅∪↿⇂↥∢⋅↓⋅ ⋯⊔∠⇂⊳⇂↓↥∪⊔⋏∙≟↓↥↿⇂⊔⊾⊔↓⋖⋅↥⇂↥⋯⇂∣⋯≱∖⋖⋅∠⇂∪⊔↻↓∐⊰≺∶↓↥⋜↧∖⇁⋖⋅⋏∙≟⊲↓∖⇁∢⊾⊔⋜↧ ower limit of ον20.5 (Aleaniz anc Lima 1999). from a theoretical viewpoint. these models are. plaguec by a wolound contradiction: in order to dominate the dvnamies of the Universe only at recent times. a very small value for," On the other hand, though the method based on OHRG have given a lower limit of $\Omega_\Lambda \geq 0.5$ (Alcaniz and Lima 1999), from a theoretical viewpoint, these models are plagued by a profound contradiction: in order to dominate the dynamics of the Universe only at recent times, a very small value for"
In cases where extremely narrow spectral lines were present (e.g. the unresolved ozone lines at nnm - see Sect. ??)),In cases where extremely narrow spectral lines were present (e.g. the unresolved ozone lines at nm - see Sect. \ref{sec:4760}) )
 up to 30 Eigenvalues could be used without amplifying the noise., up to 30 Eigenvalues could be used without amplifying the noise.
 This analysis confirmed that a Gaussian model is a valid approximation for the instrumental profile at signal-to-noise levels below ~300., This analysis confirmed that a Gaussian model is a valid approximation for the instrumental profile at signal-to-noise levels below $\sim300$.
" We have thus used a single Gaussian model for the final fitting of the telluric model, the width of the profile being the only free parameter."," We have thus used a single Gaussian model for the final fitting of the telluric model, the width of the profile being the only free parameter."
 We note that our determined FWHM are broader than expected for R=100000.," We note that our determined FWHM are broader than expected for $R=100\,000$."
" At the time most of the data presented in this paper were recorded, the focal plane array of CRIRES was slightly out of focus."," At the time most of the data presented in this paper were recorded, the focal plane array of CRIRES was slightly out of focus."
" Thus, the maximum resolving power measured in the presented dataset is R~87000."," Thus, the maximum resolving power measured in the presented dataset is $R\sim87\,000$."
 We present CRIRES observations of bright standard stars that contain clear and isolated line systems of the most abundant molecules of the Earth’s atmosphere., We present CRIRES observations of bright standard stars that contain clear and isolated line systems of the most abundant molecules of the Earth's atmosphere.
 We evaluate the performance of the telluric correction by checking the residuals for systematics by a rigid quantisation of the signal-to-noise ratio (SNR) in the corrected spectra., We evaluate the performance of the telluric correction by checking the residuals for systematics by a rigid quantisation of the signal-to-noise ratio (SNR) in the corrected spectra.
" We chose the SNR as a measure for the performance of the telluric model fit since it has the strongest practical implication for the observer, as it allows a judgement as to which precision a position or an equivalent width of a stellar line can be measured in a spectrum."," We chose the SNR as a measure for the performance of the telluric model fit since it has the strongest practical implication for the observer, as it allows a judgement as to which precision a position or an equivalent width of a stellar line can be measured in a spectrum."
" In regions of telluric contamination, the residuals from the removal of telluric lines are the main factor hindering a high signal-to-noise measurement, otherwise limited only by the source flux or the uniformity of the detector response."," In regions of telluric contamination, the residuals from the removal of telluric lines are the main factor hindering a high signal-to-noise measurement, otherwise limited only by the source flux or the uniformity of the detector response."
" In order to compare the SNR before and after the correction, we first compute the intrinsic noise level of our CRIRES spectra by adding the photon noise of the source and the measured background noise in quadrature for each pixel along the dispersion direction."," In order to compare the SNR before and after the correction, we first compute the intrinsic noise level of our CRIRES spectra by adding the photon noise of the source and the measured background noise in quadrature for each pixel along the dispersion direction."
" The background noise consists of the detector readout noise, the photon noise of the dark current and the photon noise of the thermal emission of the atmosphere and the telescope after pairwise subtraction of nodded frames."," The background noise consists of the detector readout noise, the photon noise of the dark current and the photon noise of the thermal emission of the atmosphere and the telescope after pairwise subtraction of nodded frames."
 We have validated this procedure by measuring the standard deviation in the continuum level of each spectrum., We have validated this procedure by measuring the standard deviation in the continuum level of each spectrum.
 The result was always consistent with our noise estimate., The result was always consistent with our noise estimate.
" Hence, for the input spectra we have the SNR as a function of wavelength which scales with the square root of the received flux, and thus with the square root of the atmospheric transmission T."," Hence, for the input spectra we have the SNR as a function of wavelength which scales with the square root of the received flux, and thus with the square root of the atmospheric transmission $T$."
" In addition, the telluric emission lines in the thermal infrared example (A~4700 nnm, Sec. ??))"," In addition, the telluric emission lines in the thermal infrared example $\lambda\sim4700$ nm, Sec. \ref{sec:4760}) )"
 add noise where the lower transmission of the same telluric lines in absorption degrade the SNR in comparison to the continuum level., add noise where the lower transmission of the same telluric lines in absorption degrade the SNR in comparison to the continuum level.
" The fitting procedure minimises the residuals in the observed—computed spectrum (O—-C) over the whole wavelength region of the considered spectral range of ~A/200, recorded on one CRIRES detector."," The fitting procedure minimises the residuals in the $-$ computed spectrum $-$ C) over the whole wavelength region of the considered spectral range of $\sim\lambda/200$, recorded on one CRIRES detector."
" For the removal of the telluric lines, however, the measured spectrum has to be divided by the computed model."," For the removal of the telluric lines, however, the measured spectrum has to be divided by the computed model."
" After division by the model, we measure the noise in the spectrum as the standard deviation of all pixels."," After division by the model, we measure the noise in the spectrum as the standard deviation of all pixels."
" The SNR is then simply the inverse of the measured noise, given that after the division the signal level is at unity for a normalized spectrum."," The SNR is then simply the inverse of the measured noise, given that after the division the signal level is at unity for a normalized spectrum."
" It is important to remember that even for a perfect model of the telluric lines the O/C residuals will be larger, and the SNR smaller, in comparison to the undisturbed continuum, since the spectrograph receives less flux where telluric absorption lines are present."," It is important to remember that even for a perfect model of the telluric lines the O/C residuals will be larger, and the SNR smaller, in comparison to the undisturbed continuum, since the spectrograph receives less flux where telluric absorption lines are present."
" Hence, when dividing through the telluric transmission spectrum, and thus correcting the received flux for the presence of these lines, the noise in these regions is higher than in the neighboring continuum, simply because less flux was received in the lines in the first place."," Hence, when dividing through the telluric transmission spectrum, and thus correcting the received flux for the presence of these lines, the noise in these regions is higher than in the neighboring continuum, simply because less flux was received in the lines in the first place."
" Albeit this caveat, we use the formal SNR as a direct measure of the performance of the modeling since it reflects the impact of the telluric lines on the astronomical spectrum and gives the observer a meaningful number to work with."," Albeit this caveat, we use the formal SNR as a direct measure of the performance of the modeling since it reflects the impact of the telluric lines on the astronomical spectrum and gives the observer a meaningful number to work with."
 Water vapour (H5O) is the most important absorber in the near-infrared., Water vapour $_2$ O) is the most important absorber in the near-infrared.
" In addition to the strong νι and v3 fundamental bands at nnm, several overtone and combination bands centered at 940nnm, nnm, 1380nnm, nnm, and nnm are present in near-infrared spectra."," In addition to the strong $\nu_1$ and $\nu_3$ fundamental bands at nm, several overtone and combination bands centered at nm, nm, nm, nm, and nm are present in near-infrared spectra."
" We picked a few isolated, unsaturated H5O lines at nnm in a CRIRES standard star setting from March 03, 2007."," We picked a few isolated, unsaturated $_2$ O lines at nm in a CRIRES standard star setting from March 03, 2007."
" The airmass of the observation was 1.033, the resolving power was ~80,000 at a slitwidth of0."," The airmass of the observation was 1.033, the resolving power was $\sim$ 80,000 at a slitwidth of."
"2"".. Both the GDAS model and the modified MMS model had to be scaled to a PWV of mmm for the best fit with the observation (see Fig. 4)).", Both the GDAS model and the modified MM5 model had to be scaled to a PWV of mm for the best fit with the observation (see Fig. \ref{fig:1504}) ).
 The achieved correction using the GDAS and MM5 models is generally satisfying but falls behind the expectation from the noise limit., The achieved correction using the GDAS and MM5 models is generally satisfying but falls behind the expectation from the noise limit.
 The average SNR of the measured spectrum is «200., The average SNR of the measured spectrum is $\sim$ 200.
" After the division by the model spectrum, the average"," After the division by the model spectrum, the average"
upper limit that the disk surface density cannot exceed during the evolution.,upper limit that the disk surface density cannot exceed during the evolution.
 For a lavered disk domiuated by viscous heating. Gamunic (1996) has shown that the Toome Q parameter at R is At the outer dead zone radius (Rp). X2 with the coudition that Q1. For au iaciation dominated disk. we derive a similar condition: assununes Then. Since QUA 1. Iu either case. Rp< 50 AU.," For a layered disk dominated by viscous heating, Gammie (1996) has shown that the Toomre Q parameter at R is At the outer dead zone radius $_{D}$ ), $\Sigma=\Sigma_{A}$; with the condition that $>$ 1, For an irradiation dominated disk, we derive a similar condition: assuming Then, Since $_{D}$ $>$ 1, In either case, $_{D}\lesssim$ 50 AU."
" To test this. we computed a case witha Q,92«10. Pyad ! core."," To test this, we computed a case with a $\Omega_{c}=2\times10^{-13}$ rad $^{-1}$ core."
 The core mass falls into Rojas as lige as LOO AU. but the dead zoue bevoud 30 AU is quickly cleared by the GI so that Rp is smaller than 30 AU during the T Taux phase.," The core mass falls into $_{cmax}$ as large as 400 AU, but the dead zone beyond 30 AU is quickly cleared by the GI so that $_{D}$ is smaller than 30 AU during the T Tauri phase."
 After we know the initial Rp position. the motion of Rp cau be derived by cousidering the mass conservation at Bp.," After we know the initial $_{D}$ position, the motion of $_{D}$ can be derived by considering the mass conservation at $_{D}$ ."
 For the region exteuding from Rp-AR to Rp. Assunmiug the mass flux (or eg) at Rp is zero. which will be justified in some cases later. aud considering the tine for the dead zone within A R to be depleted is Af. we eot thus where M is as equation AL.," For the region extending from $_{D}$ $\Delta R$ to $_{D}$, Assuming the mass flux (or $v_{R}$ ) at $_{D}$ is zero, which will be justified in some cases later, and considering the time for the dead zone within $\Delta$ R to be depleted is $\Delta t$, we got thus where $\dot{M}$ is as equation \ref{eq:layer4}."
" I£sve insert AL we cau derive vpxL/(,|X4).", If we insert $\dot{M}$ we can derive $_{D}$$\propto 1/(\Sigma_{a}+\Sigma_{d})$.
" Cousidering X, X, is larecr at siialler radii. the speed of p decrease with time when it is moving wards."," Considering $\Sigma_{a}$ $\Sigma_{d}$ is larger at smaller radii, the speed of $_{D}$ decrease with time when it is moving inwards."
 Finally. when Ry moves to the very iuucr radii where X429 M. Bp is almost halted.," Finally, when $_{D}$ moves to the very inner radii where $\Sigma_{d} \gg \Sigma_{a}$ , $_{D}$ is almost halted."
 Bevoud Ry the disk is purely MRI active with a coustant oa;., Beyond $_{D}$ the disk is purely MRI active with a constant $\alpha_{M}$.
" Its surface deusity evolution can be solved by with the boundary condition X5,,-X,.", Its surface density evolution can be solved by with the boundary condition $\Sigma_{R_{D}}$ $\Sigma_{A}$ .
" If j£—AR"". this solution can be suuplied by dividing it iuto two parts where the first term on the right comes from the disk evolution with a zero surface density boundary condition at Rp (Xp, 20). aud the secoud teri is the effect of the nou-zero boundary surface density at p."," If $\nu=k
R^{n}$, this solution can be simplified by dividing it into two parts where the first term on the right comes from the disk evolution with a zero surface density boundary condition at $_{D}$ $\Sigma_{R_{D}}$ =0), and the second term is the effect of the non-zero boundary surface density at $_{D}$."
 The first term bchaves like the similarity solution with a + part and a power law decrease part at larger radius. aud it eventually decreases to zero with time.," The first term behaves like the similarity solution with a $^{-1}$ part and a power law decrease part at larger radius, and it eventually decreases to zero with time."
 The surface density evolution is determined by the competition between the first aud the second term., The surface density evolution is determined by the competition between the first and the second term.
 During the imfall phase. the disk is massive. aud the first term is larger than the second terii. thus the disk. behaves like a sinularity solution.," During the infall phase, the disk is massive, and the first term is larger than the second term, thus the disk behaves like a similarity solution."
 Then after iufall stops. as the disk accretes audRy moves inwards. both of these two terms decrease. and as shown above. Ry moves invairds moreslowly.," Then after infall stops, as the disk accretes and$_{D}$ moves inwards, both of these two terms decrease, and as shown above, $_{D}$ moves inwards moreslowly."
 EventuallyRy can be considered independent of, Eventually$_{D}$ can be considered independent of
flux can be used.,flux can be used.
 The adimensional parameter is again w=WAY and the parametric equations are: being ry=C—2.5log(f.) the magnitude on an object without line and w€(—1.x).," The adimensional parameter is again $w= \ew /\Delta'_{\narrowband}$ and the parametric equations are: being $x_0=C-2.5\log\left(\fluxc\right)$ the magnitude on an object without line and $w \in (-1, \infty)$."
 The asymptotic limits correspond to a object without line («6—0.6 —ry and y—0). an object dominated by the emission line ((—x. ——o and y——2.5loge) aud an object dominated by an absorption line (——1..7 —x and y——x) In Fig. 4.," The asymptotic limits correspond to a object without line $w \to 0$, $x \to x_0$ and $y \to
0$ ), an object dominated by the emission line $w \to \infty$, $x \to
-\infty$ and $y \to -2.5\log \epsilon$ ) and an object dominated by an absorption line $w \to -1$, $x \to \infty$ and $y \to -\infty$ ) In Fig. \ref{fig:isolineas},"
 we show a magnitude-color selection diagram including example isolines for ihe WFCO9200 case (see Sect., we show a magnitude-color selection diagram including example isolines for the WFC9200 case (see Sect.
 6 [or details)., \ref{sec:examples} for details).
 The isolines of continuum flix and line flux are represented by the dashed and solid curves respectively., The isolines of continuum flux and line flux are represented by the dashed and solid curves respectively.
 We can see clearly the asymptotic behavior of both tvpes of isolines., We can see clearly the asymptotic behavior of both types of isolines.
 Although the limited numeric precision of the plot do nol show it. the curves of line flux continue up to the maximum color.," Although the limited numeric precision of the plot do not show it, the curves of line flux continue up to the maximum color."
 We have assigned an approximate redshift = to the source. but we do not know which is the redshift range z4 zs of the objects that could be selected and included in the sample.," We have assigned an approximate redshift $\bar{z}$ to the source, but we do not know which is the redshift range $z_1$ $z_2$ of the objects that could be selected and included in the sample."
 The answer (o this question is (he kev [or the incompleteness correction., The answer to this question is the key for the incompleteness correction.
 If we intend to compute a luminosity function. we need to compute (he comoving volume per object.," If we intend to compute a luminosity function, we need to compute the comoving volume per object."
 We consider that the object is at redshift = and it has a line flux. fj; and continuum flix f.., We consider that the object is at redshift $\bar{z}$ and it has a line flux $\fluxl$ and continuum flux $\fluxc$ .
 The this recovered in a given filler at redshift z (whereby (the emission line lies at wavelength A. ) is: with d; the luminosity distance., The flux recovered in a given filter at redshift $z$ (whereby the emission line lies at wavelength $\lambda_z$ ) is: with $d_L$ the luminosity distance.
 For 2=7 we recover Eq. 9.., For $z=\bar{z}$ we recover Eq. \ref{eq:flujo2}.
 Substituting the appropriate filter widths A’. we can obtain the magnitude and color of the object as a function of redshift.," Substituting the appropriate filter widths $\Delta'$, we can obtain the magnitude and color of the object as a function of redshift."
 Onlv al à range of redshifts σι.22 the object fulfills the color criterion given in Sect. 5..," Only at a range of redshifts $z_1, z_2$ the object fulfills the color criterion given in Sect. \ref{sec:dispq}."
 In Fig., In Fig.
 5 we show the reclshilt (races given bv Eq., \ref{fig:volume} we show the redshift traces given by Eq.
 59. Lor two objects in a diagram corresponding to the filter set CAFOSS200 (see Sect., \ref{eq:fluxz} for two objects in a magnitude-color diagram corresponding to the filter set CAFOS8200 (see Sect.
 6. for details)., \ref{sec:examples} for details).
 The position ofthe objects in the diagram al redshift 7 are represented as crosses on top of the traces., The position ofthe objects in the diagram at redshift $\bar{z}$ are represented as crosses on top of the traces.
 The intersection between the traces and the selection curve (drawn: with a dashed line) is labeledas τι and zs., The intersection between the traces and the selection curve (drawn with a dashed line) is labeledas $z_1$ and $z_2$ .
 The object would be selected only in the range z4<2«€ z»., The object would be selected only in the range $z_1 < z < z_2$ .
 Maps. Earth. Projectiou. Curvature Tissot (1881) indicatrices have been very useful for providing a visual presentation of local distortions in nip projectios lu a situple and compelling fashion.," Maps, Earth, Projection, Curvature Tissot (1881) indicatrices have been very useful for providing a visual presentation of local distortions in map projections in a simple and compelling fashion."
 A simall circle of tiny radius (say 0.1 degree of arc: jn raclius) is constructed at a given location. aud then eularged aud projected ou the map at that ocation.," A small circle of tiny radius (say 0.1 degree of arc in radius) is constructed at a given location, and then enlarged and projected on the map at that location."
 Tlis always produces an ellipse., This always produces an ellipse.
 Usually oue favors conformal map projections that iniiize the chages in scale factor. or equal area projections that minimize anisotropy. or recently. lap »rojections hat are ueitlier coulormal nor equal area. but which have a judicious Combination oL nminimiziug bot1 scale aucl isotropy errors (like the Wiukel-tripel used by the National Geographic Sociev for world maps).," Usually one favors conformal map projections that minimize the changes in scale factor, or equal area projections that minimize anisotropy, or recently, map projections that are neither conformal nor equal area, but which have a judicious combination of minimizing both scale and isotropy errors (like the Winkel-tripel used by the National Geographic Society for world maps)."
 Tje. Tissot elipse at a given location is specified by three parameters. the major axis. tle minor axls. aud the orientation angle 0 of the major axis of the ellipse.," The Tissot ellipse at a given location is specified by three parameters, the major axis, the minor axis, and the orientation angle $\theta$ of the major axis of the ellipse."
 Geometrically. [rom cliffereutial geoetry Cand Ceueral Relativity) we know that the measurement of local distances is measured by the metric ellsor Gah. Where a aud b cau each take the values x or y. vieldiug three inclependent components.," Geometrically, from differential geometry (and General Relativity) we know that the measurement of local distances is measured by the metric tensor $g_{ab}$, where a and b can each take the values x or y, yielding three independent components."
 Locally for two nearby. points separated by inlinitesimal map coordinate cliffereuces ον aud dy. the true cdistauce between these two points ou the globe is giveu by: The Tissot ellipse (auajor axis. minor axis. aud orientation angle. 0) can be calculated from the three components of the metric teusor.," Locally for two nearby points separated by infinitesimal map coordinate differences dx and dy, the true distance between these two points on the globe is given by: The Tissot ellipse (major axis, minor axis, and orientation angle, $\theta$ ) can be calculated from the three components of the metric tensor."
 Thus. the Tissot ellipse essentially carries the information ou tlie metric teusor for the map.," Thus, the Tissot ellipse essentially carries the information on the metric tensor for the map."
 It tells us how local tufinitesimal clistauces on the map correlate with local infinitesimal distauces ou the globe., It tells us how local infinitesimal distances on the map correlate with local infinitesimal distances on the globe.
RR Lwrae stars are low-mass stars on the horizontal branch that burn Ποια in their core.,RR Lyrae stars are low-mass stars on the horizontal branch that burn Helium in their core.
 They have contributed. to the progress in. many. fields of modern astronomy., They have contributed to the progress in many fields of modern astronomy.
 With their large amplitudes. of up to 1.5 magnitudes they have been known since the enc of the ΝΙΝ century and olfer the opportunity το study stellar evolution in real-time (LeBorgneetal2007)., With their large amplitudes of up to 1.5 magnitudes they have been known since the end of the XIX century and offer the opportunity to study stellar evolution in real-time \citep{LeBo07}.
. They obey ao period-luminosity-color relation which makes them important standard: candles., They obey a period-luminosity-color relation which makes them important standard candles.
 Also. their evolutionary status makes them useful for the study of galactic evolution.," Also, their evolutionary status makes them useful for the study of galactic evolution."
 It Lyrac stars can pulsate in the raclial fundamental mode (πα stars). in the first racial overtone (Itc stars) or in both modes. simultaneously (Ilic stars).," RR Lyrae stars can pulsate in the radial fundamental mode (RRab stars), in the first radial overtone (RRc stars) or in both modes simultaneously (RRd stars)."
 Some lut Lyrae stars. are. canclidates. for possible higher radial overtone One of the remaining unsolvecl questions in. those seemingly wellestucied stars is the centurv-old. Blazhko enigma. a modulation of the amplitude and/or. phase of the main pulsation period on time scales of typically. weeks to months.," Some RR Lyrae stars are candidates for possible higher radial overtone One of the remaining unsolved questions in those seemingly well-studied stars is the century-old Blazhko enigma, a modulation of the amplitude and/or phase of the main pulsation period on time scales of typically weeks to months."
 With the availability of new and precise data. the estimated. incidence rate of this phenomenon has recently increased dramaticallv. making it even more important to understand. the ellect.," With the availability of new and precise data, the estimated incidence rate of this phenomenon has recently increased dramatically, making it even more important to understand the effect."
 While until some vears ago. a value of 20-30 per cent was assumed for lutab stars. Ixolenbergetal.(2010a) find at least 40 per cent of the stars observed by the Kepler satellite to be modulated. and Juresiketal.(20098). give a lower limit of 50 per cent. of modulated stars in their sample and even raise the suspicion that the modulation might," While until some years ago, a value of 20-30 per cent was assumed for RRab stars, \citet{kol10a} find at least 40 per cent of the stars observed by the Kepler satellite to be modulated, and \citet{jur09} give a lower limit of 50 per cent of modulated stars in their sample and even raise the suspicion that the modulation might"
Hux ratios in Figure 3. (logfyf/f7 2) which is additional evidence for star-formation activity in these systems.,flux ratios in Figure \ref{fig_fxfo} $\log f_X/f_{opt}\approx-2$ ) which is additional evidence for star-formation activity in these systems.
 The optical spectra [or these two sources are presented in lll., The optical spectra for these two sources are presented in II.
 The remaining sources in Figure 4 either deviate from the LyLia relation of local star-forming svstems (suggesting the presence of ACN) or have X-ray luminosities well in excess of typical starburst galaxies (=10eres Aloran. Lehnert Lelanel 1999).," The remaining sources in Figure \ref{fig_lxl14} either deviate from the $L_X-L_{1.4}$ relation of local star-forming systems (suggesting the presence of AGN) or have X-ray luminosities well in excess of typical starburst galaxies $\approx10^{42}\rm\, erg\, s^{-1}$; Moran, Lehnert Helfand 1999)."
 We also [ind a double. lobe radio. source associated with diffuse X-ray cluster emission., We also find a double lobe radio source 1) associated with diffuse X-ray cluster emission.
 As discussed in Appendix ?? this cluster is likely to lie at 2=0.6.," As discussed in Appendix \ref{app1}
 this cluster is likely to lie at $z\approx0.6$."
 Powerful radio galaxies are known be good tracers of dense environments. with Fanaroll-Hillev. ἵνρο sources (FRI) associated with rich groups or clusters and Fanaroll-ltillev HE tvpe sources (EIUL) avoiding rich clusters at Iow-z (c.g. Zirbe 1997 and references therein)., Powerful radio galaxies are known to be good tracers of dense environments with Fanaroff-Rilley I type sources (FRI) associated with rich groups or clusters and Fanaroff-Rilley II type sources (FRII) avoiding rich clusters at $z$ (e.g. Zirbel 1997 and references therein).
 Recently. Zanichelli et al. (," Recently, Zanichelli et al. ("
2001) used. NVSS data to compile an intermediate redshift ἐς=0.1 1.3) cluster sample based. on. racio selection.,2001) used NVSS data to compile an intermediate redshift $z\approx0.1 - 0.3$ ) cluster sample based on radio selection.
 Follow-up optical photometry anc spectroscopy confirmed that a significant fraction of their targeted cluster candidates are real. demonstrating the strength. of racio selection in cluster studies.," Follow-up optical photometry and spectroscopy confirmed that a significant fraction of their targeted cluster candidates are real, demonstrating the strength of radio selection in cluster studies."
 Our X-ray/radio matched sample comprises a total of 3 double lobe radio sources and #112) of which only one (E11) is clearly associated with X-rav cluster emission.," Our X-ray/radio matched sample comprises a total of 3 double lobe radio sources 1, 5 and 12) of which only one 1) is clearly associated with X-ray cluster emission."
 Phe remaining two lie in either »»or or high-z clusters/groups that remain undetected. in »Àh the X-ray. ancl the optical data., The remaining two lie in either poor or $z$ clusters/groups that remain undetected in both the X-ray and the optical data.
 Indeed. as discussed in Appendix ?? we do not find evidence for optical galaxy overdensity in the vicinity of sources #55 and 112. at least o the optical magnitude limit of the SDSS.," Indeed, as discussed in Appendix \ref{app1} we do not find evidence for optical galaxy overdensity in the vicinity of sources 5 and 12, at least to the optical magnitude limit of the SDSS."
 Sources with absorption line optical spectra and. high X-ray luminosities (Ly21071075ergs ly similar to hose found in the present sample (3633. 4 in Table 1). have also been identified in previous X-ray surveys (c.g. Crilliths et al.," Sources with absorption line optical spectra and high X-ray luminosities $L_X\approx10^{42}-10^{43}\rm\, erg\,s^{-1}$ ), similar to those found in the present sample 3, 4 in Table 1), have also been identified in previous X-ray surveys (e.g. Griffiths et al."
 1995: Blair.€MM Stewart 1997: Allen. Di Matteo Fabian eio:206: Comastri ct al.," 1995; Blair, Georgantopoulos Stewart 1997; Allen, Di Matteo Fabian 2000; Comastri et al."
 as)022: Brusa οἱ al., 2002; Brusa et al.
2003)., 2003).
 Possible for the observed emission in these optically normal(e.g. no emission lines) systems include (1) BL-Lac tvpe activity (c.g. Blair et al., Possible scenarios for the observed X-ray emission in these optically normal (e.g. no emission lines) systems include (i) BL-Lac type activity (e.g. Blair et al.
 1997). (ii) Advection Dominated Accretion FlowsADAE: Naravan Yi 1995: Di Matteo et al.," 1997), (ii) Advection Dominated Accretion Flows (ADAF; Narayan Yi 1995; Di Matteo et al."
 2000). (ii) cülluse cluster of group X-ray. emission and(iv)) LLAGN with its optical signature diluted. by the ealaxv Dight (c Aloran. l'ilippenko Chornock 2002: ASevergnini ct al.," 2000), (iii) diffuse cluster of group X-ray emission and (iv) LLAGN with its optical signature diluted by the host galaxy light (e.g. Moran, Filippenko Chornock 2002; Severgnini et al."
 2003)., 2003).
The present paper is the first analytical study on the nonlinear resonant wave. where we obtain governing equations using techniques made familiar from. previous studies on nonlinear slow resonant MHD waves.,"The present paper is the first analytical study on the nonlinear resonant wave, where we obtain governing equations using techniques made familiar from previous studies on nonlinear slow resonant MHD waves."
 Before embarking on the actual derivation. let us carry out a qualitative discussion.," Before embarking on the actual derivation, let us carry out a qualitative discussion."
 First of all we should point out that in plasmas with high Reynolds numbers (as in the solar corona) efficient dissipation only operates in a thin layer embracing the resonant surface., First of all we should point out that in plasmas with high Reynolds numbers (as in the solar corona) efficient dissipation only operates in a thin layer embracing the resonant surface.
 This layer is calledlaver., This layer is called.
.. This restriction on the effect of dissipation makes the problem more tractable from a mathematical point of view. as outside the dissipative layer the dynamics of waves is described by the ideal MHD.," This restriction on the effect of dissipation makes the problem more tractable from a mathematical point of view, as outside the dissipative layer the dynamics of waves is described by the ideal MHD."
 Dissipation is a key ingredient of the problem of resonance., Dissipation is a key ingredient of the problem of resonance.
 As it was mentioned earlier dissipation removes singularities in mathematical solutions., As it was mentioned earlier dissipation removes singularities in mathematical solutions.
 From a physical point of view dissipation is important as it is the mechanism which relaxes the accumulation of energy at the resonant surface and eventually contributes to the global process of heating., From a physical point of view dissipation is important as it is the mechanism which relaxes the accumulation of energy at the resonant surface and eventually contributes to the global process of heating.
 It is important to stress that the choice of dissipation has to be related to the very physics which is described as different waves are sensitive to different dissipative mechanisms., It is important to stress that the choice of dissipation has to be related to the very physics which is described as different waves are sensitive to different dissipative mechanisms.
 Due to the dominant role of the magnetic field ir the solar corona. transport processes are highly anisotropic.," Due to the dominant role of the magnetic field in the solar corona, transport processes are highly anisotropic."
 Possible dissipation mechanisms acting in coronal structures can be described within the framework of Braginskit's theory (2). às it was shown in applications by. e.g. 222.," Possible dissipation mechanisms acting in coronal structures can be described within the framework of Braginskii's theory \citep{braginskii} as it was shown in applications by, e.g. \citet{erdelyi1995, ofman1995, mocanu2008}."
 waves are incompressible and transversal (in. polarization). therefore. it is sensible to adopt shear viscosity and magnetic resistivity.," waves are incompressible and transversal (in polarization), therefore, it is sensible to adopt shear viscosity and magnetic resistivity."
 Despite both transport processes being described by rather small coetficients. the net effect of dissipation can be increasec considerably when the dissipative coefficients are multiplied by large transversal gradients.," Despite both transport processes being described by rather small coefficients, the net effect of dissipation can be increased considerably when the dissipative coefficients are multiplied by large transversal gradients."
 The paper is organized as follows., The paper is organized as follows.
 In the next sectior we introduce the fundamental equations and discuss the mau assumptions., In the next section we introduce the fundamental equations and discuss the main assumptions.
 In Sect., In Sect.
 III. we derive the governing equation for wave dynamics inside the dissipative layer.," III, we derive the governing equation for wave dynamics inside the dissipative layer."
 Section IV is devoted to caleulating the nonlinear corrections at resonance., Section IV is devoted to calculating the nonlinear corrections at resonance.
 Finally. in Sect.," Finally, in Sect."
 V. we summarize and draw our conclusions. pointing out afew applications and further studies to be carried out in the future.," V we summarize and draw our conclusions, pointing out a few applications and further studies to be carried out in the future."
 For describing mathematically the nonlinear resonant waves We use the visco-resistive MHD equations., For describing mathematically the nonlinear resonant waves we use the visco-resistive MHD equations.
 In spite of the presence of dissipation we use the adiabatic equation as an approximation of the energy equation., In spite of the presence of dissipation we use the adiabatic equation as an approximation of the energy equation.
 Numerical studies by ? in linear MHD have shown that dissipation due to viscosity and finite electrical conductivity in the energy equation does not alter significantly the behaviour of resonant MHD waves in the driven problem.," Numerical studies by \citet{poedts1994}
 in linear MHD have shown that dissipation due to viscosity and finite electrical conductivity in the energy equation does not alter significantly the behaviour of resonant MHD waves in the driven problem."
 When the product of the ion (electron) gyrofrequency. Wie). and the ton (electron) collision time. τη. is much greater than one (as in the solar corona) the viscosity and finite electrical conductivity become anisotropic and viscosity is given by the Braginskil viscosity tensor (see Appendix A).," When the product of the ion (electron) gyrofrequency, $\omega_{i(e)}$, and the ion (electron) collision time, $\tau_{i(e)}$, is much greater than one (as in the solar corona) the viscosity and finite electrical conductivity become anisotropic and viscosity is given by the Braginskii viscosity tensor (see Appendix A)."
 The components of the viscosity tensor that remove the Alfvénn singularity are the shear components., The components of the viscosity tensor that remove the Alfvénn singularity are the shear components.
 The parallel and perpendicular components of anisotropic finite electrical conductivity only differ by a factor of 2. therefore. we will consider only one of them without loss of generality.," The parallel and perpendicular components of anisotropic finite electrical conductivity only differ by a factor of $2$, therefore, we will consider only one of them without loss of generality."
 The dynamics of waves in our model is described bythe visco-resistive MHD equations In Eqs. (1))-(5)), The dynamics of waves in our model is described bythe visco-resistive MHD equations In Eqs. \ref{eq:masscontinuity}) \ref{eq:totalpressureandsolenoid}) )
 v and B are the velocity and magnetic induction vectors. p the plasma pressure. p the density. οἱ the coefficient of magnetic diffusivity. y the adiabatic exponent. ande gio the magnetic permeability of free space.," $\mathbf{v}$ and $\mathbf{B}$ are the velocity and magnetic induction vectors, $\bar{p}$ the plasma pressure, $\bar{\rho}$ the density, $\overline{\lambda}$ the coefficient of magnetic diffusivity, $\gamma$ the adiabatic exponent, and $\mu_0$ the magnetic permeability of free space."
 In addition. V.S is the Braginskil viscosity (see Appendix A for full details).," In addition, $\nabla\cdot\mathbf{S}$ is the Braginskii viscosity (see Appendix A for full details)."
 ote that even though anisotropy has been considered the Hall term has been neglected., Note that even though anisotropy has been considered the Hall term has been neglected.
 We neglect the Hall term from the induction equation. which can be of the order of diffusion term in the solar corona. because the largest Hall terms in the perpendicular direction relative to the ambient magnetic field identically cancel.," We neglect the Hall term from the induction equation, which can be of the order of diffusion term in the solar corona, because the largest Hall terms in the perpendicular direction relative to the ambient magnetic field identically cancel."
 The components of the Hall term in the normal and parallel directions relative to the ambient magnetic field have no effect on the dynamics of waves in dissipative layers. hence these too are neglected.," The components of the Hall term in the normal and parallel directions relative to the ambient magnetic field have no effect on the dynamics of waves in dissipative layers, hence these too are neglected."
 For full details on the Hall term and the reasoning behind neglecting it. we refer to Appendix B. We adopt Cartesian coordinates x.y.z and limit our analysis to à static background equilibrium (vq= 0).," For full details on the Hall term and the reasoning behind neglecting it, we refer to Appendix B. We adopt Cartesian coordinates $x, y, z$ and limit our analysis to a static background equilibrium $\mathbf{v}_0=0$ )."
 We assume that all equilibrium quantities depend on x only., We assume that all equilibrium quantities depend on $x$ only.
 The equilibrium magnetic field. Bo. is unidirectional and lies in the yz-plane.," The equilibrium magnetic field, $\mathbf{B}_0$, is unidirectional and lies in the $yz$ -plane."
 The equilibrium quantities must satisfy the condition of total pressure balance. For simplicity we assume that the perturbations of all quantities are independent of v (ὁ/ὃν= 0).," The equilibrium quantities must satisfy the condition of total pressure balance, For simplicity we assume that the perturbations of all quantities are independent of $y$ $\da/\da y =0$ )."
 We note that since the magnetic field is not aligned with the z-axis. Alfvénn waves still exist.," We note that since the magnetic field is not aligned with the $z$ -axis, Alfvénn waves still exist."
 In linear theory of driven waves all perturbed quantities oscillate with the same frequency. c. which means that they can be Fourier-analyzed and taken tobe proportional to expG|kz— wt).," In linear theory of driven waves all perturbed quantities oscillate with the same frequency, $\omega$, which means that they can be Fourier-analyzed and taken tobe proportional to $\exp(i[kz-\omega t])$ ."
 Solutions are sought in. the form. of propagating waves., Solutions are sought in the form of propagating waves.
 All perturbations in these solutions depend on the combination 8=z— Vr. rather than z and f separately.," All perturbations in these solutions depend on the combination $\theta=z-Vt$ , rather than $z$ and $t$ separately,"
low rates. probably through a radiatively inefficient ADAF.,"low rates, probably through a radiatively inefficient ADAF."
 In this work. we show that the ADAF model developed by Li&Cao(2009) can explain the correlation between Eddington-scaled Kinetic power and bolometric luminosity found by (2007).," In this work, we show that the ADAF model developed by \citet{l2009} can explain the correlation between Eddington-scaled kinetic power and bolometric luminosity found by \citet{m2007}."
 In Fig. |l.," In Fig. \ref{zeta},"
 we show the kinetic power as functions of bolometric luminosity for different values of ¢. which describes the distribution of the magnetic field in the space above the dise.," we show the kinetic power as functions of bolometric luminosity for different values of $\zeta$, which describes the distribution of the magnetic field in the space above the disc."
 It is found that the slopes of the model calculations with different values of ¢ are almost constant. which is in good consistent with the observed correlation (see Fig. 1).," It is found that the slopes of the model calculations with different values of $\zeta$ are almost constant, which is in good consistent with the observed correlation (see Fig. \ref{zeta}) )."
 The slope changes very little with different values of «2 (see Fig. 2)., The slope changes very little with different values of $\beta$ (see Fig. \ref{beta}) ).
 This means that the slope of the observed correlation between {ωρα and Luin?Lee ds always consistent with our model calculations independent of the values of parameters adopted., This means that the slope of the observed correlation between $L_{\rm bol}/L_{\rm Edd}$ and $L_{\rm kin}/L_{\rm Edd}$ is always consistent with our model calculations independent of the values of parameters adopted.
 Figure 3 shows that a relatively high ¢ is required for low-. which means that field lines diverse rapidly if the field strength is relative high. in order to explain the observed correlation between A and Lii/Lead.," Figure \ref{beta_zeta} shows that a relatively high $\zeta$ is required for $\beta$, which means that field lines diverse rapidly if the field strength is relative high, in order to explain the observed correlation between $\lambda$ and $L_{\rm kin}/L_{\rm Edd}$."
 This provides useful clues to constructing detailed accretion disc/outflow models. which is beyond the scope of this work.," This provides useful clues to constructing detailed accretion disc/outflow models, which is beyond the scope of this work."
 We find that the internal energy is not important compared with other terms in Eq. C159)., We find that the internal energy is not important compared with other terms in Eq. \ref{Fkin}) ).
 The jet power is mainly related to the strength of the magnetic fields threading the dise and then to, The jet power is mainly related to the strength of the magnetic fields threading the disc and then to
case for most Seyfert 2s (Guainazzi Bianchi 2007).,case for most Seyfert 2s (Guainazzi Bianchi 2007).
 We therefore substituted the with many Gaussian. narrow-emission lines with the energy fixed to that of the most important transitions (see Table 2)). even if there is no direct evidence from eye inspections of the spectrum for most of them.," We therefore substituted the with many Gaussian, narrow-emission lines with the energy fixed to that of the most important transitions (see Table \ref{lines}) ), even if there is no direct evidence from eye inspections of the spectrum for most of them."
 The two power law indices were forced to be the same (no significant improvement is found letting them vary independently)., The two power law indices were forced to be the same (no significant improvement is found letting them vary independently).
 The fit 1s good (y7/d.0.f.=1.16/219). even if not as good as the one withMEKAL.," The fit is good $\chi^2_r$ /d.o.f.=1.16/219), even if not as good as the one with."
 The presence of O and radiative recombination continua (RRC) is in agreement with the assumption that the matter is indeed photoionized., The presence of O and radiative recombination continua (RRC) is in agreement with the assumption that the matter is indeed photoionized.
 The best fit values of Nj. EF. R and the iron line EW for this fit are 3.33x 107? em7?. 1.71. 1.28 and 218 eV. The spectrum and best fit models are shown in Fig.," The best fit values of $N_H$, $\Gamma$, $R$ and the iron line EW for this fit are $\times$ $^{23}$ $^{-2}$, 1.71, 1.28 and 218 eV. The spectrum and best fit models are shown in Fig."
 5 and Fig. 6..," \ref{spectrum} and Fig. \ref{model},"
 respectively. while the continuum parameters are summarized in Table 3..," respectively, while the continuum parameters are summarized in Table \ref{fitxmm}."
 From Fig., From Fig.
 6 it can be seen that the hard and soft components have similar fluxes at about 3 keV. The observed 0.5-2 keV flux is 2.7x107? erg em s. of the nuclear one. after correcting for absorption.," \ref{model} it can be seen that the hard and soft components have similar fluxes at about 3 keV. The observed 0.5-2 keV flux is $\times$ $^{-13}$ erg $^{-2}$ $^{-1}$, of the nuclear one, after correcting for absorption."
 This number ts typical of Seyfert 2 galaxies (Bianchi Guainazzi 2007)., This number is typical of Seyfert 2 galaxies (Bianchi Guainazzi 2007).
 About of the soft flux is in the continuum. and the remaining in emission lines (adopting the photoionization scenario described above).," About of the soft flux is in the continuum, and the remaining in emission lines (adopting the photoionization scenario described above)."
 Therefore. the column density of the photoionized matter. Nipcur. 1s (0.60.0337(4c/AQ. i.e. 3x 107(42/AQ) ο”. with AQ the solid angle subtended by the reflecting matter to the ionizing continuum.," Therefore, the column density of the photoionized matter, $N_{H,scatt}$, is $\times$ $\sigma_T^{-1}(4\pi/\Delta\Omega)$, i.e. $\times$ $^{22}(4\pi/\Delta\Omega)$ $^{-2}$ , with $\Delta\Omega$ the solid angle subtended by the reflecting matter to the ionizing continuum."
 In a biconical geometry. AQ=4z2(1—cos). with 8 the half-opening angle.," In a biconical geometry, $\Delta\Omega$ $\pi(1-cos\theta)$, with $\theta$ the half-opening angle."
 If 62257. as suggested by the fit with the “Ehime” model. NuNasun~ 3x107 em.," If $\theta$ $^{\circ}$, as suggested by the fit with the “Ehime” model, $N_{H,scatt}\sim3\times$ $^{23}$ $^{-2}$."
 The 0.5-10 keV XMM-Newton spectrum was fitted with a model similar to the one described in the previous paragraph. but letting the soft power-law index vary independently of the hard one.," The 0.5-10 keV $Newton$ spectrum was fitted with a model similar to the one described in the previous paragraph, but letting the soft power-law index vary independently of the hard one."
 The reason ts that the quality of the XMM-Newton spectrum is much lower than that of Suzaku due to the shorter exposure and the lack of p-n data. and the wealth of lines likely present in the soft band cannot be fitted individually.," The reason is that the quality of the $Newton$ spectrum is much lower than that of $Suzaku$ due to the shorter exposure and the lack of p-n data, and the wealth of lines likely present in the soft band cannot be fitted individually."
 Only a line at about 0.9 keV. which is apparent in the data (see Guainazzi et al.," Only a line at about 0.9 keV, which is apparent in the data (see Guainazzi et al."
 2002). is included.," 2002), is included."
 The best-fit values are summarised in Table 3.., The best-fit values are summarised in Table \ref{fitxmm}.
" The 4-10 keV flux is 9.5x107"" erg em-2 s.| corresponding to a 2-10 keV unabsorbed flux of 2.6x 107!! erg em-s- ! and luminosity of 1.1x 10? erg s'."," The 4-10 keV flux is $\times$ $^{-12}$ erg $^{-2}$ $^{-1}$ , corresponding to a 2-10 keV unabsorbed flux of $\times$ $^{-11}$ erg $^{-2}$ $^{-1}$ and luminosity of $\times$ $^{43}$ erg $^{-1}$."
 After comparing the Suzaku and XMM-Newton observations. we first note that the intrinsic column density is higher (about 3.3x107 against 1.8xI0 em-7).," After comparing the $Suzaku$ and $Newton$ observations, we first note that the intrinsic column density is higher (about $\times10^{23}$ against $\times10^{23}$ $^{-2}$ )."
" The difference is statistically significant (e.g.. fixing the Nj, in the XMM-Newton (Suzakin) observation to be equal to the Suzuki (XMM-Newton) best-fit value makes the fits much worse. Ay267 (38))."," The difference is statistically significant (e.g., fixing the $N_H$ in the $Newton$ $Suzaku$ ) observation to be equal to the $Suzaku$ $Newton$ ) best-fit value makes the fits much worse, $\Delta\chi^2$ =67 (38))."
 This suggests that the absorber is relatively close to the nucleus. varying on timescales no longer than a few years.," This suggests that the absorber is relatively close to the nucleus, varying on timescales no longer than a few years."
 The EW of the neutral iron line ts about 200 eV in the Suzaku observation. more than in XMM-Newton data (but still consistent within the errors). while the line flux is about the same. i.e 4.5x107 ph em? s! against 3.9x107 ph em sl: the two values are consistent each other within the errors. (," The EW of the neutral iron line is about 200 eV in the $Suzaku$ observation, more than in $Newton$ data (but still consistent within the errors), while the line flux is about the same, i.e $\times10^{-5}$ ph $^{-2}$ $^{-1}$ against $\times10^{-5}$ ph $^{-2}$ $^{-1}$; the two values are consistent each other within the errors. ("
"The same is true if the line is put the absorber. in the hypothesis that the absorbing matter is closer to the black hole than the line emitting matter. in which case. the line fluxes are 19x10 and 2.4x107 ph em"" s7!.)","The same is true if the line is put the absorber, in the hypothesis that the absorbing matter is closer to the black hole than the line emitting matter, in which case, the line fluxes are $\times10^{-5}$ and $\times10^{-5}$ ph $^{-2}$ $^{-1}$ .)"
 The best-fit value of R is also higher. but consistent within the errors. (," The best-fit value of $R$ is also higher, but consistent within the errors. ("
It is very poorly constrained in the XMM-Newton data due to the limited bandwidth.),It is very poorly constrained in the $Newton$ data due to the limited bandwidth.)
 The 2-10 keV XMM-Newron unabsorbed flux is about 1.5 times larger than the Suzaku one., The 2-10 keV $Newton$ unabsorbed flux is about 1.5 times larger than the $Suzaku$ one.
 As a result. we can conclude that the source was in a quite similar state during the Sucaku and XMM-Newton observations. albeit somewhat fainter in the former.," As a result, we can conclude that the source was in a quite similar state during the $Suzaku$ and $Newton$ observations, albeit somewhat fainter in the former."
" The main results from the 2007 Suzaku observation of the ""changing-look"" source. the Phoenix Galaxy (the first observation inhard X-rays. as the source was not observed by BeppoS AX). can be summarized as follows."," The main results from the 2007 $Suzaku$ observation of the “changing-look” source, the Phoenix Galaxy (the first observation inhard X-rays, as the source was not observed by $SAX$ ), can be summarized as follows."
 The source was clearly in a Compton-thinstate. as in 2001. when it was observed by XMM-Newton (Guainazzi et al.," The source was clearly in a Compton-thinstate, as in 2001, when it was observed by $Newton$ (Guainazzi et al."
 2002). and differently than in 1995. when ASCA caught it 1n a reflection—dominated state (Awaki et al.," 2002), and differently than in 1995, when $ASCA$ caught it in a reflection–dominated state (Awaki et al."
 2000. Guainazzi et al.," 2000, Guainazzi et al."
 2002)., 2002).
With the help of Fig.,With the help of Fig.
 2. we can identify differences and similarities between sources., \ref{sedtotal} we can identify differences and similarities between sources.
 Of special interest are possible shortcomings in the fits for those five sources that have two upper limits at the long-wavelength end. namely. 3€ 207. 4€ |28.45. 4€ | 34.47. ὃς 263. and 85 | 820. and the poorly constrained source 3C 245.," Of special interest are possible shortcomings in the fits for those five sources that have two upper limits at the long-wavelength end, namely, 3C 207, 4C $+$ 28.45, 4C $+$ 34.47, 3C 263, and S5 $+$ 820, and the poorly constrained source 3C 245."
 Firstly. we note that the values of the upper limits for the two sources 4€. | 34.47 and 3€ 263 are relatively low and. although located at different rest-frame wavelengths. appear to sample a similar cool blackbodsy: component.," Firstly, we note that the values of the upper limits for the two sources 4C $+$ 34.47 and 3C 263 are relatively low and, although located at different rest-frame wavelengths, appear to sample a similar cool blackbody component."
 Therefore. they are unlikely: to be far from the true values.," Therefore, they are unlikely to be far from the true values."
 Secondly. based on the similarity between the SEDs of the sources 3€ 207. 4€. | 28.45 and 85 | S20. and those of the sources LLL Zw 2. ος 390.3. and 4€ | 3447. respectively. the peak of the cool blackbody component of the former is unlikely to be overestimated by factors z-2.," Secondly, based on the similarity between the SEDs of the sources 3C 207, 4C $+$ 28.45 and S5 $+$ 820, and those of the sources III Zw 2, 3C 390.3, and 4C $+$ 34.47, respectively, the peak of the cool blackbody component of the former is unlikely to be overestimated by factors $\ga2$."
 Finally. judging [rom a comparison between the SEDs of the sources 3€ 245 and 3€ 47. whereas the peaks of their cool blackhocky components appear similar. the peak of the warm blackbodsy component of the former could be overestimated by a factor of —2.," Finally, judging from a comparison between the SEDs of the sources 3C 245 and 3C 47, whereas the peaks of their cool blackbody components appear similar, the peak of the warm blackbody component of the former could be overestimated by a factor of $\sim 2$."
 A characteristic property of interstellar dust are the spectral features centered at 9.7 pam and LS jui due to silicates., A characteristic property of interstellar dust are the spectral features centered at 9.7 $\mu$ m and 18 $\mu$ m due to silicates.
 In AGN. these features are expected το be. produced bv the dustv torus.," In AGN, these features are expected to be produced by the dusty torus."
 A pertinent problem inherent to the measurement of silicate features. remains the correct placement of the continuum. which ultimately decides. if they are seen in emission or in absorption.," A pertinent problem inherent to the measurement of silicate features remains the correct placement of the continuum, which ultimately decides if they are seen in emission or in absorption."
 This task is still challenging because the available spectra rarely cover a wavelength region large enough to see a sizeable portion of the continuum around the features. which are relatively broad. and often shifted. in wavelength. (sce Fig. 1)).," This task is still challenging because the available spectra rarely cover a wavelength region large enough to see a sizeable portion of the continuum around the features, which are relatively broad and often shifted in wavelength (see Fig. \ref{sed}) )."
 The advantage of our data set. however. is that the HUC and ALLS photometry considerably extend the wavelength range of the URS spectrum. thus allowing us to detect theoveralf continuum.," The advantage of our data set, however, is that the IRAC and MIPS photometry considerably extend the wavelength range of the IRS spectrum, thus allowing us to detect the continuum."
 In particular. our data shows that the ‘dip’ in the SED blucward of ~ 10 jan observed in most sources. which could be interpreted as blueshifted silicate absorption. is in [act the result of a warm dust component that is relatively week.," In particular, our data shows that the 'dip' in the SED blueward of $\sim$ 10 $\mu$ m observed in most sources, which could be interpreted as blueshifted silicate absorption, is in fact the result of a warm dust component that is relatively weak."
 Based on the overall continuum. we detect silicate in all our sources.," Based on the overall continuum, we detect silicate in all our sources."
 The only exception is the source 3C 245. for which we do not detect strong silicate emission and interpret the IRS spectrum as being dominated: by the warm ancl cool dust. components.," The only exception is the source 3C 245, for which we do not detect strong silicate emission and interpret the IRS spectrum as being dominated by the warm and cool dust components."
 We have measured the properties of both the LO jun and 18 μαι silicate emission features after subtracting from the I5 spectrum the overall continuum fitted in Section 4.1.1. and removing superimposed strong narrow emission lines., We have measured the properties of both the 10 $\mu$ m and 18 $\mu$ m silicate emission features after subtracting from the IRS spectrum the overall continuum fitted in Section \ref{blackbody} and removing superimposed strong narrow emission lines.
 Our results are listed in Table 7.., Our results are listed in Table \ref{silicates}.
 We detect the IS jum silicate feature in 4/8 sources. for which the HUS spectrum covers its location.," We detect the 18 $\mu$ m silicate feature in 4/8 sources, for which the IRS spectrum covers its location."
 In the remaining four sources this feature appears to be swamped by the cool blackbody component. which. based on its temperature (200 Ix). peaks around this wavelength.," In the remaining four sources this feature appears to be swamped by the cool blackbody component, which, based on its temperature $\sim 200$ K), peaks around this wavelength."
 Our first noteworthy finding is that the LO fum silicate emission feature is considerably recshiftecd in all oursources., Our first noteworthy finding is that the 10 $\mu$ m silicate emission feature is considerably redshifted in all oursources.
 Insteacl of at the. expected. rest-frame wavelength of 9.7 qun. we find its center to lie in the range of ~10.9 yam (lable 7.. column (6)).," Instead of at the expected rest-frame wavelength of $\sim 9.7~\mu$ m, we find its center to lie in the range of $\sim 10.1 - 10.9~\mu$ m (Table \ref{silicates}, column (6))."
 We note that stronely recshilted silicate emission features were observed previously in a few quasars (227).. but never consistently in an entire sample.," We note that strongly redshifted silicate emission features were observed previously in a few quasars \citep{Sieben05, Hao05, Schweitzer08}, but never consistently in an entire sample."
 On the other hand. the center of the IS jum silicate emission feature in the three objects for which it can be determined is stronely shifted in only one source (3€ S4) and in this case bluewaud.," On the other hand, the center of the 18 $\mu$ m silicate emission feature in the three objects for which it can be determined is strongly shifted in only one source (3C 84) and in this case blueward."
 From circumstellar dust. studies we know that. the exact wavelength location of the silicate emission maximum. depends on the grain size ancl composition. with larger erain sizes and [arger amounts of crystalline dust. (over amorphous dust) expected. to enhance the cmiissivity at longer wavelengths (e.g. 2)..," From circumstellar dust studies we know that the exact wavelength location of the silicate emission maximum depends on the grain size and composition, with larger grain sizes and larger amounts of crystalline dust (over amorphous dust) expected to enhance the emissivity at longer wavelengths \citep[e.g.,][]{Bou01}. ."
 In order to investigate if this, In order to investigate if this
enerev is (ALο ο M; is the irreducible mass of the hole.,"energy is $(M-M_i)c^2$, where $M_i$ is the irreducible mass of the hole."
" Setting E, equal to the spin energy ancl solving for the black hole spin. we find that where r=E,/(M). or jz(8r)? for r1."," Setting $E_*$ equal to the spin energy and solving for the black hole spin, we find that where $r \equiv E_*/(c^2M)$, or $j \approx (8r)^{1/2}$ for $r \ll 1$."
" Thus. for sources with outflows powered bv the spin energv of a black hole. the black hole spin j cau be measured when the total energv £L, of the outflow and the black hole mass Af ave known."," Thus, for sources with outflows powered by the spin energy of a black hole, the black hole spin $j$ can be measured when the total energy $E_*$ of the outflow and the black hole mass $M$ are known."
 Determinations of the black hole spin obtained using equation (1) do not depend upon the details of the energv extraction., Determinations of the black hole spin obtained using equation (1) do not depend upon the details of the energy extraction.
 The use of (his equation assumes only that the ultimate source of the energy is the black hole spin. and (hat. during the short period of spin energy. extraction. the black hole spin changes only because of the energy. extraction.," The use of this equation assumes only that the ultimate source of the energy is the black hole spin, and that, during the short period of spin energy extraction, the black hole spin changes only because of the energy extraction."
 The ratio r provides an important diagnostic of the black hole svstem., The ratio $r$ provides an important diagnostic of the black hole system.
 Studies of the ralio r provide insight into the physical state of the svstem at the lime the outflow is eeneraled irrespective of whether the outflow is powered by (he spin energy of the hole., Studies of the ratio $r$ provide insight into the physical state of the system at the time the outflow is generated irrespective of whether the outflow is powered by the spin energy of the hole.
 The spin energv of the black hole is thought to power the large-scale outflows [rom AGN (e.g. Dlandford Znajek 1977: Rees 1934: Blandford 1990: Wilson Colbert 1995: Daly 1995: Meier 1999. 2002).," The spin energy of the black hole is thought to power the large-scale outflows from AGN (e.g. Blandford Znajek 1977; Rees 1984; Blandford 1990; Wilson Colbert 1995; Daly 1995; Meier 1999, 2002)."
" The total οποιον ££, that will be channeled through a outflow can be determined for the most powerful classical double radio galaxies (e.g. Daly Guerra 2002: Wan. Daly. Guerra 2002: Daly οἱ al."," The total energy $E_*$ that will be channeled through a large-scale outflow can be determined for the most powerful classical double radio galaxies (e.g. Daly Guerra 2002; Wan, Daly, Guerra 2002; Daly et al."
 2008: ODea οἱ al., 2008; O'Dea et al.
 2003)., 2008).
" This empirically determined total enerev 2, can then be combined with the black hole mass to obtain the ratio r and a measure of or bound on the spin parameter of sources using eq. (", This empirically determined total energy $E_*$ can then be combined with the black hole mass to obtain the ratio $r$ and a measure of or bound on the spin parameter of sources using eq. (
1).,1).
 If only part of the spin energy is tapped. or if some of the tapped energy is dissipated or raciated away ancl does not reach the extremities of (he radio source. (he value of j obtained using eq. (," If only part of the spin energy is tapped, or if some of the tapped energy is dissipated or radiated away and does not reach the extremities of the radio source, the value of $j$ obtained using eq. ("
1) will be a lower bound.,1) will be a lower bound.
 If the outflows are powered by energy [rom some other source. such as an accretion disk. then the energy per unit black hole mass. r. could be used to probe the physics of the disk and. perhaps. place a limit on the black hole spin.," If the outflows are powered by energy from some other source, such as an accretion disk, then the energy per unit black hole mass, $r$, could be used to probe the physics of the disk and, perhaps, place a limit on the black hole spin."
 For mocleV. in which the large scale outflow is powered bv processes in (he accretion disk. the black hole spin may be studied by combining the beam power of the outflow and mass of the black hole.," For models in which the large scale outflow is powered by processes in the accretion disk, the black hole spin may be studied by combining the beam power of the outflow and mass of the black hole."
 There are 19 powerful radio galaxies for which we have estimates of both the black hole mass M and the total οποιον ο., There are 19 powerful radio galaxies for which we have estimates of both the black hole mass $M$ and the total energy $E_*$.
 All of these very powerful ERII radio sources have radio powers at least à [actor of ten above the classical FRI/FRID transition. ancl are referred (o," All of these very powerful FRII radio sources have radio powers at least a factor of ten above the classical FRI/FRII transition, and are referred to"
"If the parametric model, where only a handful of numbers can describe the beam, describes our best case realistic scenario, the non-parametric model, requiring no prior knowledge of the beam, represents our case.","If the parametric model, where only a handful of numbers can describe the beam, describes our best case realistic scenario, the non-parametric model, requiring no prior knowledge of the beam, represents our case."
" For the bands 100-217 GHz, which are the most important in terms of raw sensitivity to the CMB, one of the twice-annual crossings of Jupiter should yield a measurement of the beam function0."," For the bands 100-217 GHz, which are the most important in terms of raw sensitivity to the CMB, one of the twice-annual crossings of Jupiter should yield a measurement of the beam function."
396.. This compares to roughly from 5 years of Jupiter mapping from WMAP (?)., This compares to roughly from 5 years of Jupiter mapping from WMAP .
". Over the mission lifetime, Jupiter will be visible about four times."," Over the mission lifetime, Jupiter will be visible about four times."
" Errors will decrease somewhat faster than the square root of the number of observations in the non-parametric model, because filling in the plane constrains which modes can contribute, improving the overlap matrix."," Errors will decrease somewhat faster than the square root of the number of observations in the non-parametric model, because filling in the plane constrains which modes can contribute, improving the overlap matrix."
" Because we have some knowledge of the beams from ground tests and numerical models, these parametric and non-parametric cases should bracket the range of reasonable possibilities."," Because we have some knowledge of the beams from ground tests and numerical models, these parametric and non-parametric cases should bracket the range of reasonable possibilities."
" For the parametric model, we evaluate an extremely conservative method for dealing with the nonlinear gain of the HFI, excluding data where the gain deviates from linearity by more than the rms noise per sample."," For the parametric model, we evaluate an extremely conservative method for dealing with the nonlinear gain of the HFI, excluding data where the gain deviates from linearity by more than the rms noise per sample."
" In practice, the HFI analysis pipeline will correct the data for the nonlinear response."," In practice, the HFI analysis pipeline will correct the data for the nonlinear response."
" For most channels, this cut removes the peak region of Jupiter and Saturn, but keeps the central region of Mars, which provides information on the beam’s peak, so we include all three in the fit."," For most channels, this cut removes the peak region of Jupiter and Saturn, but keeps the central region of Mars, which provides information on the beam's peak, so we include all three in the fit."
" Separate amplitudes are fit for each of the planets, and are effectively marginalized out in the computation of the function."," Separate amplitudes are fit for each of the planets, and are effectively marginalized out in the computation of the function."
 Only the 353 GHz channel has significant nonlinear response at the peak of Mars; for this channel only we additionally include observations of Uranus and Neptune to aid the fitting., Only the 353 GHz channel has significant nonlinear response at the peak of Mars; for this channel only we additionally include observations of Uranus and Neptune to aid the fitting.
" The corresponding error in thewindow function is shown in figure 8,,channel."," The corresponding error in the function is shown in figure \ref{fig:tf-nonlinear},."
" Although we are including more planets, the exclusion of the peak of the beam on Jupiter, where signal-to-noise is highest, boosts the noise significantly."," Although we are including more planets, the exclusion of the peak of the beam on Jupiter, where signal-to-noise is highest, boosts the noise significantly."
" In this case the error on the function on the small scale end of the beam is raised by a factor of 2-25, with 545 GHz the least and 217 GHz the most affected."," In this case the error on the function on the small scale end of the beam is raised by a factor of 2–25, with 545 GHz the least and 217 GHz the most affected."
 Correcting for the non-linearity will allow much better performance., Correcting for the non-linearity will allow much better performance.
 Destriping to reduce noise is an important step in beam reconstruction., Destriping to reduce noise is an important step in beam reconstruction.
" We ran cases with unfiltered low frequency noise, which manifests as stripes across the face of the beam, and this will increase the noise in the beam fit or decomposition."," We ran cases with unfiltered low frequency noise, which manifests as stripes across the face of the beam, and this will increase the noise in the beam fit or decomposition."
" The non-parametric model is somewhat more sensitive to low-frequency noise than the parametric model, because the correlated noise will be mistaken for the true structure of the beam."," The non-parametric model is somewhat more sensitive to low-frequency noise than the parametric model, because the correlated noise will be mistaken for the true structure of the beam."
" Larger beams are also more affected than small beams, because they take longer to transit the planet."," Larger beams are also more affected than small beams, because they take longer to transit the planet."
" For both beam reconstruction models, low frequency noise increases thewindow function errors at 30 GHz (FWHM 32’) by nearly two orders of magnitude and at 217 GHz (FWHM 6.5’) by less than an order of magnitude."," For both beam reconstruction models, low frequency noise increases the function errors at 30 GHz (FWHM $32'$ ) by nearly two orders of magnitude and at 217 GHz (FWHM $6.5'$ ) by less than an order of magnitude."
" The parametric model errors increase at 857 GHz (FWHM 4’) by several tens of percent, but the non-parametric model’s errors at 857 GHz are driven by the poor decomposition into basis coefficients, and are not much affected by the destriping."," The parametric model errors increase at 857 GHz (FWHM $4'$ ) by several tens of percent, but the non-parametric model's errors at 857 GHz are driven by the poor decomposition into basis coefficients, and are not much affected by the destriping."
 We explore the tilt of the scalar perturbation spectrum with the likelihood in a simplified case where the likelihood function is analytic., We explore the tilt of the scalar perturbation spectrum with the likelihood in a simplified case where the likelihood function is analytic.
" We take slices through the likelihood, modified by the beam errors, and leave a fuller Markov Chain Monte Carlo of the likelihood to other work ?)."," We take slices through the likelihood, modified by the beam errors, and leave a fuller Markov Chain Monte Carlo of the likelihood to other work ."
". For a theoretical power spectrum, C;=I(l+ 1)C1/2m, given a beam deconvolved data spectrum Τι, which includes isotropic noise with power spectrum NV, the sky likelihood is We assume a flat prior on Cj, so that the posterior probability is proportional to the likelihood: P(C;|D;)ος L(D,\Ci)."," For a theoretical power spectrum, ${\cal C}_l = l(l+1)C_l/2\pi$ , given a beam deconvolved data spectrum ${\cal D}_l$ , which includes isotropic noise with power spectrum ${\cal N}_l$, the full-sky likelihood is We assume a flat prior on ${\cal C}_l$, so that the posterior probability is proportional to the likelihood: $P({\cal C}_l | {\cal D}_l ) \propto {\cal L}( {\cal D}_l | {\cal C}_l )$ ."
" For the beam deconvolved noise spectrum, we take which is exact for uniform white noise, and depends onlyonthetime domain noise variance (c2, see Table 2))"," For the beam deconvolved noise spectrum, we take which is exact for uniform white noise, and depends onlyonthetime domain noise variance $\sigma_t^2$ , see Table \ref{tab:det}) )"
Thea,"where $(n_i) = (n-1,n+1,n,n,n)$ and $(M_i)=(M_W,M_W,M_W,M_H,M_W)$."
"ssociated Fadeev-Popov Lagrangian becomes 1 447) P25(oO 4yA, 1... (or ο”)62 d PAGA OIE:+ Pig A,On) yo (3.-1) aug duo) au(290,0) 1 (--ὃ- 1 Tn +po)n cο (--ὃ-+ —0O MaysMos2974 Q EmMa; 290,629090 Moy 2g?ÜrEYshooqM,242 ao"," The potential $\bfV$ takes the elements $\begin{array}{rcl{\qquad}rcl}
\bfV_{11}^n&=&m_W^2\left(f^2-1\right) &\bfV_{12}^n&=&0\\
\bfV_{13}^n&=&\sqrt{2}m_W^{}f'&
\bfV_{14}^n&=&\displaystyle\sqrt{2}m_W^{}f
\frac{A+1}{r}\\
\bfV_{22}^n&=&\bfV_{11}^n&\bfV_{23}^n&=&\bfV_{13}^n\\
\bfV_{24}^n&=&-\bfV_{14}^n&
\bfV_{33}^n&=&\displaystyle
\frac{(A+1)^2}{r^2}+\frac{m_H^2}{2}\left(f^2-1\right)
+m_W^2\left(f^2-1\right)\\
\bfV_{34}^n&=&
\displaystyle -2\frac{A+1}{r^2}n&\bfV_{44}^n&=&
\displaystyle\frac{(A+1)^2}
{r^2}+\frac{3}{2}m_H^2\left(f^2-1\right) \\
\bfV_{55}^n&=&m_W^2\left(f^2-1\right)&\bfV_{i5}&=&0 \; .\\
\end{array}$ Choosing the dimensionless variable $M_Wr$ one realizes that the fluctuation operator depends only on the ratio $\xi=M_H/M_W$ , up to an overall factor $M_W^2$ which cancels in the ratio with the free operator."
" 2. day AT; QD D?> D? A(30D?+oCU)—M. Mix μοι- MjMas —ü?o A,Gogo"," We will need in the the Euclidean fluctuation modes $f_{n,i}^{\alpha\pm}(r,\nu^2)$, which we will denote as “mode functions” in the following."
 ο) ο ⋅ subscriptscontribution ofthe operatorsthei," They satisfy ) ] =0The superscript $\alpha$ labels $4$ linearly independent solutions, the subscript $i$ labels the $4$ components, $n$ refers to the partial wave."
nstanton the gauge-1igegssector ," The superscript $+$ denotes a solution regular (i.e. exponentially decreasing) as $r\to \infty$, the superscript $-$ denotes a solution regular at $r=0$ ."
becomes will masses(Mi. MyMy; Mg Mg)., The corresponding free $\bfV=0$ ) solutions are the modified Bessel functions $B^\pm_{n_i}(\kappa_i r)$ with _i _ir) _i_ir) where $\kappa_i=\sqrt{\nu^2+M_i^2}$ .
 Hds convenient, It is convenient to rewrite the mode functions as
field.,field.
 Effects of cose rays were reduced ou those images for which we had multiple exposures with thestsdus routinec, Effects of cosmic rays were reduced on those images for which we had multiple exposures with the routine.
rrej Tages of the nuclear region as observed through the FIGOBW. E300NV aud F139W filter with the PC are shown in Fie.," Images of the nuclear region as observed through the F160BW, F300W and F439W filter with the PC are shown in Fig."
 1., 1.
 Our values for the total flux of the uucleus in the different bands agree with those of Gordonetal.1999., Our values for the total flux of the nucleus in the different bands agree with those of \cite{gordon}.
. The radial profiles of the D. V and I data are also consistent with I&M93 and L9s.," The radial profiles of the B, V and I data are also consistent with KM93 and L98."
 As is apparent from Fie., As is apparent from Fig.
 l. the nucleus of M33. appears more concentrated in the FIGOBW filter than in the longer wavelength filters.," 1, the nucleus of M33 appears more concentrated in the F160BW filter than in the longer wavelength filters."
 Nevertheless. as shown in Fig.," Nevertheless, as shown in Fig."
 2. the profile in the FIGOBW image of the nucleus is extended. compared to the profile of the star located about NNW from the nucleus and to PSF profiles calculated. using the UST PSF generating routine TiuvTInu (Iaist&Wook 1997)).," 2, the profile in the F160BW image of the nucleus is extended compared to the profile of the star located about 1"" NNW from the nucleus and to PSF profiles calculated using the HST PSF generating routine TinyTim \cite{krist}) )."
 Further investigation of the radial structure of the UV cluission calls for deconvolution of the data taking iuto account the different iustriuneutal effects in theCamera (L98)., Further investigation of the radial structure of the UV emission calls for deconvolution of the data taking into account the different instrumental effects in the (L98).
 Since the UV image does not have a large enough signal-to-noise to allow for a proper deconvolution. we chose instead to fit convolved models to the data.," Since the UV image does not have a large enough signal-to-noise to allow for a proper deconvolution, we chose instead to fit convolved models to the data."
 Following IWAL93. we fitted radial profile models of the form: The model aud the PSF are oversampled on a txt exid for cach PC pixel.," Following KM93, we fitted radial profile models of the form: The model and the PSF are oversampled on a 4x4 grid for each PC pixel."
 For a eiven set of (7.1). the convolved model is moved on the txt exid. rebinned to the PC resolution aud compared to the data.," For a given set of $(r_{\rm o},n)$, the convolved model is moved on the 4x4 grid, rebinned to the PC resolution and compared to the data."
 The comparison with the data is performed in a 61x61 pixcl aperture (about 3x3) but we have verified that larger aud smaller apertures gave simular results., The comparison with the data is performed in a 64x64 pixel aperture (about 3”x3”) but we have verified that larger and smaller apertures gave similar results.
 We lave assumed an A type spectruii for the PSF but other choices do not affect our conchisions., We have assumed an A type spectrum for the PSF but other choices do not affect our conclusions.
 The parameter ο is varied between 0.05 aud 2 PC pixels and à» is varied between 0.5 aud 2., The parameter $r_{\rm o}$ is varied between 0.05 and 2 PC pixels and $n$ is varied between 0.5 and 2.
 The results roni the 47 ininimization routine are presented in Tab., The results from the $\chi^2$ minimization routine are presented in Tab.
 1 aud Fie., 1 and Fig.
 3., 3.
 The quoted errors correspoud to higher values than at the minunmuuu of the fit function., The quoted errors correspond to higher values than at the minimum of the fit function.
 In the roisy UV. baud only lower bouuds to the parameters could © extracted., In the noisy UV band only lower bounds to the parameters could be extracted.
" Tere we eive errors on the EWITIM instead of ou r,.", Here we give errors on the FWHM instead of on $r_{\rm o}$.
 The PFWIIM of the models decreases in accordance with he blue colour eracdient., The FWHM of the models decreases in accordance with the blue colour gradient.
 In all cases 0z 1 sugecstine hat the distribution of light at huge radi is the same in all the filters., In all cases $n\approx$ 1 suggesting that the distribution of light at large radii is the same in all the filters.
 This is consistent with the flat colour xofile« observed for 7Z9ro., This is consistent with the flat colour profiles observed for $r\gg r_{\rm o}$.
 With » fixed at 0.75 as in L98 we also find a FWITIM for the E555W data of 007., With $n$ fixed at 0.75 as in L98 we also find a FWHM for the F555W data of 07.
 This is clearly not the best solution when » varies (Tah., This is clearly not the best solution when $n$ varies (Tab.
 1)., 1).
 We find a higher ΕΠΑΕ (07009)., We find a higher FWHM 09).
 We note that INI93 find » between 0.51.3 and a EWITIML below (711., We note that KM93 find $n$ between 0.8–1.3 and a FWHM below 1.
" The simular values found for 5, in the V. LL F1012M and (to a lesser extent) B bands iudicate that their radial profiles are coniparable 1νο, the colour eracieuts are much reduced between those bands than when compared to the UV and U so that. for example. on first approximation theV-7 colour eradient is ucelicible when compared to CTYT-V)."," The similar values found for $r_{\rm o}$ in the V, I, F1042M and (to a lesser extent) B bands indicate that their radial profiles are comparable (i.e. the colour gradients are much reduced between those bands than when compared to the UV and U so that, for example, on first approximation the colour gradient is negligible when compared to )."
 Since the V. I and F1012M show very close light distributions. we investigated whether the compact enüssion from the UV and U filter could be explaince: wean underlving extended population having the V. bare distribution plus a blue point source.," Since the V, I and F1042M show very close light distributions, we investigated whether the compact emission from the UV and U filter could be explained by an underlying extended population having the V band distribution plus a blue point source."
" We fixed +,=1 PCE yxcl n=l and superposed at the center of this model a »out source of varving relative streneth."," We fixed $r_{\rm o}$ =1 PC pixel, $n$ =1 and superposed at the center of this model a point source of varying relative strength."
 For the V aac redward. filters. the best fits were consistently obtaiucc with a nil contribution from the point source.," For the V and redward filters, the best fits were consistently obtained with a nil contribution from the point source."
 However. a »oiut source of Increasing streneth was needed in D. U iux UV.," However, a point source of increasing strength was needed in B, U and UV."
 Oulv those models with fit values within the errors of he previous extended emission fits (Tab., Only those models with fit values within the errors of the previous extended emission fits (Tab.
 1) were kept ic. he fits here are as good or better than the previous ones (sce Fig., 1) were kept i.e. the fits here are as good or better than the previous ones (see Fig.
 3)., 3).
 The contribution of the point source to the otal flux within a E15 circular aperture is sunniarized in Tab., The contribution of the point source to the total flux within a 45 circular aperture is summarized in Tab.
 2., 2.
 The corresponding magnitudes are eiven in the VegaMAG system using updated tables for the zeropoints (Holtzmanetal.19953)., The corresponding magnitudes are given in the VegaMAG system using updated tables for the zeropoints \cite{holtzman}) ).
 The best fits for the D. U and UV are shown in Fig.," The best fits for the B, U and UV are shown in Fig."
 3 bv dashed lines., 3 by dashed lines.
 They are indistinguishable from the extended enission fits., They are indistinguishable from the extended emission fits.
 As a result we cannot. based ou the data in haud. distinguish between a model in which enüssion is extended in the UV but has a siualler core radius than at longer waveleneths aud a composite model consisting of a point source and an underlying distribution characterised by the visible light profile.," As a result we cannot, based on the data in hand, distinguish between a model in which emission is extended in the UV but has a smaller core radius than at longer wavelengths and a composite model consisting of a point source and an underlying distribution characterised by the visible light profile."
 The uucleus of N33 has à composite spectrum raugiug πο ATV at A-—J3800À to FSV at Ac (OConnell 1983)3., The nucleus of M33 has a composite spectrum ranging from A7V at $\lambda$$\sim$ to F5V at $\lambda$$\sim$ \cite{oconnell}) ).
 This requires at least a two component population in most models with the blue cmission being due to vouue moetalxich stars., This requires at least a two component population in most models with the blue emission being due to young metal-rich stars.
 The colour eracdient iu5-7 miplies this voung population is more ceutrallv condensed., The colour gradient in implies this young population is more centrally condensed.
 The uucleus could have beeu the site of episodic starbursts with the voungest stars being about LO Myr old. (OConnell1983.vandenBergh1991.Schinidtetal.1990 aud references therein).," The nucleus could have been the site of episodic starbursts with the youngest stars being about 10 Myr old \cite{oconnell,vand,schmidt} and references therein)."
 Recently Cordonetal.1999/ have argued that a single 7O Alyy old. starburst reproduces the UV to IR spectral cucrey distribution within 1755 of the center if dust is correctly taken intoaccount!., Recently \cite{gordon} have argued that a single 70 Myr old starburst reproduces the UV to IR spectral energy distribution within 5 of the center if dust is correctly taken into.
. The uucleus is also very similar to a elobular cluster and is likely to have undergone core-collapse (Ileruquistetal. 1991.. I&KM93).," The nucleus is also very similar to a globular cluster and is likely to have undergone core-collapse \cite{hernquist}, , KM93)."
 Blue stars formed in collisions at the ceuter nuelt explain the colour gradieut., Blue stars formed in collisions at the center might explain the colour gradient.
 This model may have difficulties accounting for the UV Iuniuositv (EHeruquisti. 1991.. LOS. Gordonetal.," This model may have difficulties accounting for the UV luminosity \cite{hernquist}, L98, \cite{gordon}) )."
 1999)). Alassevetal.1996 detected the nucleus in the UV but with 5 resolution., \cite{massey} detected the nucleus in the UV but with 5” resolution.
 Tence they had proposed that the blue component of the uncleus could be due to unresolved enüssion from a few hot stars., Hence they had proposed that the blue component of the nucleus could be due to unresolved emission from a few hot stars.
 However. the UST data shows the UV cussion is very compact with no stars of comparable brightuess within 5° of the center.," However, the HST data shows the UV emission is very compact with no stars of comparable brightness within 5” of the center."
 The total flux is about 63-10 0L tina 6/555 aperture or about 2.81079 lat S00 kpe., The total flux is about $\cdot$ $^{-15}$ in a 55 aperture or about $\cdot$ $^{38}$ at 800 kpc.
 From tlie best fit model. ~50% of the UV light comes from the imner Q'11I of the nucleus.," From the best fit model, $\sim$ of the UV light comes from the inner 14 of the nucleus."
 Auv model of the structure of the uucleus has to explain this UV eimissiou from a region ouly  0.55 pe across., Any model of the structure of the nucleus has to explain this UV emission from a region only $\sim$ 0.55 pc across.
 If a point source is prescut. thissource," If a point source is present, thissource"
To summarize. 2;Max(cs)i7di.,"To summarize, $\Delta \hat\tau_{\rm fold} \propto (\epsilon
u_2)^{3/2} \propto d_1^3$."
 Por comparison. A544xdi.," For comparison, $R_{\rm fold} \propto
d_1$ ."
 Since dy is small. a violation of the ideal relation Agua=0 ds more likely to indicate the presence of small-scale structure in the lens galaxy than would be indicated by a non-zero value of ja.," Since $d_1$ is small, a violation of the ideal relation $\Delta \hat\tau_{\rm fold} = 0$ is more likely to indicate the presence of small-scale structure in the lens galaxy than would be indicated by a non-zero value of $R_{\rm fold}$ ."
 Our analvtic expression for ρα is only valid for sources sulliciently. close to the caustic that higher-order terms are negligible., Our analytic expression for $\Delta \hat\tau_{\rm fold}$ is only valid for sources sufficiently close to the caustic that higher-order terms are negligible.
 To quantify this statement. we compare our analytic approximation with the differential time delay computed numerically from the exact. form of the lens equation.," To quantify this statement, we compare our analytic approximation with the differential time delay computed numerically from the exact form of the lens equation."
 The numerical analysis requires a specific lens model. so we consider an SLE with axis ratio g=0.5 as a representative example.," The numerical analysis requires a specific lens model, so we consider an SIE with axis ratio $q=0.5$ as a representative example."
 We use the software of Keeton(2001). to compute the astroid caustic. and then choose a point on the caustic. Far from a cusp. to serve as the originof the (Gri.d») frame.," We use the software of \citet{Keeton-gravlens} to compute the astroid caustic, and then choose a point on the caustic, far from a cusp, to serve as the originof the $(u_1,u_2)$ frame."
 For à given value of ον we solve the exact lens equation to obtain the image positions and time delay for the fold doublet.," For a given value of $u_2$, we solve the exact lens equation to obtain the image positions and time delay for the fold doublet."
 The top left panel ol Figure 2. shows AFjaq in units of 61: as a function of πο in units of 67. where θε is the Einstein angle of the lens.," The top left panel of Figure \ref{fig:checkTfold} shows $\Delta \hat\tau_{\rm fold}$ in units of $\theta_E^2$ as a function of $u_2$ in units of $\theta_E$, where $\theta_E$ is the Einstein angle of the lens."
 The analytic and numerical results are in σους agreement for sources within 0.0507 of the caustic. although the curves do begin to diverge as Πο increases.," The analytic and numerical results are in good agreement for sources within $0.05 \theta_E$ of the caustic, although the curves do begin to diverge as $u_2$ increases."
 The difference between the numerical and analvtie curves (which represents the error in the analytic approximation. denoted by =) is shown in the middle left panel.," The difference between the numerical and analytic curves (which represents the error in the analytic approximation, denoted by $\varepsilon$ ) is shown in the middle left panel."
 Phe bottom left panel shows the logarithmic slope of the error curve. d(Inz2/d(Inve).," The bottom left panel shows the logarithmic slope of the error curve, $d(\ln \varepsilon)/d(\ln u_2)$."
 Fogether. the middle and bottom left panels verify that our analytic approximation is accurate at order €*7.," Together, the middle and bottom left panels verify that our analytic approximation is accurate at order $\epsilon^{3/2}$."
 Furthermore. these panels suggest that the next non-vanishing term is of order «2 and has a positive cocllicient.," Furthermore, these panels suggest that the next non-vanishing term is of order $\epsilon^{5/2}$ and has a positive coefficient."
 The interesting implication is that the coelficient of the c? termi seems to vanish., The interesting implication is that the coefficient of the $\epsilon^2$ term seems to vanish.
 At our current order of approximation we are not able to determine whether this is rigorously true. ancl if so. how general it is: we merely olfer the remark in the hope that it may be useful for future analytic studies.," At our current order of approximation we are not able to determine whether this is rigorously true, and if so, how general it is; we merely offer the remark in the hope that it may be useful for future analytic studies."
 For now. we focus on the verification that our analytic approximation is accurate at the order to which we work.," For now, we focus on the verification that our analytic approximation is accurate at the order to which we work."
" Since the source position #2 is not observable. we compare the analytic ancl numerical time celavs as a function of the image separation d, in the right-hand. panels of Figure 2.."," Since the source position $u_2$ is not observable, we compare the analytic and numerical time delays as a function of the image separation $d_1$ in the right-hand panels of Figure \ref{fig:checkTfold}."
 Phe range of αι corresponds to that used for πο in the panels., The range of $d_1$ corresponds to that used for $u_2$ in the left-hand panels.
 For a canonical fold lens with d;=0.4607 (Ixeetonetal.2005).. our analytic expression. gives a very good approximation. indicating that our analysis can be applied in astrophysically relevant situationsnot just when the source is asvmptotically close to the caustic. but even when it lies some small but finite distance away.," For a canonical fold lens with $d_1 = 0.46 \theta_E$ \citep{Keeton-fold}, our analytic expression gives a very good approximation, indicating that our analysis can be applied in astrophysically relevant situations—not just when the source is asymptotically close to the caustic, but even when it lies some small but finite distance away."
 We now apply our perturbative method to the case of a source near a cusp point., We now apply our perturbative method to the case of a source near a cusp point.
 This approach has not been applied to cusp lenses before., This approach has not been applied to cusp lenses before.
 Appendix A of Ixeetonctal.(2003). derives the image positions and magnifications for a cusp triplet assuming a simplified form of the lens equation., Appendix A of \citet{Keeton-cusp} derives the image positions and magnifications for a cusp triplet assuming a simplified form of the lens equation.
 As we noted in Section Ἐν this simplified lens equation assumes that certain terms may be set to zero. using criteria that are less than rigorous.," As we noted in Section \ref{sec:intro}, this simplified lens equation assumes that certain terms may be set to zero, using criteria that are less than rigorous."
 We use the lens equation derived. from the fourth-orcer lens potential. and. use perturbation theory to verify the results of IxXeetonetal.(2003) and extend the analysis to a higher order of approximation.," We use the lens equation derived from the fourth-order lens potential, and use perturbation theory to verify the results of \citet{Keeton-cusp} and extend the analysis to a higher order of approximation."
 We also study time delays for à cusp lens for the first time., We also study time delays for a cusp lens for the first time.
 Our analysis does not involve simplifying assumptions. and indicates that perturbation theory is a powerful method in the study of lensing.," Our analysis does not involve simplifying assumptions, and indicates that perturbation theory is a powerful method in the study of lensing."
 We again expand the image positions. magnifications and time delays in the parameter e. but with one notable dillerence.," We again expand the image positions, magnifications and time delays in the parameter $\epsilon$, but with one notable difference."
" For à cusp oriented in the ay direction. a small “horizontal” displacement of cv, from the cusp point permits a “vertical” displacement of only c*7903 (see Fig. 1))."," For a cusp oriented in the $u_1$ direction, a small “horizontal” displacement of $\epsilon u_1$ from the cusp point permits a “vertical” displacement of only $\epsilon^{3/2} u_2$ (see Fig. \ref{fig:u-theta-def}) ),"
 since larger vertical displacements would imply a source position outside the caustic (Blandford&Naravan1986)., since larger vertical displacements would imply a source position outside the caustic \citep{Blandford_Narayan}.
. Phe lens equation is then (scePettersetal.2001.Pheoreni9.1).. where the 65 term of equation (13)) does not appear. since coss(0)=0for a cusp. corresponding to =0 in equation (11)).," The lens equation is then \citep[see][Theorem 9.1]{Petters_book}, where the $\theta_2^2$ term of equation \ref{eq:lens1-u2}) ) does not appear, since $\psi_{222}(\bmo) = 0$for a cusp, corresponding to $h = 0$ in equation \ref{eq:psiexp}) )."
 As before. we write the image positions asà series expansionin e. but now keeping an aclelitional term: (i-c.. ;Al ," As before, we write the image positions asa series expansionin $\epsilon$ , but now keeping an additional term (i.e., $\gamma_i
\epsilon^{3/2}$ )."
‘This is necessary since the vertical component of the source position enters the lens equation as €*7 wo. rather than cni as in the Fold case.," This is necessary since the vertical component of the source position enters the lens equation as $\epsilon^{3/2}
u_2$ , rather than $\epsilon u_2$ as in the fold case."
 We have The lens equation then becomes, We have The lens equation then becomes
out to be a physically more interesting case for studying recoiled BH - host interactions without compromising the general validity of our findings.,out to be a physically more interesting case for studying recoiled BH - host interactions without compromising the general validity of our findings.
" We first consider an isolated, massive, gas-rich galaxy, that we simulate at high numerical resolution."," We first consider an isolated, massive, gas-rich galaxy, that we simulate at high numerical resolution."
 The initial conditions of this galaxy are set-up as described in detail in ?.., The initial conditions of this galaxy are set-up as described in detail in \citet{Springel2005b}.
" Specifically, the total mass of the galaxy is Mooo=6.28x1013Μο (with h= 0.7), its virial velocity is v200=300kms-!, the dark matter halo concentration is c= 9, and the spin parameter is 1=0.033."," Specifically, the total mass of the galaxy is $M_{\rm 200} = 6.28 \times 10^{12}\,h^{-1}{\rm M}_\odot\,$ (with $h = 0.7$ ), its virial velocity is $v_{\rm 200}= 300\, {\rm km\,
 s^{-1}}$, the dark matter halo concentration is $c\,=\,9$ , and the spin parameter is $\lambda \,=\, 0.033$."
" The fraction of the total mass in the disc is πια= 0.041, while the bulge mass fraction amounts to mp=0.008."," The fraction of the total mass in the disc is $m_{\rm d} \,=\, 0.041$ , while the bulge mass fraction amounts to $m_{\rm
 b}\,=\, 0.008$."
 The galactic disc scale length is 4.8h~'kpc.," The galactic disc scale length is $4.8 \,h^{-1} \, {\rm
 kpc}$."
" The disc is gas-rich, with a gas fraction of fgas=0.6, while the remaining 40% are in stars."," The disc is gas-rich, with a gas fraction of $f_{\rm gas}\,=\,0.6$, while the remaining $40\%$ are in stars."
Circular velocity curves for the different galactic components are shown in Figure 1..,Circular velocity curves for the different galactic components are shown in Figure \ref{vcirc}.
" We select the numerical parameters of the multiphase model for star formation such that the gas consumption timescale is long, specifically we adopt t2=12.6Gyrs, Ao=6000, Τεν=6x105K (inthenotationof?),, and we assume a softer equation of state with qgos=0.5."," We select the numerical parameters of the multiphase model for star formation such that the gas consumption timescale is long, specifically we adopt $t_{*}^{\rm 0} = 12.6\, {\rm Gyrs}$, $A_{\rm 0} = 6000$, $T_{\rm SN}= 6 \times
10^8\, {\rm K}$ \citep[in the notation of][]{SpringelH2003}, and we assume a softer equation of state with $q_{\rm EOS} \,=\, 0.5$."
 This choice of parameters assures that during the simulated time-span (which can be a considerable fraction of the Hubble time) the galaxy stays comparatively gas-rich and the gaseous disc is stable., This choice of parameters assures that during the simulated time-span (which can be a considerable fraction of the Hubble time) the galaxy stays comparatively gas-rich and the gaseous disc is stable.
" We additionally perform runs where we keep all numerical and physical parameters the same, except for the equation of state parameter which we set to 0.05, corresponding to a nearly isothermal equation of state and resulting in a very clumpy disc."," We additionally perform runs where we keep all numerical and physical parameters the same, except for the equation of state parameter which we set to 0.05, corresponding to a nearly isothermal equation of state and resulting in a very clumpy disc."
" The galaxy is simulated with 300000 dark matter particles in the halo, 200000 star particles in the disc, 200000 gas particles in the disc and 50000 star particles belonging to the bulge."," The galaxy is simulated with $300000$ dark matter particles in the halo, $200000$ star particles in the disc, $200000$ gas particles in the disc and $50000$ star particles belonging to the bulge."
" The gravitational softening of the disc and bulge is set to 60A7!pc, and of the dark matter halo to 2h!kpc."," The gravitational softening of the disc and bulge is set to $60\,h^{-1} \, {\rm pc}$, and of the dark matter halo to $2\,h^{-1} \,
{\rm kpc}$."
" We first evolve the isolated gas-rich galaxy for 2.8x10? yrs without a central BH, until the gas disc develops a well defined spiral structure which is long lived."," We first evolve the isolated gas-rich galaxy for $2.8 \times 10^8\,$ yrs without a central BH, until the gas disc develops a well defined spiral structure which is long lived."
" Thus, during the simulated time of the central BH the gaseous disc will be in a stable, quasi stationary state."," Thus, during the simulated time of the central BH the gaseous disc will be in a stable, quasi stationary state."
" We then introduce a BH particle in the centre of the galaxy, assuming that such a gas-rich galaxy should already contain a supermassive BH."," We then introduce a BH particle in the centre of the galaxy, assuming that such a gas-rich galaxy should already contain a supermassive BH."
 To be confident that the dynamical friction force acting on the BH is numerically well resolved we set the initial mass of the BH to 5x105A7!Mg in most of ourlations!.," To be confident that the dynamical friction force acting on the BH is numerically well resolved we set the initial mass of the BH to $5 \times 10^8 \,h^{-1}{\rm M}_\odot\,$ in most of our."
". We have also investigated a lower mass BH of 5x10h! (corresponding to ~107? Mbuige), and a simulation whereMa the host galaxy is simulated at twice as high spatial resolution to verify the validity of our findings (see Appendix A))."," We have also investigated a lower mass BH of $5 \times 10^7 \,h^{-1}{\rm M}_\odot\,$ (corresponding to $\sim 10^{-3}
M_{\rm bulge}$ ), and a simulation where the host galaxy is simulated at twice as high spatial resolution to verify the validity of our findings (see Appendix \ref{appen}) )."
" Given that we want to track dynamical friction forces acting on the BH as accurately as possible, we do not perform simulations with a lower mass BHs here, because then the mass ratio between the BH and the gas or star particles would not be sufficiently large."," Given that we want to track dynamical friction forces acting on the BH as accurately as possible, we do not perform simulations with a lower mass BHs here, because then the mass ratio between the BH and the gas or star particles would not be sufficiently large."
" Also, by simulating a more massive BH we are in a more advantageous regime for spatially resolving regions close to the Bondi radius."," Also, by simulating a more massive BH we are in a more advantageous regime for spatially resolving regions close to the Bondi radius."
" Once the BH particle is introduced in the centre we let it evolve for an additional 1.4x10"" after which the BH is imparted a certain recoil velocity, with the direction either parallel to the galactic disc or perpendicular to it."," Once the BH particle is introduced in the centre we let it evolve for an additional $1.4 \times 10^7\,$ after which the BH is imparted a certain recoil velocity, with the direction either parallel to the galactic disc or perpendicular to it."
" Table 1 summarises the main parameters of our simulated galaxies, as well as initial BH masses and recoil velocities."," Table \ref{tab_kicks}
 summarises the main parameters of our simulated galaxies, as well as initial BH masses and recoil velocities."
" For comparison, we also carried out runs where the BH stays in the centre of the galaxy, and is not subject to any gravitational recoil."," For comparison, we also carried out runs where the BH stays in the centre of the galaxy, and is not subject to any gravitational recoil."
" The simulations with the stationary BH permit us to verify that the BH position always coincides with the minimum of the gravitational potential within the spatial resolution of the simulation, hence the resolution is sufficiently high to reduce two-body noise acting on the BH to negligible levels."," The simulations with the stationary BH permit us to verify that the BH position always coincides with the minimum of the gravitational potential within the spatial resolution of the simulation, hence the resolution is sufficiently high to reduce two-body noise acting on the BH to negligible levels."
" Most of our simulations were performed with BH accretion and feedback included, but we also considered cases where BH accretion, or selectively only BH feedback, was switched-off."," Most of our simulations were performed with BH accretion and feedback included, but we also considered cases where BH accretion, or selectively only BH feedback, was switched-off."
" This allows us to disentangle purely dynamical processes acting on a recoiled BH from the additional effects due to BH accretion and feedback, which can affect the properties of the surrounding gas and thus possibly alter the BH dynamics."," This allows us to disentangle purely dynamical processes acting on a recoiled BH from the additional effects due to BH accretion and feedback, which can affect the properties of the surrounding gas and thus possibly alter the BH dynamics."
 We first want to assess the importance of dynamical friction forces acting on the BH as it moves trough the simulated galaxy., We first want to assess the importance of dynamical friction forces acting on the BH as it moves trough the simulated galaxy.
" For this purpose, we consider the BH to be a collisionless particle which does not accrete, and has a constant mass of 5x10°h~'Mo."," For this purpose, we consider the BH to be a collisionless particle which does not accrete, and has a constant mass of $5 \times 10^8 \,h^{-1}{\rm
 M}_\odot\,$."
" Just before the BH is imparted the kick, we estimate the central escape velocity from the halo based on the minimum of the gravitational potential, which yields vesc(0)=V2(r=0)~1450kms !."," Just before the BH is imparted the kick, we estimate the central escape velocity from the halo based on the minimum of the gravitational potential, which yields $v_{\rm esc}(0) = \sqrt{2
 \,\Phi(r=0)} \sim 1450\,{\rm km\, s^{-1}}$ ."
" Note that the central escape velocity is rather high, vesc(0)~4.8v29o, due to the presence of a central stellar bulge."," Note that the central escape velocity is rather high, $v_{\rm esc}(0) \sim 4.8\, v_{\rm 200}$, due to the presence of a central stellar bulge."
" Using instead the oftenadopted equationΌρο= 4/2f(c)vooo,where f(c)=c(In(1+ is only a function of the concentration parameter, one would obtain a significantly lower estimate of vesc(0)~3.6 7290, as this considers only the distribution of the dark matter."," Using instead the oftenadopted equation$v_{\rm esc} = \sqrt{2\, f(c)}\, v_{\rm 200}$ ,where $f(c) = c \,(\ln(1+c) - c/(1+c))^{-1}$ is only a function of the concentration parameter, one would obtain a significantly lower estimate of $v_{\rm esc}(0) \sim 3.6 \, v_{\rm 200}$ , as this considers only the distribution of the dark matter."
yeaks in the short period regime may also be associated with as— unidentified artefacts. since they are seen in more than one object.,"peaks in the short period regime may also be associated with as--yet unidentified artefacts, since they are seen in more than one object."
 Based on this. we treat all peaks with frequencies higher han 10007/Hz with caution.," Based on this, we treat all peaks with frequencies higher than $\mu$ Hz with caution."
 On the other hand if they are confirmed © be intrinsic to the stars. then all stars presented in this paper are iybrid pulsators.," On the other hand if they are confirmed to be intrinsic to the stars, then all stars presented in this paper are hybrid pulsators."
" We have also found peaks in the ""transition region"" between he p- and g-modes (i.e. between 500 and 2000 Η2).", We have also found peaks in the “transition region” between the $p$ - and $g$ -modes (i.e. between 500 and 2000 $\mu$ Hz).
 These oaks were observed in the sdBV stars 0090100001. 22697388 by Baranetal.(2009) and. Reedetal.(2010).. respectively.," These peaks were observed in the sdBV stars 090100001, 2697388 by \cite{baran09} and \cite{reed10}, respectively."
 Their authenticity also needs confirmation., Their authenticity also needs confirmation.
 The overall survey strategy and properties of the candidates have been described in prior papers: Paper VI presents all selected targets for Q3 and Q4., The overall survey strategy and properties of the candidates have been described in prior papers; Paper VI presents all selected targets for Q3 and Q4.
 Further analysis of the pulsation spectra of these stars Gand several from the first part of the survey) is described by (Reedetal.2011.PaperVIID.. who attempt to disentangle asymptotic behaviour of g-mode pulsations in HHer class stars observer byKepler andCoRoT.," Further analysis of the pulsation spectra of these stars (and several from the first part of the survey) is described by \cite[][Paper VIII]{reed11}, who attempt to disentangle asymptotic behaviour of $g$ -mode pulsations in Her class stars observer by and."
 Data from Q3 eover 3 months. starting from BJD-2-222455092.7]23 (Q3.l) and— ending —at /2455182.0065 (Q3.3).," Data from Q3 cover 3 months, starting from 2455092.7123 (Q3.1) and ending at 2455182.0065 (Q3.3)."
 QJ covers the next 3 months from 22455|84.3679 (Q4.1) to. 2455274.7137. (Q4.3)., Q4 covers the next 3 months from 2455184.8679 (Q4.1) to 2455274.7137 (Q4.3).
. One. safe-mode event occurred. during Q3., One safe-mode event occurred during Q3.
 As a result. data for two days (between— 22455154 and 2455156) are not available.," As a result, data for two days (between 2455154 and 2455156) are not available."
 As it happened at the beginning of the third month in Q3. it did not cause any gaps in the data presented here.," As it happened at the beginning of the third month in Q3, it did not cause any gaps in the data presented here."
 However. the third month of Q3 (03.3) was truncated by those 2 days. slightly reducing the frequency resolution as well as slightly increasing the noise level in the amplitude speetrum as a result of less data (and fewer pulsation cycles).," However, the third month of Q3 (Q3.3) was truncated by those 2 days, slightly reducing the frequency resolution as well as slightly increasing the noise level in the amplitude spectrum as a result of less data (and fewer pulsation cycles)."
 Two stars in this paper are affected by the safe-mode event: 58302197 and 111558725., Two stars in this paper are affected by the safe-mode event: 8302197 and 11558725.
 Beside this safe-mode event. theKepler spacecraft. lost its fine pointing control twice.," Beside this safe-mode event, the spacecraft lost its fine pointing control twice."
 The first was a brief interval between 2245099.9] and 2455100.04. which was quickly corrected.," The first was a brief interval between 245099.91 and 2455100.04, which was quickly corrected."
 The second. longer pointing control error fell between 22453113.05 and 2455113.83. and required correction via ground command.," The second, longer pointing control error fell between 2455113.05 and 2455113.83, and required correction via ground command."
 Only data on 77668647 were affected by this fault., Only data on 7668647 were affected by this fault.
 A log of the photometric data is presented in Table I.., A log of the photometric data is presented in Table \ref{phot_log}.
 All data were obtained through the Asteroseismic Science Consortium (KASC) — see Gillilandetal.(2010)., All data were obtained through the Asteroseismic Science Consortium (KASC) – see \cite{gilliland10}.
.. The KASC data include raw as well as calibrated fluxes., The KASC data include raw as well as calibrated fluxes.
 Calibration takes into account all known instrumental issues as well as contamination from nearby objects., Calibration takes into account all known instrumental issues as well as contamination from nearby objects.
 Since spectroscopic parameters of stars in theKepler database are. generally. not precise. we believe that this contamination may not be evaluated correctly.," Since spectroscopic parameters of stars in the database are, generally, not precise, we believe that this contamination may not be evaluated correctly."
 That is why we decided to use raw fluxes for all stars in our analysis and not to account for contamination., That is why we decided to use raw fluxes for all stars in our analysis and not to account for contamination.
 This may lead to decreased amplitudes of detected modes by a factor of (1c). where ¢ is a contamination factor.," This may lead to decreased amplitudes of detected modes by a factor of $(1-c)$, where $c$ is a contamination factor."
 Future processing. supported by ground based observation of the relevant fields. will allow more accurate determination of the correction factor.," Future processing, supported by ground based observation of the relevant fields, will allow more accurate determination of the correction factor."
responsible for heating the He gas.,responsible for heating the $\alpha$ gas.
 ? and ? show that magnetic fields in these filaments can be relatively strong., \citet{fabian2008} and \citet{ferland2009} show that magnetic fields in these filaments can be relatively strong.
 Small MHD shocks or waves may also play a role in mixing the coolest X-ray and cold Ha gas phases., Small MHD shocks or waves may also play a role in mixing the coolest X-ray and cold $\alpha$ gas phases.
 ? propose that passing shock waves can accelerate the ICM past the much denser He filaments. which induces shearing and subsequently Kelvin-Helmholtz instabilities that mix the two phases.," \citet{werner2010} propose that passing shock waves can accelerate the ICM past the much denser $\alpha$ filaments, which induces shearing and subsequently Kelvin-Helmholtz instabilities that mix the two phases."
 The hot ICM electrons then ionize and heat the cool gas. which induces Ha emission.," The hot ICM electrons then ionize and heat the cool gas, which induces $\alpha$ emission."
 Shocks and sound waves have been seen in recent Chandra observations of Abell 2052 (?).., Shocks and sound waves have been seen in recent Chandra observations of Abell 2052 \citep{blanton2009}.
 Chandra images show alot of smaller shock fronts or sound waves around the edges of the bubbles. and also just in front of the blob of Ha emissio," Chandra images show a lot of smaller shock fronts or sound waves around the edges of the bubbles, and also just in front of the blob of $\alpha$ emission."
 ? base this hypothesis on the fact that most Ha regions are seen in the downstream regions of a shock., \citet{werner2010} base this hypothesis on the fact that most $\alpha$ regions are seen in the downstream regions of a shock.
 Therefore. this mechanism could explain why the X-ray gas would cool non-radiatively below 0.5 keV. Mixing and shock heating may also explain why the gas does not cool all the way to star-formitΠα» temperatures.," Therefore, this mechanism could explain why the X-ray gas would cool non-radiatively below 0.5 keV. Mixing and shock heating may also explain why the gas does not cool all the way to star-forming temperatures."
 Since observations ofM87.. for example. have show! that the Ha filaments contain dust (2).. the shocks should just gently heat the gas and not evaporate the dust.," Since observations of, for example, have shown that the $\alpha$ filaments contain dust \citep{sparks1993}, the shocks should just gently heat the gas and not evaporate the dust."
 The multi-temperature region of appears to be fully consistent with the Ha and 0.5 keV gas interactions as seen inM87., The multi-temperature region of appears to be fully consistent with the $\alpha$ and 0.5 keV gas interactions as seen in.
. However. in a recent Ha survey of 23 clusters (?).. ts one of few objects where the NUV/Hw flux ratio is consistent with shock heating.," However, in a recent $\alpha$ survey of 23 clusters \citep{mcdonald2010}, is one of few objects where the $\alpha$ flux ratio is consistent with shock heating."
 Therefore. this region may prove to be a very interesting special case.," Therefore, this region may prove to be a very interesting special case."
 More detailed Πα. CO. UV measurements are necessary to unravel the heating and feedback mechanisms operating in this filament.," More detailed $\alpha$, CO, UV measurements are necessary to unravel the heating and feedback mechanisms operating in this filament."
 Using a deep XMM-Newton observation (95 ks). we have derived 2D maps of the core of the cluster Abell 2052.," Using a deep XMM-Newton observation (95 ks), we have derived 2D maps of the core of the cluster Abell 2052."
 In the maps. we discover a cold front at a distance of ~ 130 kpe in the South-Western direction from the central cD galaxy.," In the maps, we discover a cold front at a distance of $\sim$ 130 kpc in the South-Western direction from the central cD galaxy."
 Close to the cD galaxy in the North-Western direction. we find à multi- region with cool X-ray emitting gas as low as 0.5 keV. From a careful analysis of these regions. we conclude that:," Close to the cD galaxy in the North-Western direction, we find a multi-temperature region with cool X-ray emitting gas as low as 0.5 keV. From a careful analysis of these regions, we conclude that:"
"'To correct the effects of the artifact, we use the inter-order regions to determine the spatial shape and time-variability of the residual background.","To correct the effects of the artifact, we use the inter-order regions to determine the spatial shape and time-variability of the residual background."
 This approach is very similar to that done by Starkey et al. (, This approach is very similar to that done by Starkey et al. (
"2011, in prep).","2011, in prep)."
" Because of pixel-based noise, we generally need to average BCDs together as a function of time to improve our determination of the background."," Because of pixel-based noise, we generally need to average BCDs together as a function of time to improve our determination of the background."
" At SL wavelengths the presence of the “peak up” arrays on the right-side of the BCD make the available inter-order regions very small, since scattered light from the peak up arrays affects the nearby inter-order regions."," At SL wavelengths the presence of the “peak up” arrays on the right-side of the BCD make the available inter-order regions very small, since scattered light from the peak up arrays affects the nearby inter-order regions."
 We are therefore limited to the small strip between the SL1 and SL2/SL3 orders to characterize the vertical spatial dependence of the residual background (this region is highlighted in Figure A21))., We are therefore limited to the small strip between the SL1 and SL2/SL3 orders to characterize the vertical spatial dependence of the residual background (this region is highlighted in Figure \ref{fig:vstime}) ).
 For the SL orders we determine the median BCD value as a function of time and vertical row in a box ~5 pixels wide by 11 pixels high by 3 BCDs long., For the SL orders we determine the median BCD value as a function of time and vertical row in a box $\sim$ 5 pixels wide by 11 pixels high by 3 BCDs long.
 We then create a correction image for each BCD where all columns contains this median., We then create a correction image for each BCD where all columns contains this median.
" Finally, we apply the SL flatfield to get the appropriate value within the orders."," Finally, we apply the SL flatfield to get the appropriate value within the orders."
" At LL wavelengths, the lack of peak up arrays make this process much simpler."," At LL wavelengths, the lack of peak up arrays make this process much simpler."
" In addition, there is far less time variability in the LL residual background so we use a simpler correction derived from the median BCDs shown in Figure A20.."," In addition, there is far less time variability in the LL residual background so we use a simpler correction derived from the median BCDs shown in Figure \ref{fig:llmedbcd}."
" We mask the LL orders and fit a horizontal polynomial to each row of the BCD, filling in each row of the correction image with this fit."," We mask the LL orders and fit a horizontal polynomial to each row of the BCD, filling in each row of the correction image with this fit."
 We then apply a vertical polynomial smoothing in a 10 pixel box and finally use the LL flatfield image to obtain the correction in the LL orders., We then apply a vertical polynomial smoothing in a 10 pixel box and finally use the LL flatfield image to obtain the correction in the LL orders.
 For both the SL and LL we subtract these correction images from the original BCDs., For both the SL and LL we subtract these correction images from the original BCDs.
 'The effect of the correction on several spectra can be seen in Figure A18 in black., The effect of the correction on several spectra can be seen in Figure \ref{fig:mismatch} in black.
 The agreement between orders in the overlap region is essentially always improved after the correction., The agreement between orders in the overlap region is essentially always improved after the correction.
" Most of this correction comes from removing flux in the 7—8 rregion, coincident with the 7.7 PAH feature."," Most of this correction comes from removing flux in the $-$ 8 region, coincident with the 7.7 PAH feature."
" Because of the limited information about the spatial structure of the residual background and the large effect it has on the PAH feature, we use both the corrected and uncorrected spectra to analyze the PAH band ratios in this paper."," Because of the limited information about the spatial structure of the residual background and the large effect it has on the PAH feature, we use both the corrected and uncorrected spectra to analyze the PAH band ratios in this paper."
 Future work my reveal the source of this residual background and improve our determination of the PAH feature strengths., Future work my reveal the source of this residual background and improve our determination of the PAH feature strengths.
 To conclude we make a few remarks about the sort of observations this artifact is likely to affect., To conclude we make a few remarks about the sort of observations this artifact is likely to affect.
" Generally, the residual background level is not an issue in typical IRS observations when either (1) an in-slit background can be subtracted, as in the case of point source observations; (2) the background observation is very close in time to the mapping observation, for example when a “nod” position can be used for background subtraction; or when the object being mapped is bright enough that these low-level variations are negligible."," Generally, the residual background level is not an issue in typical IRS observations when either (1) an in-slit background can be subtracted, as in the case of point source observations; (2) the background observation is very close in time to the mapping observation, for example when a “nod” position can be used for background subtraction; or (3) when the object being mapped is bright enough that these low-level variations are negligible."
 In the SMC we are observing(3) faint emission using the full IRS slit and our “off” observation occurs after ~600 pointings., In the SMC we are observing faint emission using the full IRS slit and our “off” observation occurs after $\sim$ 600 pointings.
" Thus, the oobservations present a somewhat unique situation where the residual background artifacts are important to remove."," Thus, the observations present a somewhat unique situation where the residual background artifacts are important to remove."
If we now assume a search strategy with observatious regularly spaced at interval 9/. the above distribution is modified to vield the cistributiou of age /4 at detection: In the limit of coutinuous observatious. οἱ—(0. the two expressions become identical as requirect.,"If we now assume a search strategy with observations regularly spaced at interval $\delta t$, the above distribution is modified to yield the distribution of age $t_1$ at detection: In the limit of continuous observations, $\delta t \rightarrow 0$, the two expressions become identical as required."
 Now let us integrate over redshift., Now let us integrate over redshift.
 The rate of orphan cletectious per uuit /4 becomes where nmi is the magnitude limit of the survey., The rate of orphan detections per unit $t_1$ becomes where $m$ is the magnitude limit of the survey.
 Again following Nakar et al (2002). we take n(z) to grow as n(z)x10 for 2€cU Land hold n(z) constant for z>1.," Again following Nakar et al (2002), we take $n(z)$ to grow as $n(z) \propto 10^{0.75z}$ for $z\le1$ and hold $n(z)$ constant for $z \ge 1$ ."
 We fix the cosmology to currently fashionable values. Le... Ay=70kis!MpcLAG 20.7. Oy=0.3.," We fix the cosmology to currently fashionable values, i.e., $H_0 = 70 \kmsMpc$, $\Lambda_0
= 0.7$, $\Omega_0 = 0.3$."
" Numerical inteeration OE Fu,04) was done using the ""Libeosin C library (Yoshikawa 1999).", Numerical integration of $R_{orph}(t_1)$ was done using the “Libcosm” C library (Yoshikawa 1999).
 The overall GRB rate density. was normalized to set the total ou-axis GRB rate over the redshift range 0«z5 to about 900 GRB/vear. using a lixed rest-frame /j;4=0.7 day.," The overall GRB rate density was normalized to set the total on-axis GRB rate over the redshift range $0<z<5$ to about 900 GRB/year, using a fixed rest-frame $\tjet = 0.7$ day."
 This matches reasonably well the BATSE rate., This matches reasonably well the BATSE rate.
 A more accurate normalization may be impossible eiven present observational tucertatuties., A more accurate normalization may be impossible given present observational uncertainties.
 The shape of the /4 probability distribution is indepeudent of the exact normalization of the GRB event rate., The shape of the $t_1$ probability distribution is independent of the exact normalization of the GRB event rate.
" Results of our numerical /2,,,5(/1) calculations are shown in figure L.", Results of our numerical $R_{orph}(t_1)$ calculations are shown in figure \ref{t1dist}.
 A typical rate curve will show either three or four segments., A typical rate curve will show either three or four segments.
 The first is a rising segment. determined by the jet break ime and the observational time sampling interval.," The first is a rising segment, determined by the jet break time and the observational time sampling interval."
" The second is a “plateau” where Ie,t~ —J. esseutially constant.", The second is a “plateau” where $R_{orph}(t_1)$ is essentially constant.
 This occurs when /41 exceeds both the sampling interval 9/ aud the jet break ime /;4- and simultaneously the sensitivity of the survey is sufficient to detect orphaus up to very [unLE.eh redshilt.," This occurs when $t_1$ exceeds both the sampling interval $\delta t$ and the jet break time $\tjet$, and simultaneously the sensitivity of the survey is sufficient to detect orphans up to very high redshift."
 Such coucitious are only met for comparatively seusitive surveys aud short jet break lunes., Such conditions are only met for comparatively sensitive surveys and short jet break times.
 The next segment is a falling one. whose typical time scale is determined by the magnitude unit of the survey.," The next segment is a falling one, whose typical time scale is determined by the magnitude limit of the survey."
 The last component is a low-level. late time tail comprised of the closest bursts. which cau be detected at large augles off axis but over a very siuall volume.," The last component is a low-level, late time tail comprised of the closest bursts, which can be detected at large angles off axis but over a very small volume."
" We see that the peak of /25,,5,0/4) occurs on times ranging [rom a couple of days to a couple of weeks. depending on the characteristic jet break time aud the survey maeuitucle limit."," We see that the peak of $R_{orph}(t_1)$ occurs on times ranging from a couple of days to a couple of weeks, depending on the characteristic jet break time and the survey magnitude limit."
 Moreover. it can be quite sharply peaked for surveys of relatively bright flux limits (~ 21st niaguitude) but is rather broad lor more sensitive surveys.," Moreover, it can be quite sharply peaked for surveys of relatively bright flux limits $\sim$ 21st magnitude) but is rather broad for more sensitive surveys."
 [now outline an observing strategy [or optical regimeorphan afterglow surveys that incorporates the results of precediug sections together with earlier studies., I now outline an observing strategy for optical regimeorphan afterglow surveys that incorporates the results of preceding sections together with earlier studies.
In terms of hard binaries one could argue. for the sake of semantics. that in comparison toa population drawn from a uniform distribution of periods extending from 0.1d (without restriction) our initial distributions are under-sampling the contribution of hard binaries.,"In terms of hard binaries one could argue, for the sake of semantics, that in comparison to a population drawn from a uniform distribution of periods extending from $0.1\,$ d (without restriction) our initial distributions are under-sampling the contribution of hard binaries."
 A Κον point here is (hat short-period binaries were not excluded from the primordial populations of our sinnlations by some ad-hoc process., A key point here is that short-period binaries were not excluded from the primordial populations of our simulations by some ad-hoc process.
 Instead. the distribution of orbital periods is dictated by using distributions borne from observations in combination with accounting for pre-anain sequence (AIS) evolution. — before contracting along the ILavashi track the stellar radius of a pre-MS star can be a factor of five or more greater than on the ZAAIS (Siess. Dufour Forestini 2000) and birth periods must allow for this (IXroupa 1995).," Instead the distribution of orbital periods is dictated by using distributions borne from observations in combination with accounting for pre-main sequence (MS) evolution – before contracting along the Hayashi track the stellar radius of a pre-MS star can be a factor of five or more greater than on the ZAMS (Siess, Dufour Forestini 2000) and birth periods must allow for this (Kroupa 1995)."
 Pre-MS evolution was not considered by Ivanova et al. (, Pre-MS evolution was not considered by Ivanova et al. (
2005) although thev did reject svstems where one or both stars would initially fill (heir Roche-lobes al pericentve — this was also assumed in our models.,2005) although they did reject systems where one or both stars would initially fill their Roche-lobes at pericentre – this was also assumed in our models.
 In Figure 3. we show the evolution of the core binary fraction for the four N-body sinilations introduced above., In Figure \ref{f:fig3} we show the evolution of the core binary fraction for the four $N$ -body simulations introduced above.
 Also shown is the binary fraction within the Lagrangian radius and the overall binary. fraction of the model clusters., Also shown is the binary fraction within the Lagrangian radius and the overall binary fraction of the model clusters.
 Except at late times in the IX24-50 model. when the cluster has lost more than ol ils original mass aud is nearing dissolution. we see that in each case the cluster binary fraction remains close to the primordial value.," Except at late times in the K24-50 model, when the cluster has lost more than of its original mass and is nearing dissolution, we see that in each case the cluster binary fraction remains close to the primordial value."
 Focussing on the IX100-5 simulation. Figure 4 shows (he [fractions of single stars and binaries (compared to their respective initial number) in the cluster.," Focussing on the K100-5 simulation, Figure \ref{f:fig4} shows the fractions of single stars and binaries (compared to their respective initial number) in the cluster."
 Following on from Figure 3aa the Iractions are similar at all times as expected., Following on from Figure \ref{f:fig3}a a the fractions are similar at all times as expected.
 llowever. Figure 4 also shows the fractions of single stars ancl binaries (hat have escaped the cluster and we see that from about 2 Gyr onwards the fractional escape rate of single stars is greater than that of the binaries.," However, Figure \ref{f:fig4} also shows the fractions of single stars and binaries that have escaped the cluster and we see that from about $2\,$ Gyr onwards the fractional escape rate of single stars is greater than that of the binaries."
 At the end of the simulation (20 Gyr) the difference is34%.," At the end of the simulation $20\,$ Gyr) the difference is."
.. This is offset somewhat by evolution processes (stellar and binary) that destroy binaries (see (he dotted line in Figure 4))., This is offset somewhat by evolution processes (stellar and binary) that destroy binaries (see the dotted line in Figure \ref{f:fig4}) ).
 These processes include binaries becoming unbound due {ο supernova mass-loss and/or kicks (only relevant for the first LOO Myr of evolution) aud mass (ransler-induced mergers in close binaries.," These processes include binaries becoming unbound due to supernova mass-loss and/or kicks (only relevant for the first $100\,$ Myr of evolution) and mass transfer-induced mergers in close binaries."
 The remaining clilference is balanced by the destruction of binaries in cdvuamical encounters and (his becomes more important as the cluster evolves., The remaining difference is balanced by the destruction of binaries in dynamical encounters and this becomes more important as the cluster evolves.
 We note that even though the cluster binary fraction is relatively static as the cluster evolves the characteristics of the binary population change markedly over time with hare binaries favoured at late Gimes., We note that even though the cluster binary fraction is relatively static as the cluster evolves the characteristics of the binary population change markedly over time with hard binaries favoured at late times.
 Evident [rom Figure 3. is an overall trend for the core binary fraction to increase with lime. irrespective of simulation (wpe.," Evident from Figure \ref{f:fig3} is an overall trend for the core binary fraction to increase with time, irrespective of simulation type."
 For the core binary population of the IX100-5 moclel we, For the core binary population of the K100-5 model we
As the X-ray count rate [cll during the transition to outburst. the hardness ratio might have been expected o soften to the level of the earlier outburst observation.,"As the X-ray count rate fell during the transition to outburst, the hardness ratio might have been expected to soften to the level of the earlier outburst observation."
 Llowever. this is not the case.," However, this is not the case."
 The hardness ratio appears o increase until micway through the transition. remaining iarder than the outburst observation throughout.," The hardness ratio appears to increase until midway through the transition, remaining harder than the outburst observation throughout."
 Fie., Fig.
 3 shows the X-ray light curve and hardness ratio of he transitional observation in more detail., \ref{fig:transcurve} shows the X-ray light curve and hardness ratio of the transitional observation in more detail.
 The light curve shows a decreasing count rate with irregular peaks., The light curve shows a decreasing count rate with irregular peaks.
 Apart rom the large peak in hardness ratio. there appears to be ittle correlation between the count rate and hardness ratio. with the hardness ratio remaining approximately constant hroughout.," Apart from the large peak in hardness ratio, there appears to be little correlation between the count rate and hardness ratio, with the hardness ratio remaining approximately constant throughout."
 We selected two pairs of spectra for detailed analysis., We selected two pairs of spectra for detailed analysis.
 “Phese are the SIS CIS spectra from throughout the outburst and from the entire transition observation., These are the SIS GIS spectra from throughout the outburst and from the entire transition observation.
 All the spectra have been fitted with several standard models as used. in the SSPIEC spectral fitting package. version 11.0 (CArnaud 1996).," All the spectra have been fitted with several standard models as used in the XSPEC spectral fitting package, version 11.0 (Arnaud 1996)."
 Phroughout the modelling. both SIS and GIS have been fitted simultaneously. with a energv-independent: multiplicative factor applied. between the SIS and CIS instruments as a [ree parameter.," Throughout the modelling, both SIS and GIS have been fitted simultaneously, with a energy-independent multiplicative factor applied between the SIS and GIS instruments as a free parameter."
 We started the spectral fitting of Z Cam using a single temperatureΚας model for emission [rom a hot diffuse gas (Alewe 1985. Alewe 1986. Liedahbl 1995) and to model photo-clectric absorption (Morrison AleCammon. 1983).," We started the spectral fitting of Z Cam using a single temperature model for emission from a hot diffuse gas (Mewe 1985, Mewe 1986, Liedahl 1995) and to model photo-electric absorption (Morrison McCammon, 1983)."
 This leads to acceptable fits for the transition spectrum ML=0.77. 603 degrees of freedom) but poor fits to the outburst spectrum (42=1.72. 225 d.oL).," This leads to acceptable fits for the transition spectrum $\chi^2_\nu=0.77$, 603 degrees of freedom) but poor fits to the outburst spectrum $\chi^2_\nu=1.72$, 225 d.o.f.)."
 The single temperature fits can be seen in Fig. 4..," The single temperature fits can be seen in Fig. \ref{fig:1tspectra},"
 showing excess emission. below kkeV. especially in the. outburst observation.," showing excess emission below keV, especially in the outburst observation."
 In an attempt to model the excess emission. below lkkeV a secondmekal component CK&0.6 kkeV) was added to. the. model., In an attempt to model the excess emission below keV a second component $kT\approx0.6$ keV) was added to the model.
 This component only contributes line emission around. kkeV. predominantly. iron L-shell emission. thus resulting in an improved fit.," This component only contributes line emission around keV, predominantly iron L-shell emission, thus resulting in an improved fit."
 Fhus. acceptable fits where obtained to both the outburst ancl transition spectra (see table 1)).," Thus, acceptable fits where obtained to both the outburst and transition spectra (see table \ref{table:2ts}) )."
 Not only does this two-temperature model fit. both spectra. but it is only necessary to allow one parameter to vary between the two states. namely the normalisation of the righer temperature component (47>=0.71. 829 d.oL..," Not only does this two-temperature model fit both spectra, but it is only necessary to allow one parameter to vary between the two states, namely the normalisation of the higher temperature component $\chi^2_\nu=0.71$, 829 d.o.f.,"
 fitting roth pairs of spectra simultaneously)., fitting both pairs of spectra simultaneously).
 Thus the spectra of Z Cam can be interpreted. as originating [rom two. single-emperature plasmas. with only the amount of hot gas varving between the two observations.," Thus the spectra of Z Cam can be interpreted as originating from two, single-temperature plasmas, with only the amount of hot gas varying between the two observations."
 The two-tempoerature model fits vield large values for he absorption during both the transition (Ny=11.2-1075 7) and the outburst (Ng=3.1ur ]107em 7)., The two-temperature model fits yield large values for the absorption during both the transition $N_H=11.2^{+1.7}_{-2.4}\times 10^{21}$ $^{-2}$ ) and the outburst $N_H=3.1^{+0.4}_{-0.3}\times 10^{21}$ $^{-2}$ ).
 The absorption throughout both the outburst and transition is two orders of magnitudes greater than the value of Ng=4o10% ? determined from an£02 curve-ol-growth study. of interstellar absorption lines (€. Mauche. private communication. for method see Mauche. Ravnone Cordova. 1988))," The absorption throughout both the outburst and transition is two orders of magnitudes greater than the value of $N_H=4\times 10^{19}$ $^{-2}$ determined from an curve-of-growth study of interstellar absorption lines (C. Mauche, private communication, for method see Mauche, Raymond Cordova 1988)."
 This excess over the interstellar value sugeests that a large amount of absorption is occurring close to the X-ray source. even in the earliest stages of outburst.," This excess over the interstellar value suggests that a large amount of absorption is occurring close to the X-ray source, even in the earliest stages of outburst."
 Since the excess absorption Is local to the svsteni. it is unrealistic to expect the absorbing medium to be neutral.," Since the excess absorption is local to the system, it is unrealistic to expect the absorbing medium to be neutral."
 Therefore. the neutral absorber was replaced with an ionised absorber in the single temperaturemehked model.," Therefore, the neutral absorber was replaced with an ionised absorber in the single temperature model."
 Although this again produces an acceptable fit to the transition spectra (42.=0.67. GOL doo)," Although this again produces an acceptable fit to the transition spectra $\chi^2_\nu=0.67$, 601 d.o.f.)"
 it cannot. reproduce the outburst spectra (X5=1.65. : d.o.L.).," it cannot reproduce the outburst spectra $\chi^2_\nu=1.65$, 223 d.o.f.)."
 In the previous section we have shown that the data can be modelled: with a optically-thin plasma at. two cistinct temperatures., In the previous section we have shown that the data can be modelled with a optically-thin plasma at two distinct temperatures.
 However. the X-rays are expected to originate [rom only one site. with a continuous range of teniperatures representing σας cooling onto the white «ναί," However, the X-rays are expected to originate from only one site, with a continuous range of temperatures representing gas cooling onto the white dwarf."
 In an attempt to fit our spectra with a continuous temperature distribution. we adopted the model of Done Osborne (1997). as applied to the quiescent spectrum of SS Cve (their.ρία(ρέig) model).," In an attempt to fit our spectra with a continuous temperature distribution, we adopted the model of Done Osborne (1997), as applied to the quiescent spectrum of SS Cyg (their model)."
" The model consists of neutral photo-clectric absorption to model interstellar absorption. (which we [freeze at Ny4107"" em>), ancl an ionised absorption component is required to model absorption local to the svstem."," The model consists of neutral photo-electric absorption to model interstellar absorption (which we freeze at $N_H=4\times 10^{19}$ $^{-2}$ ), and an ionised absorption component is required to model absorption local to the system."
" The emission. measure of the X-ray. emitting plasma is modelled as a power-law in temperature. (E/13,,)"" using theeceimélf model in NSPEC."," The emission measure of the X-ray emitting plasma is modelled as a power-law in temperature, $(T/T_{Max})^\alpha$ using the model in XSPEC."
 The observations of Z Cam cannot constrain the reflection continuum component due to the low number of counts at high. energies. and so this has been omitted. in our version of the model.," The observations of Z Cam cannot constrain the reflection continuum component due to the low number of counts at high energies, and so this has been omitted in our version of the model."
 However. a gaussian is include to model the kkeV fluorescent line.," However, a gaussian is included to model the keV fluorescent line."
 We shall refer to this model as the DO97 model hereafter., We shall refer to this model as the DO97 model hereafter.
 ‘The results of the modelling can be seen in Table 2. anc Pies.," The results of the modelling can be seen in Table \ref{table:done} and Figs. \ref{fig:spectra}, \ref{fig:donemodels},"
 5.6.7 9.., \ref{fig:donedem}~ \ref{fig:nhvsxi}.
 The metal abundances used throughou are those of Ancers Crevesse (1989). with the abundances frozen at solar values.," The metal abundances used throughout are those of Anders Grevesse (1989), with the abundances frozen at solar values."
 Phawing the abuncances leads to no improvement for the quiescent observation. anc margina improvement to the outburst observation (AYE= 3.0).," Thawing the abundances leads to no improvement for the quiescent observation, and marginal improvement to the outburst observation $\Delta\chi^2=3.0$ )."
 The maximum temperature of the power-law temperature enmissivitv component is poorly constrained. and so is fixec at kkeV. The DO97 model fits as well as the two-teniperature model. with v=0.64 (600 clo.)," The maximum temperature of the power-law temperature emissivity component is poorly constrained and so is fixed at keV. The DO97 model fits as well as the two-temperature model, with $\chi^2_\nu=0.64$ (600 d.o.f.)"
 during the transition and 4Z=0.68 (222 do.), during the transition and $\chi^2_\nu=0.68$ (222 d.o.f.)
 during outburst., during outburst.
 The residuals in the single temperature fits below LkkeV. (Fig. 4)), The residuals in the single temperature fits below keV (Fig. \ref{fig:1tspectra}) )
 are mocelled in the DO97 model bv a combination of increased line emission. from cool gas ancl increased: absorption., are modelled in the DO97 model by a combination of increased line emission from cool gas and increased absorption.
 We lind that both. the ionisation state ancl absorption column of the absorber remains approximately. constant in both the outburst and the transition (see Fig. 9))., We find that both the ionisation state and absorption column of the absorber remains approximately constant in both the outburst and the transition (see Fig. \ref{fig:nhvsxi}) ).
 The, The
law well at the epoch.,law well at the epoch.
 This means that although LIS'T-1 contains Complex substructures. its IEVPA still obevs the Faraday rotation law well.," This means that although HST-1 contains complex substructures, its EVPA still obeys the Faraday rotation law well."
 Figure 4. shows that RAL has the largest. value in the beginning. and gradually decreases with ποσαο present due to subllares occurred in the mean time (Llarrisetal. 2009).," Figure \ref{fig4} shows that RM has the largest value in the beginning, and gradually decreases with fluctuation present due to subflares occurred in the mean time \citep{har09}."
. This is consistent with the fact that fractional polarization is relatively low in the carly stage of the outburst. anc as expected. the RAL variability may well arise from the interior of LIST-1.," This is consistent with the fact that fractional polarization is relatively low in the early stage of the outburst, and as expected, the RM variability may well arise from the interior of HST-1."
 Faraday rotation may give rise to depolarization by a factor of sin(®)/® (Burn1966Lomanetal. 2009).. where 9=2X. if Faraday. rotation occurs internal to HS'T-1.," Faraday rotation may give rise to depolarization by a factor of $\rm{sin(\Phi)/\Phi}$ \citep{bur66, hom09}, where $\rm{\Phi=2\Delta\chi}$, if Faraday rotation occurs internal to HST-1."
 According to the resultant LAL values shown in Fig. 4.," According to the resultant RM values shown in Fig. \ref{fig4},"
 the dispersion ellect is supposed to be quite significant at S CGllz in comparison to that at the other two frequencies if Faraday rotation occurs internal to LIST-1. for example. the polarization level at S Gllz would be 4 times lower man that at 15 Cillz even when RAL took a relatively small value of 1000racm," the dispersion effect is supposed to be quite significant at 8 GHz in comparison to that at the other two frequencies if Faraday rotation occurs internal to HST-1, for example, the polarization level at 8 GHz would be 4 times lower than that at 15 GHz even when RM took a relatively small value of $\rm{1000~rad~m^{-2}}$."
 Aloreover. the relatively large beam size at low frequencies would make the fractional polarization even lower if it had an elfect on the polarization level Clavlor.1998).," Moreover, the relatively large beam size at low frequencies would make the fractional polarization even lower if it had an effect on the polarization level \citep{tay98}."
 However. one can find in the bottom panels of Fig.," However, one can find in the bottom panels of Fig."
 3. that after 2006. the polarization levels are almost comparable at the 3 frequencies.," \ref{fig3} that after 2006, the polarization levels are almost comparable at the 3 frequencies."
 This implies that. although the RAL variation is attributed to the plasma internal to HIST-1. quite a large part of Faraday rotation may occur external to H8TT-1.," This implies that, although the RM variation is attributed to the plasma internal to HST-1, quite a large part of Faraday rotation may occur external to HST-1."
 Especially in the decaving stage. the Faraday rotation may occur dominantly external to LIST-1. as is required by the comparable fractional polarization at the 3 wavebancds.," Especially in the decaying stage, the Faraday rotation may occur dominantly external to HST-1, as is required by the comparable fractional polarization at the 3 wavebands."
 This is consistent with the fact that RAL decrease gets slower with time. and more importantly. is supported by the presence of cdilfuse emission around. LIST-1 measured with VLA at low frequencies. as is shown in Fig.," This is consistent with the fact that RM decrease gets slower with time, and more importantly, is supported by the presence of diffuse emission around HST-1 measured with VLA at low frequencies, as is shown in Fig."
 1. of Harrisetal.(2008)., \ref{fig1} of \cite{har08}.
 We reduced 8 epoch polarimetric observations of AIST with VLA A array at S. 15. and 22 Cllz from 2003 to 2007.," We reduced 8 epoch polarimetric observations of M87 with VLA A array at 8, 15, and 22 GHz from 2003 to 2007."
 Due to the significant blazar-like nature shown in ΕΕ (Llarrisetal.2008).. the Lux density. spectral. polarization. and RAL variabilities are investigated in the period of a major outburst. during which a 5-rav flare is detected to be associated with it.," Due to the significant blazar-like nature shown in HST-1 \citep{har08}, the flux density, spectral, polarization, and RM variabilities are investigated in the period of a major outburst, during which a $\gamma$ -ray flare is detected to be associated with it."
 Contrary to the general case for blazars. the outburst of LIST-1 grows relatively carly at. low frequency. which is consistent with the spectral variability that the spectrum is softening in the rising stage.," Contrary to the general case for blazars, the outburst of HST-1 grows relatively early at low frequency, which is consistent with the spectral variability that the spectrum is softening in the rising stage."
 Εις implies that the newly emerging subcomponents have relatively steep spectra. and hence make the overall spectrum. softening.," This implies that the newly emerging subcomponents have relatively steep spectra, and hence make the overall spectrum softening."
 In. addition. these newly emerging subcomponents from. 2003 to 2004 are superluminal (Cheungetal.2007).. and hence expected be accelerated. in this region due to the fact that all components upstream of LST-1 are subluminal. ancl become faster with distance [rom the core.," In addition, these newly emerging subcomponents from 2003 to 2004 are superluminal \citep{che07}, and hence expected to be accelerated in this region due to the fact that all components upstream of HST-1 are subluminal, and become faster with distance from the core."
 The RALeorreeted EWPA of HST-1. shows a monotonous variation from ~90° to ~Q at the 3 wavebands over a period of more than 3 wears. and in this period the subcomponents within HIST-1 move down along the routes deviating from the overall jet. direction with the motion of one subcomponent bending from. one side of the large scale jet axis to another. and another two subcomponents deflecting significantly [rom the jet axis.," The RM-corrected EVPA of HST-1 shows a monotonous variation from $\sim-90^\circ$ to $\sim0^\circ$ at the 3 wavebands over a period of more than 3 years, and in this period the subcomponents within HST-1 move down along the routes deviating from the overall jet direction with the motion of one subcomponent bending from one side of the large scale jet axis to another, and another two subcomponents deflecting significantly from the jet axis."
 “These give quite strong evidence for the presence of a helical magnetic field around HST-1. which also. ects support from the fact that the emission features may get accelerated in this region.," These give quite strong evidence for the presence of a helical magnetic field around HST-1, which also gets support from the fact that the emission features may get accelerated in this region."
 “Phe fractional polarization varies with epochs. and is relatively low in the rising stage in comparison to that in the decaving stage.," The fractional polarization varies with epochs, and is relatively low in the rising stage in comparison to that in the decaying stage."
 Εις implies that the internal Faraday rotation anc medium. opacity play an important role to the variation of polarization level during the outburst., This implies that the internal Faraday rotation and medium opacity play an important role to the variation of polarization level during the outburst.
 Moreover. in the decaying stage alter 200(. the polarization levels are almost comparable at the 3 wavebands. indicating that Faraday rotation occurs dominantly external to HST-1 in this phase. or else it would be much lower at S Gllz than measured.," Moreover, in the decaying stage after 2006, the polarization levels are almost comparable at the 3 wavebands, indicating that Faraday rotation occurs dominantly external to HST-1 in this phase, or else it would be much lower at 8 GHz than measured."
" This work was supported in part hy EN erants 10nim 10633010 and. 10821302.InternationalIX.JCNX2-YW-TO Partiedip ronYCBSI54"" the CAS/SAPLA Creative Research Teams. and the Science and Technology Commission of Shanghai Municipality. (OOZR1437."," This work was supported in part by the grants 10625314, 10633010 and 10821302, KJCX2-YW-T03, 2007CB815405, the CAS/SAFEA International Partnership Program for Creative Research Teams, and the Science and Technology Commission of Shanghai Municipality (09ZR1437400)."
stars.,stars.
 LO would. be straightforward to obtain. spectra of higher S/N 300 ) for these stars., It would be straightforward to obtain spectra of higher S/N $\ge 300$ ) for these stars.
 Ht would also be helpful to observe additional stars without detected. planets with comparable S/N ratio., It would also be helpful to observe additional stars without detected planets with comparable S/N ratio.
 Our [findings bring to [ive the number of stellar parameters that correlate with the presence. of Doppler planets: metallicity. mass. Li abundance. vsini and Ruy. detected.," Our findings bring to five the number of stellar parameters that correlate with the presence of Doppler detected planets: metallicity, mass, Li abundance, vsini and $R^{'}_{\rm HK}$."
"Phere might also be additional secondary. parameters. such as the ""Li/'Li isotope ratio and AL/Fe and. Si/Fe abundances ratios. but. these require confirmation."," There might also be additional secondary parameters, such as the $^{6}$ $^{7}$ Li isotope ratio and Al/Fe and Si/Fe abundances ratios, but these require confirmation."
 Taken together. these five parameters should make it. possible to select stars with a high probability of hosting planets.," Taken together, these five parameters should make it possible to select stars with a high probability of hosting planets."
 We thank the anonymous reviewer for helpful comments and suggestIons., We thank the anonymous reviewer for helpful comments and suggestions.
"to scale as Llere IN, is the total number of raciating particles and dii rellects the increasing visible size (due to decrease of »aming) of the slab.",to scale as Here $N_{tot}$ is the total number of radiating particles and $\dev \mu$ reflects the increasing visible size (due to decrease of beaming) of the slab.
 “Phe two possible spectra that the code can generate are shown in fig. L..," The two possible spectra that the code can generate are shown in fig. \ref{twospectra},"
 where we used the labelling rom Granot&Sari(2002) to distinguish the cillerent power aw regimes., where we used the labelling from \citet{Granot2002} to distinguish the different power law regimes.
 In tables (1)) and (2)) we give the expressions or the absolute scalings in the cdillerent regimes 2. E. LF. C. H and the critical frequencies.," In tables \ref{scalings_table}) ) and \ref{frequencies_table}) ) we give the expressions for the absolute scalings in the different regimes $D$, $E$, $F$, $G$, $H$ and the critical frequencies."
 Scaling coefficients aside. hese equationsare similar to those given in VanderHorstetal. (2008).," Scaling coefficients aside, these equationsare similar to those given in \citet{vanderHorst2008}."
. The flux in regime D is denoted by £p. the critical peak frequency in spectrum | is denoted by i.1. he critical coolinge frequency in spectrum 1 by £524 ancl so OL.," The flux in regime $D$ is denoted by $F_D$, the critical peak frequency in spectrum 1 is denoted by $\nu_{m,1}$, the critical cooling frequency in spectrum 1 by $\nu_{c,1}$ and so on."
 The equations in the tables introduce a number of svinbols that need an explanation., The equations in the tables introduce a number of symbols that need an explanation.
" The cosmic recshilt is eiven by 2. while the luminosity distance ""ροκ Is measured in units of 1077 em."," The cosmic redshift is given by $z$ , while the luminosity distance $r_{obs,28}$ is measured in units of $10^{28}$ cm."
 £s» is the explosion encrey£ in units of 1077 ere., $E_{52}$ is the explosion energy$E$ in units of $10^{52}$ erg.
 Phe observer time in davs is denoted by furs.," The observer time in days is denoted by $t_{obs,d}$."
 Lhe characteristic distance yg we put at 1055 em ancl py and no are related via the proton mass: po=mynmo., The characteristic distance $R_0$ we put at $10^{17}$ cm and $\rho_0$ and $n_0$ are related via the proton mass: $\rho_0 = m_p n_0$.
 Ehe scaling coellicients Cp. Cr ete.," The scaling coefficients $C_D$, $C_E$ etc."
 contain a number of numerical constants (determined by fitting to output from our code) and some explicit dependencies on & and p and. are further explained in appencix C.., contain a number of numerical constants (determined by fitting to output from our code) and some explicit dependencies on $k$ and $p$ and are further explained in appendix \ref{scaling_derivation_section}.
 3efore the cooling break the scaling behaviour is dictated by the asymptotic behaviour of QUAfiyin)," Before the cooling break the scaling behaviour is dictated by the asymptotic behaviour of $Q(\nu'/\nu'_{cr,m})$."
 The steepening of the spectrum bevond the cooling breaks and the corresponding changes in the sealing behaviour are due to the fact that beyond. the cooling break [frequency the region behind the shock that still significantly. contributes to the total flux (i.e. the region)) becomes noticably smaller than the shock width., The steepening of the spectrum beyond the cooling breaks and the corresponding changes in the scaling behaviour are due to the fact that beyond the cooling break frequency the region behind the shock that still significantly contributes to the total flux (i.e. the ) becomes noticably smaller than the shock width.
 The changes in the scalings reflects the change in the size of region., The changes in the scalings reflects the change in the size of region.
 The hot region is discussed separately in appendix. D.., The hot region is discussed separately in appendix \ref{hot_region_section}.
 In simple power law model fits. the gradual. transition between regimes is often handled by a free parameter. the sharpness factor 5.," In simple power law model fits, the gradual transition between regimes is often handled by a free parameter, the sharpness factor $s$."
 In more detailed calculations like those done here the gradual transitions are included automatically and we can use this to provide the correct dependence of s on p and A., In more detailed calculations like those done here the gradual transitions are included automatically and we can use this to provide the correct dependence of $s$ on $p$ and $k$.
 Εις eliminates s as à free parameter. simplifying the fit to the data and allowing the shape of the transition to help determine whether a particular model fits the data or not.," This eliminates $s$ as a free parameter, simplifying the fit to the data and allowing the shape of the transition to help determine whether a particular model fits the data or not."
" For spectrum 1. we use the following equation to describe the Hux density near the peak break £44: where 4,4 denotes the Lux at the critical frequency My. Lor infinite sharpness 5,,4 (ic. the meeting point of the asymptotic power laws)."," For spectrum 1, we use the following equation to describe the flux density near the peak break $\nu_{m,1}$: where $F_{m,1}$ denotes the flux at the critical frequency $\nu_{m,1}$ for infinite sharpness $s_{m,1}$ (i.e. the meeting point of the asymptotic power laws)."
" When we switch off cooling in our simulation. we can determine 55,4, from fitting against the resulting spectrum while keeping the other parameters in equation (9)) fixed."," When we switch off cooling in our simulation, we can determine $s_{m,1}$ from fitting against the resulting spectrum while keeping the other parameters in equation \ref{Fsmooth_m1_equation}) ) fixed."
 The sharpness is a function mainly of p and toa lesser extent of & and the other simulation input parameters., The sharpness is a function mainly of $p$ and to a lesser extent of $k$ and the other simulation input parameters.
" Rather than attempting to include all secondary dependencies when formulating a description for 5,4. we find that the following approximation For 5,,,4 is always valid up to a few percent: When we switch on electron. cooling. the Hux is best approximated by Ifowe fit this function against simulation output using δι as a fitting parameter we find that the results are described (up to à few percent) by AX simultaneous fit using both νι and s, vields the same results."," Rather than attempting to include all secondary dependencies when formulating a description for $s_{m,1}$, we find that the following approximation for $s_{m,1}$ is always valid up to a few percent: When we switch on electron cooling, the flux is best approximated by If we fit this function against simulation output using $s_{c,1}$ as a fitting parameter we find that the results are described (up to a few percent) by A simultaneous fit using both $s_{m,1}$ and $s_{c,1}$ yields the same results."
" For spectrum 5 the order of the breaks is reversed and the smooth power law for both breaks is given hy where £5 denotes the peak fux for infinite sharpness 5,5 and the prescriptions for the sharpness are and Once again valid. up to a few percent."," For spectrum 5 the order of the breaks is reversed and the smooth power law for both breaks is given by where $F_{c,5}$ denotes the peak flux for infinite sharpness $s_{c,5}$ and the prescriptions for the sharpness are and Once again valid up to a few percent."
 Given their accuracies. all sharpness prescriptions are consistent. with CGranot&Sari (2002)..," Given their accuracies, all sharpness prescriptions are consistent with \cite{Granot2002}. ."
 Various authors have used. [lux scaling equations to derive the physical properties of GARB 070508 from afterglow data (Galamactal.L999:Granot&Sari2002:Yostet 2008)..," Various authors have used flux scaling equations to derive the physical properties of GRB 970508 from afterglow data \citep{Galama1999, Granot2002, Yost2003, vanderHorst2008}. ."
 This provides us with a context to illustrate the scaling laws derived. in section 4.., This provides us with a context to illustrate the scaling laws derived in section \ref{coefficients_section}. .
 We will use the fit parameters obtained from broadbandmodeling by VanderHorst.ctal. (2008)..., We will use the fit parameters obtained from broadbandmodeling by \citet{vanderHorst2008}. .
 They. have fit simultaneously in time and frequency. while keeping & as a fitting parameter., They have fit simultaneously in time and frequency while keeping $k$ as a fitting parameter.
 Because the only mocdol dependencies that have been introduced by this approachare thescalings of / and v (and no scaling coefficients). their fit results are still," Because the only model dependencies that have been introduced by this approachare thescalings of $t$ and $\nu$ (and no scaling coefficients), their fit results are still"
CR region. leading to its geometrical erowth.,"CR region, leading to its geometrical growth."
 Iu both dynamical aud secular phases of the bar evolution we fi that the final bar amplitude.;ioj.. correlates. with final bar (Fig.," In both the dynamical and secular phases of the bar evolution we find that the final bar amplitude, correlates with the final bar (Fig."
 12)., 12).
 So stronger bars appear to | longer as well., So stronger bars appear to be longer as well.
 Taken at face value. this property of b evolution appears to be supported by recent observatio in that the A-band tages of barred galaxies exhibit : correlation between the bar amplitude. measured by m=2 Fourier compoucut. and its size (Ehucercen et :," Taken at face value, this property of bar evolution appears to be supported by recent observations in that the $K$ -band images of barred galaxies exhibit a correlation between the bar amplitude, measured by the $m=2$ Fourier component, and its size (Elmegreen et al."
 2007)., 2007).
 Tlowever. there are caveats.," However, there are caveats."
 We find that the bar sizes are substantially siall in the eas-rich models compared to the eas-poor o1ο (Figs., We find that the bar sizes are substantially smaller in the gas-rich models compared to the gas-poor ones (Figs.
 3. 13).," 3, 13)."
 While the median bar size at the eud of the simulation in the latter models appears to be ~1.3. it is around 0.5 in the feminer ones. 1.60. 12 kpc vs 5 kpc. enmibedded iu the ideutical disks and hialos.," While the median bar size at the end of the simulation in the latter models appears to be $\sim 1.2$, it is around 0.5 in the former ones, i.e., 12 kpc vs 5 kpc, embedded in the identical disks and halos."
 This differece arises alinost cutirely during the secular plase of the uar evolution., This difference arises almost entirely during the secular phase of the bar evolution.
 More precisely. by the eud of the dynamcal stage. one can notice that the bar size sliehtlv aiticorrelates with for the eas-rich models aud ouly for =U. N6andOOSscequences," More precisely, by the end of the dynamical stage, one can notice that the bar size slightly anti-correlates with for the gas-rich models and only for $=0.016$ and 0.05 sequences."
 Phistrendissufficient lgweakandbar 15 oftheirlength., This trend is sufficiently weak and bars differ by not more than of their length.
Socssentially.thebargrowthinthedgnaimical ," So essentially, the bar growth in the dynamical phase is independent of."
What is rather striking. is the axupt change iu the bars growth ‘habit’ in the secular plae bars erow in the gas-poor models aud stagnate iu the eas-rich ones.," What is rather striking, is the abrupt change in the bar's growth `habit' in the secular phase — bars grow in the gas-poor models and stagnate in the gas-rich ones."
Searching for faint companions around bright nearby. stars is one of the most challenging goals of today’s astronomy (??)..,"Searching for faint companions around bright nearby stars is one of the most challenging goals of today's astronomy \citep{Oppenheimer09,Absil10b}."
 This quest for high-dynamie range observations has been spurred for more than a decade by the discovery of more than 500 extrasolar planets in our galactic neighbourhood., This quest for high-dynamic range observations has been spurred for more than a decade by the discovery of more than 500 extrasolar planets in our galactic neighbourhood.
 A proper characterisation of the physical and atmospheric parameters of the detected planets requires the photon emitted (or reflected) by the planets to be disentangled from those emitted by their host star., A proper characterisation of the physical and atmospheric parameters of the detected planets requires the photon emitted (or reflected) by the planets to be disentangled from those emitted by their host star.
 While a temporal separation of the photons ts possible in the particular case of transiting systems. most systems require high-dynamie range imaging with a high angular resolution in order to be characterised.," While a temporal separation of the photons is possible in the particular case of transiting systems, most systems require high-dynamic range imaging with a high angular resolution in order to be characterised."
 Adaptive optics on 10-m class telescopes. possibly aided by the use of a coronagraph and/or differential imaging techniques. is currently the most widely used technique for directly imaging extrasolar planets (e.g...???2)..," Adaptive optics on 10-m class telescopes, possibly aided by the use of a coronagraph and/or differential imaging techniques, is currently the most widely used technique for directly imaging extrasolar planets \citep[e.g.,][]{Chauvin04,Kalas08,Marois08,Lagrange09a}."
" This technique provides very high-contrast images at relatively large angular distances from the central star. up to about 1:10* beyond 1"". and can detect companions with flux ratios around 1:1000 down to about mmas separations (e.g. ?).."," This technique provides very high-contrast images at relatively large angular distances from the central star, up to about $10^6$ beyond $1\arcsec$, and can detect companions with flux ratios around 1:1000 down to about mas separations \citep[e.g.][]{Boccaletti09}."
 The upcoming generation of high-contrast imagers using extreme adaptive optics on 10-m class telescope promises to further improve these performances., The upcoming generation of high-contrast imagers using extreme adaptive optics on 10-m class telescope promises to further improve these performances.
 Combined with state-of-the-art coronagraphic devices such as vortex coronagraphs. the detection of high-contrast companions at the diffraction limit CUD. 1.ο.. about mmas at K band on a 10-m telescope) will be enabled (e.g..?)..," Combined with state-of-the-art coronagraphic devices such as vortex coronagraphs, the detection of high-contrast companions at the diffraction limit $\lambda/D$, i.e., about mas at K band on a 10-m telescope) will be enabled \citep[e.g.,][]{Mawet11}."
 Pushing the discovery space within the diffraction limit of a single aperture requires interferometric techniques., Pushing the discovery space within the diffraction limit of a single aperture requires interferometric techniques.
 On a single pupil. aperture-masking techniques have recently allowed dynamic ranges up to 1:1000 to be reached down to the diffraction limit. and permitted decent contrasts (~ 11:100) to be reached down to 2/3D (??)..," On a single pupil, aperture-masking techniques have recently allowed dynamic ranges up to 1:1000 to be reached down to the diffraction limit, and permitted decent contrasts $\sim$ 1:100) to be reached down to $\lambda/3D$ \citep{Hinkley11,Lacour11}."
 Another example is the Palomar Fiber Nuller. which also works on a single telescope. and has reached a 1:500 dynamic range around the bright star Vega in a search region extending significantly within the diffraction limit of the Palomar Hale 5-m telescope (?)..," Another example is the Palomar Fiber Nuller, which also works on a single telescope, and has reached a 1:500 dynamic range around the bright star Vega in a search region extending significantly within the diffraction limit of the Palomar Hale 5-m telescope \citep{Mennesson11}."
 To extend the search region even closer to the central star. interferometry with multiple apertures separated by more than 10mm is mandatory.," To extend the search region even closer to the central star, interferometry with multiple apertures separated by more than m is mandatory."
 Interferometry on multiple apertures has. however. not yet reached the same dynamic range as adaptive optics or even aperture-masking observations. at least when the target star is mostly unresolved (?)..," Interferometry on multiple apertures has, however, not yet reached the same dynamic range as adaptive optics or even aperture-masking observations, at least when the target star is mostly unresolved \citep{Absil10a}."
 The dynamic range can be significantly increased by observing resolved targets (e.g..22).. but such targets are scarce in the context of main sequence stars observed on 100-m class baselines.," The dynamic range can be significantly increased by observing resolved targets \citep[e.g.,][]{Duvert10,Zhao11}, but such targets are scarce in the context of main sequence stars observed on 100-m class baselines."
 Here. we present the first results of a new 4-telescope interferometric instrument called PIONIER (Precision Integrated-Optics Near-infrared Imaging ExpeRiment). which has recently been installed and commissioned at the Very Large Telescope Interferometer (2)..," Here, we present the first results of a new 4-telescope interferometric instrument called PIONIER (Precision Integrated-Optics Near-infrared Imaging ExpeRiment), which has recently been installed and commissioned at the Very Large Telescope Interferometer \citep{LeBouquin11}."
 PIONIER aims at providing enhanced imaging capabilities at the VLTI. together with an improved accuracy and stability of interferometric observables (visibility and closure phase).," PIONIER aims at providing enhanced imaging capabilities at the VLTI, together with an improved accuracy and stability of interferometric observables (visibility and closure phase)."
 One of its main scientific goals is to directly detect. faint companions around bright. stars (up to a magnitude H~ 7)., One of its main scientific goals is to directly detect faint companions around bright stars (up to a magnitude $H\sim7$ ).
 The purpose of this paper is to discuss the dynamic range of this new instrument based on the observations obtained during the commissioning phase and the first scientific observing runs., The purpose of this paper is to discuss the dynamic range of this new instrument based on the observations obtained during the commissioning phase and the first scientific observing runs.
 The observations presented m this paper were obtained during the commissioning phase and the first scientific observing runs of PIONIER at VLTL which took place in Fall 2010.," The observations presented in this paper were obtained during the commissioning phase and the first scientific observing runs of PIONIER at VLTI, which took place in Fall 2010."
 A large number of stars has been observed during several, A large number of stars has been observed during several
survey. Iaullmann et al. (,"survey, Kauffmann et al. ("
"2007) recently found. that in 1e local universe. although the AGN activity is strongly ""orrelated with the age of the stars in the bulges (see also in Wang Wer 2008: Kewley et al.","2007) recently found that in the local universe, although the AGN activity is strongly correlated with the age of the stars in the bulges (see also in Wang Wei 2008; Kewley et al."
 2006: Wild et al., 2006; Wild et al.
 2007). 10 most active AGNs are always associated with the bluest puter disks.," 2007), the most active AGNs are always associated with the bluest outer disks."
 However. not all the galaxies with blue outer lisks have an active AGN.," However, not all the galaxies with blue outer disks have an active AGN."
 This result therefore motivates 10 authors to believe that it could. be understood. if the amount of gas transported inware from disk is a variable., This result therefore motivates the authors to believe that it could be understood if the amount of gas transported inward from disk is a variable.
 In this paper. we investigate the co-evolution issue » comparing the radial distribution of the core-collapse supernovae (ecSNo) detected in AGN host galaxies with the similar distribution of the ec-SNe detected. in. starforming galaxies.," In this paper, we investigate the co-evolution issue by comparing the radial distribution of the core-collapse supernovae (cc-SNe) detected in AGN host galaxies with the similar distribution of the cc-SNe detected in starforming galaxies."
 Because cc-SNe are generally accepted: {ο be produced by the explosion of massive stars (8S. 10M.) at the end of their lifetime X10 ve (e.g. Woosley et al.," Because cc-SNe are generally accepted to be produced by the explosion of massive stars $\geq8-10M_\odot$ ) at the end of their lifetime $\la 10^{7.5}$ yr (e.g., Woosley et al."
 2002). the racial distribution of ec-SNe reasonably. represents the distribution of voung stellar populations.," 2002), the radial distribution of cc-SNe reasonably represents the distribution of young stellar populations."
 The advantage of this approach is that the result does not strongly depend on the spatial resolution of the observations., The advantage of this approach is that the result does not strongly depend on the spatial resolution of the observations.
 At first. the spectroscopic data from the SDSS Data Release 6 are cross-matched. with the SAI-SDSS image SNelog'.," At first, the spectroscopic data from the SDSS Data Release 6 are cross-matched with the SAI-SDSS image SNe."
. In our cross-matching. we require that a) the angular olfset of an individual SN measured. from the center of the host galaxy ds less than1: b) the dillerence in redshift (As) between the SN and corresponding host galaxy. is less than 0.01.," In our cross-matching, we require that a) the angular offset of an individual SN measured from the center of the host galaxy is less than; b) the difference in redshift $\Delta z$ ) between the SN and corresponding host galaxy is less than 0.01."
 The eross-matched sample in total contains 620 events. covering all three main supernova types Lai. Ib/c and Η.," The cross-matched sample in total contains 620 events, covering all three main supernova types Ia, Ib/c and II."
 Among the sample. more than of the events show As<0.003. which means that only for a few outliers. the recessional velocity dillerence. between the SNe and. their host galaxies is inferred to be larger than 900kms.+.," Among the sample, more than of the events show $\Delta z<0.003$, which means that only for a few outliers, the recessional velocity difference between the SNe and their host galaxies is inferred to be larger than $\mathrm{km\ s^{-1}}$."
 We further limit the sample to the SNe whose host galaxies have measured photometric diameters and ratios of their minor and major axes., We further limit the sample to the SNe whose host galaxies have measured photometric diameters and ratios of their minor and major axes.
 To ensure adequate sample size and minimize the bias introduced by morphology types and host galaxy luminosity. we finally restrict our sample to the SNe whose a) host galaxies show late morphology types ranging from S0/a to Sed (i.c. the T. parameter is between 0 and 6.5). b) redshifts are less than 0.045.," To ensure adequate sample size and minimize the bias introduced by morphology types and host galaxy luminosity, we finally restrict our sample to the SNe whose a) host galaxies show late morphology types ranging from S0/a to Scd (i.e., the T parameter is between 0 and 6.5), b) redshifts are less than 0.045."
 Given that the LLL are much more common than the SNellb/c. only the ILL are included in the analysis presented here.," Given that the II are much more common than the Ib/c, only the II are included in the analysis presented here."
 The spectroscopic data taken by the SDSS are then analyzed to diagnose the central power source for these galaxies., The spectroscopic data taken by the SDSS are then analyzed to diagnose the central power source for these galaxies.
 For each narrow cmiission-line galaxy. the underlvsing stellar absorption [features are first removed from the observed. spectrum by the principal. component analysis method. (see. Wang Wei 2008 for details).," For each narrow emission-line galaxy, the underlying stellar absorption features are first removed from the observed spectrum by the principal component analysis method (see Wang Wei 2008 for details)."
 The starlight-subtractecl spectrum is then used. to. measure emission line fluxes by the splof task in the package., The starlight-subtracted spectrum is then used to measure emission line fluxes by the splot task in the package.
 These emission-line galaxies are separated into ACNs and starforming galaxies according to the widely used BVT diagram (i.c. the OLLI]/LL2 vs. Να diagnostic diagram. Baldwin et al.," These emission-line galaxies are separated into AGNs and starforming galaxies according to the widely used BPT diagram (i.e., the $\beta$ vs. $\alpha$ diagnostic diagram, Baldwin et al."
 1981: Veilleux Osterbrock 1987)., 1981; Veilleux Osterbrock 1987).
 The empirical demarcation line proposed by Ixaulfmann et al. (, The empirical demarcation line proposed by Kauffmann et al. (
2003) is adopted in the classification.,2003) is adopted in the classification.
 Finally. there are in total 46 AGNs (hereafter AGN sample) and 94 starforming galaxies (hereafter SE sample) passing the above criteria.," Finally, there are in total 46 AGNs (hereafter AGN sample) and 94 starforming galaxies (hereafter SF sample) passing the above criteria."
 Note that the LL host ACINs are dominated. by. type IL AGNs that are on average a factor of LOO fainter in bolometric Luminosity than typical type lL ACINS (see recent review in Ilo 2008)., Note that the II host AGNs are dominated by type II AGNs that are on average a factor of 100 fainter in bolometric luminosity than typical type I AGNs (see recent review in Ho 2008).
 The classifications of these galaxies are illustrated in Figure 1., The classifications of these galaxies are illustrated in Figure 1.
 AGNs and starforming galaxies are svmbolized by red-open squares ancl blue-solid.: squares. respectively.," AGNs and starforming galaxies are symbolized by red-open squares and blue-solid squares, respectively."
" Following previous studies (e.g. Petrosian ""Turatto 1990: Bartunoy ct al."," Following previous studies (e.g., Petrosian Turatto 1990; Bartunov et al."
 1992: van den Bereh 1997: Petrosian et al., 1992; van den Bergh 1997; Petrosian et al.
 2005). the supernova relative distance (Reapffeoscor) measurecl [rom the center of the host galaxy is calculated for each SN by using the same method. used. in Tsvetkov et al. (," 2005), the supernova relative distance $R_{\mathrm{SN}}/R_{25,\mathrm{cor}}$ ) measured from the center of the host galaxy is calculated for each SN by using the same method used in Tsvetkov et al. ("
2004). where the projected. distance of the SN from the center of its host galaxy is {ων=(Aa|(0N8)2.,"2004), where the projected distance of the SN from the center of its host galaxy is $R_{\mathrm{SN}}=\sqrt{(\Delta\alpha)^2+(\Delta\delta)^2}$."
 Δι the direction along the position angle of the SN. Roscor. 16 projected radius of the galaxy up to surface density of 25magaresecond.7.E corrected. for the inclination of the ealaxy. is calculated as where dos is the measured galaxy photometric diameter up to surface density of 25magareseconcd7. ϐ the position angle of an individual SN. and b the ratio between major and minor axes.," At the direction along the position angle of the SN, $R_{25,\mathrm{cor}}$, the projected radius of the galaxy up to surface density of $25 \mathrm{mag\ arcsecond^{-2}}$ corrected for the inclination of the galaxy, is calculated as where $d_{25}$ is the measured galaxy photometric diameter up to surface density of $25 \mathrm{mag\ arcsecond^{-2}}$, $\theta$ the position angle of an individual SN, and $b$ the ratio between major and minor axes."
 Even though the two sub-samples are selected by taking the cuts in morphological type and redshift. the two sub-samples do not show identical distributions of their morphological νρος and redshifts.," Even though the two sub-samples are selected by taking the cuts in morphological type and redshift, the two sub-samples do not show identical distributions of their morphological types and redshifts."
 To alleviate the possible biases. we assign a weight for cach galaxy in the SE sample.," To alleviate the possible biases, we assign a weight for each galaxy in the SF sample."
 At first. he distributions of the morphological types are compared otween the two sub-samples.," At first, the distributions of the morphological types are compared between the two sub-samples."
 A weight is assigned to each galaxy in the SE sample by requiring the two distributions are identical., A weight is assigned to each galaxy in the SF sample by requiring the two distributions are identical.
 A second weight could. be obtained through he similar procedure but. for redshifts., A second weight could be obtained through the similar procedure but for redshifts.
 In cach procedure. we adopt a bin size that ensures each bin contains at least," In each procedure, we adopt a bin size that ensures each bin contains at least"
The equations presented below are implemented within the OpenMDP. parallel N--bodyv/SPLE evolution code VINE.,The equations presented below are implemented within the OpenMP parallel -body/SPH evolution code $\textsc{Vine}$.
 For all details we refer the reader to 27. and ?.., For all details we refer the reader to \citet{VINEI} and \citet{VINEII}.
" Within the SPL formulation a hyerodynamic quantity zd is interpolated by a kernel function M(r.rh) with [Mdv-- landlim,ιο=διαY) where the so called smoothing length / defines the spatial extent of the function M."," Within the SPH formulation a hydrodynamic quantity $A$ is interpolated by a kernel function $W(\textbf{r}-\textbf{r}',h)$ with $\int Wd\textbf{r}=1$ and $\lim_{h\rightarrow0}W=\delta(\textbf{r}-\textbf{r}')$, where the so called smoothing length $h$ defines the spatial extent of the function $W$."
 Γι» interpolation is then ciscretised. so that where ;/ (j) is the index of the particle at position r; (rj) and oh; GCV;) the value of the quantity zd at the position of particle £ (6j).," This interpolation is then discretised, so that where $i$ $j$ ) is the index of the particle at position $\textbf{r}_i$ $\textbf{r}_j$ ) and $A_i$ $A_j$ ) the value of the quantity $A$ at the position of particle $i$ $j$ )."
 p; and m; denote the density and mass at position of particle j. respectively.," $\rho_j$ and $m_j$ denote the density and mass at position of particle $j$, respectively."
 The VINE code uses the common W4 kernel defined by ? as where values with index 7j denote dilferences (e.g. ry= r;j)and arithmetic means (eg. fi;=0.5(hi| hj)). respectively. @=rorh. p is the number of spatial dimensions of the svstem and @ is a constant of order unity.," The $\textsc{Vine}$ code uses the common W4 kernel defined by \citet{Monaghan&Lattanzio1985} as where values with index $ij$ denote differences (e.g. $\textbf{r}_{ij}=\textbf{r}_i-\textbf{r}_j$ ) and arithmetic means (e.g. $\bar{h}_{ij}=0.5\cdot(h_i+h_j)$ ), respectively, $\varrho=|\textbf{r}-\textbf{r}'|/h$, $\nu$ is the number of spatial dimensions of the system and $\sigma$ is a constant of order unity."
" See ον, ? or ?. for more details."," See \citet{Monaghan1992}, \citet{Monaghan2001basics} or \citet{Price2005SPH} for more details."
 As long as the kernel itself is differentiable. every. function .À can be interpolatecl to a dillerentiable function. hy the procedure described above.," As long as the kernel itself is differentiable, every function $A$ can be interpolated to a differentiable function by the procedure described above."
" The most common formulation of derivatives in SPILL is (see e.g. ὃν, 2)) Using the continuity equation. the total time derivative of the density thus reads A natural ansatz to derive a conservative form of the momentum equation comprising the force due to pressure &eracients (in addition to the force due to the gravitational potential) is to use the Lagrange formalism together with the first law of thermodynamics."," The most common formulation of derivatives in SPH is (see e.g. \citealp{Monaghan1992}, \citealp{Price2005SPH}) ) Using the continuity equation, the total time derivative of the density thus reads A natural ansatz to derive a conservative form of the momentum equation comprising the force due to pressure gradients (in addition to the force due to the gravitational potential) is to use the Lagrange formalism together with the first law of thermodynamics."
 This leads to the following SPL variant of its ideal form (uy=ME E In this formulation momentum is conserved. exactly. since the contribution of particle j to the momentum of particle £ is equal and negative to the contribution of particle ido the momentum of particle j.," This leads to the following SPH variant of its ideal form $\frac{d\textbf{v}}{dt}=-\frac{\nabla P}{\rho}$ ): In this formulation momentum is conserved exactly, since the contribution of particle $j$ to the momentum of particle $i$ is equal and negative to the contribution of particle $i$ to the momentum of particle $j$."
 The change in the thermodynamical state of the eas requires an evolution equation for a state variable corresponding to the internal energy or entropy of the gas., The change in the thermodynamical state of the gas requires an evolution equation for a state variable corresponding to the internal energy or entropy of the gas.
 VINE employs an equation for the specific internal energy (un) of the gas., $\textsc{Vine}$ employs an equation for the specific internal energy $u$ ) of the gas.
 Without external heating or cooling terms. only compressional heating ancl cooling are important. and the SPLE variant of the ideal form (m=avτν) reacls: To close the set of equations. an isothermal equation of state is used throughout this paper.," Without external heating or cooling terms, only compressional heating and cooling are important and the SPH variant of the ideal form $\frac{du}{dt}=-\frac{P}{\rho}\nabla\cdot\textbf{v}$ ) reads: To close the set of equations, an isothermal equation of state is used throughout this paper."
 Artificial viscosity is required to model shocks and angular momenttun transport properly. where the latter is important to be able to simulate spiral arm formation.," Artificial viscosity is required to model shocks and angular momentum transport properly, where the latter is important to be able to simulate spiral arm formation."
 The VINE code uses the most common form of the artificial viscosity., The $\textsc{Vine}$ code uses the most common form of the artificial viscosity.
 Lt is described by the tensor IL;; as in ?.., It is described by the tensor $\mathbf{\Pi}_{ij}$ as in \cite{Monaghan1992}.
 Since the value of IL; depends on the dillerence. in velocity between the considered. particles (i.e. the velocity eradient) the viscosity increases with increasing velocity eracdient., Since the value of $\mathbf{\Pi}_{ij}$ depends on the difference in velocity between the considered particles (i.e. the velocity gradient) the viscosity increases with increasing velocity gradient.
 Moreover. the viscosity is only applied if particles are approaching each other.," Moreover, the viscosity is only applied if particles are approaching each other."
 7 suggested a viscosity limiter to avoid spurious angular momentunm and vorticity transport in gas disks., \cite{Balsara1995} suggested a viscosity limiter to avoid spurious angular momentum and vorticity transport in gas disks.
 However. a lower viscosity leads to a higher velocity dispersion of the eas and therefore to higher divergence of the velocity and magnetic Ποια.," However, a lower viscosity leads to a higher velocity dispersion of the gas and therefore to higher divergence of the velocity and magnetic field."
 As will be shown and discussed in section 5.. this higher divergence causes a more violent magnetic field amplification.," As will be shown and discussed in section \ref{DISCUSSION}, this higher divergence causes a more violent magnetic field amplification."
 The viscous terms within the momentum and energy equations read: This treatment of viscous forces allows for a sensible description of the behaviour of eas in a spiral galaxy., The viscous terms within the momentum and energy equations read: This treatment of viscous forces allows for a sensible description of the behaviour of gas in a spiral galaxy.
 Llowever. eqs.," However, eqs."
 14 and 16 are not applied when using isothermal equation of state., \ref{internalEnergy} and \ref{u_diss} are not applied when using isothermal equation of state.
 In order to follow the evolution of the magnetic field we have aclelitionally implemented equation 3. discretised as with gov denoting the spatial directions.," In order to follow the evolution of the magnetic field we have additionally implemented equation \ref{induktionsglg} discretised as with $\mu,\nu$ denoting the spatial directions."
The metal abundance of the stellar component of cluster galaxies can only be inferred from integrated spectra coupled to svnthetic stellar populations.,The metal abundance of the stellar component of cluster galaxies can only be inferred from integrated spectra coupled to synthetic stellar populations.
 Much of the stellar mass in clusters is confined to passively evolving spheroids (ellipticals and bulges) for which the iron abundance may range from ~1/3 solar to a few times solar., Much of the stellar mass in clusters is confined to passively evolving spheroids (ellipticals and bulges) for which the iron abundance may range from $\sim 1/3$ solar to a few times solar.
" Following ROT. the ratio of the iron mass in the ICM to the iron mass loked into stars and galaxies is given by: or 4.6. 2.5. 1.65. respectively lor h—(0.5. 0.75. and 1 (h=I£./100). and having adopted andZl/4,=0.3 solar and. Zf=1 solae."," Following R97, the ratio of the iron mass in the ICM to the iron mass loked into stars and galaxies is given by: or 4.6, 2.5, and 1.65, respectively for $h=0.5$, 0.75, and 1 $h=\ho/100$ ), and having adopted $\zfecm = 0.3$ solar and $\zfes =1$ solar."
 Note that with the adopted values for the quantities in equation (2) most of the cluster iron resides in the IC'M. rather than being locked into stars. especially for low values of ff...," Note that with the adopted values for the quantities in equation (2) most of the cluster iron resides in the ICM, rather than being locked into stars, especially for low values of $\ho$."
 Fherefore. it appears that there is at least as much metal mass out of cluster galaxies (in the ICMD. as there is inside them (locked into stars).," Therefore, it appears that there is at least as much metal mass out of cluster galaxies (in the ICM), as there is inside them (locked into stars)."
 Vhis sets a strong constraint by models of the chemical evolution of galaxies., This sets a strong constraint by models of the chemical evolution of galaxies.
 The constant oof clusters savs that the total mass of iron in the ICM, The constant of clusters says that the total mass of iron in the ICM
represented prominence plasma enhancements by massive line currents ancl current sheets. Low Zhang (2004) eave analytical solutions describing plasma enhancements of finite width by addressing a Iree-boundary problem.,"represented prominence plasma enhancements by massive line currents and current sheets, Low Zhang (2004) gave analytical solutions describing plasma enhancements of finite width by addressing a free-boundary problem."
 By (this approach (μον were able to describe both normal and inverse magnetic topologies., By this approach they were able to describe both normal and inverse magnetic topologies.
 This was made possible by the assumption of a circular boundary between the prominence plasma and the ambient corona and by making assumptions simplilving the polvtropie equation of state., This was made possible by the assumption of a circular boundary between the prominence plasma and the ambient corona and by making assumptions simplifying the polytropic equation of state.
 In (his paper we present classes of prominence-like solutions calculated numerically using the FINite Element Solver lor Stationary Equilibria (FINESSE. Belien et al.," In this paper we present classes of prominence-like solutions calculated numerically using the FINite Element Solver for Stationary Equilibria (FINESSE, Belien et al."
 2002) which allows the computation of stationary. eravilationally stratified. axisvimnietric or cartesian configurations.," 2002) which allows the computation of stationary, gravitationally stratified axisymmetric or cartesian configurations."
 The two codes FINESSE and PHOENIX (Blokland οἱ al., The two codes FINESSE and PHOENIX (Blokland et al.
 2007a) have already been used in tandem to determine the stability of tokamaks ancl accretion disks will steady flows (Goedbloed et al., 2007a) have already been used in tandem to determine the stability of tokamaks and accretion disks with steady flows (Goedbloed et al.
 2004a. b. Blokland et al.," 2004a, b, Blokland et al."
 2007b) but have not been adapted for the study of solar phenomena although the PARIS code (Belien et al., 2007b) but have not been adapted for the study of solar phenomena although the PARIS code (Belien et al.
 1997) which treats the eravitational stratification of axisvnunelric static loops is incorporated in the functionality ol FINESSE., 1997) which treats the gravitational stratification of axisymmetric static loops is incorporated in the functionality of FINESSE.
 Here. we use a suitably improved version of the FINESSE code to extend," Here, we use a suitably improved version of the FINESSE code to extend"
of S55Gs (in Paper I). defined in 3D redshilt space.,"of SSSGs (in Paper I), defined in 3D redshift space."
 The present paper collects optical photonmetric and spectrophotometric data lor a sample of candidate members of the 559G dominated by the elliptical NGC 1756., The present paper collects optical photometric and spectro–photometric data for a sample of candidate members of the SSSG dominated by the elliptical NGC 4756.
 The goals of the present paper are the following: 1) to obtain a redshift estimate which will provide an independent check of the svstemic velocities of the dominant NGC 4756 eroup members and to detect possible faint members 2) {ο investigate the structure of menber galaxies through a detailed surface photometry and 3) {ο analyze any possible induced activitv using medium resolution spectroscopy and diagnostic models., The goals of the present paper are the following: 1) to obtain a redshift estimate which will provide an independent check of the systemic velocities of the dominant NGC 4756 group members and to detect possible faint members 2) to investigate the structure of member galaxies through a detailed surface photometry and 3) to analyze any possible induced activity using medium resolution spectroscopy and diagnostic models.
 The study aims also to contribute to the detection and photometric characterization of the faint ealaxie population within the group area whose membership will be determined by ongoing redshift survevs., The study aims also to contribute to the detection and photometric characterization of the faint galaxie population within the group area whose membership will be determined by ongoing redshift surveys.
 The paper is organized as follows: the literature about the NGC 4756 group is presented in 2., The paper is organized as follows: the literature about the NGC 4756 group is presented in 2.
 A description of the photometric ancl spectroscopic observations. data reduction and analvsis is given in 3.," A description of the photometric and spectroscopic observations, data reduction and analysis is given in 3."
 Results about photometric and spectroscopic properties of individual bright objects. scaling relation of the candidate earlv-tvpe. from brieht to dwart. ealaxv population as well as the characteristics of (he Xrav emission in the area are presented in 4.," Results about photometric and spectroscopic properties of individual bright objects, scaling relation of the candidate early-type, from bright to dwarf, galaxy population as well as the characteristics of the X–ray emission in the area are presented in 4."
 Belore redshift survevs. NGC. 4756 ancl ils companions have been associated (o Abell 1631. a Dautz.Morgane type I cluster. within which the morphologicalex work of Dressler includes 139 ealaxies.," Before redshift surveys, NGC 4756 and its companions have been associated to Abell 1631, a Bautz–Morgan type I cluster, within which the morphological work of \citet{Dress80} includes 139 galaxies."
 Later. Dressler&Shectman(1988). compiled a catalog of radial velocities aud detected that a loreeround group. well centred on the cluster. is responsible for many of the apparently bright cluster members.," Later, \citet{Dress88} compiled a catalog of radial velocities and detected that a foreground group, well centred on the cluster, is responsible for many of the apparently bright cluster members."
 They observed that the, They observed that the
Numerous studies have shown a strong correlation. between supermassive black hole (SMBH) mass and host galaxy mass over broad mass ranges. implying that there is a close evolutionary relationship between them (Gebhardtetal.2000:Merritt&Hiring&Rix2004:Shieldsetal. 2006).,"Numerous studies have shown a strong correlation between supermassive black hole (SMBH) mass and host galaxy mass over broad mass ranges, implying that there is a close evolutionary relationship between them \citep{Gebhardt00, Merritt01, Tremaine02, Marconi03, Haring04, Shields06}."
.. Luminous quasars. at high redshifts. which represent the most massive SMBHs. are eXperiencing Vigorous accretion of a significant portion of the final SMBH mass.," Luminous quasars at high redshifts, which represent the most massive SMBHs, are experiencing vigorous accretion of a significant portion of the final SMBH mass."
 This accretion activity is believed to be triggered by global processes (galaxy mergers. interactions or perhaps secular evolution) that also trigger major episodes of star formation in the massive host galaxies. which are also rapidly being assembled at high redshifts.," This accretion activity is believed to be triggered by global processes (galaxy mergers, interactions or perhaps secular evolution) that also trigger major episodes of star formation in the massive host galaxies, which are also rapidly being assembled at high redshifts."
 Therefore. the processes of SMBH formation and growth resulting in the quasar phenomenon are directly linked to the birth of massive galaxies (Haehneltetal.1998:Richstone 2006).," Therefore, the processes of SMBH formation and growth resulting in the quasar phenomenon are directly linked to the birth of massive galaxies \citep{Haehnelt98, Richstone98, Omont01, Omont03, Beelen06, Cox06}."
. However. the nature of the relationship between quasars and galaxy formation is not well understood.," However, the nature of the relationship between quasars and galaxy formation is not well understood."
 It is widely believed that strong interactions between gas rich galaxies can result in ultra-luminous infrared galaxies (ULIRGs). defined by Lig. 102L... which have star formation rates (SFRs) of 100AZ.vr.! (Houcketal.1985:Omont2001.2003:2007:Caoetal. 2008).," It is widely believed that strong interactions between gas rich galaxies can result in ultra-luminous infrared galaxies (ULIRGs), defined by $_{\mathrm{IR}}~>$ $^{12}L_{\odot}$, which have star formation rates (SFRs) of $>~100~M_{\odot}~yr^{-1}$ \citep{Houck85, Omont01, Omont03, Flores04, Cox05, Beelen06, Daddi07, Cao08}."
. A fraction of ULIRGs have been found o contain dust-enshrouded quasars (Sanders&Mirabel1996:Lonsdaleetal. 2006).," A fraction of ULIRGs have been found to contain dust-enshrouded quasars \citep{Sanders96, Lonsdale06}."
. There is evidence at both low and high redshift that these embedded quasars are precursors to optically uminous quasars., There is evidence at both low and high redshift that these embedded quasars are precursors to optically luminous quasars.
 For example. Caoetal.(2008). find that low redshift quasars with IR. luminosities in the ULIRG range fall between optically selected PG quasars and ULIRGs (starbursts) or a variety of mid-IR spectroscopic indicators of starburst and active galactic nucleus (AGN) contributions including: polycyclic aromatic hydrocarbon (PAH) luminosities. tine structure emission ine strengths. silicate absorption strengths. spectral slope and mid-IR colour indices.," For example, \citet{Cao08} find that low redshift quasars with IR luminosities in the ULIRG range fall between optically selected PG quasars and ULIRGs (starbursts) for a variety of mid-IR spectroscopic indicators of starburst and active galactic nucleus (AGN) contributions including: polycyclic aromatic hydrocarbon (PAH) luminosities, fine structure emission line strengths, silicate absorption strengths, spectral slope and mid-IR colour indices."
 These findings suggest that IR luminous quasars could represent an intermediate stage between a dominant starburs and a dominant AGN phase., These findings suggest that IR luminous quasars could represent an intermediate stage between a dominant starburst and a dominant AGN phase.
 Quasars with high SFRs therefore may be at an intermediate stage between an embedded quasar in a star forming galaxy anc a visible quasar in a galaxy where most star formation has ceasec (see also Sandersetal. (1988)))., Quasars with high SFRs therefore may be at an intermediate stage between an embedded quasar in a star forming galaxy and a visible quasar in a galaxy where most star formation has ceased (see also \citet{Sanders88}) ).
 At low redshifts. numerous star formation indicators in quasar host galaxies are seen. e.g. high IR luminosities. PAH emission. strong (sub)-mm emission anc," At low redshifts, numerous star formation indicators in quasar host galaxies are seen, e.g. high IR luminosities, PAH emission, strong (sub)-mm emission and"
SN progenitors. and representing half of the total number of 1n concl,"SN progenitors, and representing half of the total number of supernovae (see section \ref{subsubsec:mltas}) )."
usion. the clouds forming in our simulation have solar mietallicities and velocities lower than LOO km s.+. and are therefore likely to be associated to IVCSs rather than IIVCS (2)..," In conclusion, the clouds forming in our simulation have solar metallicities and velocities lower than 100 km $^{-1}$, and are therefore likely to be associated to IVCs rather than HVCs \citep{wak08}."
 This result is in agreement with that found in Paper I when considering single fountains., This result is in agreement with that found in Paper I when considering single fountains.
 Other details of the gas dows of the reference model are given in Fig. 7..," Other details of the gas flows of the reference model are given in Fig. \ref{fig:equil},"
 and. are intended to shed light on the feedback between the disk and the hot halo., and are intended to shed light on the feedback between the disk and the hot halo.
 This energy. exchange regulates the evolution of the halo gas ancl influences the star formation history in the disk., This energy exchange regulates the evolution of the halo gas and influences the star formation history in the disk.
" Ht is also relevant in the related. issue of the so-called. ""over-cooling problem in cosmological simulations of galaxv formation. which eencrally overestimate the amount of raciatively cooled eas in galaxies."," It is also relevant in the related issue of the so-called “over-cooling” problem in cosmological simulations of galaxy formation, which generally overestimate the amount of radiatively cooled gas in galaxies."
 Energy feedback from SNe and/or AGN. albeit xoorlv understood. is thought to be a crucial ingredient o reconcile hierarchical models of galaxy formation with he observations (e.g.22).," Energy feedback from SNe and/or AGN, albeit poorly understood, is thought to be a crucial ingredient to reconcile hierarchical models of galaxy formation with the observations \citep[e.g.][]{nav97,crot06}."
 Supernovae alone are probably incapable to fully solve the problem (?).. but is nevertheless important to quantify the SN energy. which heats the hot ialo gas to correctly estimate its cooling rate and the ialo mass exchange.," Supernovae alone are probably incapable to fully solve the problem \citep{tor04}, but is nevertheless important to quantify the SN energy which heats the hot halo gas to correctly estimate its cooling rate and the disk-halo mass exchange."
 As in our model the amount of thermal energy of all he halo is much larger than the energy. injected. by the SNe. a direct evaluation of its variation due to the stellar explosions would be rather inaccurate (also because of the »erturbations generated at the ericl boundaries. as discussed in Paper D).," As in our model the amount of thermal energy of all the halo is much larger than the energy injected by the SNe, a direct evaluation of its variation due to the stellar explosions would be rather inaccurate (also because of the perturbations generated at the grid boundaries, as discussed in Paper I)."
 We thereby check the energy balance of the halo- interaction as follows., We thereby check the energy balance of the halo-fountains interaction as follows.
 In the top panel of Fig., In the top panel of Fig.
 7 we report the thermal energy of the disk gas that is lifted by the ountains and not condensed into clouds., \ref{fig:equil} we report the thermal energy of the disk gas that is lifted by the fountains and not condensed into clouds.
 This rarefied. gas endures over the disk and its thermal energy is incorporated in the hot halo., This rarefied gas endures over the disk and its thermal energy is incorporated in the hot halo.
 On the other hand. some coronal gas cools and condenses over the clouds created by the MGEs.," On the other hand, some coronal gas cools and condenses over the clouds created by the MGFs."
 The thermal energy. radiated by this cold gas is also shown in the top panel of Eig. 7.., The thermal energy radiated by this cold gas is also shown in the top panel of Fig. \ref{fig:equil}.
 At the end of the simulation. the halo has gained an energy twice larger than that raciated due to the interaction with the fountains. with à net increase of thermal energy of ~1054 of the energy. released by all the SNe.," At the end of the simulation, the halo has gained an energy twice larger than that radiated due to the interaction with the fountains, with a net increase of thermal energy of $\sim 10$ of the energy released by all the SNe."
 This value for the SNe heating cllicieney is expected to remain stable once an approximately steady regime is established between rising and falling gas., This value for the SNe heating efficiency is expected to remain stable once an approximately steady regime is established between rising and falling gas.
 The achievement of this regime is visible in the middle panel of Fig., The achievement of this regime is visible in the middle panel of Fig.
 7 that shows the fraction of rising (solid line) and. descending (dashed line) fountain gas., \ref{fig:equil} that shows the fraction of rising (solid line) and descending (dashed line) fountain gas.
 At early times only the rising gas is present., At early times only the rising gas is present.
 After nearly 40 Myr the first clouds form and fall back., After nearly 40 Myr the first clouds form and fall back.
 A dynamical equilibrium between ascending ancl descending gas is established after 130. Myr., A dynamical equilibrium between ascending and descending gas is established after 130 Myr.
 The same result is illustrated in the bottom panel in ternis of the absolute values of the rising and falling masses., The same result is illustrated in the bottom panel in terms of the absolute values of the rising and falling masses.
 As pointed out. we cannot make an exact evaluation of the energy balance between the fountains and the coronal eas.," As pointed out, we cannot make an exact evaluation of the energy balance between the fountains and the coronal gas."
 In the analvsis above (see the top panel of Fig., In the analysis above (see the top panel of Fig.
 7). we neeleeted the gain of the kinetic energy. of the hot gas interacting with the ealactic fountains ancl the change in the potential energy. of the whole halo.," 7), we neglected the gain of the kinetic energy of the hot gas interacting with the galactic fountains and the change in the potential energy of the whole halo."
 Also. we neglected the raciative losses of the hot gas. possibly enhanced by the compression due to the fountain expansion.," Also, we neglected the radiative losses of the hot gas, possibly enhanced by the compression due to the fountain expansion."
 However. the How of the halo gas remains subsonic and its cooling time is longer than the time of the simulation.," However, the flow of the halo gas remains subsonic and its cooling time is longer than the time of the simulation."
 We thus conclude that à net energy. gain is attained by the halo during its interaction with the MGE's., We thus conclude that a net energy gain is attained by the halo during its interaction with the MGFs.
 In this section we explore the influence of the Galactocentric distance on the gasdynamics of the MGESs., In this section we explore the influence of the Galactocentric distance on the gasdynamics of the MGFs.
 In particular. we have run a model similar to the RAL but with the active area centred. at A=4.5 kpc.," In particular, we have run a model similar to the RM, but with the active area centred at $R=4.5$ kpc."
 Here the gas of the disk. is clenser and has a larger effective vertical scale heigh Lap., Here the gas of the disk is denser and has a larger effective vertical scale heigh $H_{\rm eff}$.
 As a consequence. the disk-halo transition occurs at z=1.35 kpe. instead of 2=0.8 kpe as in the RAL (see also Paper 1).," As a consequence, the disk-halo transition occurs at $z=1.35$ kpc, instead of $z=0.8$ kpc as in the RM (see also Paper I)."
 The larecr density of the ambient medium makes more difficult’ the expansion of the SN. bubbles. through the plane and accelerates their backfilling., The larger density of the ambient medium makes more difficult the expansion of the SN bubbles through the plane and accelerates their backfilling.
 Phe reduced size and lifetime of the holes also reduce their probability of merging., The reduced size and lifetime of the holes also reduce their probability of merging.
 This is apparent in Fig. S.. ," This is apparent in Fig. \ref{fig:coldenism4}, ,"
where the number and the size of the holes are smaller than in the reference model (cf., where the number and the size of the holes are smaller than in the reference model (cf.
 Fig. 2))., Fig. \ref{fig:coldenism}) ).
 In Paper Lit was shown how the vertical distribution of the disk eas influences the dynamics of a SC and the shape ofthe chimney carved by the fountain., In Paper I it was shown how the vertical distribution of the disk gas influences the dynamics of a SGF and the shape of the chimney carved by the fountain.
 In particular. a single fountain located at /?=4.5 kpe gives rise to a sort of well collimated outflow that crosses the cisk/halo transition only barely.," In particular, a single fountain located at $R=4.5$ kpc gives rise to a sort of well collimated outflow that crosses the disk/halo transition only barely."
 On the contrary. given the lower density. and scale height of thelocal gas. the same fountain located at," On the contrary, given the lower density and scale height of thelocal gas, the same fountain located at"
the N-wiud outflow is maguctically connected to the disk.,the X-wind outflow is magnetically connected to the disk.
 Furthermore. the N-wind is unique in that it asstues that the accretion of material does not deposit angular monmenutun onto the star.," Furthermore, the X-wind is unique in that it assumes that the accretion of material does not deposit angular momentum onto the star."
 As mentioned in item 3 above. for most of the spec‘fie cases of our simulated winds. we found that iu he spin equilibrium state. the truncation radius was Wwαπ close to the corotation radius.," As mentioned in item 3 above, for most of the specific cases of our simulated winds, we found that in the spin equilibrium state, the truncation radius was very close to the corotation radius."
 This is simular to he prediction of the disk locking models. which inchles both the Ghosh Lamb-type models aud the N-wixl.," This is similar to the prediction of the disk locking models, which includes both the Ghosh Lamb-type models and the X-wind."
" However. in coutrast with the disk locking models. i Is not a requiremieut of APSW that R, be close to Rcus"," However, in contrast with the disk locking models, it is not a requirement of APSW that $R_{\rm t}$ be close to $R_{\rm co}$."
" Measurements of the location of the inner edge of he eas disk suggest flat Rf Re, is typically ~TO%..", Measurements of the location of the inner edge of the gas disk suggest that $R_{\rm t}$ $R_{\rm co}$ is typically $\sim 70$.
 We leave a |more detaiας colmparison between models for future work., We leave a more detailed comparison between models for future work.
 There is observational evidence that the outflow raCR from accreting voung stellar svstenis are of the order of of the accretion rates aud are therefore. accretio1 powered19903., There is observational evidence that the outflow rates from accreting young stellar systems are of the order of of the accretion rates and are therefore accretion powered.
. ITowever. it appears that a larec fraction of this flow (which is usually probed by orbidden enissiou comune from large spatial scales) origlnates in the cdisk. rather than the star006).," However, it appears that a large fraction of this flow (which is usually probed by forbidden emission coming from large spatial scales) originates in the disk, rather than the star."
. Tt is not ve clear how much of the total observed How lav originate i a stellar wind., It is not yet clear how much of the total observed flow may originate in a stellar wind.
 There is some evideuce specifically for ellar winds from CTTSs2007)... as distinct from disk winds. aud that these are accretion powered2006).. but the mass outflow rates are not vet well coustrained2005).," There is some evidence specifically for stellar winds from CTTSs, as distinct from disk winds, and that these are accretion powered, but the mass outflow rates are not yet well constrained."
". Additional work constraining the value of AA./AL,. the stellar wind driving mechanisin. and the stellar magnetic field streneth and ecometry will help to provide stringcut and quantitative tests for the APSW iodel."," Additional work constraining the value of $\dot M_{\rm w} / \dot M_{\rm a}$, the stellar wind driving mechanism, and the stellar magnetic field strength and geometry will help to provide stringent and quantitative tests for the APSW model."
 The predictions of the spin equilibrium state can also be checked observationally2002)., The predictions of the spin equilibrium state can also be checked observationally.
. This will likely require large saaples of stars. due to large uncertaiuties iun measured parameters. and since lutrinsic variability iu real systems navy onlv allow a spin equilibrium state to be achieved in a time-averaged scuse.," This will likely require large samples of stars, due to large uncertainties in measured parameters, and since intrinsic variability in real systems may only allow a spin equilibrium state to be achieved in a time-averaged sense."
 The Cohosh Laub. Newind. and APSW models all predict an equilibriuu spin rate that depends ou WV. but of these three. only the APSW inodel contains an additional dependence on the stellar wind mass outflow rate.," The Ghosh Lamb, X-wind, and APSW models all predict an equilibrium spin rate that depends on $\Psi$, but of these three, only the APSW model contains an additional dependence on the stellar wind mass outflow rate."
 For the power law fits to the simulations of Paper IL. the APSW spin equililrinim predicts a sheltly different power-law of spin V than the other models (81.1)) though the exact depeudeuce of the stellar wind torque has not been determined for all parameters.," For the power law fits to the simulations of Paper II, the APSW spin equilibrium predicts a slightly different power-law of spin $\Psi$ than the other models \ref{sec_semianalytic}) )—though the exact dependence of the stellar wind torque has not been determined for all parameters."
 Tn this series of papers. we have focused on the global woblem of calculating he magnitude of stellar wiud orques and comparing thei with other torques acting on accreting stars.," In this series of papers, we have focused on the global problem of calculating the magnitude of stellar wind torques and comparing them with other torques acting on accreting stars."
 In order to refine the APSW model tuther aud make the predictions more precise. niore work is required.," In order to refine the APSW model further and make the predictions more precise, more work is required."
" In particular. it is not vet clear how the xesence of an accretion disk will iuflucuce the stellar wind torque. aud couversely, how a stellar wind may influence the accretion process."," In particular, it is not yet clear how the presence of an accretion disk will influence the stellar wind torque, and conversely, how a stellar wind may influence the accretion process."
 Also. although it is clear hat there is enough accretion energy to power the stellar wind. it is still not known what actually drives the stellar wind aud how the acerction power may transfer to it.," Also, although it is clear that there is enough accretion energy to power the stellar wind, it is still not known what actually drives the stellar wind and how the accretion power may transfer to it."
 We suspect that a strong fiux of bydromaguetic waves can be excited near the base of the accretion shock aud cau tap the cnerev released there. which may provide an efficient driver for the APSW.," We suspect that a strong flux of hydromagnetic waves can be excited near the base of the accretion shock and can tap the energy released there, which may provide an efficient driver for the APSW."
 We defer a rigourous investigation of the APSW driving mechanisin to future work., We defer a rigourous investigation of the APSW driving mechanism to future work.
 We wish to thank many people for discussious reearding this work. including: Cabor Das. Sylvic Cabrit. Andrea Dupree. Suzan Edwards. Will Fischer. Sluvichiro Tnutsuka. Chris Johus-IErull [ανα Romanova. Frank Shu. Neivan— Stassun. Jeff Valeuti. and Syduev Wolff.," We wish to thank many people for discussions regarding this work, including: Gibor Basri, Sylvie Cabrit, Andrea Dupree, Suzan Edwards, Will Fischer, Shu-ichiro Inutsuka, Chris Johns-Krull, Marina Romanova, Frank Shu, Keivan Stassun, Jeff Valenti, and Sydney Wolff."
 We also thank the referee. Jonatlau Ferreiva. for his useful sugeestious for improving the xeper and WITP for hosting us while finishine the uanuscript.," We also thank the referee, Jonathan Ferreira, for his useful suggestions for improving the paper and KITP for hosting us while finishing the manuscript."
 This research was supported iu part by the National Scicuce Foundation uncer Craut No. PIPYO5-51161., This research was supported in part by the National Science Foundation under Grant No. PHY05-51164.
 SM is supported bv the University of Vireija hrough a Levinson/VITA Fellowship partially fuied x The Frank Levinson Εαν Foundation through he Peninsula Community Foundation., SM is supported by the University of Virginia through a Levinson/VITA Fellowship partially funded by The Frank Levinson Family Foundation through the Peninsula Community Foundation.
 REP is suppored xoa grant from NSERC., REP is supported by a grant from NSERC.
 We follow AIPO5b to calculate the location of the disk πιοο radius. A4. aud the reader will πια details in that paper.," We follow MP05b to calculate the location of the disk truncation radius, $R_{\rm t}$, and the reader will find details in that paper."
 For convenience. we list the relevant equations here.," For convenience, we list the relevant equations here."
 As iu section 3.. we adopt x.=1.," As in section \ref{sec_tdsd}, we adopt $\gamma_{\rm c} = 1$."
" To determine R,. we first find the magnetic connectivity state using the criterion where V is defined by equation (16))."," To determine $R_{\rm t}$, we first find the magnetic connectivity state using the criterion where $\Psi$ is defined by equation \ref{eqn_psi}) )."
" If condition (A1)) is satisfied. the system is in ""state 1.7"," If condition \ref{eqn_trans}) ) is satisfied, the system is in “state 1."""
 Otherwise. it is in “state 2.7," Otherwise, it is in “state 2."""
 In state 1. we determine the truncation radius using Iu state 2. we determine the truncation radius by solving," In state 1, we determine the truncation radius using In state 2, we determine the truncation radius by solving"
the larger environment of the core at r=0.1 pc.,the larger environment of the core at $r=0.1$ pc.
 The distributions of accreted material through the shells at r=R. and r=0.1 pc are different for several reasons., The distributions of accreted material through the shells at $r=R_c$ and $r=0.1$ pc are different for several reasons.
" Firstly, most of the material passing through the shell at r=0.1 pc does not reach the inner shell in the time period considered here, and so is not included in 9.."," Firstly, most of the material passing through the shell at $r=0.1$ pc does not reach the inner shell in the time period considered here, and so is not included in \ref{CDaccinner}."
" Secondly, the external molecular cloud structure is clumpy, with many voids that contain little gas, and which do not contribute to the accretion flow."," Secondly, the external molecular cloud structure is clumpy, with many voids that contain little gas, and which do not contribute to the accretion flow."
" Thirdly, the flow of material towards the cores is not necessarily purely radial."," Thirdly, the flow of material towards the cores is not necessarily purely radial."
" Instead, the accreted material often flows along the path of the high density filaments."," Instead, the accreted material often flows along the path of the high density filaments."
 10 shows that the additional accreted material is coming preferentially from just a few directions and large sections of the surface have no accreted material passing through them., \ref{CDaccouter} shows that the additional accreted material is coming preferentially from just a few directions and large sections of the surface have no accreted material passing through them.
" To get a more conceptual view of how the accretion is proceeding, let us examine the representative cores in more detail."," To get a more conceptual view of how the accretion is proceeding, let us examine the representative cores in more detail."
 Core (a) has very little accretion from its enviroment and becomes a low mass star (0.63 Mo)., Core (a) has very little accretion from its enviroment and becomes a low mass star $0.63 \: {\rm M_{\odot}}$ ).
 Core (b) undergoes a significant amount of accretion from its environment and at the end of the simulation has a mass of 11.52Mo., Core (b) undergoes a significant amount of accretion from its environment and at the end of the simulation has a mass of $11.52 \: {\rm M_{\odot}}$.
 It has a large amount of material passing through the r=0.1 pc surface from a variety of directions., It has a large amount of material passing through the $r=0.1$ pc surface from a variety of directions.
" However, once again some directions contribute more than others."," However, once again some directions contribute more than others."
 Core (c) is a clear example of accretion along a filament., Core (c) is a clear example of accretion along a filament.
" Core (d) is also accreting along a filament, but most of the accreted material comes from close to the sink."," Core (d) is also accreting along a filament, but most of the accreted material comes from close to the sink."
 We carry out a similar analysis to the covariance function analysis used in Section ?? and classify the accretion surfaces of the entire core sample., We carry out a similar analysis to the covariance function analysis used in Section \ref{sec:global} and classify the accretion surfaces of the entire core sample.
" In this case, we found it more useful to use the function Q(d0) to classify the surfaces, rather than the covariance cov(d0), as the 'accretion density' that we are dealing with here is not a real quantity, and so it makes sense to use a dimensionless function."," In this case, we found it more useful to use the function $Q(d\theta)$ to classify the surfaces, rather than the covariance $cov(d\theta)$, as the `accretion density' that we are dealing with here is not a real quantity, and so it makes sense to use a dimensionless function."
" By normalising Q(d0) by its maximum value, it is possible to make a fair comparison between the cores, regardless of how much they accrete."," By normalising $Q(d\theta)$ by its maximum value, it is possible to make a fair comparison between the cores, regardless of how much they accrete."
 The class definitions are very similar to those used previously., The class definitions are very similar to those used previously.
" For Type 0 cores, Q(d0) remains almost flat, with deviations no larger than20%,, while for Type 1 cores, it decreases monotonically as we move away from zero angular separation."," For Type 0 cores, $Q(d\theta)$ remains almost flat, with deviations no larger than, while for Type 1 cores, it decreases monotonically as we move away from zero angular separation."
" For Type 2 cores, Q(d0) shows a double-peaked structure."," For Type 2 cores, $Q(d\theta)$ shows a double-peaked structure."
" The sample of cores is smaller in this case, as the requirement of following the accretion for 2x10* years means that cores which form"," The sample of cores is smaller in this case, as the requirement of following the accretion for $2\E^4$ years means that cores which form"
should be carefully estimated.,should be carefully estimated.
 Given the fact that late-type galaxies are not considered in the SDSS-maxBCG catalogue. their contribution was computed based on available modelings of galaxy populations. in particular the distribution of galaxies as a function of spectral type and luminosity.," Given the fact that late-type galaxies are not considered in the SDSS-maxBCG catalogue, their contribution was computed based on available modelings of galaxy populations, in particular the distribution of galaxies as a function of spectral type and luminosity."
 Since our sample goes from galaxy groups to rich clusters. we have included the environment dependence of these variables.," Since our sample goes from galaxy groups to rich clusters, we have included the environment dependence of these variables."
 We based our model on the known properties of galaxies in the local universe as given by Baloghetal.(2004) who analysed the bimodal distribution of galaxies in a local (z< 0.08) sample of SDSS galaxies (DRI)., We based our model on the known properties of galaxies in the local universe as given by \cite{balogh04} who analysed the bimodal distribution of galaxies in a local $z<0.08$ ) sample of SDSS galaxies (DR1).
 These authors performed a detailed study on the color and luminosity distribution of galaxies as a function of their environment for both early and late-type galaxies., These authors performed a detailed study on the color and luminosity distribution of galaxies as a function of their environment for both early and late-type galaxies.
" For each galaxy they define a density estimator Xa which represents the local projected galaxy density (M,-— 20) in Mpe7? and they use it to classify the different environmental regimes.", For each galaxy they define a density estimator $\Sigma_5$ which represents the local projected galaxy density $M_r<-20$ ) in $^{-2}$ and they use it to classify the different environmental regimes.
 Beside mmentioned before (see Sect. 2.1).," Beside mentioned before (see Sect. \ref{ssec:maxbcg}) ),"
" the SDSS-maxBCG dataset contains also information about the number of galaxies Neal projected within a distance lower than | h-' Mpe from the BCG. with the magnitude limit of M,<16."," the SDSS-maxBCG dataset contains also information about the number of galaxies $N_{\rm gal}$ projected within a distance lower than 1 $h^{-1}$ Mpc from the BCG, with the magnitude limit of $M_r<-16$."
 Therefore for each cluster we can assume a unique value of Xa for all of its members by establishing a relation with Ng.) and. finally. with the properties of the population of late-type galaxies present.," Therefore for each cluster we can assume a unique value of $\Sigma_5$ for all of its members by establishing a relation with $N_{\rm gal}$ and, finally, with the properties of the population of late-type galaxies present."
 We proceed as follows., We proceed as follows.
 This formalism led us to derive an average fraction of about spiral galaxies in our sample of SDSS-maxBCG clusters (see Table 1))., This formalism led us to derive an average fraction of about spiral galaxies in our sample of SDSS-maxBCG clusters (see Table \ref{tab:gal_pop}) ).
" The relation between N,,, and the fraction of", The relation between $N_{\rm gal}$ and the fraction of
are cluster members according to their radial velocity.,are cluster members according to their radial velocity.
 The colour-magnitude diagram of the cluster (Fig. 1)), The colour–magnitude diagram of the cluster (Fig. \ref{fhd362}) )
 together with the above mentioned test using the TRILEGAL stellar population model suggest that all variables from our list are cluster members., together with the above mentioned test using the TRILEGAL stellar population model suggest that all variables from our list are cluster members.
 While the cluster is seen in projection on the outskirts of the Small Magellanic Cloud (SMC). any contamination of the observed upper giant branch by extragalactic stars 1s expected to be negligible due to the very low number density of SMC stars in that colour-magnitude range (cf.," While the cluster is seen in projection on the outskirts of the Small Magellanic Cloud (SMC), any contamination of the observed upper giant branch by extragalactic stars is expected to be negligible due to the very low number density of SMC stars in that colour-magnitude range (cf."
 Udalski et al. 2008))., Udalski et al. \cite{udalski08}) ).
 Seventeen red variables of NGC 2808 showed a (semi-)periodic light variation. including two variables known before (V1 and VID.," Seventeen red variables of NGC 2808 showed a (semi-)periodic light variation, including two variables known before (V1 and V11)."
 reffhd2808 shows the location of the variables in a colour—magnitude diagram of the cluster., \\ref{fhd2808} shows the location of the variables in a colour--magnitude diagram of the cluster.
 For 4 stars. all having a comparably short period. only the ANDICAM data. which show a higher sampling in time. were used for the analysis.," For 4 stars, all having a comparably short period, only the ANDICAM data, which show a higher sampling in time, were used for the analysis."
 The corresponding phased light curves using the period with the largest amplitude are shown in ref2808les.., The corresponding phased light curves using the period with the largest amplitude are shown in \\ref{2808lcs}.
 In three cases (LW6. LW9 and LWI4) we subtracted a long time trend before producing the phase diagram.," In three cases (LW6, LW9 and LW14) we subtracted a long time trend before producing the phase diagram."
 For LW2. ANDICAM and MACHO data were analysed independently and both gave a very similar period.," For LW2, ANDICAM and MACHO data were analysed independently and both gave a very similar period."
 In the plot we show only one set of ANDICAM data as there seems to be a phase shift between the older and the more recent time series., In the plot we show only one set of ANDICAM data as there seems to be a phase shift between the older and the more recent time series.
 [n addition to the phase plots shown in ref2808les.. light curves plotted against time are shown in ref28081es2. for selected stars. including those with long term trends. multiple periods or long secondary periods.," In addition to the phase plots shown in \\ref{2808lcs}, light curves plotted against time are shown in \\ref{2808lcs2} for selected stars, including those with long term trends, multiple periods or long secondary periods."
 Beside the stars shown. we found 5 targets located on the upper part of the giant branch. which clearly show variability but without a clearly expressed periodicity.," Beside the stars shown, we found 5 targets located on the upper part of the giant branch, which clearly show variability but without a clearly expressed periodicity."
 These stars are listed at the bottom of Table 3.., These stars are listed at the bottom of Table \ref{n2808vars}.
 As with the LPVs in NGC 362. irregularities in. the light change can be seen on top of the periodic variations.," As with the LPVs in NGC 362, irregularities in the light change can be seen on top of the periodic variations."
 Some stars. e.g. LW5. show clear signatures of an additional short time variation on top of the longer period of 359 days.," Some stars, e.g. LW5, show clear signatures of an additional short time variation on top of the longer period of 359 days."
 All reliably determined periods are listed in Table 3.., All reliably determined periods are listed in Table \ref{n2808vars}.
 There is no source corresponding to LWIS in the 2MASS catalogue., There is no source corresponding to LW15 in the 2MASS catalogue.
 For this star the coordinates are estimated using the SIMBAD/Aladin tool on the 2MASS J-band image., For this star the coordinates are estimated using the SIMBAD/Aladin tool on the 2MASS $J$ -band image.
 We also list in Table 3. our own K-band photometry (transfered to the 2MASS system) where available., We also list in Table \ref{n2808vars} our own $K$ -band photometry (transfered to the 2MASS system) where available.
 The given uncertainty of the A-band data includes both the measurement error and the star’s near-infrared variability during the four months of monitoring., The given uncertainty of the $K$ -band data includes both the measurement error and the star's near-infrared variability during the four months of monitoring.
" For V1Il. LW9, LWI8 and LWI9 we have only a single epoch measurement."," For V11, LW9, LW18 and LW19 we have only a single epoch measurement."
 A well expressed near- variability was detected for three stars in our sample. VI. LWI and LWIS.," A well expressed near-infrared variability was detected for three stars in our sample, V1, LW1 and LW15."
 The light curves are shown in refIR2808 together with the visual data obtained at the same time., The light curves are shown in \\ref{IR2808} together with the visual data obtained at the same time.
 It can be seen that for LW] and LW15 the optical and infrared lighteurves are nicely in phase. while for VI there is à large difference.," It can be seen that for LW1 and LW15 the optical and infrared lightcurves are nicely in phase, while for V1 there is a large difference."
 The reason for that is not clear., The reason for that is not clear.
 The slightly larger uncertainty of LW6 is most likely due to a higher photometric error rather than due to variability., The slightly larger uncertainty of LW6 is most likely due to a higher photometric error rather than due to variability.
 The star has another near-infrared source nearby so that blending becomes a problem., The star has another near-infrared source nearby so that blending becomes a problem.
 Note that for LW15 we see the dd period only as a general trend of the average brightness in refIR2808.., Note that for LW15 we see the d period only as a general trend of the average brightness in \\ref{IR2808}.
 The star obviously shows also a short period change of about 50 days length. which can be seen as some scatter in ref2808les..," The star obviously shows also a short period change of about 50 days length, which can be seen as some scatter in \\ref{2808lcs}."
 For most sources we find a good agreement of our K-band measurement with the 2MASS value taking into account the variability of the objects., For most sources we find a good agreement of our $K$ -band measurement with the 2MASS value taking into account the variability of the objects.
 One exception is V1., One exception is V1.
 Looking at the, Looking at the
The assumption that the magnetosphere is corotating has the implication (hat the inductive field is absent and hence must be screenecl,The assumption that the magnetosphere is corotating has the implication that the inductive field is absent and hence must be screened.
 With the inductive field given by equation (3)). the required current density is U.x)-—," With the inductive field given by equation \ref{Eind3}) ), the required current density is )=."
 The parallel component of the screening current follows from equation (9)): X)Y., The parallel component of the screening current follows from equation \ref{dEIp}) ): )=.
2t We emphasize (hat the current density given bv equation (26)) is required by the hypothesis (hat the magnetosphere of an oblique rotator is rigidly corotating., We emphasize that the current density given by equation \ref{scd2}) ) is required by the hypothesis that the magnetosphere of an oblique rotator is rigidly corotating.
 La order for (his to be the case. the plasma must supply the current J... given by equation (24)).," In order for this to be the case, the plasma must supply the current ${\bi J}_{\rm sc}$ , given by equation \ref{scd1}) )."
" The parallel component of Ej,4 can be screened bycharges. and the parallel component ol OE;4/0! can be screened bv the eurrent. given by equation (27))."," The parallel component of ${\bi E}_{\rm ind}$ can be screened bycharges, and the parallel component of ${\partial{\bi E}_{\rm ind}/\partial t}$ can be screened by the current given by equation \ref{scd3}) )."
 Such screening occurs provided (hat there are sufficient charges available to provide the charge and current densities., Such screening occurs provided that there are sufficient charges available to provide the charge and current densities.
 The required number of charges can be estimated as follows., The required number of charges can be estimated as follows.
 The inductive and corotational fields are proportional to «/77. and they differ in magnitude only by geometric factors.," The inductive and corotational fields are proportional to $\omega/r^2$, and they differ in magnitude only by geometric factors."
 Hence. the charge density required to screen the μμ differs from the Goldreich-Julian value. po). only by a similar geometric factor.," Hence, the charge density required to screen the $E_{{\rm ind}\parallel}$ differs from the Goldreich-Julian value, $\rho_{\rm GJ}$, only by a similar geometric factor."
 Screening of Fina occurs providedthat secondary pair creation results in a niultüiplidity. Af. greater (han unity (Deskinetal. 1993)..," Screening of $E_{{\rm ind}\parallel}$ occurs providedthat secondary pair creation results in a multiplicity, $M$, greater than unity \citep{PhysicsOfThePulsarMagnetosphere}. ."
 Specifically. if the number densiües of electrons and positrons are 0," Specifically, if the number densities of electrons and positrons are$n_\pm$ ,"
 Specifically. if the number densiües of electrons and positrons are 0.," Specifically, if the number densities of electrons and positrons are$n_\pm$ ,"
"in à given hardening interaction the binding energv of the παν typically increases by a fraction ~O.2(m/AL)=0ο AL. where the mass of the large black bole is 50M,AM. and of the smaller black holes is 10mijoM...","in a given hardening interaction the binding energy of the binary typically increases by a fraction $\sim
0.2(m/M)=0.04m_{10}M_{50}^{-1}$ , where the mass of the large black hole is $50M_{50}\,M_\odot$ and of the smaller black holes is $10m_{10}\,M_\odot$."
 The fraction of the change in binding energy. carried. away » binary recoil also depends on the mass ratio., The fraction of the change in binding energy carried away by binary recoil also depends on the mass ratio.
 For three equal-mass objects. conservation of momentum implies that he kinetic energy. of recoil of the binary is approximately a third of the change in binding energy.," For three equal-mass objects, conservation of momentum implies that the kinetic energy of recoil of the binary is approximately a third of the change in binding energy."
 LE the binary is much more massivethan the third star. AZ»m. then the kinetic energv of the binary is instead roughly mA times he change in binding energy.," If the binary is much more massivethan the third star, $M\gg m$, then the kinetic energy of the binary is instead roughly $m/M$ times the change in binding energy."
 Let e; be the escape velocity Tom the core. where typically ce:2250 kni + for a dense cluster (\Webbink 1985).," Let $v_{\rm esc}$ be the escape velocity from the core, where typically $v_{\rm esc}\approx 50$ km $^{-1}$ for a dense cluster (Webbink 1985)."
 Then the binary binding energy hat twpically produces recoil at the escape velocity is This needs to be compared with the binding energy such that gravitational raciation causes a merger before the next three-body interaction., Then the binary binding energy that typically produces recoil at the escape velocity is This needs to be compared with the binding energy such that gravitational radiation causes a merger before the next three-body interaction.
 The merger time for two stars with total mass AZ and reduced mass (iin an orbit of semimajor axis e and eecentricitv e is (Peter 1964) From numerical simulations of three equal-amass black roles (Portegics Zwart MeMillan. 2000). the typical eccentricity. distribution of binaries after the interaction of hree equalmass objects is le) (with a slight excess at high eccentricities).," The merger time for two stars with total mass $M$ and reduced mass $\mu$ in an orbit of semimajor axis $a$ and eccentricity $e$ is (Peter 1964) From numerical simulations of three equal-mass black holes (Portegies Zwart McMillan 2000), the typical eccentricity distribution of binaries after the interaction of three equal-mass objects is $P(e)=2e$ (with a slight excess at high eccentricities)."
 With this distribution. the average eccentricity weighted over merger time is οzz0.7.," With this distribution, the average eccentricity weighted over merger time is $e\approx 0.7$."
 We note. rowever. that the strong dependence of merger time on eccentricity means that deviations from this distribution can aave large effects on the rate of growth of large black holes in elobular clusters. the detectability of gravitational radiation rom these systems. and other important quantities.," We note, however, that the strong dependence of merger time on eccentricity means that deviations from this distribution can have large effects on the rate of growth of large black holes in globular clusters, the detectability of gravitational radiation from these systems, and other important quantities."
 Quinlan (1996) finds that when perturbations are small (as they are when the mass ratio is high) the eecentricity of a hard unary is increased. steadily by the hardening process until eravitational radiation. becomes. important., Quinlan (1996) finds that when perturbations are small (as they are when the mass ratio is high) the eccentricity of a hard binary is increased steadily by the hardening process until gravitational radiation becomes important.
 If so. merger imes are decreased. and conditions are more favourable for 1e growth of ~LO?AL. black holes in the centers of globular Justers.," If so, merger times are decreased and conditions are more favourable for the growth of $\sim 10^3\,M_\odot$ black holes in the centers of globular clusters."
 However. in the presence. of finite perturbations 10 outconie is not as clear.," However, in the presence of finite perturbations the outcome is not as clear."
 An important future project is jerefore a numerical study of the growth of Z250AZ. black oles in globulars. with a full treatment of the evolution of 1e eccentricitv.," An important future project is therefore a numerical study of the growth of $\gta 50\,M_\odot$ black holes in globulars, with a full treatment of the evolution of the eccentricity."
" In the meantime. we assume conservatively ju Pee)="" peearddless of the amount of hardening that mas occurred. until the point that gravitational racliation is significant."," In the meantime, we assume conservatively that $P(e)=2e$ regardless of the amount of hardening that has occurred, until the point that gravitational radiation is significant."
 The typical semimajor axis for a merger time LO? vr and ALcm ds then The binding energy. is Therefore. the ratio of the binding energy necessary to eject the binary from the cluster to the binding energy necessary for a fast moereer is 1 μιcl1. mergersὃν happen before ejection. so that. the hole ὃνgains mass and remains near the mass source.," The typical semimajor axis for a merger time $\tau_{\rm merge}= 10^6\tau_6$ yr and $M\gg m$ is then The binding energy is Therefore, the ratio of the binding energy necessary to eject the binary from the cluster to the binding energy necessary for a fast merger is If $x>1$, mergers happen before ejection, so that the hole gains mass and remains near the mass source."
 κ1. ection happens first. so that although the binary migh merge in less than a Hubble time it will do so well away from the high stellar density in the cluster. meaning that the black hole mass growth is stopped.," If $x<1$, ejection happens first, so that although the binary might merge in less than a Hubble time it will do so well away from the high stellar density in the cluster, meaning that the black hole mass growth is stopped."
 This calculation shows that if ALZD50AL. the hole will stay near the center of the potentia and grow rapidly. whereas for much smaller masses the hole will be ejected quickly without additional mass growth.," This calculation shows that if $M\gta 50\,M_\odot$ the hole will stay near the center of the potential and grow rapidly, whereas for much smaller masses the hole will be ejected quickly without additional mass growth."
 Lf the eccentricitv increases as the orbit shrinks. the critica seniümajor axis increases. and hence the ratio we increases.," If the eccentricity increases as the orbit shrinks, the critical semimajor axis increases, and hence the ratio $x$ increases."
 Therefore. eccentricity growth allows lower-mass black holes to grow by mergers without being ejected.," Therefore, eccentricity growth allows lower-mass black holes to grow by mergers without being ejected."
" The interaction cross section [or a massive black hole is dominated by gravitational focusing. so that evàary,(20M (2)."," }---The interaction cross section for a massive black hole is dominated by gravitational focusing, so that $\sigma_{\rm coll}\approx \pi r_p(2GM/v_\infty^2)$ ."
 The relevant pericenter. distance. depends on the type of encounter and the nature of the secondary object., The relevant pericenter distance depends on the type of encounter and the nature of the secondary object.
" For example. a 10437. DII companion with a tvpical eccentricity ο=0.7 and a semimajor axis ~51045. ‘em will meree via gravitational racliation within 10""z, vr."," For example, a $10\,M_\odot$ BH companion with a typical eccentricity $e=0.7$ and a semimajor axis $\sim 5\times 10^{11}\tau_6^{1/4}$ cm will merge via gravitational radiation within $10^6\tau_6$ yr."
 A third black hole passing within ~twice that distance will produce a strong interaction. πο ry~Lo12 cm.," A third black hole passing within $\sim$ twice that distance will produce a strong interaction, so $r_p\sim 10^{12}$ cm."
 In thermal equilibrium.qup the velocity. dispersion. of stars is inversely proportional to the square root of their mass. so if the average mass of a main sequence star is Q.4AZ. ancl its velocity is ey.=10μις cm then black holes have a velocity dispersion. οστzHbag11/2Efms=.219myPans cm ," In thermal equilibrium the velocity dispersion of stars is inversely proportional to the square root of their mass, so if the average mass of a main sequence star is $0.4\,M_\odot$ and its velocity is $v_{\rm ms}=10^6v_{\rm ms,6}$ cm $^{-1}$ then black holes have a velocity dispersion $v\approx {1\over 5}m_{10}^{-1/2}v_{\rm ms}=
2\times 10^5m_{10}^{-1/2}v_{\rm ms,6}$ cm $^{-1}$."
"Then ea&l0au(r,/10712em)AMzonmuo emer.2 so the time between interactions is Tat=Lffnary)c2lOoeisl,.rp1MylngLf yr. where p,=1072n,1s cni and the average in angle brackets is takenover the velocity clistribution of black holes (assumed to be Maxwellian)."," Then $\sigma_{\rm coll}\approx
10^{30}(r_p/10^{12}\,{\rm cm})M_{50}m_{10}$ $^2$, so the time between interactions is $\tau_{\rm int}=1/\langle n\sigma v_\infty\rangle\approx 2\times 10^6 v_{ms,6}
n_6^{-1}r^{-1}_{p,12}M_{50}^{-1}m_{10}^{-1/2}$ yr, where $r_p=10^{12}r_{p,12}$ cm and the average in angle brackets is takenover the velocity distribution of black holes (assumed to be Maxwellian)."
 The binding energv scales with 1α. as does the time between interactions.," The binding energy scales with $1/a$, as does the time between interactions."
 JFherefore. the orbit. starts shrinking rapidlv. but slows clown.," Therefore, the orbit starts shrinking rapidly, but slows down."
 The total time required to reach @=5.101L cm from an initially much larger semimajor axis is ~b(Mm)ng i£ a fraction 0.2(m/AL) of the binding energv is removed in each close encounter.," The total time required to reach $a=
5\times 10^{11}$ cm from an initially much larger semimajor axis is $\sim 5(M/m)\tau_{\rm int}$ if a fraction $\sim 0.2(m/M)$ of the binding energy is removed in each close encounter."
" The orbital separation. for⋅ a fixed⋅ merger time. scales as εςLe. so Tin~ and a ~SOAL. black hole will increase its mass by 1034. ina total time ~5«105,2.nm Mt"," The orbital separation for a fixed merger time scales as $M_{50}^{1/2}$, so $\tau_{\rm int}\sim M_{50}^{-3/2}$ and a $\sim 50\,M_\odot$ black hole will increase its mass by $\sim 10
\,M_\odot$ in a total time $\sim 5\times 10^7n_6^{-1}M_{50}^{-1/2}m_{10}^{-5/2}$ yr."
 l1ntegrating. after time / the mass of the black hole is Mom[(s10ηςμα.vt) .," Integrating, after time $t$ the mass of the black hole is $M_{50}\approx \left[t/\left(3\times 10^8\,n_6^{-1}m_{10}^{-5/2}\,{\rm yr}
\right)\right]^2$ ."
 Within a Hubble time there is therefore time for substantial| growth in the mass of large black holes in the centers of clusters that are at least moderately dense. with πο 20.1. .," Within a Hubble time there is therefore time for substantial growth in the mass of large black holes in the centers of clusters that are at least moderately dense, with $n_6\gta 0.1$ ."
 as argued by Quinlan (1996). three-bocly elfects. increase the eccentricity as the binary hardens. the rate of increase of black hole mass is enhanced significantly.," If, as argued by Quinlan (1996), three-body effects increase the eccentricity as the binary hardens, the rate of increase of black hole mass is enhanced significantly."
the reprocessor.,the reprocessor.
 We constructed. emissivity. profiles for the D and V bands assuming a simple ecometry of an accretion disc with constant Z£/A1. and an N-rayv source on the axis of svmmetry over the disc.," We constructed emissivity profiles for the B and V bands assuming a simple geometry of an accretion disc with constant $H/R$, and an X-ray source on the axis of symmetry over the disc."
 The intrinsic clisc accretion rate used was =UW ob the Exldington limit. a DII mass of 3.JO'AZ. anda value of 42/8=0.1.," The intrinsic disc accretion rate used was $\dot m=1$ of the Eddington limit, a BH mass of $3\times 10^{7} M_\odot$ anda value of $H/R=0.1$."
" The X-ray [lux was equal to five times the average 2109 keV. flux measured. o account for X-ray flux outside this band. (seeDeRosapoSAN. and assumed that of the impinging X-ray lux were thermalised by the disc (seedisc""ussioninArévaloetal."," The X-ray flux was equal to five times the average 2–10 keV flux measured, to account for X-ray flux outside this band \citep[see][for their X-ray broadband analysis using {\it BeppoSAX} and assumed that of the impinging X-ray flux were thermalised by the disc \citep[see discussion in][]{MR2251}."
" 2008).. Figure 10. shows the cumuative optical [Lux as a [function of radius for the intrinsic. reprocessed. and otal emission for an N-rav source heiglx of 50 /,."," Figure \ref{profile} shows the cumulative optical flux as a function of radius for the intrinsic, reprocessed and total emission for an X-ray source height of 50 $R_g$."
" The reprocessed, curve was calculated as the cUlerence between he intrinsic and total cumulative [luxes.", The reprocessed curve was calculated as the difference between the intrinsic and total cumulative fluxes.
 Evidently. the intrinsic anc reprocessecl emissivities peak at cdilferent racii with the reprocessecl emission being procuced further out.," Evidently, the intrinsic and reprocessed emissivities peak at different radii with the reprocessed emission being produced further out."
 As this is the emission. modulated. by tie variable X-ray Hux. we should. consider his outer radius to interpret the delay between X-ray ancl optical Huctuations produced. by reprocessing.," As this is the emission modulated by the variable X-ray flux, we should consider this outer radius to interpret the delay between X-ray and optical fluctuations produced by reprocessing."
 For a face-on view. the parameters used in Fie.," For a face-on view, the parameters used in Fig."
 10 oocduce a delay of ~10 jours for the X-rays to modulate up to of the reprocessed [ux in both B and V bancs. too short to explain the observations.," \ref{profile} produce a delay of $\sim 10$ hours for the X-rays to modulate up to of the reprocessed flux in both B and V bands, too short to explain the observations."
 Using a thinner disc. clown o LAR=0.001 does no have a signilicant effect on this delay. nor does reducing the impinging X.-rav [lux down to he observed 2.10 keV. value (instead of Liv0 times this value as Wels σος for the figure).," Using a thinner disc, down to $H/R=0.001$ does not have a significant effect on this delay, nor does reducing the impinging X-ray flux down to the observed 2–10 keV value (instead of five times this value as was used for the figure)."
 Varving the position of the rav source does have a strong elfect on the time delay. partly »ecause increasing the height moves the main reprocessing region further out in the disc but mostly. because. of the onger light travel time down to the disc.," Varying the position of the X-ray source does have a strong effect on the time delay, partly because increasing the height moves the main reprocessing region further out in the disc but mostly because of the longer light travel time down to the disc."
" The measured delay ofa few days would require the N-rav. sotwee to be located adoLOOOR, over the disc.", The measured delay of a few days would require the X-ray source to be located at $\sim 1000 R_g$ over the disc.
 A cilferent ecometry of the reprocessor. however. might also produce the correct time delay.," A different geometry of the reprocessor, however, might also produce the correct time delay."
 For example a lower rav source and a flared cise would present a large area of 1e disc to the X-ray. flux only at large racii. pushing the cation of reprocessecl emission further out.," For example a lower X-ray source and a flared disc would present a large area of the disc to the X-ray flux only at large radii, pushing the location of reprocessed emission further out."
 Alternatively. 10 reprocessed N-rays might come from a structure similar o the broad line region (BLIU.," Alternatively, the reprocessed X-rays might come from a structure similar to the broad line region (BLR)."
 Reverberation mapping campaigns of this source show that the broad 112 line lags Ίο UV continuum variations by 110 days. depending on yw line width (Petersonetal.2004).," Reverberation mapping campaigns of this source show that the broad $\beta$ line lags the UV continuum variations by 1–10 days, depending on the line width \citep{revmap}."
.. Therefore. he lieht ravel time from the centre of he accretion disc (probable source of optical/UV continuum) to the BLE is similar to 1¢ distance between the X-ray corona and the reprocessing material.," Therefore, the light travel time from the centre of the accretion disc (probable source of optical/UV continuum) to the BLR is similar to the distance between the X-ray corona and the reprocessing material."
 In this scenario. the broad line variability tracks 1e large. amplitude optical continuum. Huctuations. which modulate the intrinsic optical [lux originating close LOO.) to the black hole.," In this scenario, the broad line variability tracks the large amplitude optical continuum fluctuations, which modulate the intrinsic optical flux originating close $\sim 100R_g$ ) to the black hole."
 At the same time. the ines can x: modulated: by the rapid. optical [uctuations. produced » X-ray reprocessing at the same radius of he DLIt. which alreacy contains a delay to the intrinsic optical ancl X-ray emission.," At the same time, the lines can be modulated by the rapid optical fluctuations, produced by X-ray reprocessing at the same radius of the BLR, which already contains a delay to the intrinsic optical and X-ray emission."
 Recalling the three geometries. of the PCPLrocessor ISCUSSEC in Sec., Recalling the three geometries of the reprocessor discussed in Sec.
 ?? we can now judee their ability. to produce he rapid. small-amplitucie optical Ductuations only.," \ref{transfersec} we can now judge their ability to produce the rapid, small-amplitude optical fluctuations only."
 Although it is not straightforward to clisentanele long aud gajort term. [luctuations. we can use the low {lux section in 1e first half of the campaign to estimate the amplitude and timescales of what are probably signs of reprocessing.," Although it is not straightforward to disentangle long and short term fluctuations, we can use the low flux section in the first half of the campaign to estimate the amplitude and timescales of what are probably signs of reprocessing."
 The two big Lares on this segment. which show a clear lag to 10 X-rays. have timescales of 10.20 days and amplitudes of afew my in the D band.," The two big flares on this segment, which show a clear lag to the X-rays, have timescales of 10–20 days and amplitudes of a few mJy in the B band."
 Also. their relative amplitude is about hall or less of the corresponding X-ray lares.," Also, their relative amplitude is about half or less of the corresponding X-ray flares."
 “Phe reprocessing mechanism. therefore. should. produce a few mJy of optical Hux and not smooth out completely. the variations on timoe-scales of tens of davs.," The reprocessing mechanism, therefore, should produce a few mJy of optical flux and not smooth out completely the variations on time-scales of tens of days."
 The truncated. fat dise reprocessor seems unlikely. on energeties grounds because the X-ray. source sees the disc almost edge-on. given the large truncation radius needed for the lag. for anv reasonable source height.," The truncated flat disc reprocessor seems unlikely, on energetics grounds because the X-ray source sees the disc almost edge-on, given the large truncation radius needed for the lag, for any reasonable source height."
 This small solid angle subtended by the cise produces too little reprocessed Hux to make an observable signature in the optical light curves. by many orders of magnitude.," This small solid angle subtended by the disc produces too little reprocessed flux to make an observable signature in the optical light curves, by many orders of magnitude."
 This problem. is circumvented when using the IHared cise because the inclined or vertical portion of the disc can face the X-ray source directly. reaches a higher temperature ancl produces more optical Dux.," This problem is circumvented when using the flared disc because the inclined or vertical portion of the disc can face the X-ray source directly, reaches a higher temperature and produces more optical flux."
 The BLR. clouds can have large enough column densities o be optically trick to X-rays (~10τοι2008)., The BLR clouds can have large enough column densities to be optically thick to X-rays \citep[$\sim 10^{24}{.
 teprocessing in he form of optical emission lines is possible and the effect ο ‘their fluctuations on the DB and. V. band luxes can he «sstimated from spectral monitoring., Reprocessing in the form of optical emission lines is possible and the effect of their fluctuations on the B and V band fluxes can be estimated from spectral monitoring.
 The reverberation niipping campaign on »presented. by Sirpeetal.(1994).. shows that the lines in his region of the spectrum contribute around. of the DB band Lux anc vary by a maximum ratio of in flux.," The reverberation mapping campaign on presented by \citet{Stirpe}, shows that the lines in this region of the spectrum contribute around of the B band flux and vary by a maximum ratio of in flux."
 Therefore. their contribution to the B and V band variability is too small to produce the observed rapid. [iuctuations of about20'4.," Therefore, their contribution to the B and V band variability is too small to produce the observed rapid fluctuations of about."
. The BLR. however. can also produce. cülfuse continuum emission in addition to the lines. as shown bv IxXorista&CGoack (2001).," The BLR, however, can also produce diffuse continuum emission in addition to the lines, as shown by \citet{koristagoad}. ."
. These authors calculate the expected. BLE. contribution to the optical/U continuum in detail for the case of NGC 5548., These authors calculate the expected BLR contribution to the optical/UV continuum in detail for the case of NGC 5548.
 Thev show V.that, They show that
abssorption in the]Notosphlere of tjo system.,sorption in the photosphere of the system.
vandenDereh1991)) ~107 e7054 (Trinchieri.Fabbiano1991).. Trinchieri.Fabbiano&Peres1,\citealp{vanden}) $\sim 10^{39}$ $\sim 70\%$ \citep{trin}. \citealp{trin}) \citealp{market})
991)) Markert&Rallis1983))., $1500 M_{\odot}$ \citep{gebh}.
 1500... Dubusetal.(1997) ~LOG Parmarοἱal.2001)). (Takanoetal.1994:Parmar2001) ~2 e-folding ," \citet{dubus} $\sim 106$ \citealp{parmar}) \citep{takano,parmar} $\sim 2$ $e$ "
 F-test using the OGLE data alone.,$F$ -test using the OGLE data alone.
 Appeuding additional data from various sources to the OGLE data the nonlinearity of the LAIC P-L relation is still preserved from the F-test., Appending additional data from various sources to the OGLE data the nonlinearity of the LMC P-L relation is still preserved from the $F$ -test.
 Inchiding the Cepheids with period shorter than log(P)«0.1 will not alter the detection of nonlinear P-L relation eiven in this table., Including the Cepheids with period shorter than $\log(P)<0.4$ will not alter the detection of nonlinear P-L relation given in this table.
 A similar test is also done to the MACTIIO sample (Necowctal.2005).. as both the MACIIO sample alone aud the “MACTIO|SEBO” siuuples combined have returned a significant F-test result.," A similar test is also done to the MACHO sample \citep{nge05}, as both the MACHO sample alone and the “MACHO+SEBO” samples combined have returned a significant $F$ -test result."
 Therefore we believe thatsamples. as long as we have a good sample (with large nuuber of Cepheids. see next section as wel," Therefore we believe that, as long as we have a good sample (with large number of Cepheids, see next section as well)."
 It is worthwhile to poiut out that the P-L relation is believed. to be linear if the slopes of the long and short period P-L relations are similar., It is worthwhile to point out that the P-L relation is believed to be linear if the slopes of the long and short period P-L relations are similar.
" Some cases of the similar slopes can be seen frou Table 1 (οιο,, OGLE|CL91)."," Some cases of the similar slopes can be seen from Table \ref{tab1} (e.g., OGLE+CL91)."
 However noulinearity of the P-L relation will depend ou both of the slopes aud. the zero-points. aud the full form of nonlinear P-L relations is taken iuto account iu the F-test.," However nonlinearity of the P-L relation will depend on both of the slopes and the zero-points, and the full form of nonlinear P-L relations is taken into account in the $F$ -test."
 In Figure [.. we show the histograms of the F-test results from the simulated P-L relatious using the “OGLE|CL917 and “OGLE|GECGOS7 P-L relations given in Table 1. as the input P-L relations.," In Figure \ref{fighist2}, we show the histograms of the $F$ -test results from the simulated P-L relations using the “OGLE+CL91” and “OGLE+GFG98” P-L relations given in Table \ref{tab1} as the input P-L relations."
 The former P-L relatious have simular slopes while the latter P-L relations have simular zero-poiuts., The former P-L relations have similar slopes while the latter P-L relations have similar zero-points.
 From the figure the noulinear P-L relation is still detected im both cases., From the figure the nonlinear P-L relation is still detected in both cases.
 Iu the previous section the outliers were removed using the sigma-clipping aleorithii with X=2.5., In the previous section the outliers were removed using the sigma-clipping algorithm with $X=2.5$.
 Iu this section we will investigate the iuflueuce of outliors ou the detection of the uoulinear P-L relation., In this section we will investigate the influence of outliers on the detection of the nonlinear P-L relation.
 Using the six samples as given in Table 1. (again. Cepheids with log|P]«0.1 are removed). we apply the siema-clipping algorithm witli to remove outlicrs and then calculate the correspoucding £-values.," Using the six samples as given in Table \ref{tab1} (again, Cepheids with $\log[P]<0.4$ are removed), we apply the sigma-clipping algorithm with to remove outliers and then calculate the corresponding $F$ -values."
 The results are sunumuianzed in Figure 5.., The results are summarized in Figure \ref{figoutlier}.
 From this figure it can be seen that a eeueral trend exists between the adopted value of VY aud the resulting F-values for the LAIC Cepheids., From this figure it can be seen that a general trend exists between the adopted value of $X$ and the resulting $F$ -values for the LMC Cepheids.
 At high .X. the samples show a mareinal detection CF~ 3) of the nonlinearity of the P-L relation.," At high $X$, the samples show a marginal detection $F\sim3$ ) of the nonlinearity of the P-L relation."
 As the adopted value of IX. decreases to ~3 and ~2. the detection of the nonlinear P-L relation becomes significant.," As the adopted value of $X$ decreases to $\sim3$ and $\sim2$, the detection of the nonlinear P-L relation becomes significant."
 When the adopted value of IX decreases to ~1 or less. the P-L relation becomes linear as indicated by a non-significant F-test result (F< 3).," When the adopted value of $X$ decreases to $\sim1$ or less, the P-L relation becomes linear as indicated by a non-significant $F$ -test result $F<3$ )."
 The cases for using large, The cases for using large
the light-curves of GRB 071003 (Ghisellini et al.,the light-curves of GRB 071003 (Ghisellini et al.
 2009) and GRB 061126 (Ghisellint et al., 2009) and GRB 061126 (Ghisellini et al.
 2009; Nardini et al., 2009; Nardini et al.
 2010)., 2010).
 In these events. the excess of optical flux with respect to the extrapolation of the earlier light-curve appears because the standard afterglow emission is decaying as a power-law while the late prompt component ts still constant or slightly increasing and therefore it becomes dominant at late times.," In these events, the excess of optical flux with respect to the extrapolation of the earlier light-curve appears because the standard afterglow emission is decaying as a power-law while the late prompt component is still constant or slightly increasing and therefore it becomes dominant at late times."
 Because the two components originate in completely different mechanisms. a colour evolution in the optical-NIR SED during the bump is expected. similar to the one obtained for GRB 081029.," Because the two components originate in completely different mechanisms, a colour evolution in the optical-NIR SED during the bump is expected, similar to the one obtained for GRB 081029."
 A colour change related to an increasingly dominant second component is indeed observed both in GRB 071003 and in GRB 061126 (Ghisellini et al., A colour change related to an increasingly dominant second component is indeed observed both in GRB 071003 and in GRB 061126 (Ghisellini et al.
 2009. Nardini et al.," 2009, Nardini et al."
 2010)., 2010).
 On the other hand. in the late prompt model (Ghisellini et al.," On the other hand, in the late prompt model (Ghisellini et al."
" 2007). the early-time flat (or slowly increasing) evolution of the late prompt component is caused by a geometrical effect (1... increase of the visible emitting surface owing to a decrease of the Γ lorentz factor of the slower “late prompt"" shells)."," 2007), the early-time flat (or slowly increasing) evolution of the late prompt component is caused by a geometrical effect (i.e., increase of the visible emitting surface owing to a decrease of the $\Gamma$ lorentz factor of the slower “late prompt” shells)."
 For GRB 081029 the rise of the optical rebrightening is extremely steep which renders this scenario unlikely unless a delayed reactivation of the central engine ts responsible for the emission of these less energetic shells., For GRB 081029 the rise of the optical rebrightening is extremely steep which renders this scenario unlikely unless a delayed reactivation of the central engine is responsible for the emission of these less energetic shells.
 This delayed starting point would allow a new definition of the late prompt component starting time and consequently a flattening of this component's light-curve rise., This delayed starting point would allow a new definition of the late prompt component starting time and consequently a flattening of this component's light-curve rise.
 Another scenario that invokes the coexistence of two separate emitting processes. and therefore allows the observation of a discontinuity m the light-curve with a contemporaneous colour evolution. 1s the two-component jet model.," Another scenario that invokes the coexistence of two separate emitting processes, and therefore allows the observation of a discontinuity in the light-curve with a contemporaneous colour evolution, is the two-component jet model."
 The presence of a fast narrow jet. co-alligned with a wider. slower jet has been claimed to be able to deseribe the complex broad-band evolution of some GRBs (e.g.. GRB 030329 Berger et al.:," The presence of a fast narrow jet, co-alligned with a wider, slower jet has been claimed to be able to describe the complex broad-band evolution of some GRBs (e.g., GRB 030329 Berger et al.;"
 2003. GRB 080319B Racusin et al.," 2003, GRB 080319B Racusin et al."
 2009; GRB 080413B Filgas et al., 2009; GRB 080413B Filgas et al.
 2010)., 2010).
 In this scenario. a flattening or a bump in the observed light-curve can be observed when the onset of the afterglow produced by the slower (and wider) jet occurs when the afterglow emission of the faster (and narrower) Jet Is already decreasing fast because of an early-time jet break.," In this scenario, a flattening or a bump in the observed light-curve can be observed when the onset of the afterglow produced by the slower (and wider) jet occurs when the afterglow emission of the faster (and narrower) jet is already decreasing fast because of an early-time jet break."
 The rising wider jet afterglow can be characterised by a different SED in the observed bands and therefore à colour evolution can be observed during the transition., The rising wider jet afterglow can be characterised by a different SED in the observed bands and therefore a colour evolution can be observed during the transition.
 At later times the observed radiation is dominated by the wider Jet afterglow., At later times the observed radiation is dominated by the wider jet afterglow.
 The complexity of the GRB 081029 light-curve is hard to explain in the framework of this scenario., The complexity of the GRB 081029 light-curve is hard to explain in the framework of this scenario.
 The flat early-time optical light-curve observed before the rebrightening. when interpreted in the framework of a standard forward shock afterglow scenario. suggests an environment characterised by a wind-like density profile.," The flat early-time optical light-curve observed before the rebrightening, when interpreted in the framework of a standard forward shock afterglow scenario, suggests an environment characterised by a wind-like density profile."
 On the other hand. to obtain a steep rebrighening at later times. the second jet afterglow light-curve requires an ISM-like external medium.," On the other hand, to obtain a steep rebrighening at later times, the second jet afterglow light-curve requires an ISM-like external medium."
 Moreover. the optical rise that is required for two separate components is much steeper than the pre-peak rise previously observed in other GRBs (Rykoff et al.," Moreover, the optical rise that is required for two separate components is much steeper than the pre-peak rise previously observed in other GRBs (Rykoff et al."
 2009; Panaitescu Vestrand 2008)., 2009; Panaitescu Vestrand 2008).
 In the standard forward shock emission scenario for an Isotropic outflow seen face-on by the observer. during the pre-deceleration phase the optical light-curve is supposed to rise as Fx£ or P? depending on the location of the cooling frequency.," In the standard forward shock emission scenario for an isotropic outflow seen face-on by the observer, during the pre-deceleration phase the optical light-curve is supposed to rise as $F\propto t^2$ or $t^3$ depending on the location of the cooling frequency."
 A steeper Foofr! rise can be obtained for of an off-axis observer for a collimated outflow with a sharp angular boundary (Panaitescu Vestrand 2008)., A steeper $F\propto t^4$ rise can be obtained for of an off-axis observer for a collimated outflow with a sharp angular boundary (Panaitescu Vestrand 2008).
 No steeper rise. similar to the one observed for GRB 081029 can be reproduced under these standard assumptions.," No steeper rise, similar to the one observed for GRB 081029 can be reproduced under these standard assumptions."
 Panaitescu Vestrand (2008) also found that the afterglows showing a fast pre-peak rise ία.«- are characterised by a correlation between the K-corrected optical flux at the peak and the redshift-corrected peak time assuming a common redshift z=2., Panaitescu Vestrand (2008) also found that the afterglows showing a fast pre-peak rise $\alpha < -1$ ) are characterised by a correlation between the $k$ -corrected optical flux at the peak and the redshift-corrected peak time assuming a common redshift $z=2$ .
 Considering only the contribution of the second component corrected for the Galactic foreground absorption (1.e.. the triple power-law described in section 3)). the peak flux calculated in the r-band re-locating GRB 081029 at z=2 is FO=1.3 my.," Considering only the contribution of the second component corrected for the Galactic foreground absorption (i.e., the triple power-law described in section \ref{lc}) ), the peak flux calculated in the r-band re-locating GRB 081029 at $z=2$ is $F_{\rm p}^{z=2}=1.3$ mJy."
 The peak time at redshift 2 is Ko=3.5 ks., The peak time at redshift 2 is $t_{\rm p}^{\rm z=2}=3.5$ ks.
 Comparing these values with the plot shown in Fig., Comparing these values with the plot shown in Fig.
 2 of Panaitescu Vestrand (2008). the rebrightening of GRB 081029 is inconsistent with this correlation with GRB 081029 having a peak flux that is much brighter than expected for the peak time.," 2 of Panaitescu Vestrand (2008), the rebrightening of GRB 081029 is inconsistent with this correlation with GRB 081029 having a peak flux that is much brighter than expected for the peak time."
 That confirms the difficulties of explaining the fast optical bump of GRB 081029 as the signature of a second jet afterglow with the same initial time as the one responsible for the early-time light-curve., That confirms the difficulties of explaining the fast optical bump of GRB 081029 as the signature of a second jet afterglow with the same initial time as the one responsible for the early-time light-curve.
 An alternative scenario invoked for explaining late-time optical bumps is a discrete episode of energy injection into. the fireball by the late-time interaction. of slow shells with the forward shock (Jóhhannesson et al., An alternative scenario invoked for explaining late-time optical bumps is a discrete episode of energy injection into the fireball by the late-time interaction of slow shells with the forward shock (Jóhhannesson et al.
 2006: Fan Piran 2006: Covino et al., 2006; Fan Piran 2006; Covino et al.
 2008b: Rossi et al.," 2008b; Rossi et al.,"
 2011)., 2011).
 In GRB 081029. the steepness of the optical bump and the large ratio between the second componet and the extrapolation of the early broken power-law decay (between ~ 5 and 8 during the bump) imply a sudden release of a large amount of energy at late times.," In GRB 081029, the steepness of the optical bump and the large ratio between the second component and the extrapolation of the early broken power-law decay (between $\sim$ 5 and 8 during the bump) imply a sudden release of a large amount of energy at late times."
 Such a energy injection Is not supposed to produce a change in the observed synchrotron spectral slopes under the standard assumption of non-evolving micro-physical parameters in the fireball (e.g.. &. p. εν) Under this assumption. an observed spectral evolution can only be produced if one of the characteristic frequencies (e.g.. the cooling frequency v) crosses the observed bands.," Such an energy injection is not supposed to produce a change in the observed synchrotron spectral slopes under the standard assumption of non-evolving micro-physical parameters in the fireball (e.g., $\epsilon_{\rm e}$, $p$, $\epsilon_{\rm b}$ ) Under this assumption, an observed spectral evolution can only be produced if one of the characteristic frequencies (e.g., the cooling frequency $\nu_c$ ) crosses the observed bands."
 On the other hand. we know that no contemporaneous rebrightening ts observed in the X-rays.," On the other hand, we know that no contemporaneous rebrightening is observed in the X-rays."
 This evidence acts against this scenario because an episode of energy injection would merease the normalisatior of the whole synchrotron spectrum. which would be visible at all wavelengths.," This evidence acts against this scenario because an episode of energy injection would increase the normalisation of the whole synchrotron spectrum, which would be visible at all wavelengths."
 The complex light-curve of GRB 081029 is a remarkable example of how the increasing quality of optical-NIR follow- represents a hard test for the proposed afterglow emission, The complex light-curve of GRB 081029 is a remarkable example of how the increasing quality of optical-NIR follow-up represents a hard test for the proposed afterglow emission
"reffig:iraccnts shows the number counts of galaxies at 3.6, 4.5, 5.8 and8.","\\ref{fig:iraccnts} shows the number counts of galaxies at 3.6, 4.5, 5.8 and."
"0um.. The predictions of our model are compared to data from Fazio et ((2004), who combine observations from three regions ranging from an ultradeep 5’x10' field to the wide 3°x Boóttes field, in order to achieve good sampling of both the faint and bright end of the counts."," The predictions of our model are compared to data from Fazio et (2004), who combine observations from three regions ranging from an ultradeep $5\arcmin\times10\arcmin$ field to the wide $3^{\circ}\times3^{\circ}$ Boöttes field, in order to achieve good sampling of both the faint and bright end of the counts."
" We also compare to counts of GOODS-N sources by using SExtractor on the public imaging of the that field (Dickinson et al.,"," We also compare to counts of GOODS-N sources by using SExtractor on the public imaging of the that field (Dickinson et al.,"
" in prep.),"," in prep.),"
 which agree well with Fazio et al., which agree well with Fazio et al.
’s measurements.,'s measurements.
 In both datasets stars have been excluded., In both datasets stars have been excluded.
" Broadly speaking, the agreement between the model and the data is good."," Broadly speaking, the agreement between the model and the data is good."
" Although the model appears to overpredict the counts at the faint end in panels (a)-(c), the data here suffer from confusion."," Although the model appears to overpredict the counts at the faint end in panels (a)–(c), the data here suffer from confusion."
" The fit at iis poorer than in the other bands, possibly due to our somewhat unsophisticated modelling of the PAH features, which may have a more noticeable effect at"," The fit at is poorer than in the other bands, possibly due to our somewhat unsophisticated modelling of the PAH features, which may have a more noticeable effect at"
"aud respectively, where i=Vv Ll. fj denotes the tine of observation. and aj; and a7. correspondingly. the individual aud mean values of the brightucss.","and respectively, where $\mathrm{i}=\sqrt{-1}$ , $t_j$ denotes the time of observation, and $m_j$ and $\overline{m}$, correspondingly, the individual and mean values of the brightness."
 Fie., Fig.
 laa showsthe power spectrum of the V. band data (linear trend removed) for the period from 1983 to 1995., \ref{fig:1}a a showsthe power spectrum of the $V$ band data (linear trend removed) for the period from 1983 to 1998.
 The largest peak in the spectrum: corresponds to zm1100 davs ane is readily explainable: a wave with that characteristic ine scale is casily seen in Fie. 2..," The largest peak in the spectrum corresponds to $\approx 1400$ days and is readily explainable: a wave with that characteristic time scale is easily seen in Fig. \ref{fig:2},"
 superimposed on a sinooth decay in brightuess. in the optical aud infrared light curves (see also Larionov aud Clark et al. 19991).," superimposed on a smooth decay in brightness, in the optical and infrared light curves (see also Larionov \cite{L93} and Clark et al. \cite{C99}) )."
" Iun the region of the orbital period (~0.01 dot). only a faint peak corresponding to P=103"" is seen."," In the region of the orbital period $\sim 0.01~\mathrm{d}^{-1}$ ), only a faint peak corresponding to $P=103^\mathrm{d}$ is seen."
 At the frequency corresponding to a 111 day period. there is no peak at all.," At the frequency corresponding to a 111 day period, there is no peak at all."
" Iu order to judge with sufficicut confidence the reality of the orbital plus close and/or related periodicities. one should remove the louger-period componeut(s) of variability,"," In order to judge with sufficient confidence the reality of the orbital plus close and/or related periodicities, one should remove the longer-period component(s) of variability."
 Iu our case it is not sufficient to subtract the iuear trend. aud we need to take into account the slow lieht variations with a (quasi)period of z1100 davs.," In our case it is not sufficient to subtract the linear trend, and we need to take into account the slow light variations with a (quasi)period of $\approx 1400$ days."
 We lave constructed an approximating data set using the nethod of a sliding mean with the window value A. replacing raw data i; for each time f; by the weighted luean: where & is the number of data poiuts within interval A.| Al. and the weielt of jtl poiut is determuned as where of; is the time. span from the jth point to the centerof the window.," We have constructed an approximating data set using the method of a sliding mean with the window value $\Delta
$, replacing raw data $m_i$ for each time $t_i$ by the weighted mean: where $k$ is the number of data points within interval $ [t_i-\Delta, t_i+\Delta]$ , and the weight of $j$ th point is determined as where $\delta t_j$ is the time span from the $j$ th point to the centerof the window."
" The optimal value of the smoothing interval was searched by trial and error within a rauge D0AxLOO"".", The optimal value of the smoothing interval was searched by trial and error within a range $20^\mathrm{d}< \Delta < 100^\mathrm{d}$.
 The criterion for A sclection was the signal-to-noise ratio of the power-spectriun peaks in the region of interest. ic. around the orbital frequency.," The criterion for $\Delta$ selection was the signal-to-noise ratio of the power-spectrum peaks in the region of interest, i.e. around the orbital frequency."
 The raw and smoothed light curves. and also the residuals vetween them. are plotted in Fie. 2..," The raw and smoothed light curves, and also the residuals between them, are plotted in Fig. \ref{fig:2}."
 The power spectra. or the residuals injm is shown in Fie. 3.., The power spectrum for the residuals $m_j-m_j^{\prime}$ is shown in Fig. \ref{fig:3}.
" We have fouud hat. iudepeudeut of the value of the sinoothing interval adopted. the most prominent peal im the power spectra. corresponds to a period of 102383+(1, while the low-requencey coniponeuts are most effectively suppressed and he highest signal-to-noise ratio is obtaiuec when A=50"","," We have found that, independent of the value of the smoothing interval adopted, the most prominent peak in the power spectrum corresponds to a period of $102\fd 83 \pm 0\fd 1$, while the low-frequency components are most effectively suppressed and the highest signal-to-noise ratio is obtained when $\Delta =
50^\mathrm{d}$."
 We lave obtained an analogous result for the B. baud as well., We have obtained an analogous result for the $B$ band as well.
 The amplitude of the sine-wave obtained is (7015 in V and (01 in D. while in © the 103-av peak is not seen above the noise.," The amplitude of the sine-wave obtained is $0\fm 015$ in $V$ and $0\fm 01$ in $B$, while in $U$ the 103-day peak is not seen above the noise."
 The significance of he detection of the faint sine-wave superimposed on the large-zuuplitude slow variability cau be tested by adding au artificial wuplitude harmonic to the raw data., The significance of the detection of the faint sine-wave superimposed on the large-amplitude slow variability can be tested by adding an artificial low-amplitude harmonic to the raw data.
 We added a sne- with P=TG and amplitude (9015 to the raw V data. then repeated the same cleaning routine as described above (Eqs.," We added a sine-wave with $P=76^\mathrm{d}$ and amplitude $0\fm 015$ to the raw $V$ data, then repeated the same cleaning routine as described above (Eqs."
 3 and lj).," \ref{eq:1}
 and \ref{eq:2}) )."
 Fig., Fig.
 3 demonstrates hat the low-frequency filtration method described above can be usedto detect. discussed. poxiodic variability on a time scale of ~1007 with amplitude Qu].," \ref{fig:3} demonstrates that the low-frequency filtration method described above can be usedto detect, , periodic variability on a time scale of $\sim 100^\mathrm{d}$ with amplitude $\ge 0\fm 01$ ."
 Iu order to test whether the 103-day period is real or an artifact iutriusic to some part of the data set. we have split the initial data set mto two parts corresponding to," In order to test whether the 103-day period is real or an artifact intrinsic to some part of the data set, we have split the initial data set into two parts corresponding to"
de Blok 2003. and references therein).,"de Blok 2003, and references therein)."
 In the present paper. we use a suite of N-body simulations to study how halo density profiles change as a function of dark energy equation ol state. discuss our results in the context of analytic models. and discuss the observational implications of dark energy on galaxy scales.," In the present paper, we use a suite of N-body simulations to study how halo density profiles change as a function of dark energy equation of state, discuss our results in the context of analytic models, and discuss the observational implications of dark energy on galaxy scales."
 The existence of some form of dark energy is supported by a preponderance of data., The existence of some form of dark energy is supported by a preponderance of data.
 Taken together. observations of the magnitude-redshift. relation of type Ia. supernovae SNla: Perlmutter 1999: Riess 2001: Ixnop 2008: Barris 2003). the power spectrum. of cosmic microwave background (CALB) anisotropy (Spergel 2003: Tegmark 2003a) the power spectrum of galaxy clustering (Dodelson 2003: Tegmark 03b). and the luminosity function and. barvon fraction of clusters (Allen 2003). provide. nearly unimpeachable evidence For the existence of dark energy.," Taken together, observations of the magnitude-redshift relation of type Ia supernovae (SNIa; Perlmutter 1999; Riess 2001; Knop 2003; Barris 2003), the power spectrum of cosmic microwave background (CMB) anisotropy (Spergel 2003; Tegmark 2003a), the power spectrum of galaxy clustering (Dodelson 2003; Tegmark 2003b), and the luminosity function and baryon fraction of clusters (Allen 2003) provide nearly unimpeachable evidence for the existence of dark energy."
 Phe most common supposition is that the dark energy takes the form of a cosmological constant or vacuum energy., The most common supposition is that the dark energy takes the form of a cosmological constant or vacuum energy.
 In this case. the enerey density p. and pressure are related throughp=p.," In this case, the energy density $\rho$, and pressure $p$, are related through $p =
-\rho$."
 An attractive alternative candidatep. for the dark energy. is the potential energy. of a slowly-varving scalar field: ©. or “quintessence” Ratra Peebles 1988: C'aldwell. Dave.. Steinhardt 1998).," An attractive alternative candidate for the dark energy is the potential energy of a slowly-varying scalar field $\phi$, or “quintessence” Ratra Peebles 1988; Caldwell, Davé,, Steinhardt 1998)."
 A convenient parametrization of the dark energy. is through an equation of state we—pop. relating its energy density and. pressure., A convenient parametrization of the dark energy is through an equation of state $\wQ \equiv p_{\phi}/\rho_{\phi}$ relating its energy density and pressure.
 In. general. the equation of state parameter dw is time-varving. but it is useful to. moclel quintessence with a constant equation of state parameter because current observational data sets have limited: power to distinguish between a time-varving and constant equation of state Ixujat. et al.," In general, the equation of state parameter $\wQ$ is time-varying, but it is useful to model quintessence with a constant equation of state parameter because current observational data sets have limited power to distinguish between a time-varying and constant equation of state Kujat et al."
 2002)., 2002).
 Useful limits on the equation of state for the clark energy. assuming that it remains constant in time. come from SNla studies. 1.67<wos0.62 (2o: Ixnop 2003). and can be refined by combining SNla data with CMD anisotropy anc galaxy clustering statistics vielding 1.33<ww«0.79 at 2o Cleemark 2003b).," Useful limits on the equation of state for the dark energy, assuming that it remains constant in time, come from SNIa studies, $-1.67 < \wQ < -0.62$ $2\sigma$; Knop 2003), and can be refined by combining SNIa data with CMB anisotropy and galaxy clustering statistics yielding $-1.33 < \wQ < -0.79$ at $2\sigma$ (Tegmark 2003b)."
 For our simulations. we adopt an empirical view ancl study five models with constant a. that span a comparably large range of parameter space: we15.125.LO.0.75.0.5.," For our simulations, we adopt an empirical view and study five models with constant $\wQ$, that span a comparably large range of parameter space: $\wQ = -1.5, -1.25, -1.0, -0.75, -0.5$."
" Our theoretical understanding of halo profiles has improved recently largely. through numerical simulations. performed inthe context ofCDM plus cosmological constant (C XCDALD or standard CDM (SC'DAL. Ox,=1. Mo= 0) cosmologics (Navarro. Frenk. White. 1995. 1996. 1997. hereafter NEW: Ixravtsov. IxIvpin. Ixhokhlov 1997: Ghigna 1998: Jing 2000: Bullock 2001: Eke. Navarro. Steinmetz 2001: Wechsler et al."," Our theoretical understanding of halo profiles has improved recently largely through numerical simulations, performed inthe context of CDM plus cosmological constant $\Lambda$ CDM) or standard CDM (SCDM, $\Omega_{\rm M} = 1$, $\Omega_{\rm Q}=0$ ) cosmologies (Navarro, Frenk, White, 1995, 1996, 1997, hereafter NFW; Kravtsov, Klypin, Khokhlov 1997; Ghigna 1998; Jing 2000; Bullock 2001; Eke, Navarro, Steinmetz 2001; Wechsler et al."
 2002. hereafter. W02: Zhao et al.," 2002, hereafter W02; Zhao et al."
 2003: Havashi 2003: Navarro 2003: for a review. see Primack 2003).," 2003; Hayashi 2003; Navarro 2003; for a review, see Primack 2003)."
 ]t is generally. understood that the final density profiles of haloes are linked. closely to their formation histories., It is generally understood that the final density profiles of haloes are linked closely to their formation histories.
 Halo central densities are set. during an carly. rapid-accretion phase. and tend to be proportional to the density of haloes in the Universe at the time of this rapid. collapse (02: Zhao ct al.," Halo central densities are set during an early, rapid-accretion phase, and tend to be proportional to the density of haloes in the Universe at the time of this rapid collapse (W02; Zhao et al."
 2003: Tasitsiomi et al., 2003; Tasitsiomi et al.
 2003)., 2003).
 Larger values of w dead to earlier. collapse times and also to more rapid collapse of overdensities.thus we expect haloes with higher relative central densities and higher virial densities (see our discussions in 2 and 3)).," Larger values of $\wQ$ lead to earlier collapse times and also to more rapid collapse of overdensities,thus we expect haloes with higher relative central densities and higher virial densities (see our discussions in \ref{sec:Qstruc} and \ref{sec:halostruc}) )."
 In 4-6 we quantify these effects. and. test the expected scalings.," In \ref{sec:sims}- \ref{sec:DeltaVby2}
 we quantify these effects and test the expected scalings."
 We verify that the analytic technique of Bullock (2001. BOL) for predicting halo concentrations works well when applied cdirectlv to constant « cosmologies. implving that it is fairly generally applicable.," We verify that the analytic technique of Bullock (2001, B01) for predicting halo concentrations works well when applied directly to constant $\wQ$ cosmologies, implying that it is fairly generally applicable."
 AC velatecl issue ds the cllect of dark οποιον on the halo mass function., A related issue is the effect of dark energy on the halo mass function.
 Although it is not directly observable. theoretical insight into the expected halo mass function is attainable through current N-body simulations.," Although it is not directly observable, theoretical insight into the expected halo mass function is attainable through current N-body simulations."
 The effects. of dark. energv on halo mass functions have been investigated by Dartelmann. Perota. Baccigalupi (2003). Linder Jenkins (2003). Whypin (2003). Alaccio (2008). and Lokas. Bocle. llolfiman (2003) and the accurate prediction of halo mass [functions as well as accretion histories and. density structures over a [arge range of redshifts is necessary in order to take full advantage of the ability of upcoming cluster surveys. such as the Sunvaev-Zeldovich Array //astro.uchicago.edu/sza/)}). to constrain the clark energy equation of state (see Carlstrom. Holder. Reese 2001 for a review).," The effects of dark energy on halo mass functions have been investigated by Bartelmann, Perota, Baccigalupi (2003), Linder Jenkins (2003), Klypin (2003), Macciò (2003), and okas, Bode, Hoffman (2003) and the accurate prediction of halo mass functions as well as accretion histories and density structures over a large range of redshifts is necessary in order to take full advantage of the ability of upcoming cluster surveys, such as the Sunyaev-Zeldovich Array ), to constrain the dark energy equation of state (see Carlstrom, Holder, Reese 2001 for a review)."
" There appears to be general agreement wt halo mass functions can be approximated: accurately bv the ""universal"". mass function of Jenkins (2001. hereafter JOL) even in models with dark energy."," There appears to be general agreement that halo mass functions can be approximated accurately by the “universal” mass function of Jenkins (2001, hereafter J01) even in models with dark energy."
 In 5.1. we confirm the results of previous studies of quintessence cosmologics with w21. and extend this agreement to models with w<1.," In \ref{sec:mf}, we confirm the results of previous studies of quintessence cosmologies with $\wQ > -1$, and extend this agreement to models with $\wQ < -1$."
 Ixlvpin (2008) and Dolag (2003). have performed. previous numerical studies of CDAL haloes in quintessence cosmologies., Klypin (2003) and Dolag (2003) have performed previous numerical studies of CDM haloes in quintessence cosmologies.
 Where our study. overlaps with these. our results are generally in agreement.," Where our study overlaps with these, our results are generally in agreement."
 This study extends. complements. and improves upon previous studies in several wavs.," This study extends, complements, and improves upon previous studies in several ways."
 We have investigated. the ellects of normalizing thepower spectrum using clusters vs. CMD. which vield quite different results.," We have investigated the effects of normalizing thepower spectrum using clusters vs. CMB, which yield quite different results."
 We have assembled large catalogues of haloes. with masses and concentrations [or more than 1600 haloes with more than LOO particles each in cach simulation.," We have assembled large catalogues of haloes, with masses and concentrations for more than $1600$ haloes with more than $100$ particles each in each simulation."
 With this large number of haloes we are able to measure the distribution of concentrations at. [fixed mass and test the predictions of the BOL model for both the mean relation between concentration and mass and the halo-to-halo scatter., With this large number of haloes we are able to measure the distribution of concentrations at fixed mass and test the predictions of the B01 model for both the mean relation between concentration and mass and the halo-to-halo scatter.
 We also extend. the range of quintessence parameter space probed. by simulations by exploring two models with w< 1., We also extend the range of quintessence parameter space probed by simulations by exploring two models with $\wQ<-1$ .
 The format of this paper is as follows., The format of this paper is as follows.
 Alter a brief review of the basies of structure formation in models with dark energv in 2.. wegoonin 3 to describe modifications to the BOL model for halo concentrations in. quintessence cosmologies.," After a brief review of the basics of structure formation in models with dark energy in \ref{sec:Qstruc}, , we go on in \ref{sec:halostruc} to describe modifications to the B01 model for halo concentrations in quintessence cosmologies."
 In d... we describe the set of numerical simulations that we performed.," In \ref{sec:sims}, we describe the set of numerical simulations that we performed."
 In 5.2 we present the mass functions ancl concentrations of CDM haloes in these simulations and compare them to the results of the analytic DOLI model.," In \ref{sec:nbody_results}, we present the mass functions and concentrations of CDM haloes in these simulations and compare them to the results of the analytic B01 model."
 In 6.2 we briclly discuss the central density problem of CDM in light of our results.," In \ref{sec:DeltaVby2}, we briefly discuss the central density problem of CDM in light of our results."
 In 7.2 we present a summary of our findings. highlight improvements. anc dilferences to previous studies. and discuss the implications of this work.," In \ref{sec:conclusions}, we present a summary of our findings, highlight improvements and differences to previous studies, and discuss the implications of this work."
 Throughout this paper. we assume a [lat universe with present matter density relative to critical of Oi;= 0.3. and Llubble parameter 7— 0.7.," Throughout this paper, we assume a flat universe with present matter density relative to critical of $\Omega_{\rm M}=0.3$ , and Hubble parameter $h=0.7$ ."
" The quintessence energy density is then £X,=1.Ox, 0.7.", The quintessence energy density is then $\Omega_{\rm Q} = 1 - \Omega_{\rm M} = 0.7$ .
 We refer to à model witha cosmologicalconstant (w= 1) as CDM., We refer to a model witha cosmologicalconstant $\wQ=-1$ ) as $\Lambda$ CDM.
ealaxies should have experienced a last elliplical-elliplical merger (dry. merger) (?)..,galaxies should have experienced a last elliptical-elliptical merger (dry merger) \citep{2006ApJ...636L..81N}.
 Mixed nergers have not vet. been studied in details. allhough. according to ? they should be nore frequent than drv mergers (see however ο).," Mixed mergers have not yet been studied in details, although, according to \citet{2003ApJ...597L.117K} they should be more frequent than dry mergers (see however \citealp{2007arXiv0706.1243H}) )."
 Dry mergers and (heir implications Lor the formation of the red galaxy sequence have however received a lot of attention recently 299?)," Dry mergers and their implications for the formation of the red galaxy sequence have however received a lot of attention recently \citep{2005MNRAS.359.1379K,2005MNRAS.361.1043G,2005astro.ph..6044F,
2006ApJ...636L..81N,2006MNRAS.369.1081B}."
 Further refinement of theoretical models has recently been achieved by including physics in simulations of galaxy mergers and elliptical galaxy evolution., Further refinement of theoretical models has recently been achieved by including black-hole physics in simulations of galaxy mergers and elliptical galaxy evolution.
 Energetic eedback from central black holes might solve some pending problems of major merger nodels. like the suppression of late inflow of cold gas and star formation that would make ellipticals look much bluer than observed (?)..," Energetic feedback from central black holes might solve some pending problems of major merger models, like the suppression of late inflow of cold gas and star formation that would make ellipticals look much bluer than observed \citep{2005ApJ...620L..79S}."
 In summary. despite several still unsolved questions (e.g. ?2)). the major merger scenario has become a popular model in order to explain the origin of bulge-dominated. spheroidal galaxies (see e.g. ??)).," In summary, despite several still unsolved questions (e.g. \citealp{2007astro.ph..2535N}) ), the major merger scenario has become a popular model in order to explain the origin of bulge-dominated, spheroidal galaxies (see e.g. \citealp{2007arXiv0706.1246H,2007arXiv0706.1243H}) )."
 Progress in understanding galactic evolution is often driven bv strong interactions between observers and theorists., Progress in understanding galactic evolution is often driven by strong interactions between observers and theorists.
 Increasingly more sophisticated theoretical/numnerical models are confronted with contimously improving observations (hat lead to new theoretical challenges., Increasingly more sophisticated theoretical/numerical models are confronted with continuously improving observations that lead to new theoretical challenges.
 One example is the SAURON project (2)— which aims to determine the 2-dimensional structural and kinematical properties of early-iwpe galaxies using a panoramic liekl spectrograph., One example is the SAURON project \citep{2004MNRAS.352..721E} which aims to determine the 2-dimensional structural and kinematical properties of early-type galaxies using a panoramic integral-field spectrograph.
 In order (o interpret. the observations and study (he intrinsic galaxy structure. axisvinmetric Schwarzschild models are applied to the observations (?7)..," In order to interpret the observations and study the intrinsic galaxy structure, axisymmetric Schwarzschild models are applied to the observations \citep{2006MNRAS.366.1126C,2007MNRAS.379..418C}."
 The results. published so far. have revealed interesting fine structures and plywsical properties which provide new and deeper insisht into (he origin of galaxies. in particular when compared io simulation (??)..," The results, published so far, have revealed interesting fine structures and physical properties which provide new and deeper insight into the origin of galaxies, in particular when compared to simulation \citep{2000MNRAS.316..315B,2007MNRAS.376..997J}."
 In this paper we confront a recently published SAURON analvsis of prelerentially axisvimmnmnetrie elliptical galaxies (7) wilh the predictions of numerical merger simulations and cosmological models of galaxy formation., In this paper we confront a recently published SAURON analysis of preferentially axisymmetric elliptical galaxies \citep{2007MNRAS.379..418C} with the predictions of numerical merger simulations and cosmological models of galaxy formation.
 Section 2 summarizes the observations., Section 2 summarizes the observations.
 Section 3 shows that. simulations of collisionless and gaseous disk-disk mergers. neglecting star formation and stellar feedback cannot reproduce the observational results.," Section 3 shows that simulations of collisionless and gaseous disk-disk mergers, neglecting star formation and stellar feedback cannot reproduce the observational results."
 We demonstrate that star formation and stellar energv. feedback has a strong effect on the final structure of merger remnants. leading to a good agreement with the SAURON observations of Last rotating ellipticals.," We demonstrate that star formation and stellar energy feedback has a strong effect on the final structure of merger remnants, leading to a good agreement with the SAURON observations of fast rotating ellipticals."
 The origin of the round. almost isotropic. slowly rotating SAURON elliplicals is explored in section 4.," The origin of the round, almost isotropic, slowly rotating SAURON ellipticals is explored in section 4."
 Isolated. dry mergers of ellipticals that formed as discussed in seclion 3 cannot explain these objects.," Isolated, dry mergers of ellipticals that formed as discussed in section 3 cannot explain these objects."
 We show however (hat cosmological initial conditions. leading to a series of multiple major and minor mergers. coupled with local star bursts eenerale spheroidal stellar svstems in very good agreement with the observations.," We show however that cosmological initial conditions, leading to a series of multiple major and minor mergers, coupled with local star bursts generate spheroidal stellar systems in very good agreement with the observations."
 Conclusions follow in section 5., Conclusions follow in section 5.
 Either the compation WD is clestructecd aud foris a disk around the hot core. or the more assisre core is destructed (because it is hot aud large) aud forms an accretion disk around the WD.," Either the companion WD is destructed and forms a disk around the hot core, or the more massive core is destructed (because it is hot and large) and forms an accretion disk around the WD."
 Iu the later case the lower mass WD lias a shallower gravitational poteutial. hence the temperature in the accretion disk does not reacl ignition temperature.," In the later case the lower mass WD has a shallower gravitational potential, hence the temperature in the accretion disk does not reach ignition temperature."
 ashi&Soker(2011) termed this the core-degenerate (CD) scenario., \citet{KashiSoker2011} termed this the core-degenerate (CD) scenario.
 Iu tlje CD scenario the merger occurs ina CE with a massive ACB star. or shortly (<10? vr) alter the CE phase. re. during the planetary nebula phase.," In the CD scenario the merger occurs in a CE with a massive AGB star, or shortly $\la 10^5 \yr$ ) after the CE phase, i.e. during the planetary nebula phase."
 Because oL its rapid rotation the super-Chaucdraseklar WD coes not explode (e.g. Anand1965: Ostrikerlenheimer 1965: Uenishietal. 2003: Yoon&Langer 2005)).," Because of its rapid rotation the super-Chandrasekhar WD does not explode (e.g., \citealt{Anand1965}; ; \citealt{Ostriker1968}; \citealt{Uenishi2003}; \citealt{Yoon2005}) )."
 It will explode long after the mereer. only after it has spun down to allow explosion.," It will explode long after the merger, only after it has spun down to allow explosion."
 The CD scenario. inclucing fiucings frou this paper. is sumunarized sciematically iu Figure l..," The CD scenario, including findings from this paper, is summarized schematically in Figure \ref{fig:fig1}."
 This sceuario is completely differeut. from the moclel of au explosion inside a CE with a low mass evolved star as proposed by that was subsequently criticized by Applegate(1901)., This scenario is completely different from the model of an explosion inside a CE with a low mass evolved star as proposed by \citet{Hachisuetal1989} that was subsequently criticized by \citet{Applegate1991}.
. There are three key ingredieus in the CD scenario. in addition t«» the common condition that the reunant mass be Z 1.1A£..(," There are three key ingredients in the CD scenario, in addition to the common condition that the remnant mass be $\ga 1.4 M_\odot$. ("
1) The hot core is more massive than the companion cold WD.,1) The hot core is more massive than the companion cold WD.
 Iu a [utu'e paper we will conduct a »opulation svuthesis to find the expected nmuuber of such systems. (, In a future paper we will conduct a population synthesis to find the expected number of such systems. (
2) TIe merger should occur whie the core is still large. heuce hot.,"2) The merger should occur while the core is still large, hence hot."
" This limits tlie mereer to occur witLi IO""vr alter the conumor envelope phase.", This limits the merger to occur within $\sim 10^5 \yr$ after the common envelope phase.
 Ixashi&Soker(2011) showed that this condition cau be met when the AGB star is iuassive. and some of the ejected CE gas falls back. (," \citet{KashiSoker2011} showed that this condition can be met when the AGB star is massive, and some of the ejected CE gas falls back. ("
3 The delay between mereer aud explosion slould be up to 10Mvr if this scenario is to account [or SNe Ia in old-stellar populations. iu aclelition to SNe Ia in voung stellar popuatious.,"3) The delay between merger and explosion should be up to $\sim 10^{10} \yr$ if this scenario is to account for SNe Ia in old-stellar populations, in addition to SNe Ia in young stellar populations."
 This is the subject of this p:iper., This is the subject of this paper.
 In section 2 we examine tlie spinuine-dowu time by gravitationa waves., In section \ref{sec:gradiation} we examine the spinning-down time by gravitational waves.
 Contrary to previous Claims. we find that this time is very |me. aud inieht not be relevau to spin clown WDs. despite beiug very efficieut in spinnine clown neutron stars (NS).," Contrary to previous claims, we find that this time is very long, and might not be relevant to spin down WDs, despite being very efficient in spinning down neutron stars (NS)."
" In section 3) we show that it is quite plausible that magneto-dipole radiation torque cau lead to a loug «elay of the explosion. up to 10!""vr."," In section \ref{sec:long} we show that it is quite plausible that magneto-dipole radiation torque can lead to a long delay of the explosion, up to $\sim10^{10} \yr$."
 In section. { we coustrain the properties of the core-WD merger product if the CD scenario is to account [or a large [fraction of SNe Ia. Our discussion aidsummary are in section 5.., In section \ref{sec:distribution} we constrain the properties of the core-WD merger product if the CD scenario is to account for a large fraction of SNe Ia. Our discussion andsummary are in section \ref{sec:summary}. .
morphology.,morphology.
 Hence. the presence of a stellar remnant inside the molecular cavily can give rise lo the observed 5-rav. emission.," Hence, the presence of a stellar remnant inside the molecular cavity can give rise to the observed $\gamma$ -ray emission."
 There are three stellar remnant candidates in the field ol view., There are three stellar remnant candidates in the field of view.
 They are a binary svstem GX340+0 containing a neutron star (Schultz 1993).. a magnetar (CXO J164710.2-455216). and the pulsar PSR. J1648-4611. with the third one the most likely precursor.," They are a binary system GX340+0 containing a neutron star \citep{schultz93}, a magnetar (CXO J164710.2-455216), and the pulsar PSR J1648-4611, with the third one the most likely precursor."
 The first two sources are seen geometrically inside the molecular cavity. but lie ad distances different from that of the cavity. according to their present distance estimations.," The first two sources are seen geometrically inside the molecular cavity, but lie at distances different from that of the cavity according to their present distance estimations."
 The magnetar is associated with Well following the analvsis by Munoetal.(2006)... based on statistical grounds.," The magnetar is associated with Wd1 following the analysis by \citet{muno06}, based on statistical grounds."
 However. it could as well be located futher out. in the molecular cavilv.," However, it could as well be located further out, in the molecular cavity."
 For the neutron star in the binary svstem GAX340+0. a maximum clistance of 1143 kpe has been estimated based on its Iuminosity. assuming (hat it is radiating at the LEddington limit (Penninxetal.1993).," For the neutron star in the binary system GX340+0, a maximum distance of $\pm$ 3 kpc has been estimated based on its luminosity, assuming that it is radiating at the Eddington limit \citep{penninx93}."
. Yet. it is conceivable that the neutron star is not emitting with the limiting luminosity. in which ease it could be located at the distance of our cavity.," Yet, it is conceivable that the neutron star is not emitting with the limiting luminosity, in which case it could be located at the distance of our cavity."
 A thirel stellar remnant (PSR J1648-4611) is located half a degree South-East of the ring center and is seen just al the outer border of the molecular ring (see Fig., A third stellar remnant (PSR J1648-4611) is located half a degree South-East of the ring center and is seen just at the outer border of the molecular ring (see Fig.
 4)., 4).
 A distance bv dispersion measure vields a value of 55.7 kpc. depending on the electron density distribution model adopted. (Torres&Nuza2003).," A distance by dispersion measure yields a value of 5–5.7 kpc, depending on the electron density distribution model adopted \citep{Torres03}."
. That is at the same clistance of the cavitv., That is at the same distance of the cavity.
 This makes this pulsar a strong candidate for producing the οταν emission at the inner boundaries of the molecular shell., This makes this pulsar a strong candidate for producing the $\gamma$ -ray emission at the inner boundaries of the molecular shell.
 The location of the pulsar ofl-center of the cavity. could be due to a large proper motion of the pulsar.," The location of the pulsar off-center of the cavity, could be due to a large proper motion of the pulsar."
 following the SN explosion., following the SN explosion.
 It requires only a velocity as little as 8 ! for the pulsar to move from the center of the shell to ils present position. if the explosion occurred 6 Myv ago. and the pulsar is moving in the plane of the sky.," It requires only a velocity as little as 8 $^{-1}$ for the pulsar to move from the center of the shell to its present position, if the explosion occurred 6 Myr ago, and the pulsar is moving in the plane of the sky."
 Recently. the LAT instrument on board of the Fermi satellite detected unpulsed," Recently, the LAT instrument on board of the Fermi satellite detected unpulsed"
The SERs. which are taken from the literature aud based ou a classical liner conversion of the IIo huninosity iuto a SER (cee.?).. are plotted iu dependence of the total galaxy neutral gas mass iu Fie. 2..,"The SFRs, which are taken from the literature and based on a classical linear conversion of the $\alpha$ luminosity into a SFR \citep[eg.][]{kennicutt1994a}, are plotted in dependence of the total galaxy neutral gas mass in Fig. \ref{fig_sfr_m_gas_imf}."
 With decreasing galaxy neutral eas mass the relation between the traclitionally calculated SER and the ealaxy ucutral gas mass becomes increasingly stecper., With decreasing galaxy neutral gas mass the relation between the traditionally calculated SFR and the galaxy neutral gas mass becomes increasingly steeper.
" Hore we find that for massive ealaxies of our sample with neutral eas masses τον10.M. the SER scales slightly non-linear with neutral eas mass. and scales aliiost quadraticallv for less 11assive ealaxies (Maa,5ν10* AL.). This behaviour is already well kuown (e.g.??).."," Here we find that for massive galaxies of our sample with neutral gas masses $\ge 5\times10^7\;M_\odot$ the SFR scales slightly non-linear with neutral gas mass, and scales almost quadratically for less massive galaxies $M_\mathrm{gas} \le 5\times 10^7\;M_\odot$ ), This behaviour is already well known \citep[e.g.][]{karachentsev2007a,kaisin2008a}."
 The nou-linear relation between the total IIa huninosity and the ποιο SER calculated with the IGIME-theorv diverges significantly from these classical linear relations based on a constant ealaxv-wide IMP below SERs typical for SMC-tvpe galaxies., The non-linear relation between the total $\alpha$ luminosity and the underlying SFR calculated with the IGIMF-theory diverges significantly from these classical linear relations based on a constant galaxy-wide IMF below SFRs typical for SMC-type galaxies.
" For comparison. the Maa, relation iu the classical picture becomes increasingly steeper for galaxy masses less than SMC-tvpe galaxies."," For comparison, the $M_\mathrm{gas}$ relation in the classical picture becomes increasingly steeper for galaxy masses less than SMC-type galaxies."
 It can be expected that the IGIME theory revises the Ta-SFRs of galaxies with a small neutral eas content such that their lower star formation cfiicicucies will be increased., It can be expected that the IGIMF theory revises the $\alpha$ -SFRs of galaxies with a small neutral gas content such that their lower star formation efficiencies will be increased.
 The ζω relation of this ealaxy suuple as resulting from the IGIME theory is shown in Fig., The $M_\mathrm{gas}$ relation of this galaxy sample as resulting from the IGIMF theory is shown in Fig.
 3 for the standard model aud in Fig.," \ref{fig_sfr_m_gas_std-igimf}
 for the standard model and in Fig."
 1 for the mundi model.," \ref{fig_sfr_m_gas_min-igimf}
 for the minimum model."
" It can be clearly seen that for both IGIAIF models the turu-cdown evident in the classical SER-AM,. relation (Fig. 2))", It can be clearly seen that for both IGIMF models the turn-down evident in the classical $M_\mathrm{gas}$ relation (Fig. \ref{fig_sfr_m_gas_imf}) )
 disappears completely., disappears completely.
 For the standard IGINIE a bivariate regression eives a linear scaliug of the total SFR and the ealaxy neutral gas mass. For the nininimia IGIME the relation between the SER aud the galaxy neutral gas mass is Note that the main parameter in the IGIME theory is the slope of the ECAIF.," For the standard IGIMF a bivariate regression gives a linear scaling of the total SFR and the galaxy neutral gas mass, For the minimum IGIMF the relation between the SFR and the galaxy neutral gas mass is Note that the main parameter in the IGIMF theory is the slope of the ECMF."
 Both ICIAIF models. the standard model with au ECME slope of Jj=2.35 and the müniuuu model with au LECMF slope of ;j=2. cover the full range of possible ECAIFs.," Both IGIMF models, the standard model with an ECMF slope of $\beta=2.35$ and the minimum model with an ECMF slope of $\beta=2$, cover the full range of possible ECMFs."
 Therefore. the change tha emerges in the SER-M. relation when changing from the classical iuvariant IME description to the IGIME heorv is a direct result of the nature of clustered star formation.," Therefore, the change that emerges in the $M_\mathrm{gas}$ relation when changing from the classical invariant IMF description to the IGIMF theory is a direct result of the nature of clustered star formation."
 Tn conclusion. the IGIMPE theory revises the SERs of dwarf galaxies siguificautly. such tha the true SFRs may be larecr than hitherto thought by two or three orders of magnitude.," In conclusion, the IGIMF theory revises the SFRs of dwarf galaxies significantly, such that the true SFRs may be larger than hitherto thought by two or three orders of magnitude."
 The reason why this was not evident ΤΕ uow comes about because in dwarf galaxies star formation predominantly occurs in low-mass clusters which do not contain massive stars., The reason why this was not evident until now comes about because in dwarf galaxies star formation predominantly occurs in low-mass clusters which do not contain massive stars.
 The gas depletion time scale. Το. is à nieasure for how long a galaxy with current total neutral eax inass. ΝΕ cal sustain its curent SER without being refucled by accretion of fresh material.," The gas depletion time scale, $\tau_\mathrm{gas}$, is a measure for how long a galaxy with current total neutral gas mass, $M_\mathrm{gas}$, can sustain its current SFR without being refueled by accretion of fresh extra-galactic material."
 It is defined by The correspouding gas depletion time scales. based on aconstant galaxy-wide IMF. are plotted in Fig. 5..," It is defined by The corresponding gas depletion time scales, based on aconstant galaxy-wide IMF, are plotted in Fig. \ref{fig_t_gas_depl__m_gas_imf}."
" The SER-M,4. rolatious for the IMP case (eq.", The $M_\mathrm{gas}$ relations for the IMF case (eq.
 2 and 23)) cau now be converted into Peas SER rolatious. which are plotted as dashed lines in Fie. 5..," \ref{eq_sfr_m_gas_imf_high} and \ref{eq_sfr_m_gas_imf_low}) ) can now be converted into $\tau_\mathrm{gas}$ -SFR relations, which are plotted as dashed lines in Fig. \ref{fig_t_gas_depl__m_gas_imf}."
 The gas depletion time scale increases sheltly with decreasing total ealaxy neutral eas mass for galaxies more massive than σον10h M. and it increases strongly for nüasslive ealaxies with aasses less than 5Γx10* M. in agreement with other studies (77)..," The gas depletion time scale increases slightly with decreasing total galaxy neutral gas mass for galaxies more massive than $5\times 10^7$ $M_\odot$ and it increases strongly for less-massive galaxies with masses less than $5\times 10^7$ $M_\odot$, in agreement with other studies \citep*{skillman2003a,bothwell2009a}."
 This is conuuoulv interpreted as star forming dwarf galaxies having much lower star formation efficiencies. r1. than large disk ealaxies.," This is commonly interpreted as star forming dwarf galaxies having much lower star formation efficiencies, $\tau_\mathrm{gas}^{-1}$ , than large disk galaxies."
 Iu the previous section it las been shown that the SFRs of dwart galaxies are much ligher if, In the previous section it has been shown that the SFRs of dwarf galaxies are much higher if
"The fact that the siupest modern cosmological heory. standard Cold Dark. Matter (sCDAL). aliiost. fits all available data has eιοοπαρος, the search for variauts of CDM tha teca1 do better.","The fact that the simplest modern cosmological theory, standard Cold Dark Matter (sCDM), almost fits all available data has encouraged the search for variants of CDM that can do better."
 I have been asked to make the case here hat CIIDM is the )8 theory of cosmüc structi1 formation. and indeed I )dieve that it the bst of all those I have coINcredif the cosmological natter density is nea ronical (i.c.. ορzz 1) and if the expansion rate is not too arge (Le. fb—Ty(1x)xlusiAIpe1)< 0.6).," I have been asked to make the case here that CHDM is the best theory of cosmic structure formation, and indeed I believe that it the best of all those I have considered the cosmological matter density is near critical (i.e., $\Omega_0 \approx 1$ ) and if the expansion rate is not too large (i.e. $h \equiv H_0/(100 \kmsMpc) \lsim 0.6$ )."
 Bu I hits it will be helpful to discuss CIIDM toget101) with its chief competitor a110)ie CDM variauts. low-Qy CDAL with a cosmologica COSant CACDAND).," But I think it will be helpful to discuss CHDM together with its chief competitor among CDM variants, $\Omega_0$ CDM with a cosmological constant )."
 While the predictions of iorinalized. CTIDAL aud both agree reasonablv well wit ithe available data ou scales of ~10 to LOHhYMec. cach has potetial virtues auk cleects.," While the predictions of COBE-normalized CHDM and both agree reasonably well with the available data on scales of $\sim
10$ to $100 \hMpc$, each has potential virtues and defects."
 wwith 04~03 has the possibe virtue of allowing a higher expausion rate [Ty OY a given ος»nüce age fg. hit +1ο defect of predicting too ταιςji fiucnation )OW'CY Olb sla] scales.," with $\Omega_0 \sim 0.3$ has the possible virtue of allowing a higher expansion rate $H_0$ for a given cosmic age $t_0$, but the defect of predicting too much fluctuation power on small scales."
 CITDM jas less power o1 small scales. so its precictfions appear to be 1u good agreenie with «ata on f1e galaxy. distribution. although it roniadus ο seen Whether --it medics carly enough galaxy fornatioi to be conrpatible wih the latest hieryeclslift data.," CHDM has less power on small scales, so its predictions appear to be in good agreement with data on the galaxy distribution, although it remains to be seen whether it predicts early enough galaxy formation to be compatible with the latest high-redshift data."
 Also. several sorts of data suggest that neutrinos have nonzero nass. aud 1jo variait of CIIDM favorec Lby this data in which the neutrino Wass IssLALCC οποσα two species of neuTos also SCCIUS LLOLE COupatible with the large-scale structure data.," Also, several sorts of data suggest that neutrinos have nonzero mass, and the variant of CHDM favored by this data — in which the neutrino mass is shared between two species of neutrinos — also seems more compatible with the large-scale structure data."
 Except for the Z7fy problem. trere is not a shred of evideico 1n avor of a nonzero cosmological constant. only increasingly strinecut i1oper bouuds on it frou several sorts of measurements.," Except for the $H_0-t_0$ problem, there is not a shred of evidence in favor of a nonzero cosmological constant, only increasingly stringent upper bounds on it from several sorts of measurements."
 Two recent observational results particularly favor high cosmic density. and thus favor Ὁ=1 models s1ch as CIIDM over," Two recent observational results particularly favor high cosmic density, and thus favor $\Omega=1$ models such as CHDM over"
Correlation functions (2-point and. higher order ones) have proved to be powerful statistical tools in order to acdcdress the study of the galaxy. clustering (e.g.Peeblesbles.1983). and are still wielely used in both local and hieh-redshift Universe (Ciavaliscoetal.1998:Blainetal.,"Correlation functions (2-point and higher order ones) have proved to be powerful statistical tools in order to address the study of the galaxy clustering \citep[e.g.][]{peebles76,groth,peebles80, dp} and are still widely used in both local \citep{connolly, eisenstein, masjedi} and high-redshift Universe \citep{giavalisco, blain}."
.. 2004)... Studies of the two point correlation function have matured to the point that one can study how galaxies populate dark matter halos in detail (e.g...Zehavictal. Y04).. the tvpical halo masses of galaxy. populations às a —unction of redshift (e.g. Lyman breaks - Ciavaliscoctal.rm 998)). the relative clustering of different populations (e.g.. 1e tendenev of AGN to cluster like the massive galaxy population as a whole: Lictal.. 2006)). and the use of clustering measures on the smallest scale to constrain the merger history of galaxies (o.g.. Pattonetal.2002.. Belletal.. 2006.. Robainaetal.. 20102).," Studies of the two point correlation function have matured to the point that one can study how galaxies populate dark matter halos in detail \citep[e.g.,][]{zehavi}, the typical halo masses of galaxy populations as a function of redshift (e.g., Lyman breaks - \citealp{giavalisco}) ), the relative clustering of different populations (e.g., the tendency of AGN to cluster like the massive galaxy population as a whole; \citealp{li06b}) ), and the use of clustering measures on the smallest scale to constrain the merger history of galaxies (e.g., \citealp{patton}, \citealp{bell}, \citealp{robaina10}) )."
 Furthermore. the correlation function method allows us not only to study the clustering of the galaxies themselves. but also how some of their properties are clustered.," Furthermore, the correlation function method allows us not only to study the clustering of the galaxies themselves, but also how some of their properties are clustered."
 Weighted correlation functions (Boerneretal.1989) or in a eencral sense. marked statistics (Beisbart&Ixerscher.2000:al..2006:Robainaet2009) have been widely used in the last ten vears in order to study how observables depend on the separation between galaxies.," Weighted correlation functions \citep{boerner} or in a general sense, marked statistics \citep{beisbart,gott,falten,skibba06,robaina} have been widely used in the last ten years in order to study how observables depend on the separation between galaxies."
 In. particular. weighted correlation functions are frequently used to study the dependence of star formation rate (SER) on separation between ealaxies. in great. part to explore the inlluence of galaxy interactions on enhancing a galaxy pairs SER 2009)..," In particular, weighted correlation functions are frequently used to study the dependence of star formation rate (SFR) on separation between galaxies, in great part to explore the influence of galaxy interactions on enhancing a galaxy pair's SFR \citep[e.g.][]{li,robaina}. ."
"To date, the most effective means for staving off the angular momentum problem is to employ cooling suppression in the vicinity of star formation events in order to intensify the effects of feedback.","To date, the most effective means for staving off the angular momentum problem is to employ cooling suppression in the vicinity of star formation events in order to intensify the effects of feedback."
" Most of the recent researchers in this field employ it by default, while at the same time they investigate the effects of other simulation parameters on preventing the angular momentum problem (e.g.??)."," Most of the recent researchers in this field employ it by default, while at the same time they investigate the effects of other simulation parameters on preventing the angular momentum problem \citep[e.g.][]{Agertz:2010p461, Guedes:2011p1080}."
" Our results indicate that the other parameters are secondary to cooling suppression, and they do not seem to work effectively in the absence of it."," Our results indicate that the other parameters are secondary to cooling suppression, and they do not seem to work effectively in the absence of it."
" We realize that the use of cooling suppression is mostly motivated by its effectiveness in this regard, but we continue to search for other feedback parameterizations which are more physically derived."," We realize that the use of cooling suppression is mostly motivated by its effectiveness in this regard, but we continue to search for other feedback parameterizations which are more physically derived."
 There are several potential alternatives to cooling suppression as a form of feedback in cosmological hydrodynamics solutions of galaxy formation., There are several potential alternatives to cooling suppression as a form of feedback in cosmological hydrodynamics solutions of galaxy formation.
" Radiative feedback shows some promising results (?),, although scaling its effects to a large number of particle sources is a current computational challenge."," Radiative feedback shows some promising results \citep{Kim:2011p1119}, although scaling its effects to a large number of particle sources is a current computational challenge."
" Another option is that the cosmic rays produced by supernovae could be used as a means for transporting energy and momentum to the surrounding medium as an additional fluid in a simulation (e.g.,??).."," Another option is that the cosmic rays produced by supernovae could be used as a means for transporting energy and momentum to the surrounding medium as an additional fluid in a simulation \citep[e.g.,][]{Miniati:2001bz, Jubelgas:2008je}."
" Alternatively, modifying the gas in dense regions to have stiff equation of state (?) may help “puff” up early collapsinga pockets of gas."," Alternatively, modifying the gas in dense regions to have a stiff equation of state \citep{Agertz:2010p461} may help “puff” up early collapsing pockets of gas."
" Perhaps there is even a redshift-dependent feedback prescription, similar to those used by ? and ?,, where it was assumed that the IMF was top-heavy in the distant past resulting in more supernovae and a higher feedback efficiency in that epoch."," Perhaps there is even a redshift-dependent feedback prescription, similar to those used by \citet{SommerLarsen:2003p1116} and \citet{Okamoto:2005p1101}, where it was assumed that the IMF was top-heavy in the distant past resulting in more supernovae and a higher feedback efficiency in that epoch."
 We are currently investigating some of these options to be presented in a forthcoming paper., We are currently investigating some of these options to be presented in a forthcoming paper.
" We acknowledge support from NSF grants AST-0547823, AST-0908390, and AST-1008134, as well as computational resources from NASA, the NSF Teragrid, and Columbia University’s Hotfoot cluster."," We acknowledge support from NSF grants AST-0547823, AST-0908390, and AST-1008134, as well as computational resources from NASA, the NSF Teragrid, and Columbia University's Hotfoot cluster."
 We also thank the, We also thank the
period. Zi.x1/Mai.,"period $T_{\rm osc} \propto 1/M_{\rm
disc}$."
" When rey«2a. Tue decreases as ry, increases."," When $r_p \ll R_o$, $T_{\rm osc}$ decreases as $r_p$ increases."
 For larger values of ri. τω does not vary. much with this parameter.," For larger values of $r_p$, $T_{\rm osc}$ does not vary much with this parameter."
 In the simulations we have performed. Tus~107 vears.," In the simulations we have performed, $T_{\rm osc} \sim
10^5$ years."
 As this is much shorter than the disc lifetime. there is à nonzero probability that the planet is lel on a highly eccentric orbit as the disc dissipates.," As this is much shorter than the disc lifetime, there is a nonzero probability that the planet is left on a highly eccentric orbit as the disc dissipates."
 Note that the process discussed in this paper is no a mere trivial extension of the Ixozai effect., Note that the process discussed in this paper is not a mere trivial extension of the Kozai effect.
 Indeed. in the classical Ixozai ellect. the periodic behaviour of e is obtainec because the dependence. of the perturbing &ravitationa potential on the planets argument of pericentre w is through a termi proportional to cosXx.," Indeed, in the classical Kozai effect, the periodic behaviour of $e$ is obtained because the dependence of the perturbing gravitational potential on the planet's argument of pericentre $\omega$ is through a term proportional to $\cos 2 \omega$."
 When the orbit of the plane crosses the disc. the gravitational potential has a very different form. and the behaviour of e is much more dillicul to predict.," When the orbit of the planet crosses the disc, the gravitational potential has a very different form, and the behaviour of $e$ is much more difficult to predict."
 The effect. discussed. here could. be. inhibited. if other processes induced: a precession of the planets orbit. on timescales shorter than Z5...," The effect discussed here could be inhibited if other processes induced a precession of the planet's orbit on timescales shorter than $T_{\rm
osc}$."
 That may happen if other planets are present in the system. or because of clissipative tidal torques exerted by the disc. which have been ignored here (Lubow Ogilvie 2001).," That may happen if other planets are present in the system, or because of dissipative tidal torques exerted by the disc, which have been ignored here (Lubow Ogilvie 2001)."
 In the latter case. the Ixozai ellect would then work onlv for planets orbiting inside the disc inner cavity.," In the latter case, the Kozai effect would then work only for planets orbiting inside the disc inner cavity."
" When the planet crosses the disc. loss of energy. and angular momentum cireularises the orbit and aligns it with the disc plane on a timescale ~LAL,(Zr, where £8 is the planet radius (Sver et al."," When the planet crosses the disc, loss of energy and angular momentum circularises the orbit and aligns it with the disc plane on a timescale $\sim T M_p/(\Sigma(r_p) R_p^2)$, where $R_p$ is the planet radius (Syer et al."
 1991. Ivanov ct al.," 1991, Ivanov et al."
 25).1999). which can be smaller than Zi. for rj larger than a few au.," 1999), which can be smaller than $T_{\rm
osc}$ for $r_p$ larger than a few au."
 Taking this process into account together with the Ixozai effect may leac to equilibrium values of e and £., Taking this process into account together with the Kozai effect may lead to equilibrium values of $e$ and $I$.
 Fhis will be studied in a forthcoming paper., This will be studied in a forthcoming paper.
 Note that the usual tvpe LE migration mechanism that applies to planets orbiting in discs would not be relevant here. as it happens when the planet orbits in à gap that is locked in the disc evolution.," Note that the usual type II migration mechanism that applies to planets orbiting in discs would not be relevant here, as it happens when the planet orbits in a gap that is locked in the disc evolution."
 In the coplanar case. as the disc spirals toward the central star. it carries along the gap and the planet.," In the coplanar case, as the disc spirals toward the central star, it carries along the gap and the planet."
 When the planet is on an inclined orbit. it may still open up à gap but it is not locked in it. so that it is not being pushed in as the disc spirals in.," When the planet is on an inclined orbit, it may still open up a gap but it is not locked in it, so that it is not being pushed in as the disc spirals in."
 The mechanism described here relies on the planet being on an inclined. orbit. which could happen as a result. of: (i) dvnamical relaxation of a population of planets formed through fragmentation of a protostellar envelope around a star surrounced by a disc (Papaloizou Terquem 2001): (ii) mean motion resonances (Phonics Lissauer 2003. also Yu “Premaine 2001) and (ii) eravitational interactions between embryos during the planet formation stage (Levison et al.," The mechanism described here relies on the planet being on an inclined orbit, which could happen as a result of: (i) dynamical relaxation of a population of planets formed through fragmentation of a protostellar envelope around a star surrounded by a disc (Papaloizou Terquem 2001); (ii) mean motion resonances (Thommes Lissauer 2003, also Yu Tremaine 2001) and (iii) gravitational interactions between embryos during the planet formation stage (Levison et al."
 1998. Cresswell Nelson 2008).," 1998, Cresswell Nelson 2008)."
 In all cases. the process that makes the orbits inclined also makes them eccentric.," In all cases, the process that makes the orbits inclined also makes them eccentric."
 According to the results presented here. the dise could pump the eccentricities up to even larger values.," According to the results presented here, the disc could pump the eccentricities up to even larger values."
 Aleasures of the projected angle between the axes of the planets orbit ancl the stellar rotation. using the RossiterAleLaughlin ellect. are becoming available.," Measures of the projected angle between the axes of the planet's orbit and the stellar rotation, using the Rossiter--McLaughlin effect, are becoming available."
 So far. only the system NO3. which has à 12 Jupiter masses planet on à 3.19 days orbit with e=0.26. has been shown to have a spin misalignment of at least 377.3 (EHlebbrard ct al.," So far, only the system XO–3, which has a $\sim 12$ Jupiter masses planet on a 3.19 days orbit with $e=0.26$, has been shown to have a spin--orbit misalignment of at least $^{\circ}$ .3 (Hébbrard et al."
 2008. Winn et al.," 2008, Winn et al."
 2009)., 2009).
 Misalignment has also been reported for the system HD. SOGO6. which has à 4 Jupiter mass planet on a 111.44 davs orbit with e=0.93 (Moutou ct al.," Misalignment has also been reported for the system HD 80606, which has a $\sim 4$ Jupiter mass planet on a 111.44 days orbit with $e= 0.93$ (Moutou et al."
 2009. Pont et al.," 2009, Pont et al."
 2009)., 2009).
 As LIDSOGOG is a component of a binary system. the classical Ixozai elect. could. be responsible for the misalignment in this svsten.," As HD80606 is a component of a binary system, the classical Kozai effect could be responsible for the misalignment in this system."
 We thank J. Papaloizou and S. Balbus for very useful comments which have greatly improved the original version of this paper., We thank J. Papaloizou and S. Balbus for very useful comments which have greatly improved the original version of this paper.
 CT. thanks the MPLX in Heidelberg. where part of this paper was written. for hospitality and support.," C.T. thanks the MPIA in Heidelberg, where part of this paper was written, for hospitality and support."
have been obtained (see Rutener. 1988. for the detnition of the weight q).,"have been obtained (see Rufener, 1988, for the definition of the weight q)."
 The photometric reduction procedure was described by Rutener (1961. 1985): the photometric data in the Ceneva system are collected in he Ceneral Catalogue. (Bufener. 1988) aud its up-to-date database (Durli. 2006).," The photometric reduction procedure was described by Rufener (1964, 1985); the photometric data in the Geneva system are collected in the General Catalogue (Rufener, 1988) and its up-to-date database (Burki, 2006)."
 Iu the case of W Cru. the all-sky reduction was improved by the measurement of a ucielbouwing conrparison star. ITD 101021 (IK1TIID). a well measured standard. with a confirmed stabilitv: [10 measurements over a period of 20 vears. mican V. maguitude 5.1371 with a standard deviation σι = 0.0016.," In the case of W Cru, the all-sky reduction was improved by the measurement of a neighbouring comparison star, HD 101021 (K1III), a well measured standard, with a confirmed stability: 440 measurements over a period of 20 years, mean $V$ magnitude 5.1374 with a standard deviation $\sigma_{V}$ = 0.0046."
 The same value of oq can be adopted for the global precision of the micasures ou W Cru., The same value of $\sigma_{V}$ can be adopted for the global precision of the measures on W Cru.
 The data on W Cru are listed in Table 1 available at the CDS., The data on W Cru are listed in Table 1 available at the CDS.
 The aagnitudes in each of the seven filters are obtained from the visual magnitude V. aud the six colour indices in the following manner: with i represcuting one of the seven filters ἐν B. V. By. By. V4. C. Remember that the Geneva [CB] aud [BV] indices are uot normalized to zero for au AOV star as it is the case for the Joliusou UDV indices.," The magnitudes in each of the seven filters are obtained from the visual magnitude $V$ and the six colour indices in the following manner: with i representing one of the seven filters $U$ , $B$, $V$, $B_{1}$ , $B_{2}$, $V_{1}$, G. Remember that the Geneva $[U - B]$ and $[B - V]$ indices are not normalized to zero for an A0V star as it is the case for the Johnson UBV indices."
 Iu Fig., In Fig.
 l light curves in all seven photometric pass bands of the Geneva svstemi are shown., 1 light curves in all seven photometric pass bands of the Geneva system are shown.
 The following ephemeris is used for the calculations of the phases: AMGuI (ILJD)—2110731.81|198.537«E after Iohoutek (1988) and the elements are derived from alb available winnua found iu the literature.," The following ephemeris is used for the calculations of the phases: Min I $ = 2440731.84 + 
198.537 \times E$ after Kohoutek (1988) and the elements are derived from all available minima found in the literature."
 The primary nium. o which phase y=0.50 is assigned throughout iu lis paper. aud which ids caused by the eclipse of he visible component. became deeper for the shorter wavelengths.," The primary minimum, to which phase $\varphi = 0.50$ is assigned throughout in this paper, and which is caused by the eclipse of the visible component, became deeper for the shorter wavelengths."
 Iowever. light curve in the (-band has uuch larger intrinsic scatter - large changes in the width al depth of the xmuaurv müniuuni are evideut for the ight curve in the Ü-baud than for the other pass bands.," However, light curve in the $U$ -band has much larger intrinsic scatter - large changes in the width and depth of the primary minimum are evident for the light curve in the $U$ -band than for the other pass bands."
 Unfortunately. pliase coverage of the secondary wii is muuch poorer than primary aud cvcle-to-cycele variations in the secondary mini are not evident as for the primary minium.," Unfortunately, phase coverage of the secondary minimum is much poorer than primary and cycle-to-cycle variations in the secondary minimum are not evident as for the primary minimum."
 In fact. proper phase coverage in the secondary müniuuni has been secured iu a single cvele.," In fact, proper phase coverage in the secondary minimum has been secured in a single cycle."
 The width of the secdary ΙΙ when au opaque disk is eclipsing its stellar companion is larecr than the width of the primary mininuun., The width of the secondary minimum when an opaque disk is eclipsing its stellar companion is larger than the width of the primary minimum.
 At shorter wavelcuetls he secondary ninimuuni becomes not only deeper but evidently narrower., At shorter wavelengths the secondary minimum becomes not only deeper but evidently narrower.
 Moreover. some other peculiaritics πι he light curves of W Cru are clearly seen: mnaxinuuni ollowiug deeper muni is higher than preceding. there aro asvinunetnes in descending and ascending branches of both. primary aud secondary. máininia.," Moreover, some other peculiarities in the light curves of W Cru are clearly seen: maximum following deeper minimum is higher than preceding, there are asymmetries in descending and ascending branches of both, primary and secondary, minima."
 Duups aud/or inps in the lieht curves at the phases ο=0.2.0.3.0.6 are very pronounced. aud. moreover. variable from cvcle-o-cvele.," Bumps and/or humps in the light curves at the phases $\varphi = 0.2, 0.3, 0.6$ are very pronounced, and, moreover, variable from cycle-to-cycle."
 The color variatious (bottom pauncls in Fie., The color variations (bottom pannels in Fig.
 1) are indicative of temperature chanecs., 1) are indicative of temperature changes.
 In quadratures when he supereiaut is secu from the side |BVW is the largest., In quadratures when the supergiant is seen from the side $[B-V]$ is the largest.
 The amplitude of the changes in [B-V| index is about 0.25 mag., The amplitude of the changes in [B-V] index is about 0.25 mag.
 As is evident from Fig., As is evident from Fig.
 1 changes in |B] are twice aslaree then in |BV]. i.e. about 0.50 mag.," 1 changes in $[U-B]$ are twice as large then in $[B-V]$, i.e. about 0.50 mag."
 The variations are nore difficult to interpret iu [0ΠΡΙ., The variations are more difficult to interpret in $[U-B]$.
" Large aud crratic changes in [CBl avound the primary minim, are resent as consequence of a large activity in this svsteni onu tine scale comparable to svstenis orbital period.", Large and erratic changes in $[U-B]$ around the primary minimum are present as consequence of a large activity in this system on time scale comparable to system's orbital period.
 It secnus that the source of these rapid fluctuations. which jas different amplitudes from cvcle-to-cvcle is screened im he phase of the secondary ΙΙ.," It seems that the source of these rapid fluctuations, which has different amplitudes from cycle-to-cycle is screened in the phase of the secondary minimum."
 But due to lack of more intensive phase coverage our couchision 1s rather entative at this point (sce also Fie., But due to lack of more intensive phase coverage our conclusion is rather tentative at this point (see also Fig.
 3)., 3).
 The eclipseaunappiug technique introduced by Horne (1985) in combination with Doppler tomography (Marsh Iforue 1988) became an essential tool for the study of the structure and mass-flow processes in cataclysuic variables (CVs)., The eclipse-mapping technique introduced by Horne (1985) in combination with Doppler tomography (Marsh Horne 1988) became an essential tool for the study of the structure and mass-flow processes in cataclysmic variables (CVs).
 In particular. eclipsemapping is used to reconstruct the surface brightuess distribution of the accretion disks in CVs. and is often assumed to be an indirect jmagine method.," In particular, eclipse–mapping is used to reconstruct the surface brightness distribution of the accretion disks in CVs, and is often assumed to be an indirect imaging method."
 Initially. ouly the information contained iu the shape of the eclipse was used iu he assembly of the map of the accretion disk surface xiehtuess distribution.," Initially, only the information contained in the shape of the eclipse was used in the assembly of the map of the accretion disk surface brightness distribution."
 Iowever. other sources of the radiation have been completely ignored. e.g. the secondary (often cool) component.," However, other sources of the radiation have been completely ignored, e.g. the secondary (often cool) component."
 Rutten (1998 las nuaproved the echuique with the developiuneut of 3D eclipsemapping which fits he whole light-curve instead ofjust the eclipsed xut., Rutten (1998) has improved the technique with the development of 3D eclipse–mapping which fits the whole light-curve instead of just the eclipsed part.
 Iu the xeseut study we have followed Butteu's (1998) approach., In the present study we have followed Rutten's (1998) approach.
 But. unlike the case of CVs in Aleols here accretion disk is not prevailing source of radiation.," But, unlike the case of CVs in Algols here accretion disk is not prevailing source of radiation."
 Moreover. the hot component las a relatively large radius and its screcuing of the accretion disk should be aken iuto account.," Moreover, the hot component has a relatively large radius and its screening of the accretion disk should be taken into account."
 Previous studies of the elt curves of W Serpentis Dbinarics have shown that hot conrponents mueht be exposed to various deerces. Therefore. in the construction of the binary model we have started from the gencral case du. which all three main radiative sources are preseuted: hot mass-accretiueg component surrounded by an optically thick disk aud cool mass-loosine component which (canonically) is filling its Roche lobe.," Previous studies of the light curves of W Serpentis binaries have shown that hot components might be exposed to various degrees, Therefore, in the construction of the binary model we have started from the general case in which all three main radiative sources are presented: hot mass-accreting component surrounded by an optically thick disk and cool mass-loosing component which (canonically) is filling its Roche lobe."
 The stellar surfaces are approximated by equipoteutial surfaces., The stellar surfaces are approximated by equipotential surfaces.
 These surfaces are determined dy two parameters: the lass ratio 4 and the Πιο factor S (for description of the eravitational poteutial of two rotating stars in binary system see e.c. Hüküteh 2001)., These surfaces are determined by two parameters: the mass ratio $q$ and the filling factor $S$ (for description of the gravitational potential of two rotating stars in binary system see e.g. Hilditch 2001).
 We approximated anu equipotential surface with a erid of triangles., We approximated an equipotential surface with a grid of triangles.
 This was done to avoid ravtracing when calculating visibility., This was done to avoid ray–tracing when calculating visibility.
 Geometry of an accretion disk iu our model is eiven by anàdisk model. ie. the disk is deteruiued by the disk radius Ry aud an opening anele j ).," Geometry of an accretion disk in our model is given by an$\alpha$–disk model, i.e. the disk is determined by the disk radius $R_{\rm d}$ and an opening angle $\beta$ ."
0.8—1.0. as illustrated in Figure 1..,"$0.8-1.0$, as illustrated in Figure \ref{stormfig2}."
 Storm οἱ al., Storm et al.
 discussed (he source of the poor fit without coming to a conclusion and decided that the best course of action was to ignore (his phase interval in the fit., discussed the source of the poor fit without coming to a conclusion and decided that the best course of action was to ignore this phase interval in the fit.
 Thev deleted thephase interval 0.8—1.0 for all stars in their fit of equation Cr)) For kr and o., They deleted thephase interval $0.8-1.0$ for all stars in their fit of equation \ref{eq:simple}) ) for $r$ and $\phi_0$.
 some Cepheids showed a small phase shift between (he photometry ancl the radial velocities., Some Cepheids showed a small phase shift between the photometry and the radial velocities.
 The phase shift causes a loop in the upper panel ancl a displacement between the fitted curve and the points in the lower panel., The phase shift causes a loop in the upper panel and a displacement between the fitted curve and the points in the lower panel.
 For the stars showing this effect a phase shit was determined by minimizing the scatter in the upper panel., For the stars showing this effect a phase shift was determined by minimizing the scatter in the upper panel.
 The phase shift has already been imposed between the photometry and radial velocities used for Fig., The phase shift has already been imposed between the photometry and radial velocities used for Fig.
 1. and thus these effects are not seen., \ref{stormfig2} and thus these effects are not seen.
 The adopted phase shifts are listed in Table 1., The adopted phase shifts are listed in Table 1.
 There are (wo constants adopted in a surface brightness caleulation that must be the same in the linear-hisector and. Bavesian ealeulations., There are two constants adopted in a surface brightness calculation that must be the same in the linear-bisector and Bayesian calculations.
 We chose to fix the constants in the Bavesian calculation to those used in the Storm et al., We chose to fix the constants in the Bayesian calculation to those used in the Storm et al.
 paper., paper.
 These constants are the constant term in the surface brightness definition. equation (1)). and the conversion factor from angular diameter (in milliarcseconds) and distance (parsecs) to linear radius (solar radii).," These constants are the constant term in the surface brightness definition, equation \ref{eq:fparameter}) ), and the conversion factor from angular diameter (in milliarcseconds) and distance (parsecs) to linear radius (solar radii)."
 The first of these was taken to be 4.2207 and the second as 0.10727 solar radii per nias-parsec (see eq. (7)))., The first of these was taken to be 4.2207 and the second as 0.10727 solar radii per mas-parsec (see eq. \ref{eq:simple}) )).
 To perform the calculations Storm et al., To perform the calculations Storm et al.
 usec the FORTRAN subroutine SINLIN which is available from Isobe et al., used the FORTRAN subroutine SIXLIN which is available from Isobe et al.
(1990).. The computations were run on a Linux personal computer and took a fraction of a second per star., The computations were run on a Linux personal computer and took a fraction of a second per star.
" The «quantities determined in (he linear-bisector calculation (hat are of interest to us here are (he distance r. the mean linear radius δὲ, computed from the distance ancl mean angular diameter oy. and their 1 6 uncertainties."," The quantities determined in the linear-bisector calculation that are of interest to us here are the distance $r$, the mean linear radius $R$, computed from the distance and mean angular diameter $\phi_0$, and their 1 $\sigma$ uncertainties."
 These quantiües are given in Table 1., These quantities are given in Table 1.
 The Bavesian MCMC method that we applied here is described in detail by Barnes et al., The Bayesian MCMC method that we applied here is described in detail by Barnes et al.
(2003).. It is important to realize that the Bavesian calculation treats Cae in the problem asdistributions. not as specilic values to be determined.," It is important to realize that the Bayesian calculation treats the in the problem as, not as specific values to be determined."
 The goal of (he analvsis is (ο determine the probability distribution for each parameter of interest. [rom which inlerences max be drawn by appropriate means and variances.," The goal of the analysis is to determine the probability distribution for each parameter of interest, from which inferences may be drawn by appropriate means and variances."
 To compute the lor each paviuneter requires the likelihood [function on a specific model. appropriate priors. and sampling strategies for all parameters.," To compute the for each parameter requires the likelihood function on a specific model, appropriate priors, and sampling strategies for all parameters."
 The is the probability of obtaining the particular data eiven the model., The is the probability of obtaining the particular data given the model.
It does not change with distance along the plane and. as BMM explained. Is insensitive to errors in emission measure.,"It does not change with distance along the plane and, as BMM explained, is insensitive to errors in emission measure."
" The near constancy of (n)=FN, along the line of sight (see D—Οἱ) fit in Table 2)). first noted by Weisberg et al. (1980))."," The near constancy of $\avnel = \Filfac \Nec$ along the line of sight (see $D- \avnel$ fit in Table \ref{tab:2}) ), first noted by Weisberg et al. \cite{weisberg+80}) ),"
" was an early indication of an inverse relationship between Fi and N,.", was an early indication of an inverse relationship between $\Filfac$ and $\Nec$.
 We first compare our Εν Να relation with earlier determinations available in the literature. and then discuss what it tells us about the density structure in the DIG.," We first compare our $\Filfac$ $\Nec$ relation with earlier determinations available in the literature, and then discuss what it tells us about the density structure in the DIG."
" Figure 9 and Table 5. show our F,—N, relation together with earlier determinations.", Figure \ref{fig:9} and Table \ref{tab:5} show our $\Filfac$ $\Nec$ relation together with earlier determinations.
 For the same mean density in clouds. we find somewhat smaller filling factors than BMM which agree with the lower points in their distribution.," For the same mean density in clouds, we find somewhat smaller filling factors than BMM which agree with the lower points in their distribution."
 The smaller errors in the BMM relation are not only due to the larger sample. but also to the smooth distribution of 2.)=DM/D of the electron density model of Cordes Lazio (2002)) used by BMM.," The smaller errors in the BMM relation are not only due to the larger sample, but also to the smooth distribution of $\avnel = \DM /D$ of the electron density model of Cordes Lazio \cite{cordes+lazio02}) ) used by BMM."
" Berkhuijsen (20049) obtained the F,—N, relation for the DIG in the galaxy 331 from à comparison of rotation measures of the polarized continuum emission (giving 1) with magnetic field strengths of Fletcher et citefletcher+04)) and thermal radio emission (giving (n2)).", Berkhuijsen \cite{elly04}) ) obtained the $\Filfac$ $\Nec$ relation for the DIG in the galaxy 31 from a comparison of rotation measures of the polarized continuum emission (giving $\avnel$ with magnetic field strengths of Fletcher et \\cite{fletcher+04}) ) and thermal radio emission (giving $\langle\nel^2\rangle$ ).
 These data refer to the DIG in the bright emission ring in 331., These data refer to the DIG in the bright emission ring in 31.
 The good agreement between the 331 and the MW ∶−↾∧∨⋁⊏≻↖⊂≡↥↾↴↴≤↘↴⋁⊙↾⇂∖ generally valid in the DIG in galaxies., The good agreement between the 31 and the MW results suggests that the inverse correlation between $\Filfac$ and $\Nec$ is generally valid in the DIG in galaxies.
 . & (2001)) observed — | lines from the inner Galaxy at low latitudes., Roshi Anantharamaiah \cite{roshi+01}) ) observed radio recombination lines from the inner Galaxy at low latitudes.
 0.001 derived filling factors of € 0.01for extended regions of diffuse ionized gas with densities of |—IOcm-. which agree well with the extension of the curves for the DIG in Fig. 9..," They derived filling factors of $\le 0.01$ for extended regions of diffuse ionized gas with densities of $1-10\cmcube$, which agree well with the extension of the curves for the DIG in Fig. \ref{fig:9}."
 Thus the inverse correlation between FY and N. at least holds for the density range 0.03—I0 cm7., Thus the inverse correlation between $\Filfac$ and $\Nec$ at least holds for the density range $0.03-10\cmcube$ .
 Cordes et al. (1985) , Cordes et al. \cite{cordes+85}) )
estimated a filling factor of 107777? for clumps of about Ipe size causing scattering of pulsar signals., estimated a filling factor of $10^{-4.0\pm0.3}$ for clumps of about $1\pc$ size causing scattering of pulsar signals.
 Our data predict this filling factor for clumps of density N.=50cm-? and size L./m=0.9 pc. where i is the number of clumps on the line of sight.," Our data predict this filling factor for clumps of density $\Nec = 50\cmcube$ and size $L_\mathrm{e}/m
= 0.9\pc$ , where $m$ is the number of clumps on the line of sight."
 The relation of BMM gives clumps of N.2100cm for the same filling factor and sizes., The relation of BMM gives clumps of $\Nec = 100\cmcube$ for the same filling factor and sizes.
" So these clump properties also seem to follow the F,—N, relations derived for the DIG.", So these clump properties also seem to follow the $\Filfac$ $\Nec$ relations derived for the DIG.
 Furthermore. Gaustad Van Buren (1993)) obtained a similar relationship for clouds of diffuse dust within 400pe from the Sun. with mean densities of |—10em™™.," Furthermore, Gaustad Van Buren \cite{gaustad+93}) ) obtained a similar relationship for clouds of diffuse dust within $400\pc$ from the Sun, with mean densities of $1-10\cmcube$."
 Figure 9 shows that these clouds have about 3 times higher filling factors than ionized clouds of the same mean density., Figure \ref{fig:9} shows that these clouds have about 3 times higher filling factors than ionized clouds of the same mean density.
 Interestingly. their data agree very well with those of Pynzar (1993.. also shown in citeelly98)) derived from recombination lines.," Interestingly, their data agree very well with those of Pynzar \cite{pynzar93}, also shown in \\cite{elly98}) ) derived from recombination lines."
" The general validity of the inverse F,—N, relation indicates that it describes a basic property of the DIG. possibly even of the entire ISM citeelly99))."," The general validity of the inverse $\Filfac$ $\Nec$ relation indicates that it describes a basic property of the DIG, possibly even of the entire ISM \\cite{elly99}) )."
 At least two mechanisms could be responsible: thermal pressure equilibrium and turbulence causing a fractal density structure., At least two mechanisms could be responsible: thermal pressure equilibrium and turbulence causing a fractal density structure.
" Thermal pressure equilibrium leads to an inverse F,—N, relation if it were widespread. but it seems only locally valid."," Thermal pressure equilibrium leads to an inverse $\Filfac$ $\Nec$ relation if it were widespread, but it seems only locally valid."
 In the MW largefractions of gas are observedin unstable regimes and have also been found in simulations of a turbulent ISM., In the MW largefractions of gas are observedin unstable regimes and have also been found in simulations of a turbulent ISM.
 Thermal pressure equilibrium of clouds appears to be of minor, Thermal pressure equilibrium of clouds appears to be of minor
Tt is well known that the metallicity distribution of C-dwarfs in the solar neighbourhood shows a deficit: of metal-poor stars relative to the predictious of the Simple Model of chemical evolution.,It is well known that the metallicity distribution of G-dwarfs in the solar neighbourhood shows a deficit of metal-poor stars relative to the predictions of the Simple Model of chemical evolution.
 This is the so-called G-chwart problem. originally noted by van den Bergh (1962) and Schinidt (1963).," This is the so-called G-dwarf problem, originally noted by van den Bergh (1962) and Schmidt (1963)."
 AMauy solutions have been proposed. some in the frame-work of analytical models. based on the IRA.," Many solutions have been proposed, some in the frame-work of analytical models, based on the IRA."
 The problem is still preseut if we adopt the oxveen abundances in place of metallicity., The problem is still present if we adopt the oxygen abundances in place of metallicity.
 Au important result in the framework of IRA models with sas flows was pointed out by Edmunds (1990)., An important result in the frame-work of IRA models with gas flows was pointed out by Edmunds (1990).
 Ie found that the C-dwart xoblem cannot be solved by auv outflow but i is possible to solve it with particular forms of inflow., He found that the G-dwarf problem cannot be solved by any outflow but it is possible to solve it with particular forms of inflow.
" Lvudeu Bell (1975) ""best accretion model’ aud Clavtou's models (Clayton 1988) are models of just this sind.", Lynden Bell (1975) 'best accretion model' and Clayton's models (Clayton 1988) are models of just this kind.
 Of course it is also possible to reduce the uuuber of metal-poor stars with respect to that predicted by he Simple Model by asstunineg metal dependent stellar vields (Alaeder 1992) Guore metals are produced at lower uctallicity with a consequent Prompt Initial Euricliieut (P.LE.) (Travan and Cameron 1971))., Of course it is also possible to reduce the number of metal-poor stars with respect to that predicted by the Simple Model by assuming metal dependent stellar yields (Maeder 1992) (more metals are produced at lower metallicity with a consequent Prompt Initial Enrichment (P.I.E.) (Truran and Cameron 1971)).
 However. it. has con shown by several papers (Ciovaguoli aud Tosi. 1995: Cariei 1996). that in this wav shallow eracicuts along je Galactic disk are produced. which do not agree with 1ο dnost recent observational estimates (see Matteucci and Chiappini 1999 for a review).," However, it has been shown by several papers (Giovagnoli and Tosi, 1995; Carigi 1996), that in this way shallow gradients along the Galactic disk are produced, which do not agree with the most recent observational estimates (see Matteucci and Chiappini 1999 for a review)."
 Actually. this is au navoidable problem still preseut in the models discussed in this paper where we consider time-dependent IME:.," Actually, this is an unavoidable problem still present in the models discussed in this paper where we consider time-dependent IMFs."
" It is still possible to obtain a Prompt Tuitial Euricluneut Nwoassundne au PME variable with time and oeji particular very flat at carly times. in order to favour +je formation of massive stars,"," It is still possible to obtain a Prompt Initial Enrichment by assuming an IMF variable with time and in particular very flat at early times, in order to favour the formation of massive stars."
 Iu this paper we exploit the previous idea to investigate the tine behaviour of the IMFE., In this paper we exploit the previous idea to investigate the time behaviour of the IMF.
 We assume the same hypothesis of the Simple Model (Tinsley 1980) with the exception of adopting a time-depeudent IME., We assume the same hypothesis of the Simple Model (Tinsley 1980) with the exception of adopting a time-dependent IMF.
 Tn such a wav it is possible to fud an equation (Section. 2.1) which must be satisfied i order to reproduce the observed distribution of Ciodwarfs., In such a way it is possible to find an equation (Section 2.1) which must be satisfied in order to reproduce the observed distribution of G-dwarfs.
 We then use this equation to investigate the time behaviour of IMES with oue aud two slopes., We then use this equation to investigate the time behaviour of IMFs with one and two slopes.
 In Section 3 we use a ποΊσα model without IRA to test the IMEs on other observational coustraiuts., In Section 3 we use a numerical model without IRA to test the IMFs on other observational constraints.
 Tn the following we ποσα an analytical expression to approximate the data concerning the metallicity distribution of Cedwarfs in the solar neighbourhood., In the following we need an analytical expression to approximate the data concerning the metallicity distribution of G-dwarfs in the solar neighbourhood.
 The data are plotted in Fig., The data are plotted in Fig.
 1 where the differeutial distribution of oxvecn abuudances for the solar cvlinder from Rocha-Pinto aud Maciel (1996) is given (triangles)., 1 where the differential distribution of oxygen abundances for the solar cylinder from Rocha-Pinto and Maciel (1996) is given (triangles).
 We have made use of the ONVECL abundance rather than the iron abundance. du order to be cousistent with chemical evolution models which use the iustantaueous recycling approximation.," We have made use of the oxygen abundance rather than the iron abundance, in order to be consistent with chemical evolution models which use the instantaneous recycling approximation."
 Indecd. IRA is," Indeed, IRA is"
10 periodic solution (calculated via the Ewald (1921) Upproximation) ancl the isolated solution on the eric.,the periodic solution (calculated via the Ewald (1921) approximation) and the isolated solution on the grid.
 Our approach is thus related to the method described in Εις et al. (, Our approach is thus related to the method described in Huss et al. (
1997).,1997).
 Llowever. they apply the combined force calculation via a periodic PAL scheme and direct summation on only to the particle distribution in the center of their simulation volume.," However, they apply the combined force calculation via a periodic PM scheme and direct summation on only to the particle distribution in the center of their simulation volume."
 In this region. they compute the dynamical evolution of a galaxy cluster with high resolution. whereas the remaining part of the volume is treated. by a pure PM scheme ancl provides the tidal torques for the cluster in the center.," In this region, they compute the dynamical evolution of a galaxy cluster with high resolution, whereas the remaining part of the volume is treated by a pure PM scheme and provides the tidal torques for the cluster in the center."
 For this reason. they do not have to handle possible force misalignmentD in the border region5 of the simulation cube.," For this reason, they do not have to handle possible force misalignment in the border region of the simulation cube."
 In our situation. studving the evolution and fragmentation in the interior of molecular. clouds. we need high resolution and periodicity throughout the simulation volume.," In our situation, studying the evolution and fragmentation in the interior of molecular clouds, we need high resolution and periodicity throughout the simulation volume."
 Ensuring numerical stability throughout he complete volume requires additional care in combining roth methods. the PAL scheme andGRAPE.," Ensuring numerical stability throughout the complete volume requires additional care in combining both methods, the PM scheme and."
 This paper is organised as follows: The next section gives a brief overview of the main properties of the special purpose hardwareGRAPE: then follows a description. of he Ewalel method which. leads to the construction of a Green's function for a periodic. correction., This paper is organised as follows: The next section gives a brief overview of the main properties of the special purpose hardware; then follows a description of the Ewald method which leads to the construction of a Green's function for a periodic correction.
 This. Green's function is applied. in a PM like scheme. to obtain a (corrective) force term which is added to the particle forces inG," This Green's function is applied in a PM like scheme, to obtain a (corrective) force term which is added to the particle forces in."
RAPESPII. Section 4— investigates the performance of the method in selected test cases and. suggests. wavs of optimising the algorithm., Section \ref{opt-test} investigates the performance of the method in selected test cases and suggests ways of optimising the algorithm.
 The next section applies the above to a preliminary study of the fragmentation. process in the interior of molecular clouds and the section 6 finally concludes with a summary of this paper., The next section applies the above to a preliminary study of the fragmentation process in the interior of molecular clouds and the section \ref{summary} finally concludes with a summary of this paper.
 is a special purpose hardware device. which calculates the forces and the potential in the gravitational N-body problem by direct summation on a specifically designed chip with high ellicieney. (Sugimoto et al.," is a special purpose hardware device, which calculates the forces and the potential in the gravitational N-body problem by direct summation on a specifically designed chip with high efficiency (Sugimoto et al."
 1990. Ebisuzaki et al.," 1990, Ebisuzaki et al."
 1993)., 1993).
 We use the currently distributed version.SAFE. which contains 8 chips on one hoard and therefore can compute the forces on S particles in. parallel.," We use the currently distributed version, which contains 8 chips on one board and therefore can compute the forces on 8 particles in parallel."
 The board. is connected via à standard. VALUE interface to the host computer. in our case à SUN workstation.," The board is connected via a standard VME interface to the host computer, in our case a SUN workstation."
 € and FORTRAN libraries provide the software interface between the user's program and the board., C and FORTRAN libraries provide the software interface between the user's program and the board.
 The computational speed of is approximately 5 Cllops.," The computational speed of is approximately $5\,$ Gflops."
 The force law is hardwired to be a Plummer law. Llere £ is the index of the particle for which the force is calculated and |j enumerates the particles which exert the [orce: e; is the gravitational smoothing length of particle /: C and m; and mj; are Newton's constant and the particle masses. respectively.," The force law is hardwired to be a Plummer law, Here $i$ is the index of the particle for which the force is calculated and $j$ enumerates the particles which exert the force; $\epsilon_i$ is the gravitational smoothing length of particle $i$; $G$ and $m_i$ and $m_j$ are Newton's constant and the particle masses, respectively."
 To perform the force caleulation for à single particle. needs No|19 clock excles.," To perform the force calculation for a single particle, needs $N+19$ clock cycles."
 To achieve a high speed. concessions in the accuracy of the force calculations had to be made: internally works with a 20 bit fixed point. number format for particle positions. with a 56 bit fixed point number format for the forces and a 14 bit logarithmic number format for the masses (Okumura et al.," To achieve a high speed, concessions in the accuracy of the force calculations had to be made: internally works with a 20 bit fixed point number format for particle positions, with a 56 bit fixed point number format for the forces and a 14 bit logarithmic number format for the masses (Okumura et al."
 1993)., 1993).
 Conversion to and from that internal number representation is handled by the interface software., Conversion to and from that internal number representation is handled by the interface software.
 The number format limits the spacial resolution in a simulation and constrains the force accuracy., The number format limits the spacial resolution in a simulation and constrains the force accuracy.
 However. for collisionless ον svstems. the forces on a single particle typically. need not be known better than up to an error of about one »ercent.," However, for collisionless N-body systems, the forces on a single particle typically need not be known better than up to an error of about one percent."
 In that respect.GRAVE is comparable to the widely used schemes (e.g. Barnes Hut 1986).," In that respect, is comparable to the widely used schemes (e.g. Barnes Hut 1986)."
 Besides he low precision board. there is a high precision version available.LIAR. which was specifically designed for reatment of collision dominated systems (Makino. Kokubo and Taiji 1993).," Besides the low precision board, there is a high precision version available, which was specifically designed for treatment of collision dominated systems (Makino, Kokubo and Taiji 1993)."
 In addition to integrate the time evolution of an N-bodvy system by direct summation. has been incorporated into à algorithm (Makino Funato 1993) and. more recentIv. into a P3M code (DBrieu. Summers Ostriker 1995).," In addition to integrate the time evolution of an N-body system by direct summation, has been incorporated into a algorithm (Makino Funato 1993) and, more recently, into a P3M code (Brieu, Summers Ostriker 1995)."
 Besides forces ancl potential. GRAPE returns the List of neighbours for cach particle / within a specified racius 7.," Besides forces and potential, returns the list of neighbours for each particle $i$ within a specified radius $h_i$."
 This feature makes it attractive for use in smoothed particle hyvdrodynamies (Unemura 1993). where eas properties are derived by averaging over a localised ensemble of discrete particles.," This feature makes it attractive for use in smoothed particle hydrodynamics (Unemura 1993), where gas properties are derived by averaging over a localised ensemble of discrete particles."
 For particles near the surface of the integration volume. “ghost” particles are created to. correctly. extend the neighbour search. bevond the cell border: no forces are computed for these particles.," For particles near the surface of the integration volume, “ghost” particles are created to correctly extend the neighbour search beyond the cell border; no forces are computed for these particles."
 For periodic boundaries. the ehosts are replicas of particles from the opposite side of the cell.," For periodic boundaries, the ghosts are replicas of particles from the opposite side of the cell."
 Following the strategy. described in Steinmetz (1996). we have modified our SPLI code (Benz 1990. Date et al.," Following the strategy described in Steinmetz (1996), we have modified our SPH code (Benz 1990, Bate et al."
 1995) to incorporateGRAVE and furthermore added the ability to handle periodic bouncary concitions., 1995) to incorporate and furthermore added the ability to handle periodic boundary conditions.
 vote. that there also has been an attempt to treat periodic boundaries on a special purpose hardware device.WINE. using the I2wald method harcwired on a special chip (Fukushiee et al.," Note, that there also has been an attempt to treat periodic boundaries on a special purpose hardware device, using the Ewald method hardwired on a special chip (Fukushige et al."
 1993)., 1993).
 This. work has been successful., This work has been successful.
 Fukushige et al. (, Fukushige et al. (
1996) report the development of theGRAPE. which can handle an arbitrary central force law. including the Ewald method.,"1996) report the development of the, which can handle an arbitrary central force law, including the Ewald method."
 The general concept. of incorporating periodic boundaries into is the following: We use the Ewald (1921) method to optain a Cireen's function that accounts for the correction to potential and force for a svstem transformed from having vacuum boundary conditions to periodic ones., The general concept of incorporating periodic boundaries into is the following: We use the Ewald (1921) method to optain a Green's function that accounts for the correction to potential and force for a system transformed from having vacuum boundary conditions to periodic ones.
 computes the forces and potential for the isolated system., computes the forces and potential for the isolated system.
 The Green's function is then used in a PM like scheme to assign a correction term due to periodicity to eachparticle in the simulation., The Green's function is then used in a PM like scheme to assign a correction term due to periodicity to eachparticle in the simulation.
observations made with the Spitzer Space Telescope. which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under NASA contract 1107.,"observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under NASA contract 1407."
 Support for this work was provided by NASA through contract 1255091. issued bx JPL/Caltech.," Support for this work was provided by NASA through contract 1255094, issued by JPL/Caltech."
 ZB received support from UWnnearian ΟΤΙΝΑ Curauts TSOL9872 aud TO052., ZB received support from Hungarian OTKA Grants TS049872 and T049082.
"SF, composite, and AGNs galaxies.","SF, composite, and AGNs galaxies."
 In Sect., In Sect.
 4 we analyzed the evolution of the BPT diagrams., 4 we analyzed the evolution of the BPT diagrams.
 In Sect., In Sect.
" 5 we investigate the evolution of the mass-metallicity and luminosity-metallicity relations, whereas in Sect."," 5 we investigate the evolution of the mass-metallicity and luminosity-metallicity relations, whereas in Sect."
" 6 we discuss the relations between the SFR and SSFR with stellar mass and metallicity, as well as the morphology of our galaxies using colors, concentration index, and SSFRs."," 6 we discuss the relations between the SFR and SSFR with stellar mass and metallicity, as well as the morphology of our galaxies using colors, concentration index, and SSFRs."
" Finally, conclusions are given in Sect."," Finally, conclusions are given in Sect."
 7., 7.
 We selected emission line galaxies from SDSS-DR5 (Adelman-McCarthy et al., We selected emission line galaxies from SDSS–DR5 (Adelman–McCarthy et al.
 2007)., 2007).
 Data were taken with a 2.5 m telescope located at Apache Point Observatory (Gunn et al., Data were taken with a 2.5 m telescope located at Apache Point Observatory (Gunn et al.
 2006)., 2006).
" The SDSS spectra were obtained using 3 arcsec diameter fibres, covering a wavelength range of 3800-9200A, and with a mean spectral resolution 4/AA ~ 1800."," The SDSS spectra were obtained using 3 arcsec diameter fibres, covering a wavelength range of 3800-9200, and with a mean spectral resolution $\lambda$ $\Delta\lambda$ $\sim$ 1800."
 The SDSS-DRS spectroscopy database contains spectra for ~ 10° objects over ~ 5700 deg”., The SDSS–DR5 spectroscopy database contains spectra for $\sim$ $\rm{10}^6$ objects over $\sim$ 5700 $\rm{deg}^2$.
 Further technical details can be found in Stoughton et al. (, Further technical details can be found in Stoughton et al. (
2002).,2002).
" We used the SDSS-DR5 spectra from the STARLIGHTdatabase!, which were processed with the STARLIGHT spectral synthesis code, developed by Cid Fernandes and collaborators (Cid Fernandes et al."," We used the SDSS–DR5 spectra from the STARLIGHT, which were processed with the STARLIGHT spectral synthesis code, developed by Cid Fernandes and collaborators (Cid Fernandes et al."
" 2005, 2007, Mateus et al."," 2005, 2007, Mateus et al."
" 2006, Asari et al."," 2006, Asari et al."
 2007)., 2007).
" From the spectra, the STARLIGHT code subtracts the continuum, obtaining the emission lines fluxes measurements for each galaxy."," From the spectra, the STARLIGHT code subtracts the continuum, obtaining the emission lines fluxes measurements for each galaxy."
" For each emission line, the STARLIGHT code returns the rest frame flux and its associated equivalent width, linewidth, velocity displacement relative to the rest-frame wavelength and the S/N of the fit."," For each emission line, the STARLIGHT code returns the rest frame flux and its associated equivalent width, linewidth, velocity displacement relative to the rest–frame wavelength and the S/N of the fit."
" In the case of Balmer lines, the underlying stellar absorption was corrected by the STARLIGHT code using synthetic spectra obtained by fitting an observed spectrum with a combination of 150 simple stellar populations (SSPs) from the evolutionary synthesis models of Bruzual & Charlot (2003)."," In the case of Balmer lines, the underlying stellar absorption was corrected by the STARLIGHT code using synthetic spectra obtained by fitting an observed spectrum with a combination of 150 simple stellar populations (SSPs) from the evolutionary synthesis models of Bruzual $\&$ Charlot (2003)."
" From the full set of galaxies, we only consider galaxies whose spectra show in emission the Ha,Hf, [Nu], [On] A3727, ui] 44959, ur] 45007, r] 26300, [51] lines."," From the full set of galaxies, we only consider galaxies whose spectra show in emission the ], ] $\lambda$ 3727, ] $\lambda$ 4959, ] $\lambda$ 5007, ] $\lambda$ 6300, ] lines."
" We selected galaxies with a signal-to-noise ratio higher than 3c for theHa,Hf, and [Nr] lines."," We selected galaxies with a signal–to–noise ratio higher than $\sigma$ for the, and ] lines."
" In order to identify any evolution of galaxy parameters or relations, we divided our sample in four redshift intervals as follow: 0.04€z;« 0.1,0.1«τι «02,02<m03, 0.3€z3<0.4."," In order to identify any evolution of galaxy parameters or relations, we divided our sample in four redshift intervals as follow: $0.04 \leq z_0 < 0.1$, $0.1 \leq z_1 < 0.2$, $0.2 \leq z_2 < 0.3$, $0.3 \leq z_3 \leq 0.4$."
" The lower limit of zo corresponds to an aperture covering fraction of 20%, which is the minimum required to avoid domination of the spectrum by aperture effects (Kewley et al."," The lower limit of $z_0$ corresponds to an aperture covering fraction of $\%$, which is the minimum required to avoid domination of the spectrum by aperture effects (Kewley et al."
 2005)., 2005).
" This classification give us 85931 galaxies for zo, 48888 galaxies for zi, 3278 galaxies for z2, and 199 galaxies for z3."," This classification give us 85931 galaxies for $z_0$, 48888 galaxies for $z_1$, 3278 galaxies for $z_2$ , and 199 galaxies for $z_3$."
" We selected galaxies with an apparent Petrosian r magnitude of 14.5<r17.77 in the redshift samples zo, Z1, and z», which yields 82884, 44763, and 1802 galaxies, respectively, corresponding to the magnitude completeness at these redshifts (see Fig."," We selected galaxies with an apparent Petrosian $r$ magnitude of $14.5 < r < 17.77$ in the redshift samples $z_0$ , $z_1$, and $z_2$, which yields 82884, 44763, and 1802 galaxies, respectively, corresponding to the magnitude completeness at these redshifts (see Fig."
 1)., 1).
" Galaxies of the z3 sample have a different completeness range 16.9 <r 18.8, as observed in Fig."," Galaxies of the $z_3$ sample have a different completeness range 16.9 $< r <$ 18.8, as observed in Fig."
" 1, giving 119 galaxies."," 1, giving 119 galaxies."
" We used the zo and z; sample of galaxies with its respective completeness, but for galaxies of samples ζ2 and z3 we used both, those in the completeness range, and those out of the completeness range."," We used the $z_0$ and $z_1$ sample of galaxies with its respective completeness, but for galaxies of samples $z_2$ and $z_3$ we used both, those in the completeness range, and those out of the completeness range."
" The reason for this, is to improve the galaxy statistics by increasing the number of them."," The reason for this, is to improve the galaxy statistics by increasing the number of them."
" As we shown in the next sections, the main results are similar using galaxies in the magnitude completeness and galaxies of the total sample."," As we shown in the next sections, the main results are similar using galaxies in the magnitude completeness and galaxies of the total sample."
" In Section 3, to study the S2N2 diagram as a segregator of different types of galaxies, we adopt the “main galaxy sample"" (e.g. Strauss et al."," In Section 3, to study the S2N2 diagram as a segregator of different types of galaxies, we adopt the $``$ main galaxy sample"" (e.g. Strauss et al."
" 2002) with Petrosian r magnitudes in the range 14.5 «r 17.77, and the redshift interval zo, taking into account all the emission lines and the signal-to-noise ratio mentioned above, which yields 82884 galaxies."," 2002) with Petrosian $r$ magnitudes in the range 14.5 $< r <$ 17.77, and the redshift interval $z_0$, taking into account all the emission lines and the signal–to–noise ratio mentioned above, which yields 82884 galaxies."
" In this Section, we used AGN, composite, and SF galaxies."," In this Section, we used AGN, composite, and SF galaxies."
" In Section 4, we study the evolution of the three BPT diagrams using the four redshift intervals, magnitude intervals, and the signal-to-noise restrictions mentioned above."," In Section 4, we study the evolution of the three BPT diagrams using the four redshift intervals, magnitude intervals, and the signal–to–noise restrictions mentioned above."
" In this Section, we used AGN, composite, and SF galaxies."," In this Section, we used AGN, composite, and SF galaxies."
" Also,to study the metallicity evolution of the SF galaxies of the BPT diagrams (see Fig."," Also,to study the metallicity evolution of the SF galaxies of the BPT diagrams (see Fig."
" 6), weusedthe sample of Sections 5 and 6, mentioned below."," 6), weusedthe sample of Sections 5 and 6, mentioned below."
are associated with decreasing fields.,are associated with decreasing fields.
" The fact that /3,6/3; is more likelv to be negative than positive is consistent with the results of Sections 4. and 5..."," The fact that $B_l
\delta B_l$ is more likely to be negative than positive is consistent with the results of Sections \ref{fieldchanges} and \ref{fluxchanges}."
 Confirmation with a full force estimate niust await a sizable sample of good. high-cadence vector data.," Confirmation with a full force estimate must await a sizable sample of good, high-cadence vector data."
 IIudson et al. (, Hudson et al. (
2003) estimated that forces of size 107? dvnes can be important [or the physics of seismic waves.,2008) estimated that forces of size $10^{22}$ dynes can be important for the physics of seismic waves.
 Several examples in Figure 13. have force budgets of this size., Several examples in Figure \ref{forcesfig} have force budgets of this size.
 Given that (he forces calculated here involve only those pixels well modeled by Equation (1)). some of these forces might be significantly underestimated.," Given that the forces calculated here involve only those pixels well modeled by Equation \ref{atancurve}) ), some of these forces might be significantly underestimated."
 Recently. Wang Liu (2010) applied Equation (5)) to à BBSO vector data set [or the 2002 July 26 MS.7 flare and found a vertical force change of 1.6x107? dwnes.," Recently, Wang Liu (2010) applied Equation \ref{zforce}) ) to a BBSO vector data set for the 2002 July 26 M8.7 flare and found a vertical force change of $1.6\times 10^{22}$ dynes."
 Our estimated force change lor the 2002 July 26 M87 flare using Equation (6)) is 41.3x107! dvnes., Our estimated force change for the 2002 July 26 M8.7 flare using Equation \ref{losforce}) ) is $4.3\times 10^{21}$ dynes.
 The difference between the force change estimates [rom GONG longitudinal data and BBSO vector data is perhaps mostly due to the inclusion of the transverse field in Wane Liu's caleulation., The difference between the force change estimates from GONG longitudinal data and BBSO vector data is perhaps mostly due to the inclusion of the transverse field in Wang Liu's calculation.
 We note. however. that Martínnez-Oliveros Donea (2009) did not find good spatial correspondence between locations of abrupt. signilicant field changes and seismic sources in the (wo flares that thev stucied The most impressive estimated force budget in our data set is for the 2006 December 6 N6.5 flare featured in Figures 3--5..," We note, however, that Martínnez-Oliveros Donea (2009) did not find good spatial correspondence between locations of abrupt, significant field changes and seismic sources in the two flares that they studied The most impressive estimated force budget in our data set is for the 2006 December 6 X6.5 flare featured in Figures \ref{remap}- \ref{mosaic}."
 This flare has been associated with a Moreton wave. studied in detail by Balasubramaniam et al. (," This flare has been associated with a Moreton wave, studied in detail by Balasubramaniam et al. ("
2010) using lla images from the Improved solar Observing Optical Network (ISOON).,2010) using $\alpha$ images from the Improved Solar Observing Optical Network (ISOON).
 This Moreton wave traveled [rom its source al approximately SOGEG3 at about 800 km/s with an azimuthal span of about 2707., This Moreton wave traveled from its source at approximately S06E63 at about 800 km/s with an azimuthal span of about $^{\circ}$.
 A blow-up of the East-Himb portion of the full-disk magnetogram of Figure 2 is plotted in Figure 14.. showing the flaring active region. AR 10930.," A blow-up of the East-limb portion of the full-disk magnetogram of Figure \ref{fulldisk} is plotted in Figure \ref{centroids}, showing the flaring active region, AR 10930."
 Also plotted is an ISOON II-o. image taken αἱ the height of the flare showing the flare emission., Also plotted is an ISOON $\alpha$ image taken at the height of the flare showing the flare emission.
 As Figure 14. shows. the estimated central source was rather distant. 35-75 Mm distant according to Balasubramaniam et al.," As Figure \ref{centroids} shows, the estimated central source was rather distant, 35-75 Mm distant according to Balasubramaniam et al."
's estimate. [rom the centroid of RIIESSI X-rav and ISOON white light emission.,"'s estimate, from the centroid of RHESSI X-ray and ISOON white light emission."
 This flare was also well observed by the Cerro Tololo station of the GONG network., This flare was also well observed by the Cerro Tololo station of the GONG network.
 The cata set for this flare has the largest detected umber of pixels where large. clean. permanent stepwise [field changes were detected of any of the 77 Lares studied here. as well as the most impressive force budget.," The data set for this flare has the largest detected number of pixels where large, clean, permanent stepwise field changes were detected of any of the 77 flares studied here, as well as the most impressive force budget."
 The total forces involved amount to about 2x1077 dvnes., The total forces involved amount to about $2\times 10^{22}$ dynes.
 The spatial distribution of the changes occurring between 1830 UT and 1850 UT is shown in Figure 15.., The spatial distribution of the changes occurring between 1830 UT and 1850 UT is shown in Figure \ref{forcesbala}.
 The forces are organized in two regions. a small region of relatively weak forces close to the neutral line mostly directed towards (he observer (positive) and a more extended region to the south and west of somewhat larger lorces directed towards the Sun (negative). (," The forces are organized in two regions, a small region of relatively weak forces close to the neutral line mostly directed towards the observer (positive) and a more extended region to the south and west of somewhat larger forces directed towards the Sun (negative). ("
Note that in (his example (he line of sight is tilted at GO’ with respect to the local vertical.),Note that in this example the line of sight is tilted at $60^{\circ}$ with respect to the local vertical.)
 While the group of positive forces is close to the centroid of the X-ray emission. the largest forces are clustered to (he west. closer to the location of the focal point of the Moreton wave (see Figure 1H)).," While the group of positive forces is close to the centroid of the X-ray emission, the largest forces are clustered to the west, closer to the location of the focal point of the Moreton wave (see Figure \ref{centroids}) )."
 The centroids in Figure 14. show that the changes began near the sunspot and then propagated to the South aud West., The centroids in Figure \ref{centroids} show that the changes began near the negative-polarity sunspot and then propagated to the South and West.
 The field changes associated, The field changes associated
The majority of the inhomogeneities in (he chemical composition of Globular Cluster (GC) stars see ee. Snecden(1999):Grattonοἱal.(2004)-— appear due to primordial enrichment by hot-CNO eveled material processed in stars belonging to a first stellar generation.,"The majority of the inhomogeneities in the chemical composition of Globular Cluster (GC) stars —see e.g. \citet{sne99,gratton-araa2004}- — appear due to primordial enrichment by hot-CNO cycled material processed in stars belonging to a first stellar generation."
 Either massive AGB envelopes subject to hot bottom burning. (e.g.Venturaetal.2001;Ven-tura.D'Antona.&Mazzitelli2002) or the envelopes of massive lastly rotating stars (FIRMS) (Decressinetal.2007) are an ideal environment to manufacture elements through nuclear reactions in which proton captures are involved.," Either massive AGB envelopes subject to hot bottom burning, \citep*[e.g.][]{ventura2001, ventura2002}
 or the envelopes of massive fastly rotating stars (FRMS) \citep{decressin2007} are an ideal environment to manufacture elements through nuclear reactions in which proton captures are involved."
 A second stellar generation would then be born Irom the ejecta of these stars. either directly. or diluted with pristine material.," A second stellar generation would then be born from the ejecta of these stars, either directly, or diluted with pristine material."
 There are however problems with these scenario., There are however problems with these scenario.
 For the AGB progenitors. problems concern the chemistry of the anomalous stars (see Denissenkov&IHerwig2003:Venturaοἱal.2004:Fen-nerοἱal. 2004.. but consider also Ventura&D'Antona 2005)). and the requirements for (he mass budget necessary to produce a second stellar generation that. today. is generally about as numerous as the first one (D'Antona&Caloi2004:DekkiNorris2006:Prant-gos&Charbonnel 2006).," For the AGB progenitors, problems concern the chemistry of the anomalous stars (see \citealt{denis2003,ventura2004,
fenner2004}, but consider also \citealt{venturadantona2005a}) ), and the requirements for the mass budget necessary to produce a second stellar generation that, today, is generally about as numerous as the first one \citep{dantona-caloi2004, bekki-norris2006, 
prantzos2006}."
. The FRAIS model is somewhat less problematic concerning the constraints on the IME. but poses. in any case. severe problems for “normal GCs. that show a unique metallicity for all stars (with the notable exception of Cen)): this model requires (hat the envelopes of these stars contribute to the second star formation stage. bul (he supernova ejecta are expelled from (he clusters and do not enrich the second generation matter by heavy elements.," The FRMS model is somewhat less problematic concerning the constraints on the IMF, but poses, in any case, severe problems for “normal” GCs, that show a unique metallicity for all stars (with the notable exception of ): this model requires that the envelopes of these stars contribute to the second star formation stage, but the supernova ejecta are expelled from the clusters and do not enrich the second generation matter by heavy elements."
 Both scenarios predict that the stars showing chemical anomalies also must haveabundance., Both scenarios predict that the stars showing chemical anomalies also must have.
 Helium ancl the hot.CNO products in the FRAIS come to the stellar surface by means of the chemical mixing associated with the transport of angular momentum through the stellar radiative lavers. during the phase of hydrogen burning.," Helium and the hot–CNO products in the FRMS come to the stellar surface by means of the chemical mixing associated with the transport of angular momentum through the stellar radiative layers, during the phase of hydrogen burning."
 In the ease of AGBs. the helium enrichment is mainly due to the second dredge up (2DU) which precedes the thermal pulse phase.," In the case of AGBs, the helium enrichment is mainly due to the second dredge up (2DU) which precedes the thermal pulse phase."
 Also the third dredge up (23DU) contributes to helium enrichment in the envelopes. bul a further observational constraint is (hat the sum ol CNO abundancees remains remarkably constant in normal and anomalous stars hen&Meléndez 2005).," Also the third dredge up (3DU) contributes to helium enrichment in the envelopes, but a further observational constraint is that the sum of CNO abundances remains remarkably constant in normal and anomalous stars \citep[e.g.][]{cohenmelendez2005}."
. Consequently. the masses which can be responsible lor the chemical enrichment must be high enough to be subject to a few episodes of the third dredge up. that bringse into the envelope the products of helium burning.," Consequently, the masses which can be responsible for the chemical enrichment must be high enough to be subject to a few episodes of the third dredge up, that brings into the envelope the products of helium burning."
e A variation in helium content is immecdiatelv reflected on the morphology of the horizontal branch (11B). which amplifies anv evolutionary difference among the cluster stus.," A variation in helium content is immediately reflected on the morphology of the horizontal branch (HB), which amplifies any evolutionary difference among the cluster stars."
 Ilelium enrichment by à small factor (up to Y0.28— 0.30) allowed in fact an easy. interpretation ol some puzzles posed bv the HB (blue tails. gaps) (D'Antonaοἱal. 2004).," Helium enrichment by a small factor (up to $\sim 0.28 - 0.30$ ) allowed in fact an easy interpretation of some puzzles posed by the HB (blue tails, gaps) \citep{dantona2002,dantona-caloi2004}."
. Helium. variations from the MS location are much harder to be detected. but the," Helium variations from the MS location are much harder to be detected, but the"
order to search for coherent periodicities.,order to search for coherent periodicities.
 For the analysis of data. we used combined synchronized EPIC- and MOS lighteurves with 2.6 s time bins to improve sensitivity.," For the analysis of data, we used combined synchronized EPIC-pn and MOS lightcurves with 2.6 s time bins to improve sensitivity."
 We found strong peak in the Fourier spectrum at the [requeney of 1.3 10.7 Uz (Pie. 2. panel).," We found strong peak in the Fourier spectrum at the frequency of $\sim$ $\times10^{-3}$ Hz (Fig. \ref{pds_efold}, )."
 The strength of the peak in the Fourier spectrum corresponds to the period detection confidence of ~3«10.0 (Vaughanetal. 1994)., The strength of the peak in the Fourier spectrum corresponds to the period detection confidence of $\sim 3\times10^{-9}$ \citep{Vaughan94}.
. To estimate the pulsation period. more precisely. we used. an epoch folding technique. assuming no pulse period change during 2004 November 23 observation.," To estimate the pulsation period more precisely, we used an epoch folding technique, assuming no pulse period change during 2004 November 23 observation."
 The most likely value of the pulsation period. 165.6 s (Table 2)) was obtained fitting the peak in the X7 versus trial period distribution with a Gaussian.," The most likely value of the pulsation period, 765.6 s (Table \ref{timing_spec_par}) ) was obtained fitting the peak in the $\chi^{2}$ versus trial period distribution with a Gaussian."
 The period errors in Table 2. were computed following the procedure. described. in Leahy (1987)., The period errors in Table \ref{timing_spec_par} were computed following the procedure described in Leahy (1987).
 The source Lehteurves were folded. using the periods determined. from epoch folding analysis., The source lightcurves were folded using the periods determined from epoch folding analysis.
 Ehe resulting folded lighteurve in the 0.3-7 keV. enerev band during 2004. November 23 observation. is shown in Fig., The resulting folded lightcurve in the 0.3-7 keV energy band during 2004 November 23 observation is shown in Fig.
 2. panel., \ref{pds_efold} ).
 Phe source. has. quasi- pulse profile in the 0.3-7 keV energy band with a xulsed fraction of 38d3., The source has quasi-sinusoidal pulse profile in the 0.3-7 keV energy band with a pulsed fraction of $38\pm3$.
 The pulsed fraction was defined as (Luus-Llaiu)/(Gauax linia). where buns and. Lain represent source intensities at |the maximum and minimum of the pulse wolile. excluding background. photons.," The pulsed fraction was defined as $_{\rm max}$ $_{\rm min}$ $_{\rm max}$ $_{\rm min}$ ), where $_{\rm max}$ and $_{\rm min}$ represent source intensities at the maximum and minimum of the pulse profile, excluding background photons."
 ‘To investigate energy. dependence of the source pulse oolile. we extracted lieht curves in the soft. (0.3-2. keV) and hard (2-7 keV) bands for 2004 Nov. 23 observation. and folded them at the corresponding best pulsation period. (Vig. 3)).," To investigate energy dependence of the source pulse profile, we extracted light curves in the soft (0.3-2 keV) and hard (2-7 keV) bands for 2004 Nov. 23 observation, and folded them at the corresponding best pulsation period (Fig. \ref{mod_energy_depend}) )."
 Both bands show quasi-sinusoidal pulse profiles., Both bands show quasi-sinusoidal pulse profiles.
 Because of the relatively. poor statistics. we could. detect only marginal dillerence between source pulse profiles at low and high energies. which have background-corrected pulsed ractions of 34+ and 40+," Because of the relatively poor statistics, we could detect only marginal difference between source pulse profiles at low and high energies, which have background-corrected pulsed fractions of $\pm$ and $\pm$ ."
place at which the DNO Ist harmonic begins.,place at which the DNO 1st harmonic begins.
 The second harmonic starts at Z77O.S d. which is where the DNO second harmonics begin to become very prominent.," The second harmonic starts at $T \sim 0.8$ d, which is where the DNO second harmonics begin to become very prominent."
 Again we can produce a modified plot. given in Fig. H1..," Again we can produce a modified plot, given in Fig. \ref{dno4fig12},"
 in which we have shown the implied fundamental ΟΡΟ periods., in which we have shown the implied fundamental QPO periods.
 The solid line in Fig., The solid line in Fig.
 LL is a least squares fit to these points and has the equation. for the range 0.2 < Z(d) « 1.15. where the formal errors on the parameters derived from the least squares fit are given in brackets in the above equation (I5q. 2)).," \ref{dno4fig12} is a least squares fit to these points and has the equation, for the range 0.2 $<$ $T(d)$ $<$ 1.15, where the formal errors on the parameters derived from the least squares fit are given in brackets in the above equation (Eq. \ref{dno4eq2}) )."
 Equations 1 and 2 show that the ratio Z? is not constant during the factor of three increase in. fundamental perioc seen in VW Livi., Equations \ref{dno4eq1} and \ref{dno4eq2} show that the ratio $R$ is not constant during the factor of three increase in fundamental period seen in VW Hyi.
 In fact. 2 changes from 15.2 at 7 — 0.2« to 11.0 at Y= 1.15 d. which is typical of the range of values seen among other CVs. e.g. Mauche (2002) found / = 10.4and. 11.4 for two runs on SS ςνο and in Paper LE we founc R= 15 for LU Alen ane V893 Sco.," In fact, $R$ changes from 15.2 at $T$ = 0.2 d to 11.0 at $T$ = 1.15 d, which is typical of the range of values seen among other CVs, e.g., Mauche (2002) found $R$ = 10.4and 11.4 for two runs on SS Cyg and in Paper III we found $R$ = 15 for TU Men and V893 Sco."
 Values of 2 outside of this range may arise from unwittingly combining DNO ane QPO fundamentals with harmonics., Values of $R$ outside of this range may arise from unwittingly combining DNO and QPO fundamentals with harmonics.
 1n order to prepare for our suggested interpretation of the complicated behaviour of VW. Livi we recall some of the structural components Chat have been used to explain the rapid brightness modulations in CVs., In order to prepare for our suggested interpretation of the complicated behaviour of VW Hyi we recall some of the structural components that have been used to explain the rapid brightness modulations in CVs.
 There are at least two ways in which frequency. doubling of DNOs might arise., There are at least two ways in which frequency doubling of DNOs might arise.
 Given the general structure of the LIALA model. one possibility is a transition from single-role to two-pole accretion.," Given the general structure of the LIMA model, one possibility is a transition from single-pole to two-pole accretion."
 Vhis might be due simply to a change in the fickd-threacing region as Al changes. resulting in feeding of two accretion zones instead of one.," This might be due simply to a change in the field-threading region as $\dot{M}$ changes, resulting in feeding of two accretion zones instead of one."
 A transition ο quadrupole field from cipole field configuration nav occur when the inner edge of the accretion disce is pushed so close o the white dwarf surface that the less compressible higher order multipole field components determine the position of he inner edge of the disc (Lamb LOSS): but in reverse this would reduce the number of accretion zones as M. decreases and is thus unlikely to be related to frequeney doubling in VW Livi late in outburst when AZ is decreasing., A transition to quadrupole field from dipole field configuration may occur when the inner edge of the accretion disc is pushed so close to the white dwarf surface that the less compressible higher order multipole field components determine the position of the inner edge of the disc (Lamb 1988); but in reverse this would reduce the number of accretion zones as $\dot{M}$ decreases and is thus unlikely to be related to frequency doubling in VW Hyi late in outburst when $\dot{M}$ is decreasing.
 ]t is not easy to make quantitative predictions about he changes in magnetic field structure that would lead to he above ellect. but another possibility. is that frequency doubling arises through a ecometrical rearrangement. which allows us (or the reprocessing regions of the accretion disc) ο view two poles instead. of one.," It is not easy to make quantitative predictions about the changes in magnetic field structure that would lead to the above effect, but another possibility is that frequency doubling arises through a geometrical rearrangement, which allows us (or the reprocessing regions of the accretion disc) to view two poles instead of one."
 The smoothness of the ransition to à dominant doubled DNO frequency in the sense that there is no concomitant apparent change in xhaviour (coherence. jump to longer periods. amplitude) suggests that no radical restructuring has occurred in the hreacing region or in the accretion zones.," The smoothness of the transition to a dominant doubled DNO frequency – in the sense that there is no concomitant apparent change in behaviour (coherence, jump to longer periods, amplitude) – suggests that no radical restructuring has occurred in the threading region or in the accretion zones."
 We have therefore considered. models. in which either. the upper accretion zone becomes visible as the inner edge of the accretion disc retreats outwards with decreasing Ad. or the second accretion zone becomes visible.," We have therefore considered models in which either the upper accretion zone becomes visible as the inner edge of the accretion disc retreats outwards with decreasing $\dot{M}$ , or the second accretion zone becomes visible."
 Such a transition is well, Such a transition is well
Continuum emission from interstellar dust is a promising alternative way to determine the ISM content of dEs.,Continuum emission from interstellar dust is a promising alternative way to determine the ISM content of dEs.
 In particular. the ISM in the immediate. surroundings of M87 can be traced by dust emission. as strong radio continuum emission considerably reduces the ALFALFA detection sensitivity within 1° of M87 (Giovanellietal.2007)..," In particular, the ISM in the immediate surroundings of M87 can be traced by dust emission, as strong radio continuum emission considerably reduces the ALFALFA detection sensitivity within $\degr$ of M87 \citep{2007AJ....133.2569G}."
 To date. the Andromeda satellites 2205 and 1185 are the only dEs that have been detected in the far-infrared (Haas1998;Marleauetal.2006. 2009).," To date, the Andromeda satellites 205 and 185 are the only dEs that have been detected in the far-infrared \citep{Haas1998, Marleau2006, Marleau2009}."
. In particular. no dust emission has been detected from cluster dEs.," In particular, no dust emission has been detected from cluster dEs."
 While extinction features in optical images are indicative of dust in at least some cluster dEs (Ferrareseetal.2006;Lisker2006a.b).. the direct detection of dust emission has been hampered by the limited resolution and sensitivity of the previous generation of far-infrared instrumentation.," While extinction features in optical images are indicative of dust in at least some cluster dEs \citep{Ferrarese2006, Lisker2006_1, Lisker2006_2}, the direct detection of dust emission has been hampered by the limited resolution and sensitivity of the previous generation of far-infrared instrumentation."
 The advent of the Space Observatory (Pilbrattetal.2010) offers new possibilities for mapping the ISM in dEs., The advent of the Space Observatory \citep{Herschel} offers new possibilities for mapping the ISM in dEs.
 In this paper. we report on our search for far-infrared emission from dEs in the Virgo Cluster. based on the Science Demonstration Phase (SDP) observations. of the Virgo Cluster Survey (HeViCS!)).," In this paper, we report on our search for far-infrared emission from dEs in the Virgo Cluster, based on the Science Demonstration Phase (SDP) observations of the Virgo Cluster Survey )."
 In Sect.2," In Sect.,"
. we describe the observations. data reduction. and sample selection.," we describe the observations, data reduction, and sample selection."
 In Sect.3," In Sect.,"
. we present our detections and in Sect., we present our detections and in Sect.
 we discuss the results of a stacking analysis of the non-detections., we discuss the results of a stacking analysis of the non-detections.
 Sect., Sect.
 gives our summary., gives our summary.
 The central 44 deg? region of the Virgo Cluster was observed by on 29 November 2009 as part of the HeViCS SDP observations., The central $4\times4$ $^2$ region of the Virgo Cluster was observed by on 29 November 2009 as part of the HeViCS SDP observations.
 The observations were performed in parallel scai map mode. which means that both PACS (Poglitschetal.2010) and SPIRE (Griffinetal.2010) data were collected.," The observations were performed in parallel scan map mode, which means that both PACS \citep{PACS} and SPIRE \citep{SPIRE} data were collected."
" The scar speed was 60""//s and the field was covered in nominal anc orthogonal directions to minimize 1/f noise.", The scan speed was /s and the field was covered in nominal and orthogonal directions to minimize $1/f$ noise.
 The PACS anc SPIRE data were reduced with HIPE. with reduction scripts based on the standard reduction pipelines.," The PACS and SPIRE data were reduced with HIPE, with reduction scripts based on the standard reduction pipelines."
 For more details of the HeViCS SDP data reduction. we refer to Daviesetal. (2010)..," For more details of the HeViCS SDP data reduction, we refer to \citet{HeViCS-paper1}."
 We used the HeViCS SDP observations to search for dust emission from dEs in the central regions of the Virgo Cluster., We used the HeViCS SDP observations to search for dust emission from dEs in the central regions of the Virgo Cluster.
 Selecting the galaxies with morphological type dE and dSO in the Goldmine catalogue (Gavazzietal.2003) resulted in an initial sample of 370 dEs., Selecting the galaxies with morphological type dE and dS0 in the Goldmine catalogue \citep{Gavazzi2003} resulted in an initial sample of 370 dEs.
 After the removal of 16 galaxies classified as background sources and the exclusion of 115 sources with an additional optical source within a rradius. our final sample consisted of 239 dEs.," After the removal of 16 galaxies classified as background sources and the exclusion of 115 sources with an additional optical source within a radius, our final sample consisted of 239 dEs."
 Blind aperture photometry was applied to the PACS and SPIRE maps based on the positions of each of these 239 galaxies using Source Extractor (Bertin&Arnouts1996)., Blind aperture photometry was applied to the PACS and SPIRE maps based on the positions of each of these 239 galaxies using Source Extractor \citep{1996A&AS..117..393B}.
 From this sample of 239 sources. ΤΙ dEs were detected above 3c m the SPIRE 250 um image.," From this sample of 239 sources, 11 dEs were detected above $\sigma$ in the SPIRE 250 $\mu$ m image."
 Nevertheless. we only report the clear detection of 2 and951.. excluding all sources raising any kind of doubt.," Nevertheless, we only report the clear detection of 2 and, excluding all sources raising any kind of doubt."
 As such. 11502 and VCC7788 were ruled out because their FIR location differs more than FWHMyJ2 from their optical position.," As such, 1502 and 788 were ruled out because their FIR location differs more than FWHM/2 from their optical position."
 8832 was excluded because SDSS images detect à background source with substructure in the residual image. possibly indicating the presence of dust.," 832 was excluded because SDSS images detect a background source with substructure in the residual image, possibly indicating the presence of dust."
 Four other objects 7752. 8815. 11272. 11512) fall within a crowded field of sources. enhancing the chances of their being background sources.," Four other objects 752, 815, 1272, 1512) fall within a crowded field of sources, enhancing the chances of their being background sources."
 For instance. Binggelietal.(1985) reports the detection of a blue compact background galaxy. only 4 to the west of the nucleus of 8815.," For instance, \citet{Binggeli1985} reports the detection of a blue compact background galaxy, only $4\arcsec$ to the west of the nucleus of 815."
 Furthermore. 77352 and 11272 appear to have SEDs that do not correspond to a single grey body.," Furthermore, 752 and 1272 appear to have SEDs that do not correspond to a single grey body."
 This failure in fitting a SED might either be caused by poor flux determination (the error bars are large compared to the low number of counts in the detection) or be a true indication that both detections originate in sources other than dEs., This failure in fitting a SED might either be caused by poor flux determination (the error bars are large compared to the low number of counts in the detection) or be a true indication that both detections originate in sources other than dEs.
 The other two sources 11578 and 7767 are clearly detected in the SPIRE 250 uim image at 4.8 and 4c. respectively. but their high FIR flux relative to their faint optical appearance (both objects are low surface brightness [LSB] galaxies with B-band magnitudes of the order -2] mag) aroused suspicion.," The other two sources 1578 and 767 are clearly detected in the SPIRE 250 $\mu$ m image at 4.8 and $\sigma$, respectively, but their high FIR flux relative to their faint optical appearance (both objects are low surface brightness [LSB] galaxies with B-band magnitudes of the order $\sim$ 21 mag) aroused suspicion."
 Dust has already been detected in LSBs (Hinzetal.2007). but mainly in those with large amounts of star formation responsible for dust heating.," Dust has already been detected in LSBs \citep{Hinz2007}, but mainly in those with large amounts of star formation responsible for dust heating."
 A lack of evidence of recent star formation in 7767 and 11578 strongly argues against the association of the 250 jm) emission with the corresponding dEs., A lack of evidence of recent star formation in 767 and 1578 strongly argues against the association of the 250 $\mu$ m emission with the corresponding dEs.
 Furthermore. we can identify a background source in the immediate surroundings of 11578. which possibly contaminates the SPIRE 250 jm emission.," Furthermore, we can identify a background source in the immediate surroundings of 1578, which possibly contaminates the SPIRE 250 $\mu$ m emission."
 Nonetheless. these questionable detections may still be caused by dust emission from dEs.," Nonetheless, these questionable detections may still be caused by dust emission from dEs."
 Deeper fields in the future might reveal dust emission in PACS images. where the higher spatial resolution may be able to rule out possible contamination by background sources.," Deeper fields in the future might reveal dust emission in PACS images, where the higher spatial resolution may be able to rule out possible contamination by background sources."
 On the other hand. and are detected unambiguously in the SPIRE 250 pm image at 9.8 and 11.1. respectively (see Fig. 1).," On the other hand, and are detected unambiguously in the SPIRE 250 $\mu$ m image at 9.8 and $\sigma$ , respectively (see Fig. \ref{VCC781.pdf}) )."
 Adopting a fair 2c detection cutoff in the other filters. is also detected in the PACS 100 jm. PACS 160 um. SPIRE 350 jim. and SPIRE 500 pam bands.," Adopting a fair $\sigma$ detection cutoff in the other filters, is also detected in the PACS 100 $\mu$ m, PACS 160 $\mu$ m, SPIRE 350 $\mu$ m, and SPIRE 500 $\mu$ m bands."
 has additional detections only in both SPIRE 350 uim and 500 uim images., has additional detections only in both SPIRE 350 $\mu$ m and 500 $\mu$ m images.
 Unfortunately. it is impossible to measure fluxes for at 350 jim and 500 um because and its neighbouring source. SDSS J122655.56+114000.5. are not resolved at these longer wavelengths.," Unfortunately, it is impossible to measure fluxes for at 350 $\mu$ m and 500 $\mu$ m because and its neighbouring source, SDSS J122655.56+114000.5, are not resolved at these longer wavelengths."
 Since the SPIRE 250 um image contains a vast number of sources. we need to address the possibility of false detections due to background sources.," Since the SPIRE 250 $\mu$ m image contains a vast number of sources, we need to address the possibility of false detections due to background sources."
 To perform this check accurately. we search all relevant images/catalogs (SDSS. Goldmine. 2MASS. UKIDSS. Spirzer)) for sources within 18.1 “(FWHM of SPIRE 250 μπι) of the dwarf positions.," To perform this check accurately, we search all relevant images/catalogs (SDSS, Goldmine, 2MASS, UKIDSS, ) for sources within 18.1 (FWHM of SPIRE 250 $\mu$ m) of the dwarf positions."
 For1.. this identifies four SDSS sources within this PSF region that may be contaminating the dwarf emission.," For, this identifies four SDSS sources within this PSF region that may be contaminating the dwarf emission."
 The superior resolution in the PACS 160 jim image discards three of these optically identified objects. while the remaining source is rejected as a possible background candidate for the IR emission. based," The superior resolution in the PACS 160 $\mu$m image discards three of these optically identified objects, while the remaining source is rejected as a possible background candidate for the IR emission, based"
as well as a ¢epeidence upon location within the strip. a natural 'Onsequence of an effect tied to SULface eravity as weLas pulsation efficiency.,"as well as a dependence upon location within the strip, a natural consequence of an effect tied to surface gravity as well as pulsation efficiency."
 Iu Oreer to eliuuiuae that factor im characterizing οςoheid period changes. we have normalized the resulting valies of bhe light amplitude aud Pas ‘olows: (1) blue auplitudes AB were standardized hrough the ratio AB ‘A Büauax). where AB luis) Is the iaxiluiji value of AB for the stir nsation period. and (i) P was adjsted to the equivalent value for a Ceολα. with a oilsation »eriod of LOT using the eiupiricallv-obtaiied slope otted in Fie.," In order to eliminate that factor in characterizing Cepheid period changes, we have normalized the resulting values of blue light amplitude and $\dot{P}$ as follows: (i) blue amplitudes $\Delta B$ were standardized through the ratio $\Delta B$ $\Delta 
B$ (max), where $\Delta B$ (max) is the maximum value of $\Delta B$ for the star's pulsation period, and (ii) $\dot{P}$ was adjusted to the equivalent value for a Cepheid with a pulsation period of $10^{\rm d}$ using the empirically-obtained slope plotted in Fig."
 L., 4.
 Fie., Fig.
" 5 plots such data for Cepheids with wth«P10"" and increasing pulsation periods (P~ 20%.", 5 plots such data for Cepheids with $12^{\rm d} \leq P \leq 40^{\rm d}$ and increasing pulsation periods $P \simeq 20^{\rm d}$ ).
 The upper part of the diagram plots he individual data. while the middle part of he diagram plots runniug meaus for the data.," The upper part of the diagram plots the individual data, while the middle part of the diagram plots running means for the data."
 The lower part of the diagram is an alternate interpretation of the same data. as described (low.," The lower part of the diagram is an alternate interpretation of the same data, as described below."
 Similar plots are given in Fie., Similar plots are given in Fig.
 6 for Cepheids with 15<Pxs aud inercasing olsation periods (P~ 6G). aud in Fig.," 6 for Cepheids with $4^{\rm d} \leq P \leq 8^{\rm d}$ and increasing pulsation periods $P \simeq 6^{\rm d}$ ), and in Fig."
 7 for Cepheids with If<Pxsb and decreasing oulsatiou periods (2= 653., 7 for Cepheids with $4^{\rm d} \leq P \leq 8^{\rm d}$ and decreasing pulsation periods $P \simeq 6^{\rm d}$ ).
 Recall that larec values of DP shorid correspond to the hot side of the 1stability strip. and small values to the cool sido.," Recall that large values of $\dot{P}$ should correspond to the hot side of the instability strip, and small values to the cool side."
 The «ata for 20 Copheids (top portion of Fie., The data for $20^{\rm d}$ Cepheids (top portion of Fig.
 5) display a teudeucy for laree amplitude Cexieids to have rates of period iucrease typical of stars Iwing near the ceuter of the instability strip. with sanaller amplitude Cepheids falling towards the hot aud cool edges (larger aud smaller vaues of P. respectively).," 5) display a tendency for large amplitude Cepheids to have rates of period increase typical of stars lying near the center of the instability strip, with smaller amplitude Cepheids falling towards the hot and cool edges (larger and smaller values of $\dot{P}$, respectively)."
 The trend is more obvious wheu oue plots runuuiung five-poiut means of the same data. as in the middle section of Fie.," The trend is more obvious when one plots running five-point means of the same data, as in the middle section of Fig."
 5., 5.
 There are two long period Cepheids with anomalously laree values of P. SZ Cas aud AQ Pup. which are conceivably fifth crossers.," There are two long period Cepheids with anomalously large values of $\dot{P}$, SZ Cas and AQ Pup, which are conceivably fifth crossers."
 If they are omitted from the ΠοC» means and averagesOo over smiuller sunples are iucluded at the extremes of P. one obtains the results iu the lower portion of Fie.," If they are omitted from the running means and averages over smaller samples are included at the extremes of $\dot{P}$, one obtains the results in the lower portion of Fig."
 5. which are typical of uxcpendent cross-sectional wuplitude maps of the instability. strip.," 5, which are typical of independent cross-sectional amplitude maps of the instability strip."
 The scatter in P values evident iu the top portion of Fig., The scatter in $\dot{P}$ values evident in the top portion of Fig.
 5 is intrinsic to the sars. and is not the result of lavee unucertaiutfies iu the calculated values.," 5 is intrinsic to the stars, and is not the result of large uncertainties in the calculated values."
 Prestunably there are intrinsic plivsical differences from one Cepheid to arother that account for the scatter. as noted carlier.," Presumably there are intrinsic physical differences from one Cepheid to another that account for the scatter, as noted earlier."
 Differences im initial rotation velocity for the progcuiOY Inail-sequence y.ars nueht be the sole factor. even that thev would generate sufficicutly large variations in the abundances of the CNO clemens throughout the y.ar to affect the exteut of the blue loop stages (Nu&Li2001..," Differences in initial rotation velocity for the progenitor main-sequence stars might be the sole factor, given that they would generate sufficiently large variations in the abundances of the CNO elements throughout the star to affect the extent of the blue loop stages \citep{xl04}."
 The data for 6° Cepheids wit1 period increases (top porion of Fig., The data for $6^{\rm d}$ Cepheids with period increases (top portion of Fig.
 6) are more complicated., 6) are more complicated.
 Tt appears that the sample consists of two overlappingo eroups of objects. a feature that also appears iu 1C ταnume five-point means displaved iu the uiddle section of Fig.," It appears that the sample consists of two overlapping groups of objects, a feature that also appears in the running five-point means displayed in the middle section of Fig."
 6., 6.
 We assmnue that each oοτι] COsists of Cepheids displaviug an order of naenituce (factor of 10) variation in P. values Toni he hot to cool edges of the iustabilitv strip. as clisplaved bv the lone period Cepheids in Fie.," We assume that each group consists of Cepheids displaying an order of magnitude (factor of 10) variation in $\dot{P}$ values from the hot to cool edges of the instability strip, as displayed by the long period Cepheids in Fig."
 5. and use the resuts presented in the ower vortion of Fie.," 5, and use the results presented in the lower portion of Fig."
 5 as a tempate for the likely variatious in relative anplitide with P for shox )oriod Cepheids., 5 as a template for the likely variations in relative amplitude with $\dot{P}$ for short period Cepheids.
 W1011 the two eroups in Fie., When the two groups in Fig.
 6 are sermarvated iu sucl fashion aXd averages Over siialler samples are included at extreme values of P tor cach group. οιο obtaii the results in the lower vortion of Fig.," 6 are separated in such fashion and averages over smaller samples are included at extreme values of $\dot{P}$ for each group, one obtains the results in the lower portion of Fig."
 6., 6.
 The simsest explanation for the existence of two eroups among the shor period Cepheids is he existence of higher stri CLOSSILES ALLOic the stars. nauclv fifth crossine:4 for Cepheids undergoiug period increases.," The simplest explanation for the existence of two groups among the short period Cepheids is the existence of higher strip crossings among the stars, namely fifth crossings for Cepheids undergoing period increases."
 Similar results apply to the« ata for GT Cephieids with period «ecreases (Fig., Similar results apply to the data for $6^{\rm d}$ Cepheids with period decreases (Fig.
 7). when analyzed in simular fashion. vite the smaller sample size.," 7), when analyzed in similar fashion, despite the smaller sample size."
 The top yortion of Fig., The top portion of Fig.
 7 displavs excessive scatter. much Ijumτο that iu the top portion of Fig.," 7 displays excessive scatter, much like that in the top portion of Fig."
 6. with οily liusτα]. iuprovenient through ruunius five-point means (nikdle xrtion of Fie.," 6, with only marginal improvement through running five-point means (middle portion of Fig."
 T), 7).
 Restricting tιο data sets as above produces the results depiced in he lower portion of Fig., Restricting the data sets as above produces the results depicted in the lower portion of Fig.
 7. whic1 suggests au overlap beween Cepheids in second and fourth crossings of the instability strip.," 7, which suggests an overlap between Cepheids in second and fourth crossings of the instability strip."
 As noted or loig-period Cepheids. there Is an intrinsic scatter in P values that is not the restlt of lareeOo uucertauties 11i the caleulated values.," As noted for long-period Cepheids, there is an intrinsic scatter in $\dot{P}$ values that is not the result of large uncertainties in the calculated values."
 Such scatter mals sit difficult to use rate of period change for individual Cepicids to identity +reir exact location within the iustabilitv srip. althoueh approximate dlaccineuts are possib edu most cases.," Such scatter makes it difficult to use rate of period change for individual Cepheids to identify their exact location within the instability strip, although approximate placements are possible in most cases."
 The couclisions reached here. while aduüittedlv speculative. provide a reasonable explanation for," The conclusions reached here, while admittedly speculative, provide a reasonable explanation for"
surveved at 15 Gllz with the Ryle Telescope (IE) (2).. and then cach of the sources found is monitored. at 34 CGllz by ἃ source subtractor interferometer located next to the VSA.,"surveyed at 15 GHz with the Ryle Telescope (RT) \citep{waldram-02}, and then each of the sources found is monitored at 34 GHz by a source subtractor interferometer located next to the VSA."
 Alore details of the source subtraction procedure are given in Paper I., More details of the source subtraction procedure are given in Paper II.
 The source subtraction interferometer consists of two 3.7-m cliameter dishes fitted with the same receivers as the main array. and equippec with identical LF ancl correlator.," The source subtraction interferometer consists of two 3.7-m diameter dishes fitted with the same receivers as the main array, and equipped with identical IF and correlator."
 The dishes are mounted such that they track with the same rotation of parallactic angle as the main array. so that polarized. sources will be observed. in the same orientation with respect to the polarization of the feeds.," The dishes are mounted such that they track with the same rotation of parallactic angle as the main array, so that polarized sources will be observed in the same orientation with respect to the polarization of the feeds."
 “Phe antennas are housed in separate ground screens identical to the main array ground screen. to limit ground. pickup and to provide shieküing from the wind.," The antennas are housed in separate ground screens identical to the main array ground screen, to limit ground pickup and to provide shielding from the wind."
 The baseline is 9 m. giving a resolution of 4 aremin in a 9 arcmin field. which will not resolve any of the sources we observe. while completely resolving out the CMD.," The baseline is 9 m, giving a resolution of 4 arcmin in a 9 arcmin field, which will not resolve any of the sources we observe, while completely resolving out the CMB."
 The VSA is located at the Teide Observatory. Tenerife at an altitude of 2400 m. The observatory is well established and. although on-site support is always available. we are able to control both the telescope ancl observations remotely.," The VSA is located at the Teide Observatory, Tenerife at an altitude of 2400 m. The observatory is well established and, although on-site support is always available, we are able to control both the telescope and observations remotely."
 The site provides excellent. observing conditions ancl has been well-tested for atmospheric seeing over a 15-vear period., The site provides excellent observing conditions and has been well-tested for atmospheric seeing over a 15-year period.
 The inversion laver lies below the observatory for 75 per cent of the vear: the twice-daily. meteorological balloon Hights show that the precipitable water vapour is as low as 2 nim for 30 per cent of the launches., The inversion layer lies below the observatory for 75 per cent of the year; the twice-daily meteorological balloon flights show that the precipitable water vapour is as low as 2 mm for 30 per cent of the launches.
 For observations in the region 26360 Cillz. the tvpical transparency is 95 per cent.," For observations in the region $26-36$ GHz, the typical transparency is 95 per cent."
 Correlated signals due to atmospheric emission. decrease with higher angular resolution., Correlated signals due to atmospheric emission decrease with higher angular resolution.
 For example. at 33 Cillz. the S ‘Tenerife beam-switch experiment (for example. ?)) was alfected by water vapour Fluctuations for SO per cent of the time while the Jodrell Bank-LAC 33 Gllz interferometer was allected for 20 per cent of the time at ~2 resolution (6— 110) and x10 percent at ~1 resolution (£~ 210).," For example, at 33 GHz, the $8^{\circ}$ Tenerife beam-switch experiment (for example, \citet{davies-96}) ) was affected by water vapour fluctuations for 80 per cent of the time while the Jodrell Bank-IAC 33 GHz interferometer \citep{melhuish1999} was affected for 20 per cent of the time at $\sim 2^{\circ}$ resolution $\ell \sim 110$ ) and $\le10$ percent at $\sim 1^{\circ}$ resolution $\ell \sim 210$ )."
 On the VSA baselines (£z150 900) we expect even Less loss of data due to atmospheric elfects., On the VSA baselines $\ell\approx 150-900$ ) we expect even less loss of data due to atmospheric effects.
 In fact. during our first season of observations (September 2000 September 2001). the site proved to have exceptional conditions: onlv 5 percent of the data were rejected due to weather.," In fact, during our first season of observations (September 2000 – September 2001), the site proved to have exceptional conditions; only 5 percent of the data were rejected due to weather."
 An important [actor in assessing the confidence we can place in measurements as dillicult. vet as important. as the CALB power spectrum. Is the agreement between experiments with different methocs that are susceptible to dilferent systematic ellects.," An important factor in assessing the confidence we can place in measurements as difficult, yet as important, as the CMB power spectrum, is the agreement between experiments with different methods that are susceptible to different systematic effects."
 The VSA is similar in its sensitivity and angular coverage to several other CAIB experiments that have reported: results. recently: BOOMIEIanG (?).. \LANIALA (7).. DASL (7). and CBE (?)," The VSA is similar in its sensitivity and angular coverage to several other CMB experiments that have reported results recently; BOOMERanG \citep{netterfield-01}, MAXIMA \citep{hanany-2000}, DASI \citep{halverson-02} and CBI \citep{cbi}."
 llere we consider briellv the iniportant svsteniatic effects in the VSA in comparison to these experiments., Here we consider briefly the important systematic effects in the VSA in comparison to these experiments.
 BOOSXMIBanG. and. ALANIALA are. focal-plane arrays of bolometers on small. oll-axis telescopes suspended: [rom stratospheric balloons.," BOOMERanG and MAXIMA are focal-plane arrays of bolometers on small, off-axis telescopes suspended from stratospheric balloons."
 They have very. high instantaneous sensitivity. at. frequencies. where the total galactic and extragalactic foregrounds are at their minima. and they operate above most of the atmosphere. eliminating that as a source of contaminating emission.," They have very high instantaneous sensitivity at frequencies where the total galactic and extragalactic foregrounds are at their minima, and they operate above most of the atmosphere, eliminating that as a source of contaminating emission."
 Vheir main problems are concerned with calibration: the uncertainties in the shapes and areas of the beams from each detector. and their pointing on the sky. limit the ability to recover the power spectrum. at resolutions close to the beam size.," Their main problems are concerned with calibration; the uncertainties in the shapes and areas of the beams from each detector, and their pointing on the sky, limit the ability to recover the power spectrum at resolutions close to the beam size."
 Recovery of the power spectrum at very low £ is limited. by receiver stability during each scan of the telescope., Recovery of the power spectrum at very low $\ell$ is limited by receiver stability during each scan of the telescope.
 There is also an overall uncertainty in the temperature scale due to the lack of knowledge of the beam areas ancl. particularly for BOOAERanG. the lack of good. κ calibration sources observed. during the flight.," There is also an overall uncertainty in the temperature scale due to the lack of knowledge of the beam areas and, particularly for BOOMERanG, the lack of good flux calibration sources observed during the flight."
 The quoted. uncertainty for the BOOAERanG power spectrum due to the temperature calibration is 20 percent. in addition to the beam uncertainty (?).," The quoted uncertainty for the BOOMERanG power spectrum due to the temperature calibration is 20 percent, in addition to the beam uncertainty \citep{netterfield-01}."
 Foreground. contamination at 150 CGllz is small. and dominated by galactic dust. emission. the clleet of which declines at higher f.," Foreground contamination at 150 GHz is small, and dominated by galactic dust emission, the effect of which declines at higher $\ell$."
 Calibration of instruments such as the VSA is much less problematic., Calibration of instruments such as the VSA is much less problematic.
 Flux calibrators with wellestucdied: properties are available (such as Cas A or Jupiter) and these can be observed repeatedly. and eross-checked with other secondary calibrators.," Flux calibrators with well-studied properties are available (such as Cas A or Jupiter) and these can be observed repeatedly, and cross-checked with other secondary calibrators."
 Overall calibration error is ominated. by the uncertainty in the absolute temperature of Jupiter. which is about 3.5 percent. (2).. leading to a 7 percent error in the power spectrum.," Overall calibration error is dominated by the uncertainty in the absolute temperature of Jupiter, which is about 3.5 percent \citep{mason_casscal}, leading to a 7 percent error in the power spectrum."
 The VSA is much less sensitive than the xdloon experiments. however. with slightly worse random errors from 500 hours observation than MLANINLAs from 3 hours.," The VSA is much less sensitive than the balloon experiments, however, with slightly worse random errors from 500 hours observation than MAXIMA's from 3 hours."
 The main foreground contaminant is extragalactic »oint. sources. whose cllect increases as (C.," The main foreground contaminant is extragalactic point sources, whose effect increases as $\ell^2$."
 The dominant errors ancl contaminants are thus quite cdillerent from he balloon experiments. so agreement between VSA. BOOALERanG and ALANIALA would be al very good sign of the believability of the current power spectrum measurements.," The dominant errors and contaminants are thus quite different from the balloon experiments, so agreement between VSA, BOOMERanG and MAXIMA would be a very good sign of the believability of the current power spectrum measurements."
 Aluch the same arguments apply to the results from the other interferometers. DASL and CDL. which are very similar to the VSA in many respects. having similar numbers of antennas and the same observing frequency.," Much the same arguments apply to the results from the other interferometers, DASI and CBI, which are very similar to the VSA in many respects, having similar numbers of antennas and the same observing frequency."
 However. there are significant dillerences between the three interferometer experiments which also allow for independent: cross-checks of calibration and. systematics.," However, there are significant differences between the three interferometer experiments which also allow for independent cross-checks of calibration and systematics."
 Phe VSA and 1 use the same primary Bux. calibrators: DASL. located. at. the south pole. is unable to see these and instead uses an absolute measurement of system temperature based on hot ancl cold loads.," The VSA and CBI use the same primary flux calibrators; DASI, located at the south pole, is unable to see these and instead uses an absolute measurement of system temperature based on hot and cold loads."
 CBL. ancl DASL in its first season. had. no ground sereen. ancl hence sullered. from. significant pick-up of correlated emission from the ground. which hac to be removed by dillerencing fields observed at the same ground angles.," CBI, and DASI in its first season, had no ground screen, and hence suffered from significant pick-up of correlated emission from the ground, which had to be removed by differencing fields observed at the same ground angles."
 Phe VSA avoided this problem but sulfered from a spurious correlated signal related to cross-coupling between the closest-spacecl antennas. which is removed. by. Fourier filtering of the time-ordered data (see section 4.2)).," The VSA avoided this problem but suffered from a spurious correlated signal related to cross-coupling between the closest-spaced antennas, which is removed by Fourier filtering of the time-ordered data (see section \ref{filtering}) )."
 The biggest design dillerence. between the CBI and DASI on one hand and the VSA on the other is that the VS has semi-independent tracking rather than fully co-mounted antennas., The biggest design difference between the CBI and DASI on one hand and the VSA on the other is that the VSA has semi-independent tracking rather than fully co-mounted antennas.
 The resulting variation of path length to à source as the earth rotates imposes à modulation on the signal which can be used to very ellectively reject. signals [rom elsewhere in the sky. for example from the Sun or the Moon.," The resulting variation of path length to a source as the earth rotates imposes a modulation on the signal which can be used to very effectively reject signals from elsewhere in the sky, for example from the Sun or the Moon."
 This allows the VSA to observe during clavlight. improving the observing elliciency. (," This allows the VSA to observe during daylight, improving the observing efficiency. ("
Note that on a typical. baseline,Note that on a typical baseline
3a. the secoud halb of the X-ray. flare was observed. with Ixeck. aud showed evideuce of flaring at,"3a, the second half of the X-ray flare was observed with Keck and showed evidence of flaring at"
linear perturbations. and hence to cosmological parameters (Holder2006).,"linear perturbations, and hence to cosmological parameters \citep{Holder:01,Haiman:01,Weller:02,Battye:03,Majumdar:04, Younger:06}."
. ILowever. in order to exploit SZ cluster number counts one is required (o understand the selection function of these surveys.," However, in order to exploit SZ cluster number counts one is required to understand the selection function of these surveys."
 This is most easily accomplished in terms of the [ιν decrement. which depends on the svstem temperature. exposure (ime. bandwidth. ancl efficiency (Dattve&Weller2005).," This is most easily accomplished in terms of the flux decrement, which depends on the system temperature, exposure time, band-width, and efficiency \citep{Battye:05}."
. In order to obtain cosmological constraints. il is necessary to convert the observables into a limiting mass of the survey.," In order to obtain cosmological constraints, it is necessary to convert the observables into a limiting mass of the survey."
 There are (wo approaches {ο obtain (his mass limit., There are two approaches to obtain this mass limit.
 One is to start. wilh the assumption that all clusters are spherical and in hydrostatic equilibrium. and then include some nuisance parameters to. allow [or deviations from (his assumption (Verdeοἱal.2002:Dattve&Weller 2006).," One is to start with the assumption that all clusters are spherical and in hydrostatic equilibrium, and then include some nuisance parameters to allow for deviations from this assumption \citep{Verde:02,Battye:03,Younger:06}."
. Another approach is (ο use a very general parameterization of the massobservable relation. which in its most general form could easily introduce forty unknown parameters (IIu2003:Lima&IIu2005).," Another approach is to use a very general parameterization of the mass–observable relation, which in its most general form could easily introduce forty unknown parameters \citep{Hu:03a,Lima:05}."
. Currently. there is little data to constrain (he [ree parameters in either approach., Currently there is little data to constrain the free parameters in either approach.
 In future one can exploit the SZ cluster observations (hemselves. to these [ree parameters (Ila2003:Lima&Iu2005:MajumdarMohr2004: 2003).," In future one can exploit the SZ cluster observations themselves, to self-calibrate these free parameters \citep{Hu:03a,Lima:05,Majumdar:04,Battye:03}."
. ILowever if one emplovs (he most general parameterization. little power is lelt in the survevs (o constrain cosmological parameters (Ilu2003:Lima&αι2005)..," However if one employs the most general parameterization, little power is left in the surveys to constrain cosmological parameters \citep{Hu:03a,Lima:05}."
 Another possibility would be to use complementary observations. such as weak lensing to cross-calibrate the mass-observable relation (Majumdar&Mohr.2004:TheDarkEnergySurvevCollaboration2005:Seallonetal. 2006).," Another possibility would be to use complementary observations, such as weak lensing to cross-calibrate the mass-observable relation \citep{Majumdar:04,DES,Sealfon:06}."
. Α useful approach would be to have a physical parameterization of the mass-observable relation wilh some prior probability on the free parameters and (hen sell-calibrate the SZ surveys for these parameters 2006)., A useful approach would be to have a physical parameterization of the mass-observable relation with some prior probability on the free parameters and then self-calibrate the SZ surveys for these parameters \citep{Younger:06}.
. ILowever in order to obtain this prior knowledge we can not vet rely on observations because currently they are sparse. and in (he near future observations will not resolve clusters because the beams of (he instrumentis are tvpicallv larger (han 1 arcmin.," However in order to obtain this prior knowledge we can not yet rely on observations because currently they are sparse, and in the near future observations will not resolve clusters because the beams of the instruments are typically larger than 1 arcmin."
 We hence require simulations to explore the scatter in the mass-observable relation: in order to obtain realistic results. a large representative sample of galaxy. clusters is required.," We hence require simulations to explore the scatter in the mass-observable relation; in order to obtain realistic results, a large representative sample of galaxy clusters is required."
 sehgaletal.(2006) have already applied the method of (his paper to a full light-cone N-body output. (out to 2= 3) in order (ο generate and make publicly available large-area. sub-arcmimite resolution microwave sky maps.," \citet{SehgalTHB06z} have already applied the method of this paper to a full light-cone N-body output (out to $z=3$ ) in order to generate and make publicly available large-area, sub-arcminute resolution microwave sky maps."
 We intend to provide a detailed analysis of the mass-observable relation in a forthcoming paper. and will eive here only a rough qualitative discussion.," We intend to provide a detailed analysis of the mass-observable relation in a forthcoming paper, and will give here only a rough qualitative discussion."
 In particular. we can use our z=0 simulated catalog to explore how Che amount of thermal enerev in the eas will affect the SZ signal.," In particular, we can use our $z=0$ simulated catalog to explore how the amount of thermal energy in the gas will affect the SZ signal."
 One can express the strength of the integrated SZ Παν as where the integration limits are along the line of sight through the entire cluster. and," One can express the strength of the integrated SZ flux as where the integration limits are along the line of sight through the entire cluster, and"
considered to be colinear with p. aud ecual to 100m. νι(A) and Nasa(A) are tjus calculated for each spectru1.,"considered to be colinear with $\boldsymbol{\rho}$, and equal to 100m. $\Phi_{\rm syst}(\lambda)$ and $\chi_{\rm syst}(\lambda)$ are thus calculated for each spectrum."
" The ranges of values tasei for the piston 9, ard tlie parasitic""n Παν factor+ E€ are the same as hose ¢of the ""unresolved source’ As detailed i1 Fig."," The ranges of values taken for the piston $\delta_*$ and the parasitic flux factor $\epsilon""^2$ are the same as those of the 'unresolved source' As detailed in Fig."
" 1. 9, cau be seen as the reualniug piston icorrected by the Iriige tracking device."," 1, $\delta_*$ can be seen as the remaining piston uncorrected by the fringe tracking device."
 TUs 'esiduzil piston creates a Litear-like phase term. ad«ed to tl trinsic phase signature rom the ylaet.," This residual piston creates a linear-like phase term, added to the intrinsic phase signature from the planet."
 C«yusequently. a further parasiic phase term wi| be added to the parasitic pliase eru «le to ie pl:et phase sigual.," Consequently, a further parasitic phase term will be added to the parasitic phase term due to the planet phase signal."
 It is therefore important to glace this issue in the framework of a ClassiCa interferπμ ΠΙΟ since cli[Iereut contribuious from the atmosphere and e lustruleit Call a4 the measurement of the interferometric pase aud especiavdd e first-order jnear ern viat 'esidual piston (or plstoi1 error) Oy., It is therefore important to place this issue in the framework of a classical interferometric measurement since different contributions from the atmosphere and the instrument can affect the measurement of the interferometric phase and especially the first-order linear term viathe residual piston (or piston error) $\delta_*$.
 Basically witliot aly pase refereuce. ο.ler linear 1u of the astrophysical sienal is lost duriug he data reduction step estimati tid reuoviug thi 'eslcual piston.," Basically without any phase reference, the first-order linear term of the astrophysical signal is lost during the data reduction step estimating and removing this residual piston."
 This is the case with most ¢urent iuterfe'ometers., This is the case with most current interferometers.
 Finally. olal phase. including the object phase aux the parasitic coutribution. contains only the hie ‘terms.," Finally, the total phase, including the object phase and the parasitic contribution, contains only the higher order terms."
 Hence‘orthi thee[ect of the pa‘asilic interference ias [o0 be examined after removi ear conipolie uo dq.nt(A) and Nasa(A). which conseqtently become “uupistonecd Foraclear illustrat jon«of this ;»oiut. we plotted in Fie.," Henceforth the effect of the parasitic interference has to be examined after removing the linear component of $\Phi_{\rm syst}(\lambda)$ and $\chi_{\rm syst}(\lambda)$, which consequently become 'unpistoned' For a clear illustration of this point, we plotted in Fig."
 | T unpistoned qiantities. $...fA} a \svs a(A)-(A). wilh 'espect to the wavelenelh usiug he very smooth number 7 spectruim (see Tabe 1).," 4 both 'unpistoned' quantities, $\Phi_{\rm syst}(\lambda)$ and $\chi_{\rm syst}(\lambda)$ $\Phi_{\rm syst}(\lambda)$, with respect to the wavelength, using the very smooth number 7 spectrum (see Table 1)."
" Va(A) has been calculated for two \""lues of piston (A/500 and A/30) correspondiug to the tinge tracking specilicatiois. and a quite €la'ge paraslic flux factor of1056."," $\chi_{\rm syst}(\lambda)$ has been calculated for two values of piston $\lambda/500$ and $\lambda/30$ ) corresponding to the fringe tracking specifications, and a quite large parasitic flux factor of."
.. First we can note a great similarity betwee both differeutial phases. hiehlieltine the [act that he additional parasitic phase is an object-depeident quantity.," First we can note a great similarity between both differential phases, highlighting the fact that the additional parasitic phase is an object-dependent quantity."
 It is quite low all auyway below the intrinsic term in both cases., It is quite low and anyway below the intrinsic term in both cases.
" However when tle beams udergo a greater plstoi (here A/15). the Jarasitic signal exceeds the intrinsic oue aud tle phase siguature from he planet is Icleutically to the ""unresolved. sot‘ce case. we cousider the amplitude i L baud of the phase signature [rom the planet. àost=yx(cba(A)—nit(Pit OO). aud of the corresponding parasilic phase. (par=MAX\syst(A)—Pars(ÀA))mintyDiviAy=9S,4(À))."," However when the beams undergo a greater piston (here $\lambda/15$ ), the parasitic signal exceeds the intrinsic one and the phase signature from the planet is Identically to the 'unresolved source' case, we consider the amplitude in L band of the phase signature from the planet, ${a}_{\:{\rm syst}}={\rm max}(\Phi_{\rm syst}(\lambda))-{\rm min}(\Phi_{\rm syst}(\lambda))$ , and of the corresponding parasitic phase, ${a}_{\:\rm par}={\rm max}(\chi_{\rm syst}(\lambda)-\Phi_{\rm syst}(\lambda))-{\rm min}(\chi_{\rm syst}(\lambda)-\Phi_{\rm syst}(\lambda))$."
" These qlautities are calculated Or a la‘ee range of Ó, values and are represeued wth respect to the parasitic [lux factor e? in Fie.", These quantities are calculated for a large range of $\delta_*$ values and are represented with respect to the parasitic flux factor $\epsilon''^2$ in Fig.
 , 5.
The A/500 and A/30 cases give a very simiar parasitic phase aiiplitude that can only ye distiigulshed [ο| (hie spectra 5. 6 and 7.," The $\lambda/500$ and $\lambda/30$ cases give a very similar parasitic phase amplitude that can only be distinguished for the spectra 5, 6 and 7."
" Here a ¢election is considere to be possible only il. Or givei values of o, and e. the parasitic phase auplitucle lies below the horizoutal solid. line 'epreseleeig theaiiplitude of the planet sieual."," Here a detection is considered to be possible only if, for given values of $\delta_*$ and $\epsilon''^2$ , the parasitic phase amplitude lies below the horizontal solid line representing theamplitude of the planet signal."
"dissipation processes, the treatment of thermal particles in the precursor is modified as described in refEV:th..","dissipation processes, the treatment of thermal particles in the precursor is modified as described in \\ref{EV:th}."
" In the Monte Carlo particle transport routine, the escape of energetic particles upstream is modeled by removing from the simulation the particles that cross the coordinate of the free escape boundary, xo, in the upstream direction."," In the Monte Carlo particle transport routine, the escape of energetic particles upstream is modeled by removing from the simulation the particles that cross the coordinate of the free escape boundary, $x_0$, in the upstream direction."
 The value of xo is a free parameter of the model., The value of $x_0$ is a free parameter of the model.
 The energy flux carried away by the escaping particles is reflected in the value of the energy flux downstream of xg., The energy flux carried away by the escaping particles is reflected in the value of the energy flux downstream of $x_0$.
" The amount of escaping energy flux, Qesc, is calculated in the EV method as the difference between the energy flux upstream of zo and the energy flux downstream of zo (see equation (9))), while the spectral density of the flux of escaping particles can be calculated in the simulation as As a final remark, the computational time required by this Monte Carlo approach is on the order of tens to hundreds of processor hours, and parallel processing is implemented in the latest version."," The amount of escaping energy flux, $Q_{esc}$, is calculated in the EV method as the difference between the energy flux upstream of $x_0$ and the energy flux downstream of $x_0$ (see equation \ref{eflux}) )), while the spectral density of the flux of escaping particles can be calculated in the simulation as As a final remark, the computational time required by this Monte Carlo approach is on the order of tens to hundreds of processor hours, and parallel processing is implemented in the latest version."
" In order to include precursor heating, particles are introduced at an upstream location xa, which is close to the subshock, but not too close, in order that the thermal particle distribution is still isotropic at r4."," In order to include precursor heating, particles are introduced at an upstream location $\xFP$, which is close to the subshock, but not too close, in order that the thermal particle distribution is still isotropic at $\xFP$ ."
" Upstream of x4, instead of calculating F(x,p) and its moments like P,() from particle propagation, the model finds Pj(z) by solving the equation where W(x) is the same heating rate as in Eq. fl."," Upstream of $\xFP$, instead of calculating $F(x,{\bf p})$ and its moments like $\Pth(x)$ from particle propagation, the model finds $\Pth(x)$ by solving the equation where $\mathcal{W}(x)$ is the same heating rate as in Eq. \ref{wheat}."
" Then particles are introduced at a, with the temperature corresponding to P;(r4) and allowed to propagate, get injected and accelerated."," Then particles are introduced at $\xFP$ with the temperature corresponding to $\Pth(\xFP)$ and allowed to propagate, get injected and accelerated."
 'This procedure requires a differentiation between thermal and superthermal particles in order to calculate the total momentum and energy fluxes upstream (Eqs., This procedure requires a differentiation between thermal and superthermal particles in order to calculate the total momentum and energy fluxes upstream (Eqs.
 |8| and (9)., \ref{pflux} and \ref{eflux}) ).
" Having placed the shock at z=0 (with the upstream at x> 0), the definition adopted in the model is that a particle is thermal if it has never crossed the coordinate x=0 going against theflow it never got injected), and a CR particle otherwise."," Having placed the shock at $x=0$ (with the upstream at $x>0$ ), the definition adopted in the model is that a particle is thermal if it has never crossed the coordinate $x=0$ going against theflow it never got injected), and a CR particle otherwise."
" Therefore, for x>r4 Eqs."," Therefore, for $x>\xFP$ Eqs."
 , \ref{pflux} 
in comparison with (he observed spectrum (see Table 2. and Figure 2)).,in comparison with the observed spectrum (see Table \ref{tab-obsline} and Figure \ref{fig-spectra}) ).
 This indicates hieher excitation temperature and lower optical depth in L673-5MMA., This indicates higher excitation temperature and lower optical depth in L673-SMM4.
 Lo we emploved the excitation temperature of 10 Ix. the filing result was improved.," If we employed the excitation temperature of 10 K, the fitting result was improved."
 In this case. the total optical depth was derived to be 1.40. and the derived column density decreased by a factor of 1.5. 8.0x107 7.," In this case, the total optical depth was derived to be 1.40, and the derived column density decreased by a factor of 1.5, $\times10^{12}$ $^{-2}$."
 For DNC. HNMC. and ΕΟΟ. we derived the column densities using the large velocity eradient (LVG) model (Goklreich&Kwan1974)... as adopted in Hirotaetal.(2001.2003).," For DNC, $^{13}$ C, and $^{13}$ $^{+}$, we derived the column densities using the large velocity gradient (LVG) model \citep{goldreich1974}, as adopted in \citet{hirota2001, hirota2003}."
. Details of the analysis are described. elsewhere (livota&Yamamoto2006)., Details of the analysis are described elsewhere \citep{hirota2006}.
. Decause we observed only the /-—1-0 lines. the Ils density and the kinetic temperature were fixed to be Lx10 * and LO Ix. respectively.," Because we observed only the $J$ =1-0 lines, the $_{2}$ density and the kinetic temperature were fixed to be $\times$ $^{5}$ $^{-3}$ and 10 K, respectively."
" To evaluate the uncertainty in the derived column density. we made the LVG analyses by varving the Ho density [rom 5x10! to 5x10? emoM"" "," To evaluate the uncertainty in the derived column density, we made the LVG analyses by varying the $_{2}$ density from $\times$ $^{4}$ to $\times$ $^{5}$ $^{-3}$."
As a result. we confirmed that the DNC/IINE'C ratios change only by a factor of 1.2 within the above II» density range. as shown in Table 5..," As a result, we confirmed that the $^{13}$ C ratios change only by a factor of 1.2 within the above $_{2}$ density range, as shown in Table \ref{tab-dhratio}."
 Table 5. lists the DNC/IINUC and NIL/CCS ratios of a few representative sources in addition to those of CDI20-3 and L6:3-S5MMA., Table \ref{tab-dhratio} lists the $^{13}$ C and $_{3}$ /CCS ratios of a few representative sources in addition to those of CB130-3 and L673-SMM4.
 These two ratios are known to be good indicators of chemical evolutionary stages Clirota&Yamamoto2006)., These two ratios are known to be good indicators of chemical evolutionary stages \citep{hirota2006}.
. We show a plot of these ratios in Figure 4.. as prepared by IHirotaetal.(2001).," We show a plot of these ratios in Figure \ref{fig-plots}, as prepared by \citet{hirota2001}."
. Both the NIL;/CCS and deuterium fractionation ratios in CD130-3 and L6T3-9MMA are systematically lower than (hose in (wpical dark cloud cores., Both the $_{3}$ /CCSand deuterium fractionation ratios in CB130-3 and L673-SMM4 are systematically lower than those in typical dark cloud cores.
 In our previous paper (Ilirotaetal.2009).. we classified CCPRs as dense cores wilh the NIL;/CCS ratios lower Chan 10.," In our previous paper \citep{hirota2009}, we classified CCPRs as dense cores with the $_{3}$ /CCS ratios lower than 10."
 The NIL;/CC S ratios in CD130-3 and L673-S5MMA are slightly higher than the above threshold. indicating that (hese two sources are close to the evolved prestellar cores such as L1544 (see Table 5)).," The $_{3}$ /CCS ratios in CB130-3 and L673-SMM4 are slightly higher than the above threshold, indicating that these two sources are close to the evolved prestellar cores such as L1544 (see Table \ref{tab-dhratio}) )."
" The NII;/CCS ratio in L492. which is reevaluated in (his paper on (he basis of the follow-up observation of NIL, (Appendix). is almost comparable to those of CD130-3 and L673-9MMA."," The $_{3}$ /CCS ratio in L492, which is reevaluated in this paper on the basis of the follow-up observation of $_{3}$ (Appendix), is almost comparable to those of CB130-3 and L673-SMM4."
 Nevertheless. it should be emphasized that the NII;/CCS ratios in (hese cores are still lower (han those in twpical dark cloud cores.," Nevertheless, it should be emphasized that the $_{3}$ /CCS ratios in these cores are still lower than those in typical dark cloud cores."
 Note that we here emploved the NIL; data obtained with the AIPHR 100 m telescope. which has higher spatial resolution than the NRO 45 m telescope. and hence. beam dilution effects would be less significant.," Note that we here employed the $_{3}$ data obtained with the MPIfR 100 m telescope, which has higher spatial resolution than the NRO 45 m telescope, and hence, beam dilution effects would be less significant."
 In other words. the NIL;/C5 ratios obtained with (he 45 m telescope are possibly underestimated in comparison with those with the 100 m telescope.," In other words, the $_{3}$ /CCS ratios obtained with the 45 m telescope are possibly underestimated in comparison with those with the 100 m telescope."
" In fact. we confirmed that the brightness temperatures of the NIL; line toward (he reference positions(0"".. 0""7)) of CB130-3 and L673-SMMA observed with the NRO 45 m telescope were ~3 limes weaker (han those with the MPIMR 100 m telescope."," In fact, we confirmed that the brightness temperatures of the $_{3}$ line toward the reference positions, ) of CB130-3 and L673-SMM4 observed with the NRO 45 m telescope were $\sim$ 3 times weaker than those with the MPIfR 100 m telescope."
" From this result. the equivalent EWIIM source size of the NIL, cores is estimated to be ο. which is consistent with the NIL, core size derived in (he mapping observation of this study. (Sections 3.2 and 3.3: Figures 5. and 6))."," From this result, the equivalent FWHM source size of the $_{3}$ cores is estimated to be $\sim$ , which is consistent with the $_{3}$ core size derived in the mapping observation of this study (Sections 3.2 and 3.3; Figures \ref{fig-mapcb130} and \ref{fig-mapl673}) )."
mosaic maps with a regular noise: lower in the map center and higher in the outer regions.,mosaic maps with a regular noise: lower in the map center and higher in the outer regions.
 An example of our map is given in Figure 3. where we show the contour map of the central region 26 arcmin?) of the mosaic map NI., An example of our map is given in Figure \ref{contour_map} where we show the contour map of the central region $\times$ 26 $^2$ ) of the mosaic map N1.
 We analvzed the noise properties in cach mosaic maps (NI. ΝΟ. N2 Deep ancl N3).," We analyzed the noise properties in each mosaic maps (N1, N2, N2 Deep and N3)."
 As expected. NI. N2 and N2 Deep have a regular noise distribution: a circular central region with a flat noise surrounded by concentric annular region. where the noise increases for increasing distance [from the center (the oll-axis value).," As expected, N1, N2 and N2 Deep have a regular noise distribution: a circular central region with a flat noise surrounded by concentric annular region, where the noise increases for increasing distance from the center (the off-axis value)."
 No structures or irregularities were found in the rms maps., No structures or irregularities were found in the rms maps.
 In Figure Ε we plot. the variation of the rms as function of the oll-axis value for NI. N2 and N2 Deep.," In Figure \ref{rms_off} we plot the variation of the rms as function of the off-axis value for N1, N2 and N2 Deep."
 Due to the different number of pointings in N3 (five pointings instead: of seven. see Figure 1)) the mosaic map in this field has a dillerent shape and the noise distribution in the map has a semi-circular shape instead of a circular one.," Due to the different number of pointings in N3 (five pointings instead of seven, see Figure \ref{fig_elais_grids_status}) ) the mosaic map in this field has a different shape and the noise distribution in the map has a semi-circular shape instead of a circular one."
 In Figure 4dd.e.f we plot the variation of the rms value in three dillerent slices.," In Figure \ref{rms_off}d d,e,f we plot the variation of the rms value in three different slices."
 In Figure 4dd we plot the rms in the North-South direction. in Figure tec the rms in a slice that connects the center of the three northern pointings (N3 V1. N3 V2 and N3 V4). while in Figure 4ff we report the rms value in a slice that connects the center of the two souther pointings (N3 V3 anc N3 V5).," In Figure \ref{rms_off}d d we plot the rms in the North-South direction, in Figure \ref{rms_off}e e the rms in a slice that connects the center of the three northern pointings (N3 V1, N3 V2 and N3 V4), while in Figure \ref{rms_off}f f we report the rms value in a slice that connects the center of the two souther pointings (N3 V3 and N3 V5)."
 As for NI and N2. no structures or irregularities were found in the rms map of N3.," As for N1 and N2, no structures or irregularities were found in the rms map of N3."
 Using Figure 4. we determined the regions with constant rms noise to be used for the source extraction., Using Figure \ref{rms_off} we determined the regions with constant rms noise to be used for the source extraction.
 Due to the circular shape of the rms map. we searched for radio sources in à circular central region plus concentric annular regions to an olf-axis value of 42 aremin.," Due to the circular shape of the rms map, we searched for radio sources in a circular central region plus concentric annular regions to an off-axis value of 42 arcmin."
 Outside this area. the beam attenuation significantly increases the limiting Εν (sce Figure 4)).," Outside this area, the beam attenuation significantly increases the limiting flux (see Figure \ref{rms_off}) )."
 In cach region the rms usec for source extraction is the highest rms in the region., In each region the rms used for source extraction is the highest rms in the region.
 In this way we are sure that all the sources extracted are above the fixed threshold in the region (for example 5 times the rms value)., In this way we are sure that all the sources extracted are above the fixed threshold in the region (for example 5 times the rms value).
 The sizes of these regions are reported in Table 2) and in Figure 5..," The sizes of these regions are reported in Table \ref{zone_tab}
 and in Figure \ref{rms_zone}."
" Using the data reported in Table 2/— we obtained the solid angle O(5S,) as function of thepeak flux density ὃν covered by our survey.", Using the data reported in Table \ref{zone_tab} we obtained the solid angle $\Omega(S_p)$ as function of thepeak flux density $S_p$ covered by our survey.
 This ellective survey area is shown in Figure 6.., This effective survey area is shown in Figure \ref{area}. .
LFs at :—5 and 5.9 (the blue aud. exeen dashed. lines πι Figure 5) down to a depth of Mp—LF aud LS. respectively. aud found. very steep faiut-cud slopes (a~— l3) at these redshifts.,"LFs at $z=5$ and 5.9 (the blue and green dashed lines in Figure 5) down to a depth of $M_{UV}\sim-17$ and $-18$, respectively, and found very steep faint-end slopes $\alpha\sim-1.7$ ) at these redshifts."
 MeLureetal.(2009) extended the 11 saurple by inchicding brighter LBC candidates selected from the UKIDSS Ultra Deep Survey (UDS)., \citet{mcl09} extended the $HST$ sample by including brighter LBG candidates selected from the UKIDSS Ultra Deep Survey (UDS).
 They found very similar LFs at 2~5 aud 6 (the blue aud green dotted lines iu Figure 1)., They found very similar LFs at $z\sim5$ and 6 (the blue and green dotted lines in Figure 4).
 The :=6.8 LFs in Figure 5 were measured with the data from the new Πο WEC3 IR camera (Bouwenusetal.2011) aud with the deep SDF Y-baud iunagine data (Ouchi 2009)., The $z=6.8$ LFs in Figure 5 were measured with the data from the new $HST$ WFC3 IR camera \citep{bou11} and with the deep SDF $Y$ -band imaging data \citep{ouc09}.
. It is clear that our spectroscopic result is in good agreement with the trend of the LFs frou +=5.9 to 6.8 froin the previous studies., It is clear that our spectroscopic result is in good agreement with the trend of the LFs from $z=5.9$ to 6.8 from the previous studies.
 Note that k-correction within a small UV sviveleueth rauge (1300~1600 Aj) is uceleible because of the blue UV slopes., Note that k-correction within a small UV wavelength range $1300\sim1600$ ) is negligible because of the blue UV slopes.
 We paralmctrize the ealaxy UV LF at 2~6 from our spectroscopic sanuple using a Schechter function. where o is the normalization. M is the characteristic huninosity. and a is the faiut-cud slope.," We parametrize the galaxy UV LF at $z\sim6$ from our spectroscopic sample using a Schechter function, where $\phi^{\ast}$ is the normalization, $M^{\ast}$ is the characteristic luminosity, and $\alpha$ is the faint-end slope."
 Because our suuple is not deep enough to well determine a. and the measurements of à at 5xi<6 from previous studies are quite robust and consistent. we fix a to the value of 1.71 given by Bowensetal.(2007).," Because our sample is not deep enough to well determine $\alpha$, and the measurements of $\alpha$ at $5\le z\le6$ from previous studies are quite robust and consistent, we fix $\alpha$ to the value of –1.74 given by \citet{bou07}."
. We then fit our data to the above function., We then fit our data to the above function.
 The best fits are o5—(1.1512)«10?Mpe Cand A=19.97£0.32. consistent with the result of Douwoensetal.(2007) (0=Lis10?Mpe and A= 2021) and the result of MeLureetal.(2009). (05=Ls«107Mpe? aud AL= 2001) at 26.," The best fits are $\rm \phi^{\ast} = (1.7\pm1.2)\times 10^{-3}\ Mpc^{-3}$ and $M^{\ast} = -19.97\pm0.32$, consistent with the result of \citet{bou07} $\rm \phi^{\ast} = 1.4\times10^{-3}\ Mpc^{-3}$ and $M^{\ast} = -20.24$ ) and the result of \citet{mcl09} $\rm \phi^{\ast} = 1.8\times10^{-3}\ Mpc^{-3}$ and $M^{\ast} = -20.04$ ) at $z\sim6$."
 The hest-fitting M—19.97 indicates that our sample spans a large Iuniuositv rauge 0.6L~5 across the characteristic huuiuositv. L at 2~6.2.," The best-fitting $M^{\ast} = -19.97$ indicates that our sample spans a large luminosity range $0.6\, L^{\ast} \sim 5\, L^{\ast}$ across the characteristic luminosity $L^{\ast}$ at $z\sim6.2$."
 The Eresult is shown in Figure 5., The result is shown in Figure 5.
 Cosmic variance has a nunor effect on our results., Cosmic variance has a minor effect on our results.
 The six masks covered different parts of SDF. and we did not exclusively target denser regions of the LBCs candidates.," The six masks covered different parts of SDF, and we did not exclusively target denser regions of the LBGs candidates."
 We covered 79 candidates in 310 arcmin’. while we selected 196 promising candidates in the whole SDF field (~876 arcimin?).," We covered 79 candidates in 340 $^2$, while we selected 196 promising candidates in the whole SDF field $\sim 876$ $^2$ )."
 The candidate surface deusitv in the six masks (0.23= 79/310) is close to the average density of the feld (0.22= 196/876)., The candidate surface density in the six masks $0.23=79/340$ ) is close to the average density of the field $0.22=196/876$ ).
 Therefore our saluple is a representativo sample of SDF., Therefore our sample is a representative sample of SDF.
 We calculate the uncertainty originating from cosmic variance for all the SDF caudidates using the Treuti&Stiavelli(2008) calculator., We calculate the uncertainty originating from cosmic variance for all the SDF candidates using the \citet{tre08} calculator.
 Within a redshift rauge of 5.8«τς6.5 the uncertaiutv is ouly ~20%... mach smaller than statistical errors n o and AM.," Within a redshift range of $5.8<z<6.5$ the uncertainty is only $\sim20$, much smaller than statistical errors in $\phi^{\ast}$ and $M^{\ast}$."
 Our ealaxy sample is a welkdefined spectroscopic xuuple of LBGs at 2>6., Our galaxy sample is a well-defined spectroscopic sample of LBGs at $z>6$.
 The comparison in Figure 5 shows that our measurement of the UV LF at 2~6 are consistent with the results from deep JEST aud UKIDSS observations., The comparison in Figure 5 shows that our measurement of the UV LF at $z\sim6$ are consistent with the results from deep $HST$ and UKIDSS observations.
 However. our LE. hugely. depends ou the selection function described in Figure L," However, our LF largely depends on the selection function described in Figure 4."
 Our tests show that the selection function is not scusitive to the physical parameters such as age. iietallicity. aud dust extinction. but is sensitive to the distribution of EEWs.," Our tests show that the selection function is not sensitive to the physical parameters such as age, metallicity, and dust extinction, but is sensitive to the distribution of EWs."
 Because the fraction of strong LAEs among LBCs is simall. we have applied a laree correction iu the selection function.," Because the fraction of strong LAEs among LBGs is small, we have applied a large correction in the selection function."
 Figure 5 mav incicate that the distribution of sstrenpgth at :~6 is simile to that at ον~, Figure 5 may indicate that the distribution of strength at $z\sim6$ is similar to that at $z\sim3$.
 Otherwise our LE will be quite diffrent from the previous meastriucuts of the :~6 LFs., Otherwise our LF will be quite diffirent from the previous measurments of the $z\sim6$ LFs.
 Alternatively. the contaninatiou rates m previous LBC smuples should be small.," Alternatively, the contamination rates in previous LBG samples should be small."
 For example. the coutamination rate in the Dowwenusetal.(2007) sample was claimed to be as small as 95 aud iu the Dowwenusctal.(2011) sample was found to be less than 11&%.," For example, the contamination rate in the \citet{bou07} sample was claimed to be as small as $\sim3$, and in the \citet{bou11} sample was found to be less than 14."
. Tf their contamination rates are higher. the ummber deusities derived from their siuuples will be lower. in contrary to Figure 5.," If their contamination rates are higher, the number densities derived from their samples will be lower, in contrary to Figure 5."
 We also derive the LLF using the wav very similar to what we do for the UV LF., We also derive the LF using the way very similar to what we do for the UV LF.
 The lhuuiosites ZL(Lvo)) are listed in Table 1., The luminosites $L$ ) are listed in Table 1.
 We first calculate the selection function. to correct for incompleteness for our galaxies as a suple of LAEs., We first calculate the selection function to correct for incompleteness for our galaxies as a sample of LAEs.
 Suuulated galaxies are generated using the Druzual& svuthesis model as discrbied iu Section 1.2., Simulated galaxies are generated using the \citet{bru03} synthesis model as discrbied in Section 4.2.
 The only difference is the assuuption of the EEW distribution., The only difference is the assumption of the EW distribution.
 The statistics of strengths in LAEs at 7~6 has been determined in several studies., The statistics of strengths in LAEs at $z\sim6$ has been determined in several studies.
 We use the EW probability distribution estimated from a sample of photometrically selected LAEs of Ouchietal.(2008)., We use the EW probability distribution estimated from a sample of photometrically selected LAEs of \citet{ouc08}.
. Figure 6 shows the LAE selection function as a function of L(Lvo)) aud :., Figure 6 shows the LAE selection function as a function of $L$ ) and $z$.
 The coutours in the fieure are selection probabilities frou o with an interval of10:4., The contours in the figure are selection probabilities from to with an interval of.
.. The selection function is eenerallv simular to the one for LBCs iu Figure |. except hat the selection probabilities for LAEs are more than hree times higher than those for LBGs.," The selection function is generally similar to the one for LBGs in Figure 4, except that the selection probabilities for LAEs are more than three times higher than those for LBGs."
 This is because or LAEs we do not need to apply a large correction due o the sinall fraction LAEs amoug LBCs., This is because for LAEs we do not need to apply a large correction due to the small fraction LAEs among LBGs.
" The volunue densities of the LAEs are computed in our huninosity (log £) bius |12.35. 12.60]. |12.60. 12.85]. 12:50, [3.10]. and (13.10. 13.50]."," The volume densities of the LAEs are computed in four luminosity $\,L$ ) bins [42.35, 42.60], [42.60, 42.85], [42.85, 43.10], and [43.10, 43.50]."
 The lower limit 12.35 of the first bin is chosen to be our detection limit (probability in Figure 6 below 5543)., The lower limit 42.35 of the first bin is chosen to be our detection limit (probability in Figure 6 below ).
 The brightest bin contains onlv oue galaxy aud the bin size is arbitrary. so we do not include this bin iu the following analysis.," The brightest bin contains only one galaxy and the bin size is arbitrary, so we do not include this bin in the following analysis."
 We estimate the LLF using a Schechter function. Most galaxies iu our suuple cover the faint eud of the LLF. so we are not able to well determine £.," We estimate the LF using a Schechter function, Most galaxies in our sample cover the faint end of the LF, so we are not able to well determine $\,L^{\ast}$."
 Therefore we fix logL to be the value of 13.0 (e.g.IKashikawactetal.2011) when we fit our data to the above function.," Therefore we fix $\,L^{\ast}$ to be the value of 43.0 \citep[e.g.][]{kas06,hu10,cas11,kas11} when we fit our data to the above function."
 The best fits are o(S2dx33).10°Mpe? aud a=ol6Tx0.51., The best fits are $\rm \phi^{\ast} = (8.2\pm3.3)\times 10^{-5}\ Mpc^{-3}$ and $\alpha = -1.67\pm0.54$.
 Figure 7 shows our measurements of the densities (filled circles with error bars) aud the best inodel fit (solid curve)., Figure 7 shows our measurements of the densities (filled circles with error bars) and the best model fit (solid curve).
 As we mentioned. the selection function is sensitive to the EEW distribution. especially in the faint eud of the xuuple.," As we mentioned, the selection function is sensitive to the EW distribution, especially in the faint end of the sample."
 The filled triangles in Figure 7 represent the volunue densitics if the asstmed EW clistribution is increased by 50 ((to roughlv iuateh the EW distribution for the spectroscopic saniple showin iu Ouchietal. (2008)))., The filled triangles in Figure 7 represent the volume densities if the assumed EW distribution is increased by 30 (to roughly match the EW distribution for the spectroscopic sample shown in \citet{ouc08}) ).
 The density iu the faüiutest biu is iucreased by a factor of two., The density in the faintest bin is increased by a factor of two.
 Iu Figure 7 we also show the comparison with the, In Figure 7 we also show the comparison with the
"Even for the extreme case that the magnetic field touches the disk at the inner boundary. (he internal torque g anc the radiation flux F' of the disk are also zero at the inner boundary: bevond the marginally stable orbit. as r decreases. g aud F increase but suddenly drops to zero al r—ry,,.","Even for the extreme case that the magnetic field touches the disk at the inner boundary, the internal torque $g$ and the radiation flux $F$ of the disk are also zero at the inner boundary: beyond the marginally stable orbit, as $r$ decreases, $g$ and $F$ increase but suddenly drops to zero at $r=r_{ms}$."
 We emphasize that this feature differs from the case of a disk magnetically connected to the material in the transition region. where the magnetic stress is demonstrated to be non-zero at the inner boundary. ancl extends into the transition region Ixrolik2000:LawleyandIxvolik2001 ).," We emphasize that this feature differs from the case of a disk magnetically connected to the material in the transition region, where the magnetic stress is demonstrated to be non-zero at the inner boundary and extends into the transition region \citep{kro99,ago00,haw00,kro01}."
. If the black hole rotates faster (han the disk. (he black hole pumps energy. and angular momentum into the disk through magnetic coupling. which increases the efficiency. of the disk.," If the black hole rotates faster than the disk, the black hole pumps energy and angular momentum into the disk through magnetic coupling, which increases the efficiency of the disk."
 Most inlerestinely. with the existence of the magnetic coupling. a disk can radiate without accretion: all the power of the the disk comes [rom (he rotational energv of the black hole. such a non-accretion disk has an infinite efficiency.," Most interestingly, with the existence of the magnetic coupling, a disk can radiate without accretion: all the power of the the disk comes from the rotational energy of the black hole, such a non-accretion disk has an infinite efficiency."
 We have discussed in detail a simple but important case: the magnetic field touches the non-accretion disk αἱ a single radius ry or equivalently. (he magnetic field is distributed on the disk within a narrow region around ry.," We have discussed in detail a simple but important case: the magnetic field touches the non-accretion disk at a single radius $r_0$ – or equivalently, the magnetic field is distributed on the disk within a narrow region around $r_0$."
 For this simple case. the radiation [ας has a sharp peak al rj: itis zero lor r«ry and decreases quickly for r>ry.," For this simple case, the radiation flux has a sharp peak at $r_0$: it is zero for $r<r_0$ and decreases quickly for $r>r_0$."
 Al large radii. the radiation [lux approaches £xrDu)E ," At large radii, the radiation flux approaches $F\propto r^{-3.5}$."
This behavior of the radiation flux is very. different from that of a standard accretion disk. where the radiation flux spreads widely over radii and approaches Foxr? at large radii.," This behavior of the radiation flux is very different from that of a standard accretion disk, where the radiation flux spreads widely over radii and approaches $F\propto r^{-3}$ at large radii."
 We have compared the radiation flix of a non-accretion disk with the magnetic field touching the disk at the inner boundary (1.6. Che marginally stable orbit). with the radiation flux of a standard accretion disk in detail.," We have compared the radiation flux of a non-accretion disk with the magnetic field touching the disk at the inner boundary (i.e. the marginally stable orbit), with the radiation flux of a standard accretion disk in detail."
 The results are summarized in Fig., The results are summarized in Fig.
 4 Fig. 6.., \ref{fig4} – Fig. \ref{fig5}.
 Clearly. the radiation profile of the non-accretion disk is very different from that of the stanclarel accretion disk: (he non-accretion disk magnetically coupled to a rapidly rotating black hole has a bigger emissivity index throughout the disk and a radiation region closer to the center ol the disk.," Clearly, the radiation profile of the non-accretion disk is very different from that of the standard accretion disk: the non-accretion disk magnetically coupled to a rapidly rotating black hole has a bigger emissivity index throughout the disk and a radiation region closer to the center of the disk."
 We have also shown that the magnetic field lines touching the disk at the inner boundary is most efficient in extracting enerev from the black hole. (hough the power of the black hole peaks al O5=04/2.," We have also shown that the magnetic field lines touching the disk at the inner boundary is most efficient in extracting energy from the black hole, though the power of the black hole peaks at $\Omega_D = \Omega_H /2$."
 A non-accrelion state can exist only if the internal torque of the disk exactly balances the external torque on the disk produced by the magnetic coupling., A non-accretion state can exist only if the internal torque of the disk exactly balances the external torque on the disk produced by the magnetic coupling.
 If the disk has so much internal torque that cannot be balanced by (the external torque. the excess internal torque will produce accretion.," If the disk has so much internal torque that cannot be balanced by the external torque, the excess internal torque will produce accretion."
 For an accretion disk with (he magnetic field touching the disk at a single racius. a specific feature is (hat the radiation flux of the disk may have (wo maxima (Fig.," For an accretion disk with the magnetic field touching the disk at a single radius, a specific feature is that the radiation flux of the disk may have two maxima (Fig."
 ὃ and Fig. 9))., \ref{fig7} and Fig. \ref{fig8}) ).
" When the black hole rotates faster than the disk. the black hole pumps energy to the disk aud the second maximum in the radiation flux is produced by adding a bright ""bump to the standard radiation flux (Fig. 3))."," When the black hole rotates faster than the disk, the black hole pumps energy to the disk and the second maximum in the radiation flux is produced by adding a bright “bump” to the standard radiation flux (Fig. \ref{fig7}) )."
 When the black hole rotates slower than the disk. the black hole extracts energv [rom the disk. in such a case the efficiency. of the disk is decreased by (he magnetic coupling anc (the second maximum in (he radiation flux," When the black hole rotates slower than the disk, the black hole extracts energy from the disk – in such a case the efficiency of the disk is decreased by the magnetic coupling – and the second maximum in the radiation flux"
form more reaclily (e.g. Brom Loeb 2003).,form more readily (e.g. Bromm Loeb 2003).
 To characterize (his Gransition we define f (o be a complementary extreme value or Fisher Tippett distribution (see equation (6) below)., To characterize this transition we define f to be a complementary extreme value or Fisher Tippett distribution (see equation (6) below).
 Because the present context of the model is different. from that of HOA. we redefine the variables and rewrite (he equations governing (he chemical evolution.," Because the present context of the model is different from that of H04, we redefine the variables and rewrite the equations governing the chemical evolution."
 The conserved quantity is (he total barvon mass. AI).," The conserved quantity is the total baryon mass, $_{t}$."
" It is made up of 4 components: the gas contributing to star formation M, which initially is the same as M). the mass in long-lived stars and/or remnants AL. the gas lost due to supernova driven winds idenüfied with the dilhuse intergalactic medium Mogjé;3;. and the gas left to form later generations λος."," It is made up of 4 components: the gas contributing to star formation $_{g}$ which initially is the same as $_{t}$, the mass in long-lived stars and/or remnants $_{s}$, the gas lost due to supernova driven winds identified with the diffuse intergalactic medium $_{dIGM}$, and the gas left to form later generations $_{LG}$."
" From 04. with the oxvgen abundance O replacing Z. Alweay replacing Ημ. aud Alpe; replacing M,,; the equations become: and where the complementary extreme value function is written For this particular problem the initial conditions are M,—M;. M.—0. ΩΩ Να. and O20."," From H04, with the oxygen abundance O replacing Z, $_{dIGM}$ replacing $_{WHIM}$, and $_{LG}$ replacing $_{ml}$ the equations become: and where the complementary extreme value function is written For this particular problem the initial conditions are $_{g}$ $_{t}$, $_{s}$ =0, $_{dIGM}$ =0, $_{LG}$ =0, and $_{\rm 0}$ =0."
" The parameters ΟΡ and W in equation (6) represent the point of inflection and the ""vidih of the transition lrom one to zero of the extreme value distribution (see ee, http://matliworld.wollram.com/IExtremeValueDistribution.htul).", The parameters $_{f}$ and W in equation (6) represent the point of inflection and the `width' of the transition from one to zero of the extreme value distribution (see e.g. http://mathworld.wolfram.com/ExtremeValueDistribution.html).
LETC observation gives the highest iutriusic soft N-ray absorption. we use that valuc OY coniparisou with he extinction values derived frou hne optica] observations.,"LETG observation gives the highest intrinsic soft X-ray absorption, we use that value for comparison with the extinction values derived from the optical observations."
 The photou index in the 0.3.7.| keV. band is 2.7+(j2, The photon index in the 0.3–7.0 keV band is $\rm 2.7 \pm 0.2$.
" The isotropic huuinositv iu the (0.3 to 7.0 keV. hate ls 2.1-10Perest,", The isotropic luminosity in the 0.3 to 7.0 keV band is $\rm 2.1 \cdot 10^{43}\ erg\ s^{-1}$.
 The strougest absorption e«eoo frou neutral oxy:exl at O.537 keV is clearly visible in the LETC spectruu., The strongest absorption edge from neutral oxygen at 0.537 keV is clearly visible in the LETG spectrum.
. To constrain the oxveen abuudaice relative to hydrogen. we have fitted the sSpectruni bv leaving the oxvee- abundance free.," To constrain the oxygen abundance relative to hydrogen, we have fitted the spectrum by leaving the oxygen abundance free."
 The spectral fi is still acceptable but does not significatly iaprove the spectral fitting results., The spectral fit is still acceptable but does not significantly improve the spectral fitting results.
 The oxveen abundance is founL to be consistent with solu values., The oxygen abundance is found to be consistent with solar values.
 Otwer clement amudauces can not be constrained with the preseut statistics of he Chancra LETC observation, Other element abundances can not be constrained with the present statistics of the Chandra LETG observation.
" Tf we assunie hat (1) the empirical relalon between iuterstelar N-rav absorption and optical extinction of ETE5-1072Ng (Gorenstein L975) appies to LES 1927|65I: (i) that the dust egeraius are opically thi ο N-ravs. Le, that the amount of N-rav asorption he cus erains is not much different froli hat bv eqivalent mass of material in gaseous for (ii) t he N-rav and optical raciation travel tlux11ο it1C Sade natter: (iv) that the obscuration is not a stro18o ‘ction of iue and (v) a Galactic eas to dust ra10. then the itral hydrogen column deusities derive from the X-Dav spectra can be converted into N-rav Ay vahes."," If we assume that (i) the empirical relation between interstellar X-ray absorption and optical extinction of $A_V = 4.5 \cdot 10^{-22} N_H$ (Gorenstein 1975) applies to 1ES 1927+654; (ii) that the dust grains are optically thin to X-rays, i.e. that the amount of X-ray absorption by the dust grains is not much different from that by an equivalent mass of material in gaseous form; (iii) that the X-ray and optical radiation travel through the same matter; (iv) that the obscuration is not a strong function of time and (v) a Galactic gas to dust ratio, then the neutral hydrogen column densities derived from the X-ray spectra can be converted into X-ray $\rm A_V$ values."
 The aN intrinsic N-rav she derived from the ROSAT alc Chandra observations is 0.323 (followiic Coreusteiu 1975) or 0.58 (following Precdehl and Sclunitt 1995)., The maximum intrinsic X-ray $A_V$ derived from the ROSAT and Chandra observations is 0.33 (following Gorenstein 1975) or 0.58 (following Predehl and Schmitt 1995).
 The optical spectu of LES 1927|651 was taken in June 2001 using the Carelec spectrograph with the 1.9314 telescope at the Haute-Provence observatory., The optical spectrum of 1ES 1927+654 was taken in June 2001 using the Carelec spectrograph with the 1.93m telescope at the Haute-Provence observatory.
 The spectral resolution is 6 (3 pixels) and the signal-to-noise ratio is about 30., The spectral resolution is 6 (3 pixels) and the signal-to-noise ratio is about 30.
 The optical spectrt in Fig., The optical spectrum in Fig.
 10 shows narrow emission lines and several eaaxv absorption liue features of Ca | Iv AA3933/3968. CII C AI30L. MeI ASL?. aud Na I ADS93 indicating the oresence of a ost galaxy contium," 10 shows narrow emission lines and several galaxy absorption line features of Ca $+$ K $\lambda\lambda3933/3968$, CH G $\lambda4304$, Mg I $\lambda5173$, and Na I $\lambda5893$ indicating the presence of a host galaxy continuum."
 The observed PWIA of most emission lines is about 330 luus ft. which comes]λος to the iistruienutal resolutk=-," The observed FWHM of most emission lines is about 330 km $^{-1 }$, which corresponds to the instrumental resolution."
 The lower FITWM value of 27d20 kn | for IT? line is probably chwe to the preseuce of an uudoerlvi absorption component. which can barely be seen when the spectrum is displaved at ful resolujon.," The lower FHWM value of $270\pm20$ km $^{-1 }$ for the $\beta$ line is probably due to the presence of an underlying absorption component, which can barely be seen when the spectrum is displayed at full resolution."
 The flux ratios of [ο Ts=11ο. ο Πα=0.2. L/Tlo.=0.6. aux [S Πα=0.6 poiut to a Seyfert 2 classification using the standard diagnosics of Veilleux Osterbrock (1987).," The flux ratios of [O $\beta=14.6$, [O $\alpha=0.2$, [N $\alpha=0.6$, and [S $\alpha=0.6$ point to a Seyfert 2 classification using the standard diagnostics of Veilleux Osterbrock (1987)."
 The rest frame euvaleut widths of he absorption lines CIT € A1304 ane Me I A5173 of 0.15£0.03 aud 0.1E250.03.A. respectively. have beeu usec ο decompose the spectrum iuto t1 power-law featureless continu and the host galaxy contribution.," The rest frame equivalent widths of the absorption lines CH G $\lambda4304$ and Mg I $\lambda5173$ of $0.18\pm0.03$ and $0.14\pm0.03$, respectively, have been used to decompose the spectrum into the power-law featureless continuum and the host galaxy contribution."
 Assuming an optical spectral iudex of 1.unfoptSs2.0 for AGN aid typical rest frame equivaleu witis of 3.3 to ha or CII € A130 and of 2.1 to La for Me I ASL? (see Goodrich Osterbrock 1983) we lave determine he fracion of the power-law featureless coutiniin to be between 9? and 97 at Lso0A., Assuming an optical spectral index of $-1.5 \le \alpha_{opt} \le 2.0$ for AGN and typical rest frame equivalent widths of 3.3 to 5.8 for CH G $\lambda4304$ and of 2.4 to 4.8 for Mg I $\lambda5173$ (see Goodrich Osterbrock 1983) we have determined the fraction of the power-law featureless continuum to be between 92 and 97 at 4800.
. The optica continuuni enissio1 of 1 ES 1927|651 secius therefore to be dominaed by he AGN., The optical continuum emission of 1 ES 1927+654 seems therefore to be dominated by the AGN.
 From the Πα to IL? flux ratio we have determined the optical extinction Ay-=2.0 in the narrow liue region. assuniue Case D recombination (see Velleux et al.," From the $\alpha$ to $\beta$ flux ratio we have determined the optical extinction $_V$ =2.0 in the narrow line region, assuming Case B recombination (see Veilleux et al."
 L997)., 1997).
" The actual visible extinction may in fact be even lower by a few percent. baring iu munud that the IL} cmussion liue is affected by an mucdering absorption. which cau uot be precisely iieasured Dere,"," The actual visible extinction may in fact be even lower by a few percent, baring in mind that the $\beta$ emission line is affected by an underlying absorption, which can not be precisely measured here."
 Wule the the optical spectra does rot clearly indicate t1C X»esence o a significant broad Ila «Οροςit. we have nevertheless tiec to fit the Ila regioji with an additioua Cpoet (see Fig.," While the the optical spectrum does not clearly indicate the presence of a significant broad $\alpha$ component, we have nevertheless tried to fit the $\rm \alpha$ region with an additional component (see Fig."
 ). lower ue) to estimate its posside strengh.," 10, lower panel), to estimate its possible strength."
 The adcition of he broad couponent does rot sienificautv improve the fit. indicating that tlie presence of a oad II component Is LO required by the data.," The addition of the broad component does not significantly improve the fit, indicating that the presence of a broad $\rm H\alpha$ component is not required by the data."
" The line flux of any possible road comporent. has to be lower tiu 5r, per cent of the warow Io compoucut."," The line flux of any possible broad component, has to be lower than 5 per cent of the narrow $\alpha$ component."
 The equivaleit width ofthe LALLOW Πα line is 2.1 A..The central waveleneth of this )ossible xoad Ta enüssion component is aso slightly süfted to he red., The equivalent width ofthe narrow $\alpha$ line is 2.4 .The central wavelength of this possible broad $\alpha$ emission component is also slightly shifted to the red.
 We note that the lower sigjal-to-nolse spectrum, We note that the lower signal-to-noise spectrum
"Carlesian coordinates /—1.2.3.1 <7x the particle number. and time steps σηxn,l.","Cartesian coordinates $k = 1,2,3$ , $1\leq i \leq$ the particle number, and time steps $1\leq n \leq n_x+1$."
" The expansion parameter ed,y/o and lime /,,το are evaluated hall-wav between time steps n and natlL. producing a leaplrog approximation to the equation of motion."," The expansion parameter $a_{n+1/2}$ and time $t_{n+1/2}$ are evaluated half-way between time steps $n$ and $n+1$, producing a leapfrog approximation to the equation of motion."
" In the action method the present positions .r;4,, ave given and the initial condition in equation(2)) is represented bv setting dy=0 so that equation (5)) αἱ n—1 is 0 -— = ina) + + Ηδη) When the free coordinates i;i, al 1<n<n, are adjusted to make the same number of 5;,, all vanish then the 2;,, are a numerical solution in leaplrog approximationto the equation of motion (4)) for given final positions $;4,,4 and (he initial condition in equation (2))."," In the action method the present positions $x_{i,k,n_x+1}$ are given and the initial condition in equation\ref{eq:initialcondition}) ) is represented by setting $a_{1/2}=0$ so that equation \ref{eq:discrete_eom}) ) at $n=1$ is 0 = = ) + + H_o^2 When the free coordinates $x_{i,k,n}$ at $1\leq n\leq n_x$ are adjusted to make the same number of $S_{i,k,n}$ all vanish then the $x_{i,k,n}$ are a numerical solution in leapfrog approximationto the equation of motion \ref{eq:eqofm}) ) for given final positions $x_{i,k,n_x+1}$ and the initial condition in equation \ref{eq:initialcondition}) )."
 One approach to τν=0 is to walk down the the gradient of S to a minimum (Peebles 1989). but that misses maxima ancl saddle points.," One approach to $S_{i,k,n}=0$ is to walk down the the gradient of $S$ to a minimum (Peebles 1989), but that misses maxima and saddle points."
" The 5;,, are driven to zero al thesestationary points of (he action by iterativelv applving the coordinate shilts 9:7;;.,, [rom the solution to + = (."," The $S_{i,k,n}$ are driven to zero at thesestationary points of the action by iteratively applying the coordinate shifts $\delta x_{i,k,n}$ from the solution to + = 0."
" Computing the drjy,, bx matrix inversion. as in Peebles et al. ("," Computing the $\delta x_{i,k,n}$ by matrix inversion, as in Peebles et al. ("
2001). is exceedingly slow for an interesting number of time steps.,"2001), is exceedingly slow for an interesting number of time steps."
 The NNAÀM described in the Appendix speeds (his up by making use of the fact that in equation (5)) the matrix 05)4.)/O.0)47). Is nonzero only near the diagonal.," The NNAM described in the Appendix speeds this up by making use of the fact that in equation \ref{eq:discrete_eom})) the matrix $\partial S_{i,k,n}/\partial x_{j,k',n'}$ is nonzero only near the diagonal."
 This analvsis of the motion of LAIC also uses numerical integration of the equation of motion (4)) forward in(nme Irom initial positions and proper peculiar velocities taken froma NNAAI solution., This analysis of the motion of LMC also uses numerical integration of the equation of motion \ref{eq:eqofm}) ) forward intime from initial positions and proper peculiar velocities taken froma NNAM solution.
 These initial conditions are — origa)/2. = =," These initial conditions are = )/2, = = ."
reconstructed function and the fiducial model.,reconstructed function and the fiducial model.
 The reconstruction with 10 PCs encloses the fiducial model within 2c confidence levels in the whole redshift range covered by the data., The reconstruction with 10 PCs encloses the fiducial model within $2\sigma$ confidence levels in the whole redshift range covered by the data.
" Considering that initially we had 34 parameters ;, this represents a reduction of ~70% in the parameter space dimensionality."," Considering that initially we had 34 parameters $\beta_i$, this represents a reduction of $\approx 70\%$ in the parameter space dimensionality."
The Keck AO/OSIRIS integral field spectra of HI 1348A/B were obtained with the svstem aligned along the long side of the 0756x2724 FOV.,The Keck AO/OSIRIS integral field spectra of HII 1348A/B were obtained with the system aligned along the long side of the $0\farcs56\times2\farcs24$ FOV.
 We obtained five 300 sec exposures al Lfbb ancl four 300 sec exposures al Jhb. wilh an extra 300 sec exposure on sky al each band for sky subtraction.," We obtained five 300 sec exposures at $Hbb$ and four 300 sec exposures at $Jbb$, with an extra 300 sec exposure on sky at each band for sky subtraction."
 An AQ star. ILD 24899. was observed immediately. afterwards in each band for telluric calibration.," An A0 star, HD 24899, was observed immediately afterwards in each band for telluric calibration."
 The data were reduced with the OSIRIS data reduction pipeline(DRP)!.. following the procedure described in ?..," The data were reduced with the OSIRIS data reduction pipeline, following the procedure described in \citet{mcelwain2007}."
 We used the default 7 pix (07245) radius aperture to extract the spectra of ΠΠ 1348A and D. where the center of the aperture was varied with wavelength to account for differential refraction arising both from the Earth's atmosphere and from the dichroie in the heck AO system.," We used the default 7 pix $\farcs$ 245) radius aperture to extract the spectra of HII 1348A and B, where the center of the aperture was varied with wavelength to account for differential refraction arising both from the Earth's atmosphere and from the dichroic in the Keck AO system."
 Based on the input set of individual exposures. the OSIRIS DRIP? produces a single mecdian-combined. wavelength-calibrated spectrum for each extracted object.," Based on the input set of individual exposures, the OSIRIS DRP produces a single median-combined, wavelength-calibrated spectrum for each extracted object."
 A collapsed ΠΡΟ OSIRIS data cube of HII 1348A/B is shown in Figure 2.., A wavelength-collapsed $Hbb$ OSIRIS data cube of HII 1348A/B is shown in Figure \ref{fig:hii1348_h_osiris}.
 The final Jbb and IHbb spectra of HIT 1348B are shown in Figure 5..., The final $Jbb$ and $Hbb$ spectra of HII 1348B are shown in Figure \ref{fig:hii1348b_ifu}.
 Comparison spectra of M and L. dafs are [rom the IRTF Spectral (??)..," Comparison spectra of M and L dwarfs are from the IRTF Spectral \citep{cushing2005, rayner2009}."
 To assess the best-fit spectral tvpes of the Jbb and {100 spectra. we followed a \? minimization when comparing the target spectra and the comparison standards.," To assess the best-fit spectral types of the $Jbb$ and $Hbb$ spectra, we followed a $\chi^2$ minimization when comparing the target spectra and the comparison standards."
 The 4? fitting was limited to the pam and pam wavelength regions. respectively. to avoid regions of high noise.," The $\chi^2$ fitting was limited to the $\mu$ m and $\mu$ m wavelength regions, respectively, to avoid regions of high noise."
 The spectral fitting vields best-fit spectral tvpes of M9. and AIO/MS for the J56 and {00 spectra. respectively.," The spectral fitting yields best-fit spectral types of M9 and M7/M8 for the $Jbb$ and $Hbb$ spectra, respectively."
 While III] 1348À and HIT 1348B likely eonstitute a physically bound svstem. (here is still a possibilitv that HIE 1348D may be an unrelated Pleiades brown dwarf.," While HII 1348A and HII 1348B likely constitute a physically bound system, there is still a possibility that HII 1348B may be an unrelated Pleiades brown dwarf."
 We estimate the probability for such a spurious alienment by considering the surface density of Pleiades brown cwarfs and the total skv area surveved bv 7. in their AO survey of the Pleiades. (, We estimate the probability for such a spurious alignment by considering the surface density of Pleiades brown dwarfs and the total sky area surveyed by \citet{bouvier1997} in their AO survey of the Pleiades. (
HII 1348B was the only potential substellar companion found by that survey. and we targeted it [or,"HII 1348B was the only potential substellar companion found by that survey, and we targeted it for"
thermal plasma component to improve the fit and obtained αν = 134/401 which was good enough for our analysis.,thermal plasma component to improve the fit and obtained $\chi^{2}$ $\nu$ = 1.34/401 which was good enough for our analysis.
 1n the spectral fitting. the parameters of the spectral models were tied between dilferent instrument spectra but et to vary free. axwt from the CGaussian line energy at 6.4 keV and the line width c which were fixed.," In the spectral fitting, the parameters of the spectral models were tied between different instrument spectra but let to vary free, apart from the Gaussian line energy at 6.4 keV and the line width $\sigma$ which were fixed."
 In order ο estimate the abundances. the abundance parameter of he models was le to vary free for most cata sets.," In order to estimate the abundances, the abundance parameter of the models was let to vary free for most data sets."
 For hose sources for which abundance was significantly higher han the solar value. it was fixed at 1.0.," For those sources for which abundance was significantly higher than the solar value, it was fixed at 1.0."
" An example of à source with a super-solar abundance is Z Cam for which. the obtained: abundance was 0bο Zo+ with: the xwiial covering | shotoclectric absorption combined with he cooling Llow model when the abundance was let to vary ree,", An example of a source with a super-solar abundance is Z Cam for which the obtained abundance was $^{+0.34}_{-0.19}$ $Z_{0}$ with the partial covering + photoelectric absorption combined with the cooling flow model when the abundance was let to vary free.
 The measured. equivalent widths of the 6.4 keV line for each source are given in Table 3.. and the results of the spectral fitting for the optically thin thermal plasma and the cooling How mocels are given in Fable 4 and 5.. respectively.," The measured equivalent widths of the 6.4 keV line for each source are given in Table \ref{ev}, and the results of the spectral fitting for the optically thin thermal plasma and the cooling flow models are given in Table \ref{mekal} and \ref{mkcflow}, respectively."
 These results show that in general. better AZ £6 values are achieved with the cooling How. model.," These results show that in general, better $\chi^{2}_{\nu}$ $\nu$ values are achieved with the cooling flow model."
 For example. the improvement with the cooling How model was statistically significant for SW. UMa ane T Leo.," For example, the improvement with the cooling flow model was statistically significant for SW UMa and T Leo."
 Fig., Fig.
 1 illustrates the X-ray spectra of the new NIS observations. i.e. VY Aqr. SW UMa. BZ UAla. SS Aur. απά ASAS J0025. which have been fitted with the cooling [ow modcl absorbed by photoelectric absorption with an added 6.4 keV Gaussian line component.," \ref{spectra} illustrates the X-ray spectra of the new XIS observations, i.e. VY Aqr, SW UMa, BZ UMa, SS Aur, and ASAS J0025, which have been fitted with the cooling flow model absorbed by photoelectric absorption with an added 6.4 keV Gaussian line component."
 Most. of the N-ray. spectra show that the clearest. discrete emission feature seen in the spectra of our source sample is the iron Fe Ixo complex. except in GW Lib. for which the signal-to-noise at ~ 6 keV is too low for a reliable measurement.," Most of the X-ray spectra show that the clearest, discrete emission feature seen in the spectra of our source sample is the iron Fe $\alpha$ complex, except in GW Lib, for which the signal-to-noise at $\sim$ 6 keV is too low for a reliable measurement."
 Since the studied sources are all within  200 pe. i-c.. within the solar neighbourhood. the effect of interstellar absorption should be negligible.," Since the studied sources are all within $\sim$ 200 pc, i.e., within the solar neighbourhood, the effect of interstellar absorption should be negligible."
 Thus. high measured absorption columns would mainky be cue to intrinsic absorption. associated with the sources.," Thus, high measured absorption columns would mainly be due to intrinsic absorption, associated with the sources."
" For most of the sources. the measurecl absorption columns were typically of the order of afew 107"" em.7. or even lower (107 em.2 7) which indicate low intrinsic absorption."," For most of the sources, the measured absorption columns were typically of the order of a few $\times$ $^{20}$ $^{-2}$, or even lower $^{19}$ $^{-2}$ ) which indicate low intrinsic absorption."
 The highest intrinsic absorption columns are found in Vso3 Sco. Z Cam and. WP Cas when compared. to the rest of the source sample.," The highest intrinsic absorption columns are found in V893 Sco, Z Cam and HT Cas when compared to the rest of the source sample."
" ALL these three sources have partial covering absorbers ma, with values of the order of 1077 1077  depending on the model.", All these three sources have partial covering absorbers $n_{H_{2}}$ with values of the order of $^{21}$ – $^{22}$ $^{-2}$ depending on the model.
 In. addition to the partial covering absorber. Z Cam also has a simple," In addition to the partial covering absorber, Z Cam also has a simple"
The central galaxies of many clusters of galaxies are surrounded by filaiieuts observed in optical euission lines to extend up to many teus of kiloparsecs (6. ον ?7)).,"The central galaxies of many clusters of galaxies are surrounded by filaments observed in optical emission lines to extend up to many tens of kiloparsecs (e. g. \citealt{Craw99, Cons01}) )."
 The filaments in NGC 1275. the ceutral ealaxy of the Perseus cluster. are typically of the order of 100 pe thick.," The filaments in NGC 1275, the central galaxy of the Perseus cluster, are typically of the order of 100 pc thick."
 ? sugeested that the filaments originate in interactions vetween the hot iutracluster medimm aud the relativistic asma associated with the ACN-driven radio features ound in and around the ceutral galaxies., \citet{McNa96} suggested that the filaments originate in interactions between the hot intracluster medium and the relativistic plasma associated with the AGN-driven radio features found in and around the central galaxies.
 ? modelled he filaments of NGC 1275 as material ablated from. clouds and dragged by a inore diffuse outflow moving away your NGC 1275 at more than 10° lan +., \citet{Pope08b} modelled the filaments of NGC 1275 as material ablated from clouds and dragged by a more diffuse outflow moving away from NGC 1275 at more than $10^3$ km $^{-1}$.
 ? sueeested hat dissipation occumiug during the ποιο trausfer roni the diffuse flow to the filaments might power the filamentary enmüssion. au idea related o that of ? that urbuleut musing lavers plav a role in generating the cnussion.," \citet{Pope08a} suggested that dissipation occurring during the momentum transfer from the diffuse flow to the filaments might power the filamentary emission, an idea related to that of \citet{Craw92} that turbulent mixing layers play a role in generating the emission."
 ??| have constructed detailed models of the optical and infrared dine ciission of filamentary easons 16atec by sinilar dissipation aud by cosmic ray induced ionisatiou.," \citet{Ferl08, Ferl09} have constructed detailed models of the optical and infrared line emission of filamentary gas heated by similar dissipation and by cosmic ray induced ionisation."
 They were motivated iu par by the difficulties encountered in attempts to match the obscrvec Spectra with the results of models in which xXotoionisation by stellar radiation is assumed to be the dominant heating mechanism (e. e. ?2))., They were motivated in part by the difficulties encountered in attempts to match the observed spectra with the results of models in which photoionisation by stellar radiation is assumed to be the dominant heating mechanism (e. g. \citealt{John07}) ).
" lufrared IT, enussion lines trace the optical eiissiou of the Bkuneuts (2)..", Infrared $_{2}$ emission lines trace the optical emission of the filaments \citep{Jaff05}.
 2? obtained CO(1-0). and. CO(2-1) sinele dish data on NGC 1275 aud ifs surrounding filamcuts and they observed HCN(G-2) a the position of 3C8L the radio source ocated at the centre of the galaxy NGC 1275.," \citet{Salo06, Salo08b} obtained CO(1-0), and CO(2-1) single dish data on NGC 1275 and its surrounding filaments and they observed HCN(3-2) at the position of 3C84, the radio source located at the centre of the galaxy NGC 1275."
 These two sets of observations established that both reeious contain cold molecular eas, These two sets of observations established that both regions contain cold molecular gas.
",ous, 72""7? have obtained iuterferometrie observatious of the 1) cnussion of some of he NGC 1275 filaments."," \citet{Lim08, Salo08a,
 Ho09} have obtained interferometric observations of the CO(2-1) emission of some of the NGC 1275 filaments."
 The required heating rates per particle due to dissipation and/or cosmic ray induced ionisation found by 7? are several orders of magnitude ligher than those in the nearest Mülkv. Way Cüant Molecular Clouds. aud the required thermal pressures identified bv 2 (see also 77)) are about au order of magnitude higher than those of the envelopes of dense cores associated with low lass star formation.," The required heating rates per particle due to dissipation and/or cosmic ray induced ionisation found by \citet{Ferl09} are several orders of magnitude higher than those in the nearest Milky Way Giant Molecular Clouds, and the required thermal pressures identified by \citet{Ferl09} (see also \citealt{Sand04, Sand07}) ) are about an order of magnitude higher than those of the envelopes of dense cores associated with low mass star formation."
 Though the heating aud ionisation rates in he cold molecular regions of the NGC 1275 fibuneuts may )o lower than those in the corresponding optical emission regions. they still may greatly exceed those relevaut to the rearby Milky: Way deuse cores.," Though the heating and ionisation rates in the cold molecular regions of the NGC 1275 filaments may be lower than those in the corresponding optical emission regions, they still may greatly exceed those relevant to the nearby Milky Way dense cores."
 If these rates are relatively veh. the chemical composition iu the NGC 1275 filaments should differ from that in a typical nearby dense core aud nay provide a diagnostic of the conditions iu the filaueuts.," If these rates are relatively high, the chemical composition in the NGC 1275 filaments should differ from that in a typical nearby dense core and may provide a diagnostic of the conditions in the filaments."
" ILleuce. we have examined the chemistry of molecular regions subject to leating rates per volume and cosmic rav Ponisation rates per particle in ranges that exteud up ο those approaching those that. according to ?.. obtain in the optical enüssiou line regions of the NOC 1275 flaicuts,"," Hence, we have examined the chemistry of molecular regions subject to heating rates per volume and cosmic ray ionisation rates per particle in ranges that extend up to those approaching those that, according to \citet{Ferl09}, obtain in the optical emission line regions of the NGC 1275 filaments."
 Our ourpose is to ideutify those molecular species that may be detectable and to determunme which ofthose species nav act as diagnostics of the dissipation ieating and of the cosmic ray iutensitv iu the filbunents., Our purpose is to identify those molecular species that may be detectable and to determine which of those species may act as diagnostics of the dissipation heating and of the cosmic ray intensity in the filaments.
 We shall consider a rauge of parameters up to those indicated bv the observations: however. it ds nof our ourpose here to uodel the filaments in NGC 1275.," We shall consider a range of parameters up to those indicated by the observations; however, it is not our purpose here to model the filaments in NGC 1275."
 Section 7? contaius details of the physical aud chemical models that we adopted to describe the uolecular component of the filaments., Section \ref{sec:model} contains details of the physical and chemical models that we adopted to describe the molecular component of the filaments.
 Iu section ?? we xeseut our results. and section ??. euds the paper with a discussion aud sole conclusions.," In section \ref{sec:resu} we present our results, and section \ref{sec:conclu} ends the paper with a discussion and some conclusions."
 We compute the filament chemistry dy assis that the flament may be represented by a steady-state photon-donunated region (a PDR) with constant thermal pressure., We compute the filament chemistry by assuming that the filament may be represented by a steady-state photon-dominated region (a PDR) with constant thermal pressure.
 To solve for the chemistry iu the filament woe enplov the PPDR time-dependent code. as," To solve for the chemistry in the filament we employ the PDR time-dependent code, as"
‘orm of the foreground correlations.,form of the foreground correlations.
 ?. have made the first detailed oreground model and have simulated the maps that include both he diffuse emission from our Galaxy and extragalactic sources (radio galaxies and clusters)., \citet{jelic08} have made the first detailed foreground model and have simulated the maps that include both the diffuse emission from our Galaxy and extragalactic sources (radio galaxies and clusters).
 ? have also studied both galactic and extragalactic foregrounds., \citet{gleser08} have also studied both galactic and extragalactic foregrounds.
 ? has used all publicly available tota yower radio surveys to obtain all-sky Galactic maps at the desirec Tequeney range and ? has studied foreground contamination in the context of the power spectrum estimation., \citet{deoliveira08} has used all publicly available total power radio surveys to obtain all-sky Galactic maps at the desired frequency range and \citet{bowman09} has studied foreground contamination in the context of the power spectrum estimation.
 Recently. a Galactic 3D emission model has been developec oy ??? (the code). derived from a 3D distribution of the Galactic therma electrons. cosmic-ray electrons and magnetic fields.," Recently, a Galactic 3D emission model has been developed by \citet{sun08, waelkens09, sun09}
 (the code), derived from a 3D distribution of the Galactic thermal electrons, cosmic-ray electrons and magnetic fields."
 The code is able to reproduce all-sky or zoom-in maps of the Galactic emission over a wide frequency range., The code is able to reproduce all-sky or zoom-in maps of the Galactic emission over a wide frequency range.
 In addition to simulations. a number of observational projects have given estimates of Galactic foregrounds in small selected areas.," In addition to simulations, a number of observational projects have given estimates of Galactic foregrounds in small selected areas."
 ?— have used 153 MHz observations with GMRT to characterize the visibility correlation function of the foregrounds., \citet{ali08} have used 153 MHz observations with GMRT to characterize the visibility correlation function of the foregrounds.
 ? have measured the spectral index of thediffuse radio background between 100 and 200, \citet{rogers08} have measured the spectral index of thediffuse radio background between 100 and 200.
MHz. ? have set an upper limit to the diffuse polarized Galactic emission: and 2? obtained the most recent and comprehensive targeted observations with the Westerbork Synthesis Radio Telescope (WSRT).," \citet{pen09}
 have set an upper limit to the diffuse polarized Galactic emission; and \citet{bernardi09a,bernardi10} obtained the most recent and comprehensive targeted observations with the Westerbork Synthesis Radio Telescope (WSRT)."
 However. current observations are not able to fully constrain the foregrounds. especial the Galactic polarized synchrotron emission. as required by EoR experiments.," However, current observations are not able to fully constrain the foregrounds, especial the Galactic polarized synchrotron emission, as required by EoR experiments."
 The importance of the polarized foreground stems from the fact that the LOFAR instrument. in common with all current= interferometric EoR experiments. has an instrumentally polarized response.," The importance of the polarized foreground stems from the fact that the LOFAR instrument, in common with all current interferometric EoR experiments, has an instrumentally polarized response."
 An improper polarization calibration will give rise to leakages of the complex polarized signal to the total signal., An improper polarization calibration will give rise to leakages of the complex polarized signal to the total signal.
 Since the Galactic polarized emission is quite structured along the frequency direction. the leakage of polarized intensity will have similar structures along the frequency and will mimic the EoR signal.," Since the Galactic polarized emission is quite structured along the frequency direction, the leakage of polarized intensity will have similar structures along the frequency and will mimic the EoR signal."
 Therefore. for reliable detection of the EoR signal it is essential at this stage: (i) to simulate the polarized foregrounds. and (1) to test the influence of the leakages on the extraction of the EoR signal.," Therefore, for reliable detection of the EoR signal it is essential at this stage: (i) to simulate the polarized foregrounds, and (ii) to test the influence of the leakages on the extraction of the EoR signal."
 This paper focuses on both aspects., This paper focuses on both aspects.
 In our previous foreground model (?).. the total intensity Galactic emission maps were obtained from three Gaussian random fields.," In our previous foreground model \citep{jelic08}, the total intensity Galactic emission maps were obtained from three Gaussian random fields."
 The first two were for the amplitudes of synchrotron and free-free emission and the third was for the spectral index of the synchrotron emission., The first two were for the amplitudes of synchrotron and free-free emission and the third was for the spectral index of the synchrotron emission.
 The polarized maps were simulated in a similar way but with added multiple 2D Faraday screens along the line of sight., The polarized maps were simulated in a similar way but with added multiple 2D Faraday screens along the line of sight.
" Despite the ability of that model to simulate observed characteristics of the Galactic emission (e.g. spatial and frequency variations of brightness temperature and its spectral index). the model had some disadvantages: e.g. the Galactic emission was derived and depolarization effects were not taken into account,"," Despite the ability of that model to simulate observed characteristics of the Galactic emission (e.g. spatial and frequency variations of brightness temperature and its spectral index), the model had some disadvantages: e.g. the Galactic emission was derived and depolarization effects were not taken into account."
 Our new foreground model. presented in this paper. is an extension of our previous foreground model (2)..," Our new foreground model, presented in this paper, is an extension of our previous foreground model \citep{jelic08}."
 It simulates both Galactic synchrotron and free-free emission in total and polarized intensity. but in a more realistic way and with some additiona benefits.," It simulates both Galactic synchrotron and free-free emission in total and polarized intensity, but in a more realistic way and with some additional benefits."
 The Galactic emission in our current model is derivec from the physical quantities and 3D characteristics of the Galaxy (thecosmicrayandthermalelectrondensity.themagneticfield:e.g.2??.andreferences therein).," The Galactic emission in our current model is derived from the physical quantities and 3D characteristics of the Galaxy \citep[the cosmic ray and thermal electron density, and the magnetic field; e.g.]
[and references therein]{beck96, berkhuijsen06, sun08}."
" In addition. the mode has the flexibility to simulate any peculiar case of the Galactic emission including very complex 3D polarized structures produced by ""Faraday sereens” and depolarization due to Faraday thick layers."," In addition, the model has the flexibility to simulate any peculiar case of the Galactic emission including very complex 3D polarized structures produced by “Faraday screens” and depolarization due to Faraday thick layers."
 Our Galactic emission model has some similarities with the model. but the difference between the two is the main purpose of the simulations.," Our Galactic emission model has some similarities with the model, but the difference between the two is the main purpose of the simulations."
 The simulation is based on a very complex Galactic model with aim to reproduce the observed all-sky maps of the Galactic emission., The simulation is based on a very complex Galactic model with aim to reproduce the observed all-sky maps of the Galactic emission.
 Because of its complexity. the high resolution zoom-in maps require a lot of computing power and time (2)..," Because of its complexity, the high resolution zoom-in maps require a lot of computing power and time \citep{sun09}."
 In contrast. our model is restricted to produce fast and relatively small maps of Galactic emission. which are then used as a foreground template for the LOFAR-EoR end to end simulation.," In contrast, our model is restricted to produce fast and relatively small maps of Galactic emission, which are then used as a foreground template for the LOFAR-EoR end to end simulation."
 Since the foreground subtraction is usually done along the frequency direction. our model also includes 3D spatial variations of the spectral index of the Galactic synchrotron radiation.," Since the foreground subtraction is usually done along the frequency direction, our model also includes 3D spatial variations of the spectral index of the Galactic synchrotron radiation."
 The paper is organized as follows., The paper is organized as follows.
 Section 22. gives a brief theoretical overview of the Galactic emission and Faraday rotation., Section \ref{sec:theory} gives a brief theoretical overview of the Galactic emission and Faraday rotation.
 The observational constrains of the Galactic emission are presented in Sec. ??.., The observational constrains of the Galactic emission are presented in Sec. \ref{sec:obs}.
 The simulation algorithm is described in Sec. ??.. ," The simulation algorithm is described in Sec. \ref{sec:sim}, ,"
while a few examples of simulated maps of the Galactic emission are presented in Sec. ??.., while a few examples of simulated maps of the Galactic emission are presented in Sec. \ref{sec:ex}.
 Section ?? also give a quantitative comparison of our simulations with the observations.," Section \ref{sec:ex}
 also give a quantitative comparison of our simulations with the observations."
 Section ?? deseribes the EoR signal and instrumental response simulations of the LOFAR-EoR pipeline., Section \ref{sec:lofar-eor} describes the EoR signal and instrumental response simulations of the LOFAR-EoR pipeline.
 We discuss the influence of the polarized foregrounds on the extraction of the EoR in Sec. ??.., We discuss the influence of the polarized foregrounds on the extraction of the EoR in Sec. \ref{sec:calibration}.
" The paper concludes with summary and conclusions (Sec. ??)),", The paper concludes with summary and conclusions (Sec. \ref{sec:so}) ).
 In radio astronomy. at frequencies where the Rayleigh-Jeans law is applicable. the radiation intensity. / (energy emitted per unit time per solid angle and per unit area and unit frequency). at the frequency is commonly expressed in terms of the brightness temperature C75: where c is the speed of light and 4g Boltzmann's constant.," In radio astronomy, at frequencies where the Rayleigh-Jeans law is applicable, the radiation intensity, $I$ (energy emitted per unit time per solid angle and per unit area and unit frequency), at the frequency $\nu$ is commonly expressed in terms of the brightness temperature $T_b$ ): where $c$ is the speed of light and $k_B$ Boltzmann's constant."
" The emission coefficient. 7 (energy emitted per unit time per solid angle and per unit volume). at a certain frequency can also be expressed in terms of the unit temperature. j,(v7)=azlv}. 80 that: where the integral is taken along the line of sight (LOS)."," The emission coefficient, $j$ (energy emitted per unit time per solid angle and per unit volume), at a certain frequency can also be expressed in terms of the unit temperature, $j_{b}(\nu)=\frac{c^2}{2
k_B \nu^2}j(\nu)$, so that: where the integral is taken along the line of sight (LOS)."
 In the following subsection we will give a brief theoretical overview of the Galactic synchrotron and free-free emission. as well as Faraday rotation. that will be used later in the simulation.," In the following subsection we will give a brief theoretical overview of the Galactic synchrotron and free-free emission, as well as Faraday rotation, that will be used later in the simulation."
" The Galactic emission will be expressed in terms of j, and 75.", The Galactic emission will be expressed in terms of $j_b$ and $T_b$.
 Synchrotron emission originates from the interaction between relativistically moving charges and magnetic fields., Synchrotron emission originates from the interaction between relativistically moving charges and magnetic fields.
 In our own galaxy. synchrotron emission arises from cosmic ray (CR) electrons produced mostly by supernova explosions and. the Galactic magnetic field.," In our own galaxy, synchrotron emission arises from cosmic ray (CR) electrons produced mostly by supernova explosions and the Galactic magnetic field."
 A fairly complete exposition of the synchrotronemission theory is presented in e.g. ? and ?.., A fairly complete exposition of the synchrotronemission theory is presented in e.g. \citet{pacholczyk70} and \citet{rybicki86}. .
 Here we only give a simple description of the emission., Here we only give a simple description of the emission.
Collisionless shocks are widely believed. to he the primary acceleration iechanisin eiving rise to the ubiquitous existence of energetic particles i space.,Collisionless shocks are widely believed to be the primary acceleration mechanism giving rise to the ubiquitous existence of energetic particles in space.
 The theory of diffusive shock acceleration (DSA) was proposed some 30 vears ago (Axfordetal.1978:Bell1975:Dluudford&Ostriker1978:INirxauskv1977) and ds currently believed to be the most important miechanisni for a variety of astrophysical environments. for example. interplanetarv shocks. the heliosphieric termination shock. aud shocks associated witli/ supernova remnants.," The theory of diffusive shock acceleration (DSA) was proposed some $30$ years ago \citep{Axford1978,Bell1978MNRAS,Blandford1978,Krymsky1977} and is currently believed to be the most important mechanism for a variety of astrophysical environments, for example, interplanetary shocks, the heliospheric termination shock, and shocks associated with supernova remnants."
 The theory predicts a universal power-law enerev flux spectrum dJ/dExEi for strong shocks with a deusity compression ratio of 1., The theory predicts a universal power-law energy flux spectrum $dJ/dE\propto E^{-1}$ for strong shocks with a density compression ratio of 4.
 The physical iuechanisu by which particles are accelerated from thermal energies to auch ligher energies where DSA is presumed to be applicable (the injection problemi) has received some recent attentions but has not reached a commuon conscusus explanation., The physical mechanism by which particles are accelerated from thermal energies to much higher energies where DSA is presumed to be applicable (the injection problem) has received some recent attentions but has not reached a common consensus explanation.
 Many acceleration theories. for example. shock drift acceleration (SDA) (seereviews.Armstrongetal.1985:Decker1988) and shock surfing acceleration (Sagdoeev1966:Leeetal.1996:Zakot1996) have been proposed.," Many acceleration theories, for example, shock drift acceleration (SDA) \citep[see
reviews,][]{Armstrong1985,Decker1988SSRv} and shock surfing acceleration \citep{Sagdeev1966RvPP,Lee1996JGR,Zank1996JGR} have been proposed."
 The acceleration of low cnerey protons iu the shocks containing large-scale pre-existing magnetic fluctuatious is very efficient (Cdacalone2005a.b).. which suggests that there may not be an injection problem.," The acceleration of low energy protons in the shocks containing large-scale pre-existing magnetic fluctuations is very efficient \citep{Giacalone2005ApJa,Giacalone2005ApJb}, which suggests that there may not be an injection problem."
 It is generally thought that. pre-accelerated particles will interact resonantly with maeuetic turbulence which results iu isotropizatiou and diffusion.," It is generally thought that, pre-accelerated particles will interact resonantly with magnetic turbulence which results in isotropization and diffusion."
" Αα previous works cousideriug magnetic turbulence focused ou the ions (οι,Bell1978:Cüacaloneetal.1992:Ne2003:Ciacalone2001. 2005a.b)."," Many previous works considering magnetic turbulence focused on the ions \citep[e.g.,][]{Bell1978MNRAS,Giacalone1992GeoRL,Ng2003ApJ,Giacalone2004ApJ,Giacalone2005ApJa,Giacalone2005ApJb}."
. Tlowever. the acceleration of electrous is less well uuderstood since for clectrous whose evroradii are very small. the evelotrou resonauce condition is not easilv satisfied thus they caunot interact resonantly with large-scale ambicut turbulence on ion-scale.," However, the acceleration of electrons is less well understood since for electrons whose gyroradii are very small, the cyclotron resonance condition is not easily satisfied thus they cannot interact resonantly with large-scale ambient turbulence on ion-scale."
 While the scattering provided by whistler waves (Shimadactal.1999) is one possibility. Jokipi&Ciüacalone(2007) proposed au attractive solution to the iujectiou problem that docs not require pitch-auele scattering. be. couscrving the first adiabatic invariant.," While the scattering provided by whistler waves \citep{Shimada1999} is one possibility, \citet{Jokipii2007ApJ} proposed an attractive solution to the injection problem that does not require pitch-angle scattering, i.e., conserving the first adiabatic invariant."
 The idea is that the low-rigklitv particles. especially clectrous. cau move rapidly along meandering magnetic field lines aud thus travel back aud forth between shock front.," The idea is that the low-rigidity particles, especially electrons, can move rapidly along meandering magnetic field lines and thus travel back and forth between shock front."
 The particles eaim cuerev from the difference between upstream aud downstream flow velocities., The particles gain energy from the difference between upstream and downstream flow velocities.
 Encrectic clectrous are often observed to be associated with collisionless shocks., Energetic electrons are often observed to be associated with collisionless shocks.
 Accelerated electrons. are thought to produce type II radio bursts in the solar corona and interplanetary space., Accelerated electrons are thought to produce type II radio bursts in the solar corona and interplanetary space.
 Andersonetal.(1979) reported spacecraft measurements of wpstream electrous (>16 keV) of the Earth's bow shock that originate from a thin regiou close to the point of taugeucyv between interplanetary magnetic field lines aud the shock surface., \citet{Anderson1979GeoRL} reported spacecraft measurements of upstream electrons $>16$ keV) of the Earth's bow shock that originate from a thin region close to the point of tangency between interplanetary magnetic field lines and the shock surface.
" Psurutani&Lin(1985). showed observations of energetic electrons associated with interplauetary shocks showing “spike-like” fux cuhancements for energies &2 keV. The spike events were observed at quasi-perpendicular shocks with 05,τοῦ, where Op, ds the angle between upstream magnetic field) aud shock normal."," \citet{Tsurutani1985JGR} showed observations of energetic electrons associated with interplanetary shocks showing ""spike-like"" flux enhancements for energies $\gtrsim 2$ keV. The spike events were observed at quasi-perpendicular shocks with $\theta_{Bn}
\gtrsim 70^\circ$, where $\theta_{Bn}$ is the angle between upstream magnetic field and shock normal."
" Some shock crossings had no culancements of energetic electrous which were reported to be associated with low shock speeds and sinall 05,. Siun", Some shock crossings had no enhancements of energetic electrons which were reported to be associated with low shock speeds and small $\theta_{Bn}$.
cttctal.(2005). preseuted data which shows energetic electrons are accelerated close to shock frout., \citet{Simnett2005} presented data which shows energetic electrons are accelerated close to shock front.
 They also showed some accelerated electrons can escape far upstreani of quasi-perpendicular interplauctary shocks., They also showed some accelerated electrons can escape far upstream of quasi-perpendicular interplanetary shocks.
 The clear, The clear
 (Ichikietal.2006).. (Langer," \citep{ichiki06}. \citep{langer03},"
"etal.2003).. ~LO!!C. ~LO*—I0G η&10em7. (e.g...Masunaga&Inutsuka2000) (e.g..Omukai&Nishi19983). 10 n,z10!em.*. i,cI0!em.*. (Qi.c9LOMem.7). n,c10?!em.7. ~200—300necθαὉ(Onmkaietal."," $\sim 10^{-11}\ {\rm G}$ $\sim 10^{-7}-10^{-8}\ {\rm G}$ $n \sim 10^3 \cm$ \citep[e.g.,][]{masunaga00} \citep[e.g.,][]{omukai98} $\sim10$ $\nc \lesssim 10^{11}\cm$ $\nc \simeq 10^{11}\cm$ $\nc \simeq 10^{16}\cm$ $\nc \simeq 10^{21}\cm$ $\sim200-300$$\nc \simeq 10^3\cm$\citep{omukai05,bromm02}."
 5~1.1 , $\gamma\simeq 1.1$ 
"2010). so 232.24kms' was added to the published heliocentric v, velocities.",", so $232.24\kms$ was added to the published heliocentric $v_\phi$ velocities."
" Given that the sample lies near to themid-plane. where a=2 would apply. we adopted a=1.5. and we also set the local scale-height to fy=200pe: when a and 7j, are allowed to vary when fitting to the data. they prove to be strongly correlated in the sense that low values ofa enhance the contribution of populations with smaller A, and thus favour larger values of /n for balance."," Given that the sample lies near to themid-plane, where $\alpha=2$ would apply, we adopted $\alpha = 1.5$, and we also set the local scale-height to $h_0 = 200 \pc$: when $\alpha$ and $h_0$ are allowed to vary when fitting to the data, they prove to be strongly correlated in the sense that low values of $\alpha$ enhance the contribution of populations with smaller $\Rg$ and thus favour larger values of $h_0$ for balance."
 However. the ditferences between the residuals of the various best fits are not statistically significant.," However, the differences between the residuals of the various best fits are not statistically significant."
 shows two typical tits performed on the region 150.<« 250. which demonstrate how nicely and naturally the formula reproduces the non-Gaussianity of the azimuthal velocity distribution.," shows two typical fits performed on the region $150 <
v_\phi/\kms < 250$ , which demonstrate how nicely and naturally the formula reproduces the non-Gaussianity of the azimuthal velocity distribution."
" The fits are for scale-lengths A.=7.Skpe (blue dashed line) and &,,=Skpe (red solid line).", The fits are for scale-lengths $R_{\sigma}=7.5\kpc$ (blue dashed line) and $R_{\sigma}=5\kpc$ (red solid line).
 The shorter length provides the better fit at low v; and the worse fit to the high- tail., The shorter scale-length provides the better fit at low $v_\phi$ and the worse fit to the high-velocity tail.
 The shorter length scale also yields the smaller value of the velocity-dispersion parameter. 7)=22.90+O45kms| versus συ=24.57xO48kms|.," The shorter length scale also yields the smaller value of the velocity-dispersion parameter, $\sigma_0 = 22.90 \pm 0.45 \kms$ versus $\sigma_0= 24.57 \pm 0.48
\kms$."
 In the plane the core of the velocity distribution is dominated by the youngest part of the thin disc. while the wings of the distribution will be dominated by the thick dise. and as we proceed from the core to the wings of the distribution stars of ever increasing age will grow in importance.," In the plane the core of the velocity distribution is dominated by the youngest part of the thin disc, while the wings of the distribution will be dominated by the thick disc, and as we proceed from the core to the wings of the distribution stars of ever increasing age will grow in importance."
 Hence the true distribution in v; reflects the entire star-formation history of the Galaxy and we cannot expect to obtain a perfect fit to it by adjusting a single velocity-dispersion parameter. ση.," Hence the true distribution in $v_\phi$ reflects the entire star-formation history of the Galaxy and we cannot expect to obtain a perfect fit to it by adjusting a single velocity-dispersion parameter, $\sigma_0$."
 loreover. the statistics of the GCS catalogue to some extent reflect he complex selection biases involved in the catalogue's formation. and we have made no attempt to replicate these biases.," Moreover, the statistics of the GCS catalogue to some extent reflect the complex selection biases involved in the catalogue's formation, and we have made no attempt to replicate these biases."
 Neglect of hese complexities is presumably why our fit is less good than that obtained by Schónrich.Binney&Dehnen(2010)., Neglect of these complexities is presumably why our fit is less good than that obtained by \cite{SBD}.
. The presence of kinematically hotter objects is contirmed by the estimated value for cy drifting to higher values when we expand the velocity interval on which we perform the fits., The presence of kinematically hotter objects is confirmed by the estimated value for $\sigma_0$ drifting to higher values when we expand the velocity interval on which we perform the fits.
 A couple of thin-dise scale-heights above the plane the population will be more homogeneous. being dominated by the thick disc. and it should be possible to obtain better fits by adjusting c7.," A couple of thin-disc scale-heights above the plane the population will be more homogeneous, being dominated by the thick disc, and it should be possible to obtain better fits by adjusting $\sigma_0$."
 Interestingly. on testing for a systematic shift in v; the formula recovered to 2kms! the expected local standard of rest both from the GCS data and from the velocity distributions of the torus models at low altitudes.," Interestingly, on testing for a systematic shift in $v_\phi$ the formula recovered to $2 \kms$ the expected local standard of rest both from the GCS data and from the velocity distributions of the torus models at low altitudes."
 At higher altitudes the performance deteriorates due to the higher uncertainties., At higher altitudes the performance deteriorates due to the higher uncertainties.
 We found that for local stars AE. does not play a significant role. although it gives a bias of order |kms|.," We found that for local stars $\Delta E_z$ does not play a significant role, although it gives a bias of order $1 \kms$."
 Ina forthcoming re-determination of the local standard of rest we will take this effect into account., In a forthcoming re-determination of the local standard of rest we will take this effect into account.
 Fig., Fig.
 9 compares the GCS data with the vj distributions that follow from the fits to ry., \ref{fig:GCSfitU} compares the GCS data with the $v_R$ distributions that follow from the fits to $v_\phi$ .
 It is clear that the formula under-estimates the width of the ry distribution., It is clear that the formula under-estimates the width of the $v_R$ distribution.
 In Section ??> we found that in the case of the torus model at small [ο]. the fitted vj distribution," In Section \ref{sec:VR}
 we found that in the case of the torus model at small $|z|$ , the fitted $v_R$ distribution"
at radii where the pair production is significant. leading to a brightening of the optical peak emission. (Figure 3) if the ejecta Lorentz factor is less than X300 (Ligure 4).,"at radii where the pair production is significant, leading to a brightening of the optical peak emission (Figure 3) if the ejecta Lorentz factor is less than $\siml 300$ (Figure 4)."
 A higher Lorentz [actor suppresses Πο pair-wind acceleration/enrichment and its brightening o ‘the afterglow emission until μμ=(1|νι(17)—1M/300)7 s. out. leaves unaltered the afterglow emission :ul later times. vast Che deceleration timescale.," A higher Lorentz factor suppresses the pair-wind acceleration/enrichment and its brightening of the afterglow emission until $t_{pair} = (1+z) r_{pair}/(c \Gamma^2) \sim 10\, (\Gamma/300)^{-2}$ s, but leaves unaltered the afterglow emission at later times, past the deceleration timescale."
 For conservative values of the afterglow kinetic energy and clectron/magnetic field parameters. the uoper limits set x ROVSE and LOTIS on the prompt optical emission from. several GRBs at 10510? s. ranging [rom 32 to 16 mag. require that the density of the cireumburst nrecliuma within OL pe of the GRB progenitor is less than 1tWen.," For conservative values of the afterglow kinetic energy and electron/magnetic field parameters, the upper limits set by ROTSE and LOTIS on the prompt optical emission from several GRBs at $10^2-10^3$ s, ranging from 13 to 16 mag, require that the density of the circumburst medium within 0.01 pc of the GRB progenitor is less than $100\,\cm3$."
 This imit is Consistent with the typical densities of Ql30em* inferred at à1011 em from fits to the broadband emission of several GRB afterglows Panaitescu Ixumar 2002).," This limit is consistent with the typical densities of $0.1-30 \,\cm3$ inferred at $r > 10^{17}$ cm from fits to the broadband emission of several GRB afterglows Panaitescu Kumar 2002)."
 Moreover. our results show that a wind-like meciunm around je COD. progenitor. corresponding to a mass-loss to wind Esoed ratio above 10.7AL./vr/(10?m/s) viclels optical 'ounterparts brighter than 2=13 up to almost 107 s. Thus. 1e wind of a Woll-Ttavet star is too dense ancl incompatible =vith the ROVSE and LOTIS upper limits.," Moreover, our results show that a wind-like medium around the GRB progenitor, corresponding to a mass-loss to wind speed ratio above $10^{-5}\, M_\odot/{\rm yr}/ (10^3\, {\rm km/s})$ yields optical counterparts brighter than $R = 13$ up to almost $10^3$ s. Thus, the wind of a Wolf-Rayet star is too dense and incompatible with the ROTSE and LOTIS upper limits."
 Figures 3. 4. 6. and 7 show that the |xightening of 1e pair-wind of the carly afterglow emission is much more ominent for a wind-like external medium.," Figures 3, 4, 6, and 7 show that the brightening of the pair-wind of the early afterglow emission is much more prominent for a wind-like external medium."
 This feature is esent within the first 100 seconds of the optical afterglow and. provides a test for the type of cireumburst environment. using future Swift observations.," This feature is present within the first 100 seconds of the optical afterglow and provides a test for the type of circumburst environment, using future Swift observations."
" For nz10""em all photons with energy above the threshold for e production are transformed into pairs and the GRB spectrum is severely modified."," For $n \simg 10^6\, \cm3$ all photons with energy above the threshold for $e^\pm$ production are transformed into pairs and the GRB spectrum is severely modified."
 In. this case the pair-wind becomes optically thick to Thomson scattering and the burst may not be detectable., In this case the pair-wind becomes optically thick to Thomson scattering and the burst may not be detectable.
 However. an optical counterpart resulting from the interaction of the ejecta with the pair-wind would be very bright. (2—10). unless it is attenuated. by the dust. within the dense medium.," However, an optical counterpart resulting from the interaction of the ejecta with the pair-wind would be very bright $R \sim 10$ ), unless it is attenuated by the dust within the dense medium."
 Such Hashes. if produced. would be missed bx. ROWSE and LOTIS. in the absence of the GRB emission.," Such flashes, if produced, would be missed by ROTSE and LOTIS, in the absence of the GRB emission."
 We acknowledge many interesting discussions with Dohcdan Paczviisski. who suggested that we investigate the elfect of pair-wind on GRB optical afterglows.," We acknowledge many interesting discussions with Bohdan Paczyńsski, who suggested that we investigate the effect of pair-wind on GRB optical afterglows."
 We are indebted to the referee. Andrei Beloborodov. for a number of helpful sugeestions that have improved the paper significantly.," We are indebted to the referee, Andrei Beloborodov, for a number of helpful suggestions that have improved the paper significantly."
Dwarf novae (Νο) are mass-exchanging binaries in which a low-mass main sequence-like star (the secondary) fills its ltoche lobe and loses hydrogen-rich matter toa white dwarl (WD) star (the primary).,Dwarf novae (DNe) are mass-exchanging binaries in which a low-mass main sequence-like star (the secondary) fills its Roche lobe and loses hydrogen-rich matter to a white dwarf (WD) star (the primary).
 The mass transfer is regulated by an accretion disc. which undergoes evclic changes of the mass aceretion rate AJ.," The mass transfer is regulated by an accretion disc, which undergoes cyclic changes of the mass accretion rate $\dot{M}$."
" The low mass accretion rate (ALz10.TAAL L ) quiescent. stage is ""Minterrupted intermittentlyη. every few weeks (to months) bv a ügh mass accretion rate (Alz=1O“Al. 4). the DN outburst stage which lasts days (to weeks)."," The low mass accretion rate $\dot{M} \approx 10^{-11} M_{\odot}$ $^{-1}$ ) quiescent stage is interrupted intermittently every few weeks (to months) by a high mass accretion rate $\dot{M} \approx 10^{-8} M_{\odot}$ $^{-1}$ ), the DN outburst stage which lasts days (to weeks)."
 At 65 pe VW Ilviisone of thecosests (Waner LOST) and xightest example of the SU UMa sub-class of DNe. which uncergo both normal DN outbursts and superoutbursts.," At 65 pc VW Hyi is one of the closests (Waner 1987) and brightest example of the SU UMa sub-class of DNe, which undergo both normal DN outbursts and superoutbursts."
" The mass of the accreting WD was estimated to be 0.63 A. (Schoembs Vogt 1981). but more recently a eravitational redshift determination vielded a larger mass M,=O.SGAL- (Sion et al."," The mass of the accreting WD was estimated to be 0.63 $M_{\odot}$ (Schoembs Vogt 1981), but more recently a gravitational redshift determination yielded a larger mass $M_{wd}=0.86 M_{\odot}$ (Sion et al."
 1997)., 1997).
 The inclination of the svstem is 7c60 degrees (Schocmbs Vogt. 1981) and. with an orbital »iod of 107 minutes. it [ies below the CV period. gap. where gravitational wave emission is thought to drive mass ransfer. resulting in very low accretion rates curing dwarl nova quiescence.," The inclination of the system is $i \approx 60$ degrees (Schoembs Vogt 1981) and, with an orbital period of 107 minutes, it lies below the CV period gap, where gravitational wave emission is thought to drive mass transfer, resulting in very low accretion rates during dwarf nova quiescence."
 In addition. it lies along a line of sight with an ideally low interstellar column of z6101 em? Polidan. Mauche Wade 1990).," In addition, it lies along a line of sight with an ideally low interstellar column of $\approx 6 \times 10^{17}$ $^{-2}$ (Polidan, Mauche Wade 1990)."
 As a consequence. VW Lyi has been observed in nearly all wavelength ranges. and -- is. therefore. one of the best-stuclicc systems.," As a consequence, VW Hyi has been observed in nearly all wavelength ranges, and it is, therefore, one of the best-studied systems."
 Llowever. there have been some discrepancies between the observed X-ray luminosity and the expected boundary —aver (BL) luminosity in VW. Hi :uxl other svstenis.," However, there have been some discrepancies between the observed X-ray luminosity and the expected boundary layer (BL) luminosity in VW Hyi and other systems."
. We gsugeest here that the missing boumacy laver luminosity in quiescence is released mainly in the ultraviolet as the second FUY component commonly. identified. as the “accretion belt”.," We suggest here that the missing boundary layer luminosity in quiescence is released mainly in the ultraviolet as the second FUV component commonly identified as the ""accretion belt""."
 In the next section We review multi-waveleneth observations of VW. Livi and. identivw each components in the system., In the next section we review multi-wavelength observations of VW Hyi and identify each components in the system.
 In particular. in section 3. we suggest that the," In particular, in section 3, we suggest that the"
which has been detected in a few clusters thanks to the sensitivity aud wide spectral coverage of aud )).,which has been detected in a few clusters thanks to the sensitivity and wide spectral coverage of and ).
 A HAR excess has been measured in the Coma cluster by aud (Fusco-Femiano 1999: Rephaeli. Gruber. Blanco 1999) and recently confirmed by combining the spectrum of the first //Phoswich Detection System (PDS: Frontera 19912) observation with the spectrum obtained with a second deeper observation 2001).," A HXR excess has been measured in the Coma cluster by and (Fusco-Femiano 1999; Rephaeli, Gruber, Blanco 1999) and recently confirmed by combining the spectrum of the first /Phoswich Detection System (PDS; Frontera 1997a) observation with the spectrum obtained with a second deeper observation (Fusco-Femiano 2004)."
 The results of these detectious have been challenged by the data analysis ol Rossetti Moleudi (2001: therealter RMOL)., The results of these detections have been challenged by the data analysis of Rossetti Molendi (2004; thereafter RM04).
 The origin of this cilference is currently uuder investigation., The origin of this difference is currently under investigation.
 However. in Sec.," However, in Sec."
 2 we examine the systematic ellects reported in their paper and the possible systematic error in the net count rates. discussed in Nevalainen (20014: thereafter NEO). due to uuresolved poiut sources present in the field of view (FOV) of the PDS.," 2 we examine the systematic effects reported in their paper and the possible systematic error in the net count rates, discussed in Nevalainen (2004; thereafter NE04), due to unresolved point sources present in the field of view (FOV) of the PDS."
 The flux derived from the combined spectrum of the Coma cluster. in the 20-80 keV energy range. is consistent. with the value of (1.2 ∐⋗⊔≺↵↥⋅∑≟∢∙⊔−⊳∖↓∐↕≺↵⋜↕⊳∖⇂⊔⋅≺↵≺⊔≻⊽∖⇁↕∐≺↵∫≀≽⊀∖⊤∫≜⊔∐≺↵⊳∖⋜⋃∐≺↵≺↵∐≺↲↓⋅∑≟⊽∖⇁∣≻⋜↕∐⇂⋜↕⋯⇂∢∙∩↥∐∐⋅⋯≺↵≺⇂∣≻⊽∖⊽⋜↕⊳∖≺↲∢∙∩∐≺⇂ 1 ⋅ ≺⇂≺↵≺↵↥↽≻≺↵↕⋅∩∣≻⊳∖≺↵⋅∖⇁⋜↕∏∩∐↸⋮∐≺↵↥↽≻∐⋜↕≺↵∐∙∖↽∁↽," The flux derived from the combined spectrum of the Coma cluster, $(1.5\pm0.5)\times 10^{-11}\erg$ in the 20–80 keV energy range, is consistent with the value of $(1.2\pm0.3)\times
10^{-11}\erg$ measured by the in the same energy band and confirmed by a second deeper observation (Rephaeli Gruber 2002)."
∶↓⋅⋃∣≻≺↵↕⋅⊇∪∪⊇⋝⋅↥∖⊽∩∐−↕∐≺↲↓⋅⋯⋜↕↥↓⋅⋜↕≺∐⋜↕↕∎∩∐↥⋜↕⊳∖∣⋈↵≺↵∐≼⇂≺↵↕≺↵∢∙↕≺↲≺⇂⋜↕↥⊳∖∩↥∐ ∣≻⊽∖⇁∫≀≽⊸∖≓⊤∫≜↴↸⋮∁∶↕⋅⋃∣∥↵↕⋅∙∖↽∐≺↵↥↽≻∐⋜↕≺↵∐⊇≼⋊≻⊇⋝⋅⇀−∖↕⋜↕↥∩∖∖↽≺↵⋅∢, Non-thermal radiation has been detected also in A2256 by and (Fusco-Femiano 2000; Rephaeli Gruber 2003) and A2319 by (Gruber Rephaeli 2002).
∙∩↥∐∎∐⇂≺↲∐∢∙≺↵↥≺↵∖⊽≺↵↥⋅∖∖↽∐⊔⋅≺↵⊳∖↽≻≺↲∢∙↕↕∩∊↽⊲∩⋯⋜≹⋜↕⋯∟−∖⊇⊇⊐≺∃⋅ ∐∩⋯∐≺↵↓⋅⋯⋜↕↥∐⇀∖≟∐↕⋅⋯∐⋜↕∏∩↕∐⋜↕⊳∖∣≻≺↵≺↵∐⊓↵↕↩∢∙↕≺↲≺⊔≻⊽∖⇁≵↗↗∤∣↗∣↗↙⋔⊽," At a lower confidence level, with respect to Coma and A2256, HXR radiation has been detected by in A754 (Fusco-Femiano 2003)."
↲⊸∖≓↥∐⇀−∖⊤⊐↓↸⋮∌⊲⋃⊳∖∢∙∩−∌⊲≺↵⋯↥⋜↕∐∩∤∣↙∣∣⋅ ⊇∪∪∶⋝⋝⋅⇀−∖∐⋃↥↽≻↥↽≻↩↓⋅∐∩∐↕↕≺↵↕⋅⋯⋜↕↥∐∟∖↥∎⋯∐∐⋜↕⊳∖≻↩↩∐↕⋅≺↲↥↽≻∩↕⋅↕≺↵≺⇂↥∐⇀−∖∶⋝≺∃≺∃⊤↸∫∌⊲⇂," An upper flux limit has been reported in A3667 (Fusco-Femiano 2001), A119 (Fusco-Femiano 2003) and A2163 (Feretti 2001)."
⊳∖∢∙∩−∌⊲≺↲∐↥⋜↕∐∩∤∣↙∣∣⋅⊇∪∪↽⊔⋅ ≺⇂≺↵↕≺↵∢⋅↕≺↵≺⇂⋜↕↕⋜↕∿⊇⊔↥≺↵∖⇁≺↵⊔↥∿⋅↴∪↸⋰∣⋰⊒↙∩⋅∐≺↲∐∩∐−⊳∖↕∑∸↥∏∐∢↝⋜↕∐↕↥⊽∖⊽⇀−∖⊂∶↓∖⊽⊣⋅∩∐↕⋜↕∐∐∐⋜⋯↵≺⇂∢↝↥⇂⊳∖↕≺↵↕⋅⋟∖, A component is detected at a $\sim$ $\sigma$ level in $\sim$ of the non-significantly AGN-contaminated clusters observed by (NE04).
∩∣≻↜∖≺↲↕⋅∖⊽≺↵≺⇂∣≻⊽∖⊽ ⊺↥↩⋯∩⊳∖↕∐⊆≺↵↥⊽∖⇁∐∐≺↵⋅↥↽∐⋅↩↕⋜↕∏∩↥∩∎↕∐↩∐∩∐↕∐≺↵↥⋅⋯⋜↕↥∐⇀↖∐≹↕⋅⋜↕≺∐⋜↕∏∩↥↥⊳∖↥∐∖⇁≺↲⋅⊳∖≺↲⊂⊲∩⋯↥↽≻↕∩∐⋖∫∐↽⊲⋟ ≺↵∐∏⊳∖⊳∖↕∩∐∣≻⊽∖⇁↕∐≺↵⊳∖⋜⋃∐≺↵⋅⋜↕≺∐∩⊳∖⊽∖⊽∐∢∙∐↕⋅∩↕⋅∩≺↵↥≺↵∢∙⋃⋅∩∐⊳∖↓⋅≺↵⊳∖↥↽≻∩∐⊳∖∐≻↥≺↵↥∎≺∐⋅↕∐≺↵≺↵⊸∖↕≺↵∐⇂≺↵≺⊔⋅⋯⇂∎∩≺↲∐∐⊳∖⊳∖↥∩∐↥↽≻↓⋅≺↵⊳∖≺↵⋯ ↥∐⋜↕∐∩∎↕∐↩∐↕↩⊔↥∩∐↩≼⇂∢∙∐⇂⊳∖↕≺↵↕⋅⊳∖⋅⊳∖∢∙⋜↕↕≺," The most likely interpretation of the nonthermal HXR radiation is inverse Compton (IC) emission by the same radio synchrotron electrons responsible for the extended radio emission present in all of the mentioned clusters, scattering the cosmic microwave background (CMB) photons."
↵⋅↥∐∑≟↕∐↩∢∙∩⊳∖∐∏∢∙∐∐∢∙↕⋅∩∖∖↽⋜↕∖⇁≺↵∣≻⋜↕∢∙⊆∑≟↕⋅⋯∐∐⊔∫⊂⊲↓⊟⊓↽≻∐∩↕∩∐⊳∖⋅∐↥⊳∖ ∖∖↽≺↵∐↥⊆∐∩∖∖↽↕∐⋜↕↕∐≺↵⋜↕↕≺↵↕⋅∐⋜↕∏∖⊽≺↵∐∐≺↵⋅↥↽≻⋅≺↵↕⋜↕∏∩∐∣≻⋜↕⊳∖≺↲≼⇂∩∐∐∩∐↕∐≺↵⋅⋯⋜↕↥∐⋅≺↵∐⊳∖⊳∖⋃⋅⋜↕∐↥⋯∑∸≺↵∐∐⊳∖⊳∖↥∩∐∐⋅∩⋯ stuprathermal electrons (Ixaastra. Bleeker. Mewe 1098: Eusslin. Lieu. Bierman 1999: Sarazin ]xempener 2000) has remarkable energetic problems (Petrosian 2002).," It is well know that the alternative interpretation based on bremsstrahlung emission from suprathermal electrons (Kaastra, Bleeker, Mewe 1998; Ensslin, Lieu, Biermann 1999; Sarazin Kempener 2000) has remarkable energetic problems (Petrosian 2002)."
 Another possible explanation for the detected emission may be due to a significant contamination by obscured) AGNs (Matt 1999: Hasiuger 2001)., Another possible explanation for the detected emission may be due to a significant contamination by obscured AGNs (Matt 1999; Hasinger 2001).
 However. the observatory that is sensitive enough to probe a sieuilicaut fraction of the obscured AGN has not detected such sources in several clusters (e.g. Molnar 2002).," However, the observatory that is sensitive enough to probe a significant fraction of the obscured AGN has not detected such sources in several clusters (e.g. Molnar 2002)."
 Besicles. the co-addecl spectrum of the sample of clusters in NEO! gives indication Lor an exteuded distribution of the emission against a significant contamination [rom obscured ACN.," Besides, the co-added spectrum of the sample of clusters in NE04 gives indication for an extended distribution of the emission against a significant contamination from obscured AGN."
 Iu this Letter. we present the combined PDS spectrum of Abell 2256 obtained by summiug," In this Letter, we present the combined PDS spectrum of Abell 2256 obtained by summing"
of rare bright objects).,of rare bright objects).
 Indeed. Takeuchi ct al.," Indeed, Takeuchi et al."
 attribute the bright end. bump! seen in both our samples to. the increasinglv dominant GN contribution at. luminosities Lin>5.0-107L.;. (the point at which their power law LE diverges [rom our Schechter LE). which we have removed as per §38.1..," attribute the `bright end bump' seen in both our samples to the increasingly dominant AGN contribution at luminosities $\mathrm{L}_{\mathrm{IR}} > 5.0 \times 10^{11} \mathrm{L}_{\sun}$ (the point at which their power law LF diverges from our Schechter LF), which we have removed as per \ref{agn}."
 Apart from this difference. the LE of t1c parent sample is essentially identical to our UV-matched subsample. suggesting that we have not introduced: undue bias by insisting on à UV observation for inclusion in the sample (see 83.2.1. below).," Apart from this difference, the LF of the parent sample is essentially identical to our UV-matched subsample, suggesting that we have not introduced undue bias by insisting on a UV observation for inclusion in the sample (see \ref{non} below)."
 Fig., Fig.
 3. shows the FUY luminosity [function for t1ο large UV sample (blue triangles) and the LVL data (green squares)., \ref{fig:uv_lf} shows the FUV luminosity function for the large UV sample (blue triangles) and the LVL data (green squares).
 The fitted black ancl blue curves are again. respectively. the least-squares fit to the L/Vinas data points. and the maximum likelihood fit.," The fitted black and blue curves are again, respectively, the least-squares fit to the $_{\mathrm{max}}$ data points, and the maximum likelihood fit."
 Being the sample from Duat et al. (, Being the sample from Buat et al. (
2007). the UV. LE is identical to theirs.,"2007), the UV LF is identical to theirs."
 The dashed line is the field. FUY LE. from GALEN. presented. by ? (again corrected for the more recent cosmology).," The dashed line is the field FUV LF from GALEX, presented by \cite{2005ApJ...619L..15W} (again corrected for the more recent cosmology)."
 Again. the close similarity between our sample and the canonical field? LE suggested. that. our data do provide a representative sample of star forming galaxies in the local Universe.," Again, the close similarity between our sample and the canonical `field' LF suggested that our data do provide a representative sample of star forming galaxies in the local Universe."
 We note however that there is a discrepancy with tw LE of Weder et al. (, We note however that there is a discrepancy with the LF of Wyder et al. (
2005) at the UV-bright end.,2005) at the UV-bright end.
 Again. the two LE derivation methods produce similar fits - the faint. end slope is 1.23d:0.09 for the μμ method. 1.31+0.09 for maximum likelihood).," Again, the two LF derivation methods produce similar fits - the faint end slope is $1.23 \pm 0.09$ for the $_{\mathrm{max}}$ method, $1.31 \pm 0.09$ for maximum likelihood)."
 The construction of all the luminosity functions ane star formation distribution functions was carried. out both by removing all galaxies with only upper flux limits (UV flux limits in the LReselected sample. and vice-versa). and. by treating the upper Dux limit as a detection equal to the limiting flux magnitude - as in 2.. the results in all cases do not depend on whether the galaxies are included. so we opted to include them.," The construction of all the luminosity functions and star formation distribution functions was carried out both by removing all galaxies with only upper flux limits (UV flux limits in the IR-selected sample, and vice-versa), and by treating the upper flux limit as a detection equal to the limiting flux magnitude - as in \cite{2005ApJ...619L..59M}, the results in all cases do not depend on whether the galaxies are included, so we opted to include them."
 ‘This insensitivity to the inclusion of non-detections does not extend to the distribution of obscured star formation. as traced by the ratio Lin/Lrev.," This insensitivity to the inclusion of non-detections does not extend to the distribution of obscured star formation, as traced by the ratio $\mathrm{L}_{\mathrm{IR}} / \mathrm{L}_{\mathrm{FUV}}$."
 Phis parameter is highly sensitive to low flux values. such as those assumed from an upper limit (being a ratio. low UV Iluxes lead to high values of Lin/Lpovy. ). and is discussed more in re[secrx..," This parameter is highly sensitive to low flux values, such as those assumed from an upper limit (being a ratio, low UV fluxes lead to high values of $\mathrm{L}_{\mathrm{IR}} / \mathrm{L}_{\mathrm{FUV}}$ ), and is discussed more in \\ref{sec:irx}."
 We also have το consider the οσο of. sample incompleteness on the form of our star formation rate function., We also have to consider the effect of sample incompleteness on the form of our star formation rate function.
 The Large IR and UM datasets will suller incompleteness starting at ~OAL:vr (note the downturns the the LR and UV samples shown in Figs.," The large IR and UV datasets will suffer incompleteness starting at $\sim 0.1\, \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}$ (note the downturns the the IR and UV samples shown in Figs."
 2 and 5 staring at LOL. )., \ref{fig:tir_lf} and \ref{fig:uv_lf} starting at $\sim 10^8 \mathrm{L}_{\sun}$ ).
 Ες is. hewever. the star formation rate at which the LVL data are highly complete. and hence our resultant. function will remain statistically robust here.," This is, however, the star formation rate at which the LVL data are highly complete, and hence our resultant function will remain statistically robust here."
 A the faintest end of our function. however. we run into LVL incompleteness which cannot be ignored.," At the faintest end of our function, however, we run into LVL incompleteness which cannot be ignored."
 ο calculate the LVL dataset to be complete tologSER=2:5M.vr| based on the resultant Ho Dux limits of the sample.," \cite{2009ApJ...692.1305L} calculate the LVL dataset to be complete to $\log \, \mathrm{SFR} = -2.5\, \mathrm{M}_{\sun}\; \mathrm{yr}^{-1}$ based on the resultant $\alpha$ flux limits of the sample."
 Aolow this level. too. small number statistics lead to derived star [formation rates being somewjw uncertain.," Below this level, too, small number statistics lead to derived star formation rates being somewhat uncertain."
 We do not truncate our sample below this SER. but it must be noted that values of & calculated at the faint. end are inherently uncertain. and for reference we mark this regime on our star formation function shown in Fig. 4..," We do not truncate our sample below this SFR, but it must be noted that values of $\Phi$ calculated at the faint end are inherently uncertain, and for reference we mark this regime on our star formation function shown in Fig. \ref{fig:sfr_tot}."
 We note. however. that calculations of the faint end. slope do not depend. on this uncertain regime. and all derived: parameters do not change i we fit only to the statistically complete cata.," We note, however, that calculations of the faint end slope do not depend on this uncertain regime, and all derived parameters do not change if we fit only to the statistically complete data."
 ]t is worth noting that the volumes probed by our field IR and UV. datasets are large enough to avoid. bias due to large scale structure., It is worth noting that the volumes probed by our field IR and UV datasets are large enough to avoid bias due to large scale structure.
" We can define a ""correlation volume. derived from the correlation length. which is an estimate of the absolute minimum volume vou need to probe in order to be reasonably free of bias clue to structure."," We can define a `correlation volume', derived from the correlation length, which is an estimate of the absolute minimum volume you need to probe in order to be reasonably free of bias due to structure."
 We estimate this to be 500 Alpe? (approximating the correlation radius as 5 Alpe: 7)., We estimate this to be 500 $^3$ (approximating the correlation radius as 5 Mpc; \citealt{Zehavi:2005aa}) ).
 For a Dux limited survey. the volume probed will vary as a function of luminosity: at a luminosity of 107L.: the lower limit of our field cata our IR and UV datasets respectively probe volumes of 20 and SO times this minimum ‘correlation volume’.," For a flux limited survey, the volume probed will vary as a function of luminosity; at a luminosity of $10^8 \;\mathrm{L}_{\sun}$ – the lower limit of our field data – our IR and UV datasets respectively probe volumes of 20 and 80 times this minimum `correlation volume'."
 For each galaxy in our sample. we derived a star formation rate as follows.," For each galaxy in our sample, we derived a star formation rate as follows."
 The LRA. defined as is used as a measure of the internal dust. absorption.," The IRX, defined as is used as a measure of the internal dust absorption."
 This then can be used to accurately estimate the UV. attenuation: it is known that the value of IRA is a good tracer of UV attenuation. and importantly reniains robust independent of the deals behind the extinction. such as &cometry and dust. properties (see 2)).," This then can be used to accurately estimate the UV attenuation; it is known that the value of IRX is a good tracer of UV attenuation, and – importantly – remains robust independent of the details behind the extinction, such as geometry and dust properties (see \citealt{Meurer:1999aa}) )."
 Conversions between IRN and A(GEUV) int 1e literature are given by ?:: 2... 9.. and ?.., Conversions between IRX and A(FUV) in the literature are given by \cite{Calzetti:2000aa}; \cite{2004MNRAS.349..769K}; \cite{2005ApJ...619L..51B}; ; and \cite{Cortese:2008aa}.
 For the purposes of this paper. we use the IRNACEUVM) relation derived by ? AGEUV)2 — 0028267 | 0.392. where wis the ΗΝ defined above.," For the purposes of this paper, we use the IRX–A(FUV) relation derived by \cite{2005MNRAS.360.1413B} : (FUV) = -0.028 x^3 + 0.392 x^2 + 1.094 x + 0.546, where $x$ is the IRX defined above."
 Once the observed UV luminosity has been corrected [ου dust attenuation using the above expression. we converted to a star formation rate as per the method described by. οτι This calibration uses à Salpeter IME. from. 0.1 to 100 Abi. was derived specifically for the GALLEN bands. and is similar to other estimators in the literature.," Once the observed UV luminosity has been corrected for dust attenuation using the above expression, we converted to a star formation rate as per the method described by \cite{2006ApJS..164...38I}: This calibration uses a Salpeter IMF from 0.1 to 100 $_{\sun}$, was derived specifically for the GALEX bands, and is similar to other estimators in the literature."
 For comparison the ? conversion factor: when converted from a monochromatic UV luminosity - taking the elfective central wavelength of the GALEN FUY filter to be 1532 - gives logSERAL.vr1]=log(breviesasb]0.56: —12% lower than our value.," For comparison the \cite{1998ARA&A..36..189K} conversion factor: when converted from a monochromatic UV luminosity - taking the effective central wavelength of the GALEX FUV filter to be 1532 - gives $ \log \mathrm{SFR} \; [\mathrm{M}_{\sun} \; \mathrm{yr}^{-1}] = \log (\mathrm{L}_{\mathrm{FUV, corr}}\; [\mathrm{L}_{\sun}]) - 9.56$: $\sim 12\%$ lower than our value."
" We adopt the former conversion. as it was derived specifically for the GALLEN filters by ον,"," We adopt the former conversion, as it was derived specifically for the GALEX filters by \cite{2006ApJS..164...38I}."
 We note that this method does carry some implicit assumptions. such as à constant star formation history. and a lack of stochasticity in populating the upper (UV-emitting) regions of the LME.," We note that this method does carry some implicit assumptions, such as a constant star formation history, and a lack of stochasticity in populating the upper (UV-emitting) regions of the IMF."
 As noted above. for the purpose of calculating ον we include non detections. hy assigning a Εαν equal to thedetection limit the overall results do not change if we instead discard. all non-detected. galaxies.," As noted above, for the purpose of calculating SFRs we include non detections, by assigning a flux equal to thedetection limit – the overall results do not change if we instead discard all non-detected galaxies."
 We then generated a ‘resultant’ sample from our large, We then generated a `resultant' sample from our large
image giant planets around young nearby star. contrasts of ~10°°—107 are requested.,"image giant planets around young nearby star, contrasts of $\sim 10^{-6}-10^{-7}$ are requested."
 To this aim. speckle noise has to be reduced by a further factor of 10 - 100.," To this aim, speckle noise has to be reduced by a further factor of 10 - 100."
 This 1s done by applying some differential imagingσι analysis techniques to the final datacube extracted from the IFS data. such as the simultaneous spectral differential imaging (S-SDI) etal.2000) and the spectral deconvolution (SD) (Thatteetal. 2007).," This is done by applying some differential imaging analysis techniques to the final datacube extracted from the IFS data, such as the simultaneous spectral differential imaging (S-SDI) \citep{Ma00} and the spectral deconvolution (SD) \citep{Th07}."
. A natural evolution of the S-SDI technique. using an IPS. has been defined during our work and is described in detail in Section. ??..," A natural evolution of the S-SDI technique, using an IFS, has been defined during our work and is described in detail in Section \ref{mdi}."
 Another possible technique. normally applied with other analysis techniques. is angular differential imaging (ADI) (Maroisetal.20006).," Another possible technique, normally applied with other analysis techniques, is angular differential imaging (ADI) \citep{Ma06}."
. In this section we briefly present the algorithms we used to implement these methods., In this section we briefly present the algorithms we used to implement these methods.
 As previously said. this techique is an extension to more spectral channels of the previous S-SDI techniques. such as the single differential imaging for two channels and the double differential imaging for three channels (Maroisetal.2000).," As previously said, this techique is an extension to more spectral channels of the previous S-SDI techniques, such as the single differential imaging for two channels and the double differential imaging for three channels \citep{Ma00}."
. The final result of the CSP simulation code consists of a datacube composed of 33 (for the Y-J-mode) or 38 (for the Y- monochromatic images., The final result of the CSP simulation code consists of a datacube composed of 33 (for the Y-J-mode) or 38 (for the Y-H-mode) monochromatic images.
 On these images we apply the following steps: There are three critical issues in this procedure: This method was proposed for the first time in Sparks&Ford(2002) and further developed in Thatteetal...(2007)., On these images we apply the following steps: There are three critical issues in this procedure: This method was proposed for the first time in \citet{Sparks02} and further developed in \citet{Th07}.
 It exploits that speckles are expected to change regularly with wavelength., It exploits that speckles are expected to change regularly with wavelength.
 Outside à given separation. defined as the bifurcation. radius (BR). the speckle spatial excursion over," Outside a given separation, defined as the bifurcation radius (BR), the speckle spatial excursion over"
Our paper “Testing Magnetic Star Formation Theory” (Crutcher.Hakobian.&Troland2008) (hereafter CLIT) has been severely criticized bv Mousehovias&Tassis(2008) Cherealter MT).,Our paper “Testing Magnetic Star Formation Theory” \citep{CHT08} (hereafter CHT) has been severely criticized by \citet{MT08} (hereafter MT).
 Although all of their eriticisms were addressed implicitly in our paper. in the sense that we provided rationale for the observational program aud (he analvsis of those observations. the paper did not explicitly discuss in detail the matters (μον claimed (o be flaws in our analvsis.," Although all of their criticisms were addressed implicitly in our paper, in the sense that we provided rationale for the observational program and the analysis of those observations, the paper did not explicitly discuss in detail the matters they claimed to be flaws in our analysis."
 Here we briefly do that: our brevity. requires reading. CIT aud. AIT before this note., Here we briefly do that; our brevity requires reading CHT and MT before this note.
 The ambipolar diffusion (AD) theory that the mass-to-fhix ratio (AL/@) increase in the central region. [rom an initial subcriGical value that supports the entire cloud against gravitv., The ambipolar diffusion (AD) theory that the mass-to-flux ratio $M/\Phi$ ) increase in the central region from an initial subcritical value that supports the entire cloud against gravity.
 Previous measurements of M/d toward cloud cores have nol definitivelv distinguished between AD and turbulence driving of star formation. e.g.. (2008).," Previous measurements of $M/\Phi$ toward cloud cores have not definitively distinguished between AD and turbulence driving of star formation, e.g., \citet{TC08}."
. The goal of our experiment was to attempt a more definitive test., The goal of our experiment was to attempt a more definitive test.
 In order to carry out a ceuantitative and unbiased test. it is necessary {ο have detailed. «quantitative AD models with which to compare the data.," In order to carry out a quantitative and unbiased test, it is necessary to have detailed, quantitative AD models with which to compare the data."
 We designed our experiment to test the models that have been published by the AD theorists., We designed our experiment to test the models that have been published by the AD theorists.
 Specilicallv. in CHT we discuss three such AD models: (1) a dimensionless parameter model (Ciolek&Mouschovias1994) for which," Specifically, in CHT we discuss three such AD models: (1) a dimensionless parameter model \citep{CM94} for which"
exceeds the level of the statistical uncertainty for narrower uidpasses (13«S MlIz).,"exceeds the level of the statistical uncertainty for narrower bandpasses $B <
8~{\rm MHz}$ )."
 We find that the sensitivity of a ow-lrequency array to the spherically averaged. BAO. scale decreases with increasing redshift. becoming 2 per cent NL.," We find that the sensitivity of a low-frequency array to the spherically averaged BAO scale decreases with increasing redshift, becoming $> 2$ per cent by $z \sim 8$."
 The details of the quantitative results presented in this oper are sensitive to the reionisation model assumed., The details of the quantitative results presented in this paper are sensitive to the reionisation model assumed.
 More eenerallv. the constraints available at a particular redshift will be sensitive to the details of the reionisation history.," More generally, the constraints available at a particular redshift will be sensitive to the details of the reionisation history."
 The rest constraints are likely to be derived. from observations of the Universe when the IGM is ~50 per cent ionised - at which time the Iuctuations are maximised (2) - but the bias is relatively weakly. dependent on scale at the scale of the BAO signal (and hence does not introduce a large systematic error)., The best constraints are likely to be derived from observations of the Universe when the IGM is $\sim50$ per cent ionised - at which time the fluctuations are maximised \citep{lidz2008} - but the bias is relatively weakly dependent on scale at the scale of the BAO signal (and hence does not introduce a large systematic error).
 As a result. i£ reionisation completes earlier than z 6. the most precise constraints will be obtained at. higher redshift than reported. here.," As a result, if reionisation completes earlier than $z \sim 6$ , the most precise constraints will be obtained at higher redshift than reported here."
 However. as the [oregrouncs become brighter when observing at higher redshift. an early reionisation scenario would decrease the maximal precision with which the BAO scale could be measured.," However, as the foregrounds become brighter when observing at higher redshift, an early reionisation scenario would decrease the maximal precision with which the BAO scale could be measured."
 A detailed analysis of the corresponding constraints on the dark energy equation of state is non-trivial and depends on the functional form. assumed (2).., A detailed analysis of the corresponding constraints on the dark energy equation of state is non-trivial and depends on the functional form assumed \citep{wlg2008}.
 We do not. explore the translation of BAO scale to cosmological constraints in this paper., We do not explore the translation of BAO scale to cosmological constraints in this paper.
 Such an analysis should include independent complementary. constraints (c.g. from the CAIB). and take account of uncertainties in. for example. the horizon scale (e.g.27)..," Such an analysis should include independent complementary constraints (e.g. from the CMB), and take account of uncertainties in, for example, the horizon scale \citep[e.g.][]{mcquinn2006,mao_tegmark2008}."
 Simple estimates of the cosmological constraints from a high redshift BAO measurement were presented in ?, Simple estimates of the cosmological constraints from a high redshift BAO measurement were presented in \citet{wlg2008}.
 Ina recent paper 2? propose a parametrization of the 3d 2leem PS including peculiar velocity cllects which is motivated by the generic form expected. for the ionisation PS and ionisation-matter PS., In a recent paper \citet{mao_tegmark2008} propose a parametrization of the 3d cm PS including peculiar velocity effects which is motivated by the generic form expected for the ionisation PS and ionisation-matter PS.
 Our analysis of the accuracy with which the BAO scale may be recovered in the presence ol à scale dependent cem bias is complementary to their detailed. discussion of the cosmological constraints available via the redshifted. cem PS., Our analysis of the accuracy with which the BAO scale may be recovered in the presence of a scale dependent cm bias is complementary to their detailed discussion of the cosmological constraints available via the redshifted cm PS.
 ?/— demonstrate that their »weanmetrization can recover the ionisation PS constructed rom simulations of reionisation and therefore can be used Oo constrain cosmological parameters from the 21eem Ps., \citet{mao_tegmark2008} demonstrate that their parametrization can recover the ionisation PS constructed from simulations of reionisation and therefore can be used to constrain cosmological parameters from the cm PS.
 Whilst they mention that their parametrization strugeles to describe the high-/ ionisation PS by a redshift of z=7 when he ionisation PS become significantly scale dependent. this is not found to be a problem up to the μις=0.2 MMpce1 hat we consider in this paper.," Whilst they mention that their parametrization struggles to describe the $k$ ionisation PS by a redshift of $z=7$ when the ionisation PS become significantly scale dependent, this is not found to be a problem up to the $k_{\rm max} =
0.2$ $^{-1}$ that we consider in this paper."
 It would be interesting to see row well such a parametrization would perform at the end of reionisation in the wavelength range that we have used Or our fits., It would be interesting to see how well such a parametrization would perform at the end of reionisation in the wavelength range that we have used for our fits.
 Our calculations of the constraints on the BAO scale from 2leen intensity Ductuations are promising when compared to the expected. performance of future, Our calculations of the constraints on the BAO scale from cm intensity fluctuations are promising when compared to the expected performance of future
strength of magnetic field in the interstellar medium is about 3 7G. There are three components of background photons in the Galaxy. which interact with electrons: thev are relic. infrared. and optical photons.,"strength of magnetic field in the interstellar medium is about 3 $\mu$ G. There are three components of background photons in the Galaxy which interact with electrons: they are relic, infrared and optical photons."
 In Fig., In Fig.
 1 we present the spectra of background photons at the position of 47 Tuc aud Terzan 5 in the Galaxy which were obtained with the GALPROP code (Strong Moskalenko 1993)., \ref{photon} we present the spectra of background photons at the position of 47 Tuc and Terzan 5 in the Galaxy which were obtained with the GALPROP code (Strong Moskalenko 1998).
 IIowever inside GC's we have an additional component of optical photons which are emitted by stars of the cluster., However inside GCs we have an additional component of optical photons which are emitted by stars of the cluster.
 Their spatial distribution is strongly noununilorm., Their spatial distribution is strongly nonuniform.
" It may reach the value about eo,=300 eV  for 47 Tue and about «7,=100 eV for Terzan 5 in the cluster center but it decreases rapidly with the distance from the GC center.", It may reach the value about $w^0_{op}=300$ eV $^{-3}$ for 47 Tuc and about $w^0_{op}=100$ eV $^{-3}$ for Terzan 5 in the cluster center but it decreases rapidly with the distance from the GC center.
 Thus. for 47 Tuc the spatial distribution οἱ optical photons was derived by Michie (1963) and Ixuranov Postuov (2006) which is: where r= 0.5 pc. ry=50 pe and ry=y2riaj/3.," Thus, for 47 Tuc the spatial distribution of optical photons was derived by Michie (1963) and Kuranov Postnov (2006) which is: where $r_c = $ 0.5 pc, $r_t = 50$ pc and $r_h = \sqrt{2r_c r_t/3}$."
 The diffusion coefficient inside clusters (r« 7.) is supposed to be smaller than in the surrounding medium. that seems (o be reasonable since GCs have high densities of stars with their winds which can create turbulence in the medium inside GCs.," The diffusion coefficient inside clusters $r<r_c$ ) is supposed to be smaller than in the surrounding medium, that seems to be reasonable since GCs have high densities of stars with their winds which can create turbulence in the medium inside GCs."
 Therefore. the diffusion coefficient was taken in the form where 9(r) is the Heaviside (step) function.," Therefore, the diffusion coefficient was taken in the form where $\theta(r)$ is the Heaviside (step) function."
 The values of Dy and Dy are derived [rom the data (see below)., The values of $D_0$ and $D_1$ are derived from the data (see below).
 The value of rj=100pe~2r. is also derived [rom the spatial distribution of gamma-ray emission., The value of $r_D = 100~\mbox{pc} \simeq 2r_c$ is also derived from the spatial distribution of gamma-ray emission.
 We assume (hat pulsars inject a monoenergetic spectrum of electrons in the form, We assume that pulsars inject a monoenergetic spectrum of electrons in the form
"based on a functional form given by: where re and fee""Tug) represent (he observed ancl model [Iuxes. respectively. at the wavelength: band represented by. index A: σι represents [αν uncertaintv: ry represents the absorption in band A relative to sly. based on the reddening law from Allersetal.(2006).","based on a functional form given by: where $f_\lambda^{\rm obs}$ and $f_\lambda^{\rm mod}(T_{\rm eff})$ represent the observed and model fluxes, respectively, at the wavelength band represented by index $\lambda$; $\sigma_\lambda$ represents flux uncertainty; $r_\lambda$ represents the absorption in band $\lambda$ relative to $A_V$ , based on the reddening law from \citet{allers06}."
 The quantity 9 is a constant whose significance we now discuss., The quantity $\beta$ is a constant whose significance we now discuss.
 The 2nd term on the right haud side of (1)) is designed to compensate lor spurious biases towards low Tig values resulting [rom the potential presence of long-waveleneth excesses due to cireumstellar dust., The 2nd term on the right hand side of \ref{eq1}) ) is designed to compensate for spurious biases towards low $T_{\rm eff}$ values resulting from the potential presence of long-wavelength excesses due to circumstellar dust.
 It is necessitated by the increasing degeneracy which exists between Ziy and “ly as the temperature increases and the Bavleighi-Jeans regime is approached., It is necessitated by the increasing degeneracy which exists between $T_{\rm eff}$ and $A_V$ as the temperature increases and the Rayleigh-Jeans regime is approached.
 Because of this degeneracy. the SED of a high-temperature object wilh large extinction can be reproduced equally well by a model based on a low-temperature object seen (hough small extinction.," Because of this degeneracy, the SED of a high-temperature object with large extinction can be reproduced equally well by a model based on a low-temperature object seen though small extinction."
 For such an object. the presence of even a relatively small lone-waveleneth excess due to circumstellar dust can favor a spurious low-temperature solution.," For such an object, the presence of even a relatively small long-wavelength excess due to circumstellar dust can favor a spurious low-temperature solution."
 To compensate for this effect. we have imposed a penalty against low values οἱ ely. Whose severity is controlled by the value of ο.," To compensate for this effect, we have imposed a penalty against low values of $A_V$, whose severity is controlled by the value of $\beta$."
 We have optimized the value of ;? using published photometry for a large number of spectroscopicallv-confinned brown dwarls of a variely of ages. as discussed in Appenclix 1: the optimal value was found to be 0.7.," We have optimized the value of $\beta$ using published photometry for a large number of spectroscopically-confirmed brown dwarfs of a variety of ages, as discussed in Appendix 1; the optimal value was found to be 0.7."
 This optimization incorporated (he results of our own spectroscopic observations of 7 of the objects in our p Oph source extraction list (Marshetal.2010)., This optimization incorporated the results of our own spectroscopic observations of 7 of the objects in our $\rho$ Oph source extraction list \citep{mar10}.
 In evaluating o. the model flixes were caleulated by logarithmic interpolation of the Labulated model [Inxes with respect to Zar and wavelength.," In evaluating $\phi$, the model fluxes were calculated by logarithmic interpolation of the tabulated model fluxes with respect to $T_{\rm eff}$ and wavelength."
 The wavelength interpolation was necessitated by the mismatch between the observational wavebands (the 3.6 jm and 4.5 jan bands) and the £’ and M. bands for which the model photometry was tabulated., The wavelength interpolation was necessitated by the mismatch between the observational wavebands (the 3.6 $\mu$ m and 4.5 $\mu$ m bands) and the $L'$ and $M$ bands for which the model photometry was tabulated.
 Although the wavelength mismatches are relatively small. some errors in theinterpolation," Although the wavelength mismatches are relatively small, some errors in theinterpolation"
"iotometry is 1716 PSP in ας (CAdehuan-ALcC""thvetal. 2006).",photometry is 4 PSF in $r$ -band \citep{sdss}.
. Thus some systems that cannot be distinguished iu the SDSS data cau be identified iu the DLS., Thus some systems that cannot be distinguished in the SDSS data can be identified in the DLS.
" A total of eight pairs in the DLS are nüssiug frou SDSS,", A total of eight pairs in the DLS are missing from SDSS.
 Objects with stellar light profiles. ie. ACN i liigh redshift. may be prefercutially excluded roni our sample because DLS objects with stellar ieht profiles are not targeted for spectroscopy.," Objects with stellar light profiles, i.e. AGN at high redshift, may be preferentially excluded from our sample because DLS objects with stellar light profiles are not targeted for spectroscopy."
 However. with. the DLS i29 mae 2 surface. »ielatuess Lindt aud the 0799 resolution. we should rc alde to detect bright host galaxy bulees to the iuit ποτ=0.376.," However, with the DLS 29 mag $^{-2}$ surface brightness limit and the 9 resolution, we should be able to detect bright host galaxy bulges to the limiting $z=0.376$."
 We discuss ACN detection iu LOL! detail in refaen.., We discuss AGN detection in more detail in \\ref{agn}.
" T1ο Hectospec observations identify a spectroscopic xür Of emission line galaxies at R.A.= 9h(ΠΡΟ, decl."," The Hectospec observations identify a spectroscopic pair of emission line galaxies at R.A.= $9^{\mbox{\small h}}16^{\mbox{\small m}}58_{.}^{\mbox{s}}903$, decl."
 σποτ 120001., = $^{\circ}$ 597 (J2000).
 The nage. whic1 shows a clouble-uuclei exteuded object. aud he spectra are available on our.," The image, which shows a double-nuclei extended object, and the spectra are available on our."
 We do iot include this object in our analysis because the selection of spectroscopic pairs is not complete or ΠΠ, We do not include this object in our analysis because the selection of spectroscopic pairs is not complete or uniform.
 Οµ]ν 230 ealaxics in our sample lack spectra (2.3%)., Only 230 galaxies in our sample lack spectra $2.3\%$ ).
 Of these. ~21 galaxies (1054) may be members of pairs that satisfy our criteria for analysis refsaunple)).," Of these, $\sim 24$ galaxies $10\%$ ) may be members of pairs that satisfy our criteria for analysis \\ref{sample}) )."
 The fraction of galaxies in pairs in our sample as a whole is ~1054 refsaiple))., The fraction of galaxies in pairs in our sample as a whole is $\sim10\%$ \\ref{sample}) ).
 hnuage processing for the DLS uses t16 IRAF package to correct for bias. to Hat field. to andobtain sec astromoetrie calibratio1.," Image processing for the DLS uses the IRAF package to correct for bias, to flat field, and to obtain basic astrometric calibration."
 Wittmanetal.(2006) construct a sacked DLS inage of each subfield iu the R-baud. correcting for cosmic raves and saturated objects. fixing the astrometric calibration. and correcting the mage shape aud photomerv for optical distortion.," \citet{wittman06}
 construct a stacked DLS image of each subfield in the $R$ -band, correcting for cosmic rays and saturated objects, fixing the astrometric calibration, and correcting the image shape and photometry for optical distortion."
 The uncertaiutyv in the R-band magnitude ix~0.05 mas., The uncertainty in the $R$ -band magnitude is $\sim 0.05$ mag.
" We derive the galaxy catalog usiic SExtractor (Bertin&Α1101ts1996) on the final stackec Πασος,", We derive the galaxy catalog using SExtractor \citep{sextract} on the final stacked images.
 We rounove objects within the radius of the diffraction pattern around bright stars., We remove objects within the radius of the diffraction pattern around bright stars.
 The total area excluded around --778 stars is: 0.215ET deg?2 (5.1%H of thο SUTVOYV)., The total area excluded around 778 stars is 0.215 $^2$ $5.4\%$ of the survey).
 The absolute maguitide Ay calculation includes both k- aud evolutionary (ο) corrections., The absolute magnitude $M_R$ calculation includes both k- and evolutionary (e) corrections.
 We use the ο correction determined by Ας(2001) using the Pegase cocο (LeBorene&Rocca-Voliucrage2002)., We use the k+e correction determined by \citet{annis01} using the Pegase code \citep{pegase}.
.. The kle correction requires classification as one of niue galaxy types (bright cluster galaxy. elliptical.S0. Sa. Sb. She. Sc. Sd. and meegulu).," The k+e correction requires classification as one of nine galaxy types (bright cluster galaxy, elliptical,S0, Sa, Sb, Sbc, Sc, Sd, and irregular)."
 We classify the galaxies according o the SDSS (gy7) color aud galaxy redshift., We classify the galaxies according to the SDSS $g-i$ ) color and galaxy redshift.
 We apply the SDSS ας k|6 correction to t DLS R-banud photometry becase the shape of the filter loa passes are simular and the difference iu correctiois for the different xuads is negligible., We apply the SDSS $r$ -band k+e correction to the DLS $R$ -band photometry because the shape of the filter band passes are similar and the difference in corrections for the different bands is negligible.
 We estimate the uncertainty iu the ke correction roni the difference in the correction for adjacent ealaxy types., We estimate the uncertainty in the k+e correction from the difference in the correction for adjacent galaxy types.
 The uncertain ranges from a uaxiunun of ~0.02 mae at :05 to 0.03 123 mag at 0.38. depen i the galaxy ype.," The uncertainty ranges from a maximum of $\sim0.02$ mag at $z=0.08$ to 0.03 – 0.13 mag at $z=0.38$, depending on the galaxy type."
 There are 5l eaaxies dà our i=0.0803706 seunpe with a Dectospec redshift but no SDSS photometry for ealaxy classification aud ο correction. including three galaxies iu major udis refinajpair]).," There are 51 galaxies in our $z=0.080-0.376$ sample with a Hectospec redshift but no SDSS photometry for galaxy classification and k+e correction, including three galaxies in major pairs \\ref{majpair}) )."
 lu 31 of the 51 cases. we use the DLS (V 7) color to estimate the (y7) color for the -e correction.," In 34 of the 51 cases, we use the DLS $V-R$ ) color to estimate the $g-i$ ) color for the k+e correction."
 For the l7 cases where we lack voth DLS (V... RH) aud SDSS data (including two major pairs). we assume the ealaxy type is Sa.," For the 17 cases where we lack both DLS $V-R$ ) and SDSS data (including two in major pairs), we assume the galaxy type is Sa."
 je ο correction for the galaxy. type Sa is in he middle of the rauge for the differeut galaxy vpes: choosing a different k|e correction lias no sieuificaut effect on the results., The k+e correction for the galaxy type Sa is in the middle of the range for the different galaxy types; choosing a different k+e correction has no significant effect on the results.
 At the redshift += 13. the apparent R= 20.3 Init corresponds to a ste corrected absolute magnitude Mj=21.1 or a typical late-tvpe galaxy. aud to Mg=21.8 or a typical early-type galaxy.," At the redshift $z=0.3$ , the apparent $R=20$ .3 limit corresponds to a k+e corrected absolute magnitude $M_R = -21.4$ for a typical late-type galaxy, and to $M_R =-21.8$ for a typical early-type galaxy."
 The Iectospee spectrograph 270 liue 1 erating vields ~6 pixel| dispersion over the wavelength range 3700-9300 (Fabricautctal. 2008).," The Hectospec spectrograph 270 line $^{-1}$ grating yields $\sim 
6$ $^{-1}$ dispersion over the wavelength range 3700-9300 \citep{fabricant08}. ."
. The 300 optical fibers with 1755 diameter are placed racially within a 1 diameter field., The 300 optical fibers with 5 diameter are placed radially within a $^{\circ}$ diameter field.
 About 30 of the fibers measure sky backeround durinug cach poiutiug., About 30 of the fibers measure sky background during each pointing.
 We reduce Iectospec data with the standard Tarvard-Sinithsonian Center for Astroplivsics, We reduce Hectospec data with the standard Harvard-Smithsonian Center for Astrophysics
It is useful to beein bw considering theoretical expectations for fluid inotious in a stably stratified medium.,It is useful to begin by considering theoretical expectations for fluid motions in a stably stratified medium.
 We first review the case for a hvdrodynaiiic fiuid. which im clusters is stabilized by a strong cutropy eradieunt.," We first review the case for a hydrodynamic fluid, which in clusters is stabilized by a strong entropy gradient."
 We then cousider the ΑΠΟ case where anisotropic conduction alters the picture., We then consider the MHD case where anisotropic conduction alters the picture.
 Tn the purely hwdrodywuanüc case. the imediu is convectively unstable ifgluS/dlur<<(Q0.," In the purely hydrodynamic case, the medium is convectively unstable if$d\, {\rm ln} S/d\, {\rm ln} r < 0$."
 Ou the other haud. if the amedimm is stably stratified with dluS/dlur20. a fluid clement which is adiabatically displaced about its equilibriumn position will experieuce a restoring force f=phραlnSfdr)Av. assunune it remus du pressure balance with its sumroundiues.," On the other hand, if the medium is stably stratified with $d\, {\rm ln} S/d\, {\rm ln} r > 0$, a fluid element which is adiabatically displaced about its equilibrium position will experience a restoring force $f = \rho \ddot{r}= - \rho g (3/5) (d\, {\rm ln} S/d r) \Delta r$, assuming it remains in pressure balance with its surroundings."
 This is the equation for simple hinuonuic motion. at frequency: the frequency.," This is the equation for simple harmonic motion, at frequency: the frequency."
 A formal WIND. analysis reveals that the dispersion relation for such g-1iodos. where eravity is the restoring force. is: 3 ⋅↖↖⇁∐↸∖↥⋅↸∖∕≍⋮−∶∕≍⋮−⊥↖∕≍⋮⊼∙⊺∐↴∖↴∐⊔∐∐∖≼," A formal WKB analysis reveals that the dispersion relation for such $g$ -modes, where gravity is the restoring force, is: where $k^{2}=k_{\perp}^{2} + k_{r}^{2}$."
∐⋜↧↑↸∖↕⋅↖↽↑∐↸∖∐ ↕∐∏≻∐↸∖↴∖↴↑∐⋜↧↑↕⋟∪↥⋅∪↴∖↴↸⊳↕∐⋜↧↑∪↥⋅⋅↖↽⋯∪∩∖↴∖↴↖↖↽↕↑∐↥⋅↸∖⋜↧↕∕↖⋮∣⋅∙⊔∩ ↙⊨⋅∐∖↽∶∪↑∐↸∖↥⋅↖↖⇁↕↴∖↴↸∖↑↕∐∖↕⊔∪≼∐∖↴∖↴∐⋜↕↖↽↸∖↕∐↓⋜↕∶↴∙⊾↕∐⋜∐⋅⋅↖⇁↥⋅⋜∥∐⋜↧↕ ↖↖↽⋜↧↖↽↸∖↕∐⋯∐⋈∖↥⋅⋜⋯≺↧⋜⋯∖↸∖↖⇁⋜⋯↸∖↴∖↴↸⊳↸∖∐∙↕⋟∐⋅↖↽," This immediately then implies that for oscillatory modes with real $k_{r}$, $\omega < \omega_{\rm BV}$ ; otherwise the modes have imaginary radial wavenumber and are evanescent."
↴∖↴↕↸⊳⋜↧∐⋅↖⇁∙∪∐↸∖⋯∐ ⋜↧↕↖↖⇁⋜↧⋅↖⇁↴∖↴⋜↧↸⊳↕∐↸∖↖↽↸∖⋜↧↕∪↖↖↽↕≯↥⋅↸∖≺∣⋯∖∐↸⊳⋅↖⇁↥⋅↸∖↴∖↴↻∪∐↴∖↴↸∖↴⋝⋅↖↽⋯⋜∐↘↽↕∐∶↴⋁↑∐↸∖ uode progressively more tangential. but it is impossible o drive the system at frequencies higher than the uaximuuni respouse frequency of wpy. correspondius o completely vertical oscillations.," Physically, one can always achieve a low frequency response by making the mode progressively more tangential, but it is impossible to drive the system at frequencies higher than the maximum response frequency of $\omega_{\rm BV}$, corresponding to completely vertical oscillations."
 This thus implies hat waves driven at frequencies wo<<wpy can be resonantly excited. aud will be trapped. reflected aud ocused imside the resonance radius where wpy=uw.," This thus implies that waves driven at frequencies $\omega < \omega_{\rm BV}$ can be resonantly excited, and will be trapped, reflected and focused inside the resonance radius where $\omega_{\rm BV}=\omega$."
 This las duplications for how stirring can be amplified iuto voluine-filling turbulence., This has implications for how stirring can be amplified into volume-filling turbulence.
 For our purposes. an even more interesting aspect of such ganodes is that they tend to be preferentially tanecutial. as has been widely confirmed in calculations of stellar pulsation (?).. ealaxy clusters (?).. aud the earth's atmosphere and oceans (7)..," For our purposes, an even more interesting aspect of such $g$ -modes is that they tend to be preferentially tangential, as has been widely confirmed in calculations of stellar pulsation \citep{cox80}, galaxy clusters \citep{lufkin95}, and the earth's atmosphere and oceans \citep{riley00}."
 Let us consider a osimple order of imaguitude arguueut for wl lis should be the case., Let us consider a simple order of magnitude argument for why this should be the case.
 For simplicity. cousider an isothermal atmosphere.," For simplicity, consider an isothermal atmosphere."
 Iu steady state. aud adopting he Boussinesq approximation (such that V:v= 0). he coutinuity equation reads v:Vp= 0.," In steady state, and adopting the Boussinesq approximation (such that $\nabla \cdot {\bf v} =0$ ), the continuity equation reads ${\bf v} \cdot \nabla \rho =0$ ."
" Since the cekeround denusitv p, does not vary horizoutallw. the »erturbed continuity equation is v:Vo!|e;dp,/d:=0."," Since the background density $\rho_{o}$ does not vary horizontally, the perturbed continuity equation is ${\bf v} \cdot \nabla \rho^{\prime} + v_{z} d\rho_{o}/dz = 0$."
 If v|eC. [eu~IV. and £ is a characteristic lengthscale. hen: The buovancy forces frou stratification dnipose additional coustraits.," If $|{\bf v}| \sim U$, $|{ v_{z}}| \sim W$, and $L$ is a characteristic lengthscale, then: The buoyancy forces from stratification impose additional constraints."
 The vorticity evolution —form of ↑∐↸∖⋯∪⋯↸∖∐↑∏⋯↸∖≺∣∏⋜↧⊓∪∐−↕∪↥⋅∙↙∕≓↖↖⇁⋜↧↖↽↸∖↴∖↴⋖↕⋅⋜∖⋅∙⋜↧↴∖↴↴∖⋯⋯∐∶↴∙⊾ ⋅⋅≻⊳ ⋅ ⋅ which gives to order of iiagnitudo. Combining equation (5)) aud (7)). we obtain: where Fr ois the Froude number. which m this case conrpares the characteristic frequencies of inertial aud gravitational (buovaucv) forces.," The vorticity evolution form of the momentum for $g$ -waves (i.e., assuming $\delta P/P \ll \delta \rho/\rho$ ) is \citep{lufkin95}: which gives to order of magnitude, Combining equation \ref{eq:OOM_continuity}) ) and \ref{eq:OOM_vorticity}) ), we obtain: where $Fr$ is the Froude number, which in this case compares the characteristic frequencies of inertial and gravitational (buoyancy) forces."
 Ri~Εν is often known as the Richardson umber., $Ri\sim 1/Fr^{2}$ is often known as the Richardson number.
 Equation (8)) reveals two important properties., Equation \ref{eq:froude}) ) reveals two important properties.
 When turbulence is abseut or verv weak. Fr&1. motion is fimdameutally horizoutal. WU: strong restoring buovancy forces oppose lotion in the vertical direction.," When turbulence is absent or very weak, $Fr \ll 1$, motion is fundamentally horizontal, $W \ll U$: strong restoring buoyancy forces oppose motion in the vertical direction."
 Ilowever. stronger stirring. FrzQU. can overwhelu such forces and re-establish isotropic turbulence.," However, stronger stirring, $Fr \gsim \mathcal{O}(1)$, can overwhelm such forces and re-establish isotropic turbulence."
 Wile all of the preceding discussion hold ouly for short wavelength modes in the WISB approximation. uuncerical simulations of non-linear evolution of global ganodes iu ealaxy clusters excited dy orbiting galaxies show that hey can still be very usefully interpreted in teris of the near WISB analvsis (?)..," While all of the preceding discussion hold only for short wavelength modes in the WKB approximation, numerical simulations of non-linear evolution of global $g$ -modes in galaxy clusters excited by orbiting galaxies show that they can still be very usefully interpreted in terms of the linear WKB analysis \citep{lufkin95}."
 Tow do these expectations change iu the presence of anisotropic conduction?, How do these expectations change in the presence of anisotropic conduction?
 Tere. behavior is governed w temperature rather than cutropy eradicuts. and the atmosphere is buovautlv mustable regardless of the sigu of VT.," Here, behavior is governed by temperature rather than entropy gradients, and the atmosphere is buoyantly unstable regardless of the sign of $\nabla T$."
 For the ΠΟΙ. the instability saturates once the nagnetic field is perpendicular to VT.," For the HBI, the instability saturates once the magnetic field is perpendicular to $\nabla T$."
 Displaced fluid clemeuts now experience a restoring force f=prpatdlutfdirjAr (?).. and hence oscillate about their equilibrimn position at a frequency: which is typically ~2 times smaller than wp.," Displaced fluid elements now experience a restoring force $f = \rho \ddot{r} = - \rho g (d\, {\rm ln} T/d r) \Delta r$ \citep{sharma09a}, and hence oscillate about their equilibrium position at a frequency: which is typically $\sim$ 2 times smaller than$\omega_{\rm BV}$ ."
 Note that this asstaues fogo1Kt.~Lies (where 8~ ds the conductivity coefficient) so that adisplaced blob's temperature is determined by conductive rather than adiabatic cooling., Note that this assumes $t_{\rm cond} \sim L^{2}/\kappa \ll t_{\rm sc} \sim L/c_{s}$ (where $\kappa \sim v_{e} l_{e}$ is the conductivity coefficient) so that adisplaced blob's temperature is determined by conductive rather than adiabatic cooling.
" This requires: which in any case is required by the WEKD assiuption. Ler, "," This requires: which in any case is required by the WKB assumption, $L \ll r$ ."
The dispersion relation is (?)::, The dispersion relation is \citep{quataert08}: :
to obtain a system of n equations and à unknowns.,to obtain a system of $n$ equations and $n$ unknowns.
 We can numerically solve this svstem of » equations to obtain n unknown parameters (for our models n» is less than 10)., We can numerically solve this system of $n$ equations to obtain $n$ unknown parameters (for our models $n$ is less than 10).
 Each result. of the solution of the svstem will be a set of n values., Each result of the solution of the system will be a set of $n$ values.
 When we go to solve the system. we do not obtain a single solution. but a set of solutions. since the system is highly non linear: these solutions have to be selected. by means of some selection criteria in order to obtain values of parameters which give rise to physically plausible mocdeLg," When we go to solve the system, we do not obtain a single solution, but a set of solutions, since the system is highly non linear; these solutions have to be selected by means of some selection criteria in order to obtain values of parameters which give rise to physically plausible models."
 The search for the roots of the system begins with a choice of a range of parameter values: one must be sure that there are no roots outside the chosen range. and. also avery large interval does not necessarily include all roots. increasing CPU time.," The search for the roots of the system begins with a choice of a range of parameter values: one must be sure that there are no roots outside the chosen range, and also a very large interval does not necessarily include all roots, increasing CPU time."
 We choose a physically acceptable interval for parameters. giving Av random starting points o the algorithm. where A. is a number fixed by theuser’.," We choose a physically acceptable interval for parameters, giving $\mathcal{N}$ random starting points to the algorithm, where $\mathcal{N}$ is a number fixed by the."
 From these starting points the CPU tries to converge to he solutions. using the Newton's method.," From these starting points the CPU tries to converge to the solutions, using the Newton's method."
 In. particular. o assign the same probability to the starting values of a »vwameter in the relative range. we generate these values uniformly. cüstributed. also to avoid introducing any bias in he search.," In particular, to assign the same probability to the starting values of a parameter in the relative range, we generate these values uniformly distributed, also to avoid introducing any bias in the search."
 We have developed. a simple code. namedLePRe*.. Lensing Parameters Reconstruction. written [orMathematica.," We have developed a simple code, named, Lensing Parameters Reconstruction, written for."
 Generally. solving the equation svstem vields AZUN. solutions. because for some values of the starting points Newton's method. fails to converge to a solution.," Generally, solving the equation system yields $\mathcal{M} <
\mathcal{N}$ solutions, because for some values of the starting points Newton's method fails to converge to a solution."
 These © solutions are not all physically acceptable and to sort among these we have to impose a set of selection criteria. that constrains parameter values to. physically plausible ranges.," These $\mathcal{M}$ solutions are not all physically acceptable and to sort among these we have to impose a set of selection criteria, that constrains parameter values to physically plausible ranges."
 Schematicallv. we can deseribe these selection criteria as following.," Schematically, we can describe these selection criteria as following."
which are interpreted as a local manifestation of a more global phenomenon.,which are interpreted as a local manifestation of a more global phenomenon.
 In this paper we ask whether we can obtain better-fitting catacubes by replacing the kinematic moclel of MPL with the dynamical model that FBOG and ΕΟΝ developed: to account. for. observations of external galaxies., In this paper we ask whether we can obtain better-fitting datacubes by replacing the kinematic model of MF11 with the dynamical model that FB06 and FB08 developed to account for observations of external galaxies.
" In Section 2 we describe the model we will be using. which is an updated version of that described. in I""BOs. and explain how we evaluate the fit between a mocel datacube and that of the LAB survey."," In Section \ref{model} we describe the model we will be using, which is an updated version of that described in FB08, and explain how we evaluate the fit between a model datacube and that of the LAB survey."
 In Section 3.1. we investigate whether the LIL thick disc of the Milkv Way can oe produced. by a pure galactic fountain., In Section \ref{pureFountain} we investigate whether the $\hi$ thick disc of the Milky Way can be produced by a pure galactic fountain.
 Then in Section 3.2. we present evidence that this fountain interacts with he ambient coronal gas., Then in Section \ref{accretingFountain} we present evidence that this fountain interacts with the ambient coronal gas.
 In Section 77. we use our model to infer the large-scale structure of the Galaxys ihalo. and argue that there is a fundamental distinction oetween LIVCs. which have an extragalactic origin. and Intermediate Velocity Clouds. which are dominated by ountain gas.," In Section \ref{sec:layer} we use our model to infer the large-scale structure of the Galaxy's halo, and argue that there is a fundamental distinction between HVCs, which have an extragalactic origin, and Intermediate Velocity Clouds, which are dominated by fountain gas."
 In Section ?? we show that our model's »uwameter values are remarkably consistent with results rom) earlier studies of (a) small-scale hyclrodyvnamical simulations of turbulent mixing. and. (b) analysis of data or external galaxies.," In Section \ref{discussion} we show that our model's parameter values are remarkably consistent with results from earlier studies of (a) small-scale hydrodynamical simulations of turbulent mixing, and (b) analysis of data for external galaxies."
" Section 727. sums up and looks to the νο,", Section \ref{conclusions} sums up and looks to the future.
 We set up a galactic fountain model for the Milky Way by integrating orbits of gas clouds in the Galactic potential., We set up a galactic fountain model for the Milky Way by integrating orbits of gas clouds in the Galactic potential.
 The latter has been constructed. using the. standard decomposition into a bulge. stellar ancl gaseous discs. and a dark-matter halo.," The latter has been constructed using the standard decomposition into a bulge, stellar and gaseous discs, and a dark-matter halo."
 The physical parameters and functional forms for the various components have been taken from ModelLL.p. 113..," The physical parameters and functional forms for the various components have been taken from \citet[][Model II, p.~113]{GD2}."
 The clouds are ejected. from. the disc into the halo. and we follow their orbits until they return to the disc as described in EBO06.," The clouds are ejected from the disc into the halo, and we follow their orbits until they return to the disc as described in FB06."
 The probability of ejection at speed ο at angle @ with respect to the vertical direction is where fy is a characteristic velocity and D is a constant hat determines the extent to which clouds are. ejected »erpendieular to the disc., The probability of ejection at speed $v$ at angle $\theta$ with respect to the vertical direction is where $h_{\rm v}$ is a characteristic velocity and $\Gamma$ is a constant that determines the extent to which clouds are ejected perpendicular to the disc.
 Here we treat fe as à. free xwameter but adopt P—5 from BOG such a Large value of E implies strong collimation of ejecta towards the normal o the plane by cool gas near the plane., Here we treat $h_{\rm v}$ as a free parameter but adopt $\Gamma=5$ from FB06 – such a large value of $\Gamma$ implies strong collimation of ejecta towards the normal to the plane by cool gas near the plane.
 The integration is performed in the (2.2) plane and then spread. through he Galaxy by assuming azimuthal svmmetry.," The integration is performed in the $(R,z)$ plane and then spread through the Galaxy by assuming azimuthal symmetry."
 Phe number of clouds ejected as a function of radius is proportional to he SER. density at that racius (see Section 2.2))., The number of clouds ejected as a function of radius is proportional to the SFR density at that radius (see Section \ref{SFlaw}) ).
 Other xwameters and properties of the model are as described in FBOG except for some updates to the model that are specified in the following subsections., Other parameters and properties of the model are as described in FB06 except for some updates to the model that are specified in the following subsections.
 We follow the orbits of gas that leaves the disc at relatively low temperatures E 10!IX) after. being swept up in the expansion of a superbubble that is powered. by multiple supernova explosions., We follow the orbits of gas that leaves the disc at relatively low temperatures $\lsim 10^4 \K$ ) after being swept up in the expansion of a superbubble that is powered by multiple supernova explosions.
 Ehe virial-temperature gas filling the bubble is presumed to contain negligible mass and to merge with the pre-existing coronal eas., The virial-temperature gas filling the bubble is presumed to contain negligible mass and to merge with the pre-existing coronal gas.
 In the Milky. Way. emission rom the ihalo occurs mainly at negative line-of-ight. velocities (7)..," In the Milky Way, emission from the halo occurs mainly at negative line-of-sight velocities \citep{vanWoerden+85}."
 ‘This observation suggests that the neutral gas is seen mainly as it descends to the plane. so a significant. fraction. of he ejected gas must be ionised. and. indeed from their 3D kinematic mocdels of the galactic halo METL estimated that he rising clouds are 60 per cent ionised.," This observation suggests that the neutral gas is seen mainly as it descends to the plane, so a significant fraction of the ejected gas must be ionised, and indeed from their 3D kinematic models of the galactic halo MF11 estimated that the rising clouds are $\sim60$ per cent ionised."
 LE we are to compare our models with the LAB survey of neutral eas. we ave to find a wav to estimate the fraction of a cloud that is ionisecl at cach part of its orbit.," If we are to compare our models with the LAB survey of neutral gas, we have to find a way to estimate the fraction of a cloud that is ionised at each part of its orbit."
 FBOG built models with and without phase-change. the atter having the whole orbit. “visible” as geas. the former only the descending part.," FB06 built models with and without phase-change, the latter having the whole orbit “visible” as gas, the former only the descending part."
 Here we refine his treatment as follows., Here we refine this treatment as follows.
 We assume that a cloud ejected rom the cise with a kick velocity ey will become visible nneutral) only when where ον ds the vertical component of the clouc’s velocity. 6o2Mekcos9 and OSfiiysLis a free parameter that regulates the “visibilitw” of the cloud.," We assume that a cloud ejected from the disc with a kick velocity $v_{\rm kick}$ will become visible neutral) only when where $v_z$ is the vertical component of the cloud's velocity, $v_{z,0}\!=\!v_{\rm kick}\times\cos\theta$ and $0\!\le\!f_{\rm
ion}\!\le\!1$ is a free parameter that regulates the “visibility” of the cloud."
 LW fi=1. the fountain cloud is neutral only for negative values of e; iin the descending part of the orbit). while if fi=0. the cloud will always be visible.," If $f_{\rm ion}\!=\!1$, the fountain cloud is neutral only for negative values of $v_z$ in the descending part of the orbit), while if $f_{\rm ion}=0$, the cloud will always be visible."
 We infer the correct value of fieun rom the data cube (see Sections 3.1. and 3.2))., We infer the correct value of $f_{\rm ion}$ from the data cube (see Sections \ref{pureFountain} and \ref{accretingFountain}) ).
 We assume that the strength of the supernova feedback. tthe number of eas clouds ejected. per surface clement. is proportional to the local SER.," We assume that the strength of the supernova feedback, the number of gas clouds ejected per surface element, is proportional to the local SFR."
 Given that our model is axisvmametric. we require an estimate of the SER as a function of /?.," Given that our model is axisymmetric, we require an estimate of the SFR as a function of $R$."
 In BOG and FBOS that was estimated from the surface density of neutral plus molecular gas using the Sehmiicdt-Wennicutt law (??)..," In FB06 and FB08 that was estimated from the surface density of neutral plus molecular gas using the Schmidt-Kennicutt law \citep{Schmidt59, Kennicutt98}."
 The SER was then set. to zero bevoncl a certain cut-oll racius., The SFR was then set to zero beyond a certain cut-off radius.
 Here. we refine. this procedure as follows.," Here, we refine this procedure as follows."
 We use a sample of galaxies. with known racial trends of gas surface density and SER. density and derive à star-formation law for molecular gas., We use a sample of galaxies with known radial trends of gas surface density and SFR density and derive a star-formation law for molecular gas.
 We then use this law to obtain the Milky Ways SER as a function of & from the axisvnunetrisecl surface density of molecular eas given in ?.., We then use this law to obtain the Milky Way's SFR as a function of $R$ from the axisymmetrised surface density of molecular gas given in \cite{bm98}.
 Fie., Fig.
 1 isa plot of SER density versus the molecular gas density., \ref{fSFlaw} is a plot of SFR density versus the molecular gas density.
 The filled triangles show data points for galaxies in the sample of ? that have stellar masses above 10A7.., The filled triangles show data points for galaxies in the sample of \citet{Leroy+08} that have stellar masses above $10^9 M_{\odot}$.
 Each point refers to an azimuthal average at the given radius., Each point refers to an azimuthal average at the given radius.
 The scatter in the relation between gas density and SER is remarkably small considering that the points come from very dillerent. galaxies., The scatter in the relation between gas density and SFR is remarkably small considering that the points come from very different galaxies.
 We fitted these points with a power law of the form: where Mspr and My are. respectively. the SET and the molecular gas density.," We fitted these points with a power law of the form: where $\Sigma_{\rm SFR}$ and $\Sigma_{\rm H_2}$ are, respectively, the SFR and the molecular gas density."
 We find The slope of this relation is very dillerent. from. the standard: Schimidt-Ixennicutt law. mainlv because we are," We find The slope of this relation is very different from the standard Schmidt-Kennicutt law, mainly because we are"
adopting the values 0.2. 0.45. 0.6. 0.85 ancl 1 for fy.,"adopting the values 0.2, 0.45, 0.6, 0.85 and 1 for $f_{\rm d}$ ."
 Poy is the stellar potential with the maximal stellar mass-to-lieht ratio and. Py. is the potential of the dark halo which is constrained by the observed. rotation curve., $\Phi_{\rm{stellar}}$ is the stellar potential with the maximal stellar mass-to-light ratio and $\Phi_{\rm{halo}}$ is the potential of the dark halo which is constrained by the observed rotation curve.
 For a fair comparison. all the simulations in this series of increasing disk [fraction were run for the same amount of total time. namely ~ 1.6 Cir.," For a fair comparison, all the simulations in this series of increasing disk fraction were run for the same amount of total time, namely $\sim$ 1.6 Gyr."
 A logarithmic plot (figure 8)) of the gas density for the sequence of simulations. shows that the density. contrast of the non-axisvmummetric features in the gas increases as the disk fraction increases.," A logarithmic plot (figure \ref{morph_diskfrac}) ) of the gas density for the sequence of simulations, shows that the density contrast of the non-axisymmetric features in the gas increases as the disk fraction increases."
 Fhis corroborates Figure 21. in this paper ancl Figure 8 of Paper E which. quantified. this trend by taking the average of the amplitude of the velocity deviations from axisvmnmetry. and found that this average increases more or less linearly with the disk fraction.," This corroborates Figure \ref{wiggle_amplitude} in this paper and Figure 8 of Paper I which quantified this trend by taking the average of the amplitude of the velocity deviations from axisymmetry, and found that this average increases more or less linearly with the disk fraction."
 In addition to the change in the density. contrast. figure S shows that features become more “angular” with increasing disk fraction.," In addition to the change in the density contrast, figure \ref{morph_diskfrac} shows that features become more “angular” with increasing disk fraction."
 For example. in the lowest disk fraction case (20 per cent disk). the spiral arms appear to be rounded and smooth in their curvature.," For example, in the lowest disk fraction case (20 per cent disk), the spiral arms appear to be rounded and smooth in their curvature."
 As the disk fraction is increased to 85 per cent. the lower spiral armi develops a squareness which is even more pronounced Lor the LOO per cent. disk [raction case.," As the disk fraction is increased to 85 per cent, the lower spiral arm develops a squareness which is even more pronounced for the 100 per cent disk fraction case."
 ‘To display these changes in amplitude and morphology in more detail. figure 9.— plots the density profile along an azimuthal cut through the disk at a position angle 0=202.57.," To display these changes in amplitude and morphology in more detail, figure \ref{denscuts_all} plots the density profile along an azimuthal cut through the disk at a position angle $\vartheta = \rm 202.5^{\circ}$."
 The top panel of this figure shows that the smaller amplitude density. peak (lower spiral arm) located abd? 6 kpe moves outwards with increasing clisk fraction sullering larger shifts in position with changing clisk fraction than the larger amplitude peaks at smaller. radii (2< 2 kpc)., The top panel of this figure shows that the smaller amplitude density peak (lower spiral arm) located at $R \approx$ 6 kpc moves outwards with increasing disk fraction suffering larger shifts in position with changing disk fraction than the larger amplitude peaks at smaller radii $R < $ 2 kpc).
 To further quantify this. we plot the gas density ancl velocity amplitude as a function of azimuth for f?στ 3 kpc (figure 10)).," To further quantify this, we plot the gas density and velocity amplitude as a function of azimuth for $R \approx$ 3 kpc (figure \ref{shock_az_fd}) )."
 In agreement with Woodward's (1975) analytic calculations (his figure 7). we find that the location of the density peak moves towards larger azimuths with increasing perturbation strength.," In agreement with Woodward's (1975) analytic calculations (his figure 7), we find that the location of the density peak moves towards larger azimuths with increasing perturbation strength."
 Figure 9 also indicates that as the disk fraction increases. regions of increasingly lower gas density appear immediately. adjacent to the gas density peaks in the arms.," Figure \ref{denscuts_all} also indicates that as the disk fraction increases, regions of increasingly lower gas density appear immediately adjacent to the gas density peaks in the arms."
 Hence to differentiate between models with cifferent fy a code has to perform well in the low density regions. which is one of the strengths of eric codes over particle codes in general and of the BCUN scheme in particular.," Hence to differentiate between models with different $f_{\rm d}$ a code has to perform well in the low density regions, which is one of the strengths of grid codes over particle codes in general and of the BGK scheme in particular."
 Alorphological changes with increasing disk fraction are not limited to spiral arms: they are also present in the inner region of the galaxy., Morphological changes with increasing disk fraction are not limited to spiral arms: they are also present in the inner region of the galaxy.
 Even at the resolution of our study. the erev-scale maps of the inner regions of the disk (right column of figure 8)). reveal the development of oll-axis shocks. the emergence of an oval ring around the inner region. and the strengthening of the 4/1 shockat the endpoints of this oval ring.," Even at the resolution of our study, the grey-scale maps of the inner regions of the disk (right column of figure \ref{morph_diskfrac}) ), reveal the development of off-axis shocks, the emergence of an oval ring around the inner region, and the strengthening of the 4/1 shockat the endpoints of this oval ring."
 Since we are ultimately. comparing the kinematical, Since we are ultimately comparing the kinematical
compressible plasma with a non-zero plasma pressure differ by terms of order (Κ.Α). even when we allow MHD radiation.,"compressible plasma with a non-zero plasma pressure differ by terms of order $(k_z R)^2$, even when we allow MHD radiation."
 If we neglect contributions proportional to (Κ.Α) then the simple conclusion is that the frequency of the kink wave and its damping due to resonant absorption are the same in the three cases that we have considered., If we neglect contributions proportional to $(k_z R)^2$ then the simple conclusion is that the frequency of the kink wave and its damping due to resonant absorption are the same in the three cases that we have considered.
 Again. as in the two previous cases. the eigenfunctions in the thin dissipative layer can be described by the functions £(7) and G(r).," Again, as in the two previous cases, the eigenfunctions in the thin dissipative layer can be described by the functions $\tilde{F}(\tau)$ and $\tilde{G}(\tau)$."
 Again the conclusion is that kink MHD waves are highly Alfvénnie in the dissipative layer., Again the conclusion is that kink MHD waves are highly Alfvénnic in the dissipative layer.
 This paper has examined the nature of MHD kink waves., This paper has examined the nature of MHD kink waves.
 This was done by determining the frequency. the damping rate and in particular the eigenfunctions of MHD kink waves for three widely different MHD waves cases: a compressible pressureless plasma. an incompressible plasma and a compressible plasma with non-zero plasma pressure which allows for MHD radiation.," This was done by determining the frequency, the damping rate and in particular the eigenfunctions of MHD kink waves for three widely different MHD waves cases: a compressible pressureless plasma, an incompressible plasma and a compressible plasma with non-zero plasma pressure which allows for MHD radiation."
 The overall conclusion ts that kink waves are very robust and do not care about the details of the MHD wave environment., The overall conclusion is that kink waves are very robust and do not care about the details of the MHD wave environment.
 In all three cases the frequency and the damping rate are for practical purposes the same as they differ at most by terms proportional to (k-R)-., In all three cases the frequency and the damping rate are for practical purposes the same as they differ at most by terms proportional to $(k_z R)^2$.
 In the magnetic flux tube the kink waves are in all three cases. to a high degree of accuracy incompressible waves with negligible pressure perturbations and with mainly horizontal motions.," In the magnetic flux tube the kink waves are in all three cases, to a high degree of accuracy incompressible waves with negligible pressure perturbations and with mainly horizontal motions."
 The main restoring force of kink waves in the magnetised flux tube is the magnetic tension force., The main restoring force of kink waves in the magnetised flux tube is the magnetic tension force.
 The gradient pressure force cannot be neglected except when the frequency of the kink wave is equal or slightly differs from the local Alfvénn frequency. t.e. in the resonant layer.," The gradient pressure force cannot be neglected except when the frequency of the kink wave is equal or slightly differs from the local Alfvénn frequency, i.e. in the resonant layer."
 The adjective fast is not the correct adjective to characterise kink waves., The adjective fast is not the correct adjective to characterise kink waves.
 If an adjective is to be used it should be Alfvénnic., If an adjective is to be used it should be Alfvénnic.
 However. it is better to realize that kink waves have mixed properties and cannot be put in one single box.," However, it is better to realize that kink waves have mixed properties and cannot be put in one single box."
caaunot be the exclusive source of the afterglow cnussion.,annot be the exclusive source of the afterglow emission.
evidence. Surma&Bender(1995) suggested Chat a metal-rieh (up to 2Z.) and ~7 Gyr old stellar population witha mass of (8.1+2.5)x10?M... exists in the central regions of NGC! 4365.,"evidence, \citet{sb95} suggested that a metal-rich (up to $2 \, Z_{\odot}$ ) and $\sim7$ Gyr old stellar population witha mass of $(8.1\pm2.5)\times10^9\, \msun$ exists in the central regions of NGC 4365."
 However. new results [rom integral lield spectroscopy of NGC 4365 (Davies and comparison with more recent SSP models indicate a Iuminositv-weiehted age for the stellar population in the galaxy of about 14 Gas over a continuous central region covering 0.5 effective radii.," However, new results from integral field spectroscopy of NGC 4365 \citep{dav01} and comparison with more recent SSP models indicate a luminosity-weighted age for the stellar population in the galaxy of about 14 Gyrs over a continuous central region covering 0.5 effective radii."
 NGC 4365 does not show any of the usual signs of merger activity within the last few Gvrs tidal tails. disturbed morphology ete.).," NGC 4365 does not show any of the usual signs of merger activity within the last few Gyrs tidal tails, disturbed morphology etc.)."
 Although it does have a kinematically decoupled core (ΝΟ). Daviesetal.(2001). found no age differences between the KDC and the remainder of the galaxy. and both components appear to have the same elevated Meg-to-Fe ratio.," Although it does have a kinematically decoupled core (KDC), \citet{dav01} found no age differences between the KDC and the remainder of the galaxy, and both components appear to have the same elevated Mg-to-Fe ratio."
" An intermediate age population maa still exist within the central e1"" of the nucleus. but its huminosity and associated mass are probably much smaller than required to account for the intermediate-age GC's (Carolloetal.1997)."," An intermediate age population may still exist within the central $\sim1\arcsec$ of the nucleus, but its luminosity and associated mass are probably much smaller than required to account for the intermediate-age GCs \citep{car97}."
". Even if the KDC indicates à merger history. it is therefore not clear that the event which produced the KDC is related to the origin of the intermediate-age GCs,"," Even if the KDC indicates a merger history, it is therefore not clear that the event which produced the KDC is related to the origin of the intermediate-age GCs."
 In present-day starbursts. formation of clusters is generally accompanied by formation of field stars. so it would be natural to expect a stellar population with Che same age ancl metallicity as the GCs.," In present-day starbursts, formation of clusters is generally accompanied by formation of field stars, so it would be natural to expect a stellar population with the same age and metallicity as the GCs."
 Indeed. the metallicity of the metal-rich GCs in many systems (ends to correlate with that of the parent galaxy (e.g.Forbes.Brodie.&Grillmair1997).. suggesting a common origin.," Indeed, the metallicity of the metal-rich GCs in many systems tends to correlate with that of the parent galaxy \citep[e.g.][]{fbg97}, suggesting a common origin."
 Η the old age of the general stellar population in NGC 4365 and the intermediate age of the metalrich GC's are both confirmed. an interesting question will be whether or not there are anv field stars associated with the intermediate-age GCs at all.," If the old age of the general stellar population in NGC 4365 and the intermediate age of the metal-rich GCs are both confirmed, an interesting question will be whether or not there are any field stars associated with the intermediate-age GCs at all."
 One possibility is that some intermecliale-age field stars do exist in NGC 4365. but constitute a too small fraction of the total population to strongly influence the integrated light.," One possibility is that some intermediate-age field stars do exist in NGC 4365, but constitute a too small fraction of the total population to strongly influence the integrated light."
 llow many voung stars could be “hidden” in NGC 4365 without significantly affecting (he integrated light?, How many young stars could be “hidden” in NGC 4365 without significantly affecting the integrated light?
 A detailed answer to this question is bevond the scope of the present paper. bul we can make some rough estimates.," A detailed answer to this question is beyond the scope of the present paper, but we can make some rough estimates."
 First. it is important to specify (he age of the voung population. because the Balmer line strengths and Iuminositv per unit mass both increase strongly al vounger ages. causing even a small number of very voung stars to have a slrong impact on the integrated spectrum.," First, it is important to specify the age of the young population, because the Balmer line strengths and luminosity per unit mass both increase strongly at younger ages, causing even a small number of very young stars to have a strong impact on the integrated spectrum."
 In a realistic modelling. the effects of different metallicities also need to be considered.," In a realistic modelling, the effects of different metallicities also need to be considered."
 However. here we are mainlv interested in the in Dalmer line strengths in (he presence of voung stars ancl for this purpose the metallicity effects can. al least as a first approximation. be considered second-order.," However, here we are mainly interested in the in Balmer line strengths in the presence of young stars and for this purpose the metallicity effects can, at least as a first approximation, be considered second-order."
 In the following we will simply assume Solar metallicity., In the following we will simply assume Solar metallicity.
 The Maraston SSP models that we have used in previous sections do not tabulate mass-to-light ratios. so for this purpose we use models by Druzual Charlot (2000: ccommn.).," The Maraston SSP models that we have used in previous sections do not tabulate mass-to-light ratios, so for this purpose we use models by Bruzual Charlot (2000; comm.)."
 These models eive M/L ratios for various filters. of which we will use the B-band numbers as the best approximation to the region around 112.," These models give M/L ratios for various broad-band filters, of which we will use the $B$ -band numbers as the best approximation to the region around $\beta$ ."
 For 2-Gyr ancl 5-Gyr populations. the D-band Iuminosities per unit mass," For 2-Gyr and 5-Gyr populations, the $B$ -band luminosities per unit mass"
temperatures less than 10 keV because the SZ first order correction to the Kompaneets approximation is sufficient in this temperature range to calculate the SZ etfect with a good precision (i.e. the SZ second order correction is very small).,temperatures less than 10 keV because the SZ first order correction to the Kompaneets approximation is sufficient in this temperature range to calculate the SZ effect with a good precision (i.e. the SZ second order correction is very small).
 The formalism based on an extension of the Kompaneets equation is obtained by solving the Boltzmann equation. using the limitations of the single scattering approximation., The formalism based on an extension of the Kompaneets equation is obtained by solving the Boltzmann equation using the limitations of the single scattering approximation.
 Thus. the Boltzmann equation (see Challinor Lasenby 1998) can only be applied to describe the SZ effect in optically thin clusters.," Thus, the Boltzmann equation (see Challinor Lasenby 1998) can only be applied to describe the SZ effect in optically thin clusters."
 The single scattering approximation is usually sufficient to calculate the SZ ettect in galaxy clusters., The single scattering approximation is usually sufficient to calculate the SZ effect in galaxy clusters.
 We have checked that the single scattering. approximation is valid for the considered simulated cluster., We have checked that the single scattering approximation is valid for the considered simulated cluster.
 To derive the temperature variance we propose to measure the SZ ettect at two frequencies for disentangling the SZ relativistο correction from the SZ contribution which can be described by the Kompaneets approximation., To derive the temperature variance we propose to measure the SZ effect at two frequencies for disentangling the SZ relativistic correction from the SZ contribution which can be described by the Kompaneets approximation.
 The easiest way to do this is to choose one frequency at 217 GHz (i.e. the crossover frequency of the SZ effect in the Kompaneets approximation) and another frequency at which the SZ etfect is described by the Kompaneets approximation., The easiest way to do this is to choose one frequency at 217 GHz (i.e. the crossover frequency of the SZ effect in the Kompaneets approximation) and another frequency at which the SZ effect is described by the Kompaneets approximation.
" The SZ relativistic corrections are very small compared with the total SZ signal for relatively cool clusters at frequencies which are not close to the crossover frequency and are not very high (we have checked that the SZ relativistic corrections at frequencies of y>S540 GHz are higher than of the total thermal SZ effect for a galaxy cluster with 4,7.=5 keV).", The SZ relativistic corrections are very small compared with the total SZ signal for relatively cool clusters at frequencies which are not close to the crossover frequency and are not very high (we have checked that the SZ relativistic corrections at frequencies of $\nu \gtrsim 540$ GHz are higher than of the total thermal SZ effect for a galaxy cluster with $k_{\mathrm{b}}T_{\rm e}=5$ keV).
 At very high frequencies the contribution of SZ relativistic corrections is significant., At very high frequencies the contribution of SZ relativistic corrections is significant.
 Thus. we choose the second frequency at 150 GHz. which is a good choice since the modern experiments such as the Atacama Cosmology Telescope!.. Atacama Pathfinder Experiment telescope*.. and South Pole Telescope make SZ observations at this frequency.," Thus, we choose the second frequency at 150 GHz, which is a good choice since the modern experiments such as the Atacama Cosmology Telescope, Atacama Pathfinder Experiment telescope, and South Pole Telescope make SZ observations at this frequency."
 The CMB intensity change produced by the SZ etfect from a galaxy cluster with inhomogeneous density and temperature distributions in the formalism based. on an extension of the Kompaneets equation. correct to first order in the expansion parameter O. can be written as where οι) is the spectral function taken from Eq. (," The CMB intensity change produced by the SZ effect from a galaxy cluster with inhomogeneous density and temperature distributions in the formalism based on an extension of the Kompaneets equation, correct to first order in the expansion parameter $\Theta$, can be written as where $g_{1}(x)$ is the spectral function taken from Eq. ("
28) of Challinor Lasenby (1998) and is given by Since the optical depth of the electron plasma r can be derived from X-ray observations or SZ observations. we rewrite Eq. (6)),"28) of Challinor Lasenby (1998) and is given by Since the optical depth of the electron plasma $\tau$ can be derived from X-ray observations or SZ observations, we rewrite Eq. \ref{form2}) )"
" as where (4,7)=[nkralindi is the temperature averaged along the line-of-sight and GRPoy)=πο...—[nudi is the squared. temperature averaged along the line-of-sight."," as where $\langle k_{\mathrm{b}} T_{\mathrm{e}}\rangle=\int
n_{\mathrm{e}}k_{\mathrm{b}}T_{\mathrm{e}} dl/\int n_{\mathrm{e}} dl
$ is the temperature averaged along the line-of-sight and $\langle(k_{\mathrm{b}} T_{\mathrm{e}})^2\rangle=\int n_{\mathrm{e}}
(k_{\mathrm{b}}T_{\mathrm{e}})^2 dl/\int n_{\mathrm{e}} dl $ is the squared temperature averaged along the line-of-sight."
" As shown by Challinor Lasenby (1998). for 4,7.=5 keV. the second-order effects are very small over the entire intensity spectrum compared with the total thermal SZ ettect."," As shown by Challinor Lasenby (1998), for $k_{\mathrm{b}}
T_{\mathrm{e}}=5$ keV, the second-order effects are very small over the entire intensity spectrum compared with the total thermal SZ effect."
 Therefore. we can use Eq.(8)) to calculate approximately the SZ effect for the simulated cluster.," Therefore, we can use \ref{aver}) ) to calculate approximately the SZ effect for the simulated cluster."
" The SZ intensity at 150. GHz (x=2.64) is approximately given by the term of Eq.(8)) proportional to (A,7). since the SZ relativistic corrections at this frequency are very small compared with the total thermal SZ signal. i.e. The SZ intensity at 217 GHz (x=3.83) is approximately given by the term of Eg.(8)) proportional to (4,7,17). because the SZ relativistic corrections are a dominant contribution at this frequency. i.e. We note that both the temperature variance along the sight. which is given by and the standard temperature deviation. which is given by can be derived for the simulated cluster from SZ observations at frequencies of 150 GHz (x=2.64) and 217 GHz (x=3.83)."," The SZ intensity at 150 GHz (x=2.64) is approximately given by the term of \ref{aver}) ) proportional to $\langle k_{\mathrm{b}}
T_{\mathrm{e}}\rangle$, since the SZ relativistic corrections at this frequency are very small compared with the total thermal SZ signal, i.e. The SZ intensity at 217 GHz (x=3.83) is approximately given by the term of \ref{aver}) ) proportional to $\langle(k_{\mathrm{b}}
T_{\mathrm{e}})^2\rangle$, because the SZ relativistic corrections are a dominant contribution at this frequency, i.e. We note that both the temperature variance along the line-of-sight, which is given by and the standard temperature deviation, which is given by can be derived for the simulated cluster from SZ observations at frequencies of 150 GHz (x=2.64) and 217 GHz (x=3.83)."
 The temperature variance and standard. temperature deviation for a homogeneous gas temperature distribution equal zero and. therefore. both the temperature variance and standard temperature deviation along the line-of-sight are a measure of inhomogeneity of the gas temperature along the line-of-sight.," The temperature variance and standard temperature deviation for a homogeneous gas temperature distribution equal zero and, therefore, both the temperature variance and standard temperature deviation along the line-of-sight are a measure of inhomogeneity of the gas temperature along the line-of-sight."
 In terms of the SZ intensities. the temperature variance along the line-of-sight is given by where Ais and AP are the SZ intensities at frequencies of 1560. GHz (x=2.64) and 217. GHz (x=3.83). respectively.," In terms of the SZ intensities, the temperature variance along the line-of-sight is given by where $\Delta I_{\mathrm{150GHz}}$ and $\Delta I_{\mathrm{217GHz}}$ are the SZ intensities at frequencies of 150 GHz (x=2.64) and 217 GHz (x=3.83), respectively."
 The standard temperature deviation can be derived from SZ intensities at frequencies of 1560 GHz and 217 GHz by using Eq. ¢11)), The standard temperature deviation can be derived from SZ intensities at frequencies of 150 GHz and 217 GHz by using Eq. \ref{eqvar}) )
 and the relation between the temperature variance and standard temperature deviation. var2c7.," and the relation between the temperature variance and standard temperature deviation, $\rm var = \sigma^2$."
 We culeulate the map of the standard temperature deviation along the line-of-sight using Eqs. (999. 103).," We calculate the map of the standard temperature deviation along the line-of-sight using Eqs. \ref{varexpr}) ), \ref{sigmaexpr}) ),"
 and (HE) and the SZ intensity maps at frequencies of 150 GHz and 217 GHz derived in the framework of the Wright formalism in Sect., and \ref{eqvar}) ) and the SZ intensity maps at frequencies of 150 GHz and 217 GHz derived in the framework of the Wright formalism in Sect.
 2., 2.
 The map of the standard temperature deviation along the line-of-sight is shown in Fig., The map of the standard temperature deviation along the line-of-sight is shown in Fig.
 2 (middle left-hand pane, \ref{Fig1} (middle left-hand panel).
 We note that the values of the standard temperature deviation are high in regions of high temperature on the SZ-temperature map. which is shown in Fig.," We note that the values of the standard temperature deviation are high in regions of high temperature on the SZ-temperature map, which is shown in Fig."
 2. (top left-hand panel)., \ref{Fig1} (top left-hand panel).
" The standard temperature deviation map contains high temperature ""arc-Ike"" structures. which are also seen in the temperature map."," The standard temperature deviation map contains high temperature “arc-like” structures, which are also seen in the temperature map."
 As noted by Prokhorov et al. (, As noted by Prokhorov et al. (
"2010b). the origin of these ""arc-like"" structures is the major merger of two galaxy clusters occurring at z=0.8.","2010b), the origin of these “arc-like” structures is the major merger of two galaxy clusters occurring at z=0.8."
 Dubois et al. (, Dubois et al. (
2010) demonstrates that this major merger drives the cluster gas to temperatures twice the virial temperature thanks to violent merger shock waves.,2010) demonstrates that this major merger drives the cluster gas to temperatures twice the virial temperature thanks to violent merger shock waves.
 The middle left-hand xnel in Fig., The middle left-hand panel in Fig.
 2. shows that the standard, \ref{Fig1} shows that the standard
suppression in source counts in ACS fields.,suppression in source counts in ACS fields.
 An enhancement would imply that bias wins over negative feedback in these overdense regions., An enhancement would imply that bias wins over negative feedback in these overdense regions.
" They found the ACS populations to be overdense in two fields, underdense in two field, and equally dense as the GOODS populations in one field."," They found the ACS populations to be overdense in two fields, underdense in two field, and equally dense as the GOODS populations in one field."
" Somewhat surprisingly, they did not find a clear correlation between density of i775 dropouts and the region’s overdensity."," Somewhat surprisingly, they did not find a clear correlation between density of $i_{775}$ dropouts and the region's overdensity."
" In this paper, we use semi-analytic models to study reionization within overdense regions."," In this paper, we use semi-analytic models to study reionization within overdense regions."
 The main aim of this work is to quantify the effects of enhancement in the number of sources and radiative feedback within such regions and explore the possibility whether the galaxy luminosity function in overdense regions can be used as potential probe of feedback and reionization history., The main aim of this work is to quantify the effects of enhancement in the number of sources and radiative feedback within such regions and explore the possibility whether the galaxy luminosity function in overdense regions can be used as potential probe of feedback and reionization history.
 It is known that if ionization feedback is the main contributor to the suppression of star formation in low mass haloes then one can distinguish between early and late reionization histories by constraining the epoch at which feedback-related low-luminosity flattening occurs in the galactic luminosity function., It is known that if ionization feedback is the main contributor to the suppression of star formation in low mass haloes then one can distinguish between early and late reionization histories by constraining the epoch at which feedback-related low-luminosity flattening occurs in the galactic luminosity function.
 The effect of reionization feedback on the high redshift galaxy luminosity function was first demonstrated using semi-analytic models by ?.., The effect of reionization feedback on the high redshift galaxy luminosity function was first demonstrated using semi-analytic models by \citet{2007MNRAS.377..285S}.
 We apply their method to study the luminosity function in overdense regions., We apply their method to study the luminosity function in overdense regions.
 We describe our model for reionization and give details about various parameters and their calibration in $2., We describe our model for reionization and give details about various parameters and their calibration in 2.
 We explain our method of identifying biased regions and outline modification to our reionization model in such regions in $3., We explain our method of identifying biased regions and outline modification to our reionization model in such regions in 3.
 Details of our luminosity function calculation appear in $4., Details of our luminosity function calculation appear in 4.
 We present and discuss our results in 85 and $6., We present and discuss our results in 5 and 6.
" Throughout the paper, we use the best-fit cosmological parameters from the 7-year WMAP data (?)., ie., aflat universe with Q,4,=0.26, Q4=0.73, Ok=0.044 and Qy=0.045, and h=0.713."," Throughout the paper, we use the best-fit cosmological parameters from the 7-year WMAP data \citep{2010arXiv1001.4635L}, i.e., aflat universe with $\Omega_m =
0.26$, $\Omega_{\Lambda} = 0.73$, $\Omega_K = 0.044$ and $\Omega_b = 0.045$, and $h=0.713$."
" The parameters defining the linear dark matter power spectrum are og=0.80, ns=0.96, dns/dink=—0.034."," The parameters defining the linear dark matter power spectrum are $\sigma_8=0.80$, $ n_s=0.96$, ${\rm d} n_s/{\rm d} \ln k=-0.034$."
" In this section, we first summarise the basic features of the semi-analytic model used for studying the globally averaged reionization history."," In this section, we first summarise the basic features of the semi-analytic model used for studying the globally averaged reionization history."
 We then describe in detail the modifications made to this model in order to study reionization in biased regions., We then describe in detail the modifications made to this model in order to study reionization in biased regions.
 Our model for reionization and thermal history of the average IGM is essentially that developed in ? (CF05)., Our model for reionization and thermal history of the average IGM is essentially that developed in \citealt{2005MNRAS.361..577C} (CF05).
 The main features of this model are as follows., The main features of this model are as follows.
" The model accounts for IGM inhomogeneities by adopting a lognormal distribution with the evolution of volume filling factor of ionized hydrogen (H11)) regions Qui(2) being calculated according to the method outlined in ?;; reionization is said to be complete once all the low-density regions (say, with overdensities"," The model accounts for IGM inhomogeneities by adopting a lognormal distribution with the evolution of volume filling factor of ionized hydrogen ) regions $Q_{\rm HII}(z)$ being calculated according to the method outlined in \citet{2000ApJ...530....1M}; ; reionization is said to be complete once all the low-density regions (say, with overdensities"
Observations of nearby galaxies provide convincing evidence for the existence of supermassive black holes.,Observations of nearby galaxies provide convincing evidence for the existence of supermassive black holes.
 Although most such galaxies exhibit little or no nuclear activity. dynamical arguments based. on the observed: stellar. and eas distributions firmly imply the presence of supermassive compact objects in their cores IxXormendy lüchstone 1995. Magorrian 1998. Llo 1998. van der Alarel 1999).," Although most such galaxies exhibit little or no nuclear activity, dynamical arguments based on the observed stellar and gas distributions firmly imply the presence of supermassive compact objects in their cores Kormendy Richstone 1995, Magorrian 1998, Ho 1998, van der Marel 1999)."
 Interestingly. these studies show that virtually all earlv-tvpe galaxies host black holes with masses in the range 107 a few 100AL.," Interestingly, these studies show that virtually all early-type galaxies host black holes with masses in the range $10^8-$ a few $10^9$."
.. In disk galaxies. however. the (rather sparser) observations indicate that the black hole niasses are typically of order 10*AL. SerA* in our own Galaxy. M31. and the maser galaxy NGC 4258). consistent with the black hole masses in galaxies being roughly proportional to the masses of the spheroidal components.," In disk galaxies, however, the (rather sparser) observations indicate that the black hole masses are typically of order $10^7 \Msun$ $^*$ in our own Galaxy, M31, and the maser galaxy NGC 4258), consistent with the black hole masses in galaxies being roughly proportional to the masses of the spheroidal components."
" Although the central black bole masses in early tvpe galaxies are large enough to be the remnants. of OSO phenomena (with luminosities ranging [rom 107"" to 10eres 1) anc giant ellipticals are known to. host radio galaxies and raclio-loucl quasars at high recshilts. in nearby ellipticals. only low-luminosity radio cores are commonly observed (Sadler. Jenkins Ixotanji O89: Wrobel IEleershen. 1991)."," Although the central black hole masses in early type galaxies are large enough to be the remnants of QSO phenomena (with luminosities ranging from $10^{46}$ to $10^{48} \ergps$ ) and giant ellipticals are known to host radio galaxies and radio-loud quasars at high redshifts, in nearby ellipticals, only low-luminosity radio cores are commonly observed (Sadler, Jenkins Kotanji 1989; Wrobel Heershen 1991)."
 Spiral galaxies. in contrast. frequently exhibit. nuclear activitv Sevlert nuclei). with optical/UV emission. strong. power-Iaw X-ray continua (with a typical intrinsic photon index. E— 2) and complex variability observed.," Spiral galaxies, in contrast, frequently exhibit nuclear activity Seyfert nuclei), with optical/UV emission, strong, power-law X-ray continua (with a typical intrinsic photon index, $\Gamma \sim 2$ ) and complex variability observed."
 Such phenomena are usually attributed to standard. thin-clisk accretion. coupled with a hot coronal plasma with high radiative elliciencies.," Such phenomena are usually attributed to standard, thin-disk accretion, coupled with a hot coronal plasma with high radiative efficiencies."
 Postulating that the earlv-type systems are simply πιάνο of fuel to power their activity appears implausible since these galaxies possess extensive hot. gaseous halos," Postulating that the early-type systems are simply `starved' of fuel to power their activity appears implausible since these galaxies possess extensive hot, gaseous halos"
shows the scenario for which the orbit of the outer planet is inclined by 135° to that of the inner.,shows the scenario for which the orbit of the outer planet is inclined by $^{\circ}$ to that of the inner.
 That scenario displays a remarkable lack of stability across the entire 3-0 region of plausible orbits for the outer planet., That scenario displays a remarkable lack of stability across the entire $\sigma$ region of plausible orbits for the outer planet.
 The only scenario in which orbits within the σ uncertainty range for the outer planet display strong stability is shown in the lower-right hand panel of Fig., The only scenario in which orbits within the $\sigma$ uncertainty range for the outer planet display strong stability is shown in the lower-right hand panel of Fig.
" (à =180°), when the orbit is both retrograde and coplanar."," \ref{sixplot} $i=$ $^{\circ}$ ), when the orbit is both retrograde and coplanar."
" In that scenario, almost all orbits for the outer planet not crossing that of the inner are dynamically stable on long timescales."," In that scenario, almost all orbits for the outer planet not crossing that of the inner are dynamically stable on long timescales."
" In that extreme scenario, even some configurations in which the orbits of the two planets cross show significant stability."," In that extreme scenario, even some configurations in which the orbits of the two planets cross show significant stability."
" Those scenarios are the only ones of our entire suite of almost 105,000 test integrations in which orbits of the HU Aqr planets situated within the 1-c error bounds stipulated in Qianetal.(2011) can display any long-term stability."," Those scenarios are the only ones of our entire suite of almost 105,000 test integrations in which orbits of the HU Aqr planets situated within the $\sigma$ error bounds stipulated in \citet{Qian2011} can display any long-term stability."
 We have investigated the stability of the recently discovered planetary system around HU Aqr across the =+3-0 uncertainty ranges for the orbit of the outer planet., We have investigated the stability of the recently discovered planetary system around HU Aqr across the $\pm$ $\sigma$ uncertainty ranges for the orbit of the outer planet.
" In the simplest, coplanar, case, we find that the system is unstable on timescales «5x10? years, except for scenarios in which the mutual separation between the two planets was greater than 5 at the outer planet’s periastron."," In the simplest, coplanar, case, we find that the system is unstable on timescales $<~5\times10^{3}$ years, except for scenarios in which the mutual separation between the two planets was greater than 5 at the outer planet's periastron."
" When the orbit of the outer planet is set such that it reaches periastron between 3 and 5 beyond the orbit of the inner planet, the orbital stability is somewhat greater than scenarios where the minimum separation between the two planets is smaller then 3Rg,, but the planetary system still falls apart on timescales too short to inspire any confidence that the tested orbits represent the true state of the HU Aqr system."," When the orbit of the outer planet is set such that it reaches periastron between 3 and 5 beyond the orbit of the inner planet, the orbital stability is somewhat greater than scenarios where the minimum separation between the two planets is smaller then 3, but the planetary system still falls apart on timescales too short to inspire any confidence that the tested orbits represent the true state of the HU Aqr system."
" We see a steady increase in the median lifetime of the system across much of the tested phase-space as the mutual inclination of the orbits climbs from the coplanar case, through 5° to 15° to 45°, as expected from the resulting reduction in the amount of time the two planets spend in close proximity."," We see a steady increase in the median lifetime of the system across much of the tested phase-space as the mutual inclination of the orbits climbs from the coplanar case, through $^{\circ}$ to $^{\circ}$ to $^{\circ}$, as expected from the resulting reduction in the amount of time the two planets spend in close proximity."
" However, even in the most extreme prograde case investigated (i— 45?), the lifetime of systems within the l-c parameter space did not exceed 10° years."," However, even in the most extreme prograde case investigated $i = 45^{\circ}$ ), the lifetime of systems within the $\sigma$ parameter space did not exceed $10^{5}~$ years."
" This is much shorter than the expected system age, suggesting that we have either caught the system during a period of significant dynamical instability in which the planets are"," This is much shorter than the expected system age, suggesting that we have either caught the system during a period of significant dynamical instability in which the planets are"
is the electron Iuid. vorticity.,is the electron fluid vorticity.
 Phe evolution equation for the electron fluid vorticity (eq. (8))), The evolution equation for the electron fluid vorticity (Eq. \ref{ep_vort}) ))
 shows that the only sources of vorticity are the magnetic field and the photon vorticity., shows that the only sources of vorticity are the magnetic field and the photon vorticity.
 This means that if the vorticities were zero in the initial condition. as is the case with initial zero-vorticity conditions we consider here. none of these quantities can be generated for any scale in the linear regime.," This means that if the vorticities were zero in the initial condition, as is the case with initial zero-vorticity conditions we consider here, none of these quantities can be generated for any scale in the linear regime."
 In particular we can conclude that no magnetic field is generated in linear order for scalar perturbations., In particular we can conclude that no magnetic field is generated in linear order for scalar perturbations.
 Lt should be noted that using Eq. CX8)), It should be noted that using Eq. \ref{boltz_eq}) )
 and Eq. (A9)), and Eq. \ref{vor_sou}) )
 allows us to follow moces at which photons are frec-streaming at any given epoch., allows us to follow modes at which photons are free-streaming at any given epoch.
 Therefore the above conclusion holds for all scales larger than the scales at which electrons ancl protons can be treated as Εις»., Therefore the above conclusion holds for all scales larger than the scales at which electrons and protons can be treated as fluids.
 There are various terms which have to be included in going to second order in perturbation theory., There are various terms which have to be included in going to second order in perturbation theory.
 Second order terms can arise [rom treating metric perturbations to second order (Martinez-CGonzalez. Sanz Silk 1992) or bv including the scconc order terms in the cleetron-photon scattering (Vishniac 1987. Jubas Docdelson 1995. Lu. Scott. Silk 1994).," Second order terms can arise from treating metric perturbations to second order (Martinez-Gonzalez, Sanz Silk 1992) or by including the second order terms in the electron-photon scattering (Vishniac 1987, Jubas Dodelson 1995, Hu, Scott, Silk 1994)."
 Vishniac (1987) adopted the simple procedure of including the spatial dependence of densities to include the second order effects., Vishniac (1987) adopted the simple procedure of including the spatial dependence of densities to include the second order effects.
 Detailec analyses (Jubas Docdelson 1995. Hu. Scott. Silk 1994) showed that Vishniac's procedure gives the most important scconc order effect in the electron-proton scattering for sub-horizon scales.," Detailed analyses (Jubas Dodelson 1995, Hu, Scott, Silk 1994) showed that Vishniac's procedure gives the most important second order effect in the electron-proton scattering for sub-horizon scales."
 Phis allows us to study scales smaller than the horizon at the last scattering surface. 44.+c100Mpe.," This allows us to study scales smaller than the horizon at the last scattering surface, $H^{-1} \simeq 100 \, 
\rm Mpc$."
 At larger scales other second. order elfects from electron-photon scattering zux the second order metrice perturbations might be comparable or dominate., At larger scales other second order effects from electron-photon scattering and the second order metric perturbations might be comparable or dominate.
" We adopt Vishiniac's procedure here and obtain the second order term from treating the spatial dependence of densities Le. in 2. τον and p, in the source term for magnetic fiel generation (Eq. (12)))."," We adopt Vishiniac's procedure here and obtain the second order term from treating the spatial dependence of densities i.e. in $R$, $\tau_{\gamma e}$ and $\rho_e$ in the source term for magnetic field generation (Eq. \ref{source_b}) ))."
 To get estimates of the generated: magnetic field we solve for the dilference in photon and. barvonic bulk velocity in the tight-coupling approximation., To get estimates of the generated magnetic field we solve for the difference in photon and baryonic bulk velocity in the tight-coupling approximation.
 We argue below that the main contribution to the source S(k.5g) for any scale comes from epochs at which the tight-coupling approximation is valid.," We argue below that the main contribution to the source $S({\bf k},\eta)$ for any scale comes from epochs at which the tight-coupling approximation is valid."
 Phe source term is evaluated in the tight-coupling approximation in Appendix A and given by Eq. (AL4))., The source term is evaluated in the tight-coupling approximation in Appendix A and given by Eq. \ref{mag_gen_eq}) ).
 We argue there that for adiabatie evolution. the only term that can source the magnetic field generation is given by Eq. (AL5)).," We argue there that for adiabatic evolution, the only term that can source the magnetic field generation is given by Eq. \ref{mag_gen_eq1}) )."
 This allows us to solve the evolution of the generated magnetic Ποια at any scale: Note that e? is independent. of time., This allows us to solve the evolution of the generated magnetic field at any scale: Note that $a \bar R$ is independent of time.
 Eq. X15)), Eq. \ref{mag_gen_eq1}) )
 can be solved using linear theory evolution of density and. velocity perturbations for each scale from initial conditions at the time at which the scale is super-horizon to the epoch of recombination. Meee [σαι CX21))).," can be solved using linear theory evolution of density and velocity perturbations for each scale from initial conditions at the time at which the scale is super-horizon to the epoch of recombination, $\eta_{\rm rec}$ (Eq. \ref{den_vel_evo}) ))."
 At the epoch of recombination the photons decouple from the barvons ancl therefore the source for magnetic ield generation vanishes., At the epoch of recombination the photons decouple from the baryons and therefore the source for magnetic field generation vanishes.
 We do not attempt an explicit solution here but seek an approximate understanding of the generated magnetic field., We do not attempt an explicit solution here but seek an approximate understanding of the generated magnetic field.
 We first justify our use of the tight-coupling approximation., We first justify our use of the tight-coupling approximation.
 As discussed in Appendix A. for cach scale the ight-coupling approximation is valid for epochs before the Silk damping regime (Iq. CX19))).," As discussed in Appendix A, for each scale the tight-coupling approximation is valid for epochs before the Silk damping regime (Eq. \ref{delta_sol_silk}) ))."
 From Eq. (14)).," From Eq. \ref{mag_field_gen2}) ),"
 the magnetic ield at a given scale K gets contribution from density ancl velocity. perturbations at all scales., the magnetic field at a given scale ${\bf k}$ gets contribution from density and velocity perturbations at all scales.
 Le should however be noted hat if Kk corresponds to à scale at which the density ancl velocity perturbations are in damping regime the source of magnetic ield is also in the damping regime. (, It should however be noted that if ${\bf k}$ corresponds to a scale at which the density and velocity perturbations are in damping regime the source of magnetic field is also in the damping regime. (
More precisely one is interested in the power spectrum of the magnetic field. which. rom Eq. (14)).,"More precisely one is interested in the power spectrum of the magnetic field, which, from Eq. \ref{mag_field_gen2}) ),"
 is a four-point function containing density. fields., is a four-point function containing density fields.
 For magnetic field at scale k. the integrand of the source is x{ΑΚ—K'D. here PG) is the power spectrum of the density field: JkK/|& if both & and A are not in the damping or frec-streaming regime.)," For magnetic field at scale ${\bf k}$, the integrand of the source is $\propto P(k')P(|{\bf k} - {\bf k'}|)$, here $P(k)$ is the power spectrum of the density field; $|{\bf k} - {\bf k'}| \simeq k$ if both $k$ and $k'$ are not in the damping or free-streaming regime.)"
 This means that most of the contribution to the magnetic field.comes from epochs at which the tight-coupling approximation is valid., This means that most of the contribution to the magnetic fieldcomes from epochs at which the tight-coupling approximation is valid.
" More precisely much of the contribution to magnetic field at a scale & comes [rom epochs 7X;ga22,7ni (Eq.i CX20))).", More precisely much of the contribution to magnetic field at a scale $k$ comes from epochs $\eta \la \eta_d \simeq \omega_d^{-1}$ (Eq. \ref{damp_sc}) )).
 For scales that are not in the dampingin] regimeIn at me the upper limit of the integralo in Eq. (14)), For scales that are not in the damping regime at $\eta_{\rm rec}$ the upper limit of the integral in Eq. \ref{mag_field_gen2}) )
 ds tee., is $\eta_{\rm rec}$ .
 For smaller scales. the upper Limit is 2yy.," For smaller scales, the upper limit is $\simeq \eta_d$."
 Having identified some generic features of the magnetic field source terms we can give an order-of-magnitude estimate of the generated. field. from leq. X15))., Having identified some generic features of the magnetic field source terms we can give an order-of-magnitude estimate of the generated field from Eq. \ref{mag_gen_eq1}) ).
 For all scales a reasonable upper limit on the generated magnetic field for the eurrentIv-Davoured ACDM model. at a scale Lobt. is: Llere app=me for scales that are not in Silk damping regime at the epoch of last scattering and ay2ye for smaller scales.," For all scales a reasonable upper limit on the generated magnetic field for the currently-favoured $\rm \Lambda CDM$ model, at a scale $L \simeq k^{-1}$, is: Here $\eta_1 = \eta_{\rm rec}$ for scales that are not in Silk damping regime at the epoch of last scattering and $\eta_1 \simeq \eta_d$ for smaller scales."
 In Eq. (15)), In Eq. \ref{gen_mag_field}) )
 we have used the fact that once the sources of generating magnetic fields vanish. the magnetic field evolves such that 7B remains constant (see e.g. Wasserman 19758).," we have used the fact that once the sources of generating magnetic fields vanish, the magnetic field evolves such that $a^2 B$ remains constant (see e.g. Wasserman 1978)."
 In Eq. (153).," In Eq. \ref{gen_mag_field}) ),"
 e(L)—GP)77. (be matter power spectrum is given e.g. by Bardeen et al., $v_e(L) \simeq (k P(k))^{1/2}$ (the matter power spectrum is given e.g. by Bardeen et al.
" 1986) and 9,91) is the IMS of the density field.", 1986) and $\delta_e(\eta_1)$ is the RMS of the density field.
 One can compare the magnetic fields at different scales by evaluating the sources at the epoch of recombination ne., One can compare the magnetic fields at different scales by evaluating the sources at the epoch of recombination $\eta_{\rm rec}$.
 As seen from Iq. (A21)), As seen from Eq. \ref{den_vel_evo}) )
 the density field is either non-evolving or in the oscillatory phase for much of the period prior to recombination (except for modes QS&S(nv in the matter-clominatecl epoch)., the density field is either non-evolving or in the oscillatory phase for much of the period prior to recombination (except for modes $\eta^{-1} \la k \la \left( {\eta / \sqrt{3}} \right )^{-1}$ in the matter-dominated epoch).
 The velocity field however grows x [ο super-horizon scales in the radiation dominated cra ancl [or AX)V3) Lin the matter dominated era., The velocity field however grows $\propto \eta$ for super-horizon scales in the radiation dominated era and for $k \la (\eta/\sqrt{3})^{-1}$ in the matter dominated era.
 Therefore if Ίσα. (15)), Therefore if Eq. \ref{gen_mag_field}) )
 is evaluated at ος the smallscale magnetic field is," is evaluated at $\eta_{\rm rec}$ , the smallscale magnetic field is"
though (he uncertainties are large.,though the uncertainties are large.
 Detailed modeling of the disk properties awaits additional millimeter aud subamnillimeter observations wilh sufficient angular resolution to separate the disk and envelope components., Detailed modeling of the disk properties awaits additional millimeter and submillimeter observations with sufficient angular resolution to separate the disk and envelope components.
 The similar contimuum flixes of IID 100546. and TW Ilva suggest that the disk in ihe more distant and warmer ID 100546 svstem is substantially more massive (han the disk around TW Ilva., The similar continuum fluxes of HD 100546 and TW Hya suggest that the disk in the more distant and warmer HD 100546 system is substantially more massive than the disk around TW Hya.
 Since — J=1-0 emission is not detected from HD 100546 with comparable sensitivity to the TW IIva detection. it is likely that the abundance reduction is even more extreme.," Since $^+$ J=1-0 emission is not detected from HD 100546 with comparable sensitivity to the TW Hya detection, it is likely that the abundance reduction is even more extreme."
 Perhaps photodissociation of CO in the disk atmosphere is enhanced in the environment of this more massive. hotter star.," Perhaps photodissociation of CO in the disk atmosphere is enhanced in the environment of this more massive, hotter star."
 Grady et al. (, Grady et al. (
1997) 8uggest Chat much of the accreting gas seen in the optical along the line-o[-sight to IID 100546 is associated with stu-grazing planetesimals. such as comets or asteroids.,"1997) suggest that much of the accreting gas seen in the optical along the line-of-sight to HD 100546 is associated with star-grazing planetesimals, such as comets or asteroids."
 This may explain both the lack of observed bouncdary-laver emission. which mieht be expected [rom accretion of mostly gaseous material. and sienatures of bipolar outflow in forbidden line emission (hat is characteristic of an optically (hick. gaseous accretion disk.," This may explain both the lack of observed boundary-layer emission, which might be expected from accretion of mostly gaseous material, and signatures of bipolar outflow in forbidden line emission that is characteristic of an optically thick, gaseous accretion disk."
 This fact. together with the non-detection of emission and the highly processed nature of the dust around ILD 100546. perhaps indicates that the gas is nearly fully depleted in the disk. in which case the planet building process may be well advanced.," This fact, together with the non-detection of $^+$ emission and the highly processed nature of the dust around HD 100546, perhaps indicates that the gas is nearly fully depleted in the disk, in which case the planet building process may be well advanced."
 We have made high resolution observations of the nearby voung stars TW Ilva aud IID 100546 with the ATCA using three antennas equipped at 89 Gllz., We have made high resolution observations of the nearby young stars TW Hya and HD 100546 with the ATCA using three antennas equipped at 89 GHz.
 These observations indicate the promise for future millimeter-wave observations of protoplanetary disks located deep in the southern sky., These observations indicate the promise for future millimeter-wave observations of protoplanetary disks located deep in the southern sky.
 The observations of the TW Iva disk resolve emission [rom the J=10 line. and the measured extent aud brightiess support disk models with hieh depletion. [actors inferred from previous. spatially unresolved. measurements of higher J transitions.," The observations of the TW Hya disk resolve emission from the $^+$ J=1–0 line, and the measured extent and brightness support disk models with high depletion factors inferred from previous, spatially unresolved, measurements of higher J transitions."
 The detailed modeling of the TW Iva disk represents a first step towards probing the gas content of (his important svstem. whose dispersal and chenistry max be similar {ο the early Solar Svstem.," The detailed modeling of the TW Hya disk represents a first step towards probing the gas content of this important system, whose dispersal and chemistry may be similar to the early Solar System."
 A more complete understanding of the details of disk chemistry will require spatially resolved. images with higher sensitivity. eventually sampling species (hat probe all of the relevant chemical families.," A more complete understanding of the details of disk chemistry will require spatially resolved images with higher sensitivity, eventually sampling species that probe all of the relevant chemical families."
 Thermal dust continuum emission from these sources was easily detected and promises to be resolved with observations using longer baselines., Thermal dust continuum emission from these sources was easily detected and promises to be resolved with observations using longer baselines.
 The available ATCA antenna stations will offer baselines up to 1.5 km in length. corresponding to an angular resolution better (han 075 at this frequency. which will be adequate to probe the disk surface density and to start to," The available ATCA antenna stations will offer baselines up to 1.5 km in length, corresponding to an angular resolution better than $0\farcs5$ at this frequency, which will be adequate to probe the disk surface density and to start to"
In order to explain Magellanic Stream celetections. our model favours7z2.8.,"In order to explain Magellanic Stream detections, our model favours$\tau \approx 2.8$ ."
he FOSb;«5 sample aud the histogram of those data.,the $_{i>5}$ sample and the histogram of those data.
 As can be seen. the LF is well described by a yower-law for in(I)οS.25.8.," As can be seen, the LF is well described by a power-law for $m(R)\lesssim 25.8$."
 The lack of £ünt objects requires that a break must exist at Dy3:100 Sih)., The lack of faint objects requires that a break must exist at $D_b\lesssim 100$ km.
 The faint-object slope is highly uncertain. with a l-siema upper-limit of a»~0.6.," The faint-object slope is highly uncertain, with a 1-sigma upper-limit of $\alpha_2\sim0.6$."
 This fi lowever is consistent. with that of the FOsb sample best-fit at the L-sigma level., This fit however is consistent with that of the F08b sample best-fit at the $\sim1$ -sigma level.
 Though tlT inclinations of the detections we present in this survey are inaccurate. they lint that the break iu he LF for the FOsb;.5 sample is not as sharp as the fit would suggest.," Though the inclinations of the detections we present in this survey are inaccurate, they hint that the break in the LF for the $_{i>5}$ sample is not as sharp as the fit would suggest."
MD The detectious from tliW. survey are consistent with the same break diameter (36 kin) that breaks to a much flatter faint object slope similar to the 1-sigiua upper-limit of the fit (a20)., The detections from this survey are consistent with the same break diameter (36 km) that breaks to a much flatter faint object slope similar to the 1-sigma upper-limit of the fit $(\alpha_2\sim0)$.
 The fit to the FOSb;.5 samp7 however. constrains the break to be brightward of a(t)<.26.8.," The fit to the $_{i>5}$ sample however, constrains the break to be brightward of $m(R)\lesssim 26.8$."
 Additional data brightward of> m(R)~2T.5 with accurate incliuations is required. before the break location cau be accurately constrained., Additional data brightward of $m(R)\sim27.5$ with accurate inclinations is required before the break location can be accurately constrained.
 The low-inclination LF is sulliciently described by a power-law., The low-inclination LF is sufficiently described by a power-law.
 The best-fit power-law is au acceptable description of the sample. with parameters (a.125)=(0.59.2LO) and an Darling statistic. P(A>X4)=0.15.," The best-fit power-law is an acceptable description of the sample, with parameters $(\alpha,m_o)=(0.59,24.0)$ and an Anderson-Darling statistic, $P(\Delta>\Delta_{obs})=0.15$."
 Indeed the best-fit of Equation | is found when the break occurs at the bright-eud limit of the data. suggesting there is uo stroug evideuce lor a break in the FOSb;<5 data.," Indeed the best-fit of Equation \ref{eq:Fraser2008} is found when the break occurs at the bright-end limit of the data, suggesting there is no strong evidence for a break in the $_{i<5}$ data."
 Presented iu Figures 12. and 19 are the likelihood contours of the fit to the low-incliuation LF. and the histogram of the observations.," Presented in Figures \ref{fig:amo_lowi} and \ref{fig:diff_lowi} are the likelihood contours of the fit to the low-inclination LF, and the histogram of the observations."
 As can be seen. there is no apparent evidence for a break in the magnitude range of the observatious.," As can be seen, there is no apparent evidence for a break in the magnitude range of the observations."
 To test whether or uot the excited and cokl samples have differeut. LEs. we utilized the Andersou-Darling test described above (see Section L.3)) to determine whether or not the broken LF fit of the FOSb;.5 sample is an adequate description of the FOSb;.5 sample.," To test whether or not the excited and cold samples have different LFs, we utilized the Anderson-Darling test described above (see Section \ref{sec:analysis}) ) to determine whether or not the broken LF fit of the $_{i>5}$ sample is an adequate description of the $_{i<5}$ sample."
 We however. set a»—0 rather than the best-fit value. as this is a more realistic [aint object slope than that of the best-fit of the FOSb;«s sample (a2~—2€). as suggested by our survey.," We however, set $\alpha_2=0$ rather than the best-fit value, as this is a more realistic faint object slope than that of the best-fit of the $_{i>5}$ sample $(\alpha_2\sim-\infty)$, as suggested by our survey."
 This slope is still within the ]-sigima credible region of the fit (see Section 5.2))., This slope is still within the 1-sigma credible region of the fit (see Section \ref{sec:highi_fit}) ).
 We louud that for the FüSb;«5 sample. au LF with (αι.ασ.Ην.logA)=(0.7.0.26.2.22.7) has PCA>No.)=0.L. aud is a sufficient description ol the cold sample LE.," We found that for the $_{i<5}$ sample, an LF with $(\alpha_1,\alpha_2,m_b,\log A) =(0.7,0,26.2,22.7)$ has $P(\Delta>\Delta_{obs})=0.4$, and is a sufficient description of the cold sample LF."
 We therefore couclude that we have uo evideuce that those IXBOs with i>5 have a different LF than those with 7<5., We therefore conclude that we have no evidence that those KBOs with $i>5$ have a different LF than those with $i<5$.
 ? found that the LFs of the cold ancl excited populations were different., \citet{Bernstein2004} found that the LFs of the cold and excited populations were different.
 ? [ouud similar results for their survey alone. auc concluded that the low inclination population LF exhibits a bright-eud slope a0.5—1.5. steeper than ay~0.7 fouud when they considered objects of all 7 in their survey.," \citet{Fuentes2008} found similar results for their survey alone, and concluded that the low inclination population LF exhibits a bright-end slope $\alpha\sim0.8-1.5$, steeper than $\alpha_1\sim0.7$ found when they considered objects of all $i$ in their survey."
 When they considered all available data - the FOsSb sample. but not including the survey presented bere. auc includiug their own seareli - they found that the low-inclinatiou group exhibited a significantly steeper bright-object slope than the high-inclination sample.," When they considered all available data - the F08b sample, but not including the survey presented here, and including their own search - they found that the low-inclination group exhibited a significantly steeper bright-object slope than the high-inclination sample."
 This finding, This finding
thus preventing a evrochronological age determinatiou (Barnes2007)..,thus preventing a gyrochronological age determination \citep{2007ApJ...669.1167B}.
 The svuchronization timescale 19077) for the star is 0.23£0.02 Coy., The synchronization timescale \citep{1977A&A....57..383Z} for the star is $0.23 \pm 0.02$ Gyr.
 We searched within oof the sky position of WASP-30 for common proper motion stars., We searched within of the sky position of WASP-30 for common proper motion stars.
 The V213.6 star USNO-BI.O Os00-O67L908 is 13/009 away and appears to be comoving with WASP-30 (Table 23... bottom panel).," The $V = 13.6$ star USNO-B1.0 0800-0674908 is 09 away and appears to be comoving with WASP-30 (Table \ref{tab:mcmc}, bottom panel)."
 Using 2N[ASS photometry. we coustructed a colom-maguitude diagram.," Using 2MASS photometry, we constructed a colour-magnitude diagram."
 A distance modulus of 8.50+0.05 (500+10 pe) was required to place the comoving star on the nuadu-sequence and sugeests that WASP-30 is ~2.3-Cayr old (Figure 2.. upper pauclj," A distance modulus of $8.50 \pm 0.05$ $500 \pm 10$ pc) was required to place the comoving star on the main-sequence and suggests that WASP-30 is $\sim$ 2.3-Gyr old (Figure \ref{fig:colmag-mr}, upper panel)."
 IHowever. the appareut magnitude and spectral type of WASP-30 suggest a stnaller distance modulus of 7.9τε0.2 (366477 pe).," However, the apparent magnitude and spectral type of WASP-30 suggest a smaller distance modulus of $7.9 \pm 0.2$ $366 \pm 77$ pc)."
 As the comoving star niav be a mere line-of-sight neielibour. this age determination should be treated with caution.," As the comoving star may be a mere line-of-sight neighbour, this age determination should be treated with caution."
 Iu the lower panel of Figure 2.. WASDP-30b is plotted in a dgnasseradius diagram toecther with isochrones of isolated BDs from models with dusty atmospheres (DUSTYO0.Chabrieretal.20000) anc models with ust-free atmospheres (CONDOS.Baraffeetal.2003)..," In the lower panel of Figure \ref{fig:colmag-mr}, WASP-30b is plotted in a mass-radius diagram together with isochrones of isolated BDs from models with dusty atmospheres \citep[DUSTY00,][]{2000ApJ...542..464C} and models with dust-free atmospheres \citep[COND03,][]{2003A&A...402..701B}."
 By an age of 1 Cr. an isolated BD with the mass of WASP-30b is expected to have cooled to ~ 1700 K. aud the transition from dusty L-dwarts to dust-free T-dwarts is expected to take place at == 13001700 Is (Baratteetal.2003)..," By an age of 1 Gyr, an isolated BD with the mass of WASP-30b is expected to have cooled to $\sim$ 1700 K, and the transition from dusty L-dwarfs to dust-free T-dwarfs is expected to take place at = 1300–1700 K \citep{2003A&A...402..701B}."
 ConsideringTig this and the fact that WASP-30b is highly imadiated. it is ikelv that the dusty atimosphere models of Chabrieretal.(2000b) ave more representative.," Considering this and the fact that WASP-30b is highly irradiated, it is likely that the dusty atmosphere models of \citet{2000ApJ...542..464C} are more representative."
 Depending on the ransitiou temperature. WASP-30) may remain a dusty -dxvarf for the lifetime of its host star. or it may at some voit transition to a dust-free T-dwart.," Depending on the transition temperature, WASP-30b may remain a dusty L-dwarf for the lifetime of its host star, or it may at some point transition to a dust-free T-dwarf."
 WASP-30b has a mass of 0.05819+0.00081 aand a radius of 0.0011+0.0022A8., WASP-30b has a mass of $0.05819 \pm 0.00084$ and a radius of $0.0914 \pm 0.0022$.
.. The DUSTY00 nodels predict the radius of an isolated. BBD to be 0.119. 0.101. 0.00E and 0.08 aat ages of. respectively. 0.1. 0.5. 1 and 5 Cyr.," The DUSTY00 models predict the radius of an isolated, BD to be 0.149, 0.104, 0.094 and 0.084 at ages of, respectively, 0.1, 0.5, 1 and 5 Gyr."
 From a smiple linear interpolation of the DUSTYO0 inodel values. the measured radius of WASP-30b aud its 1-0 nucertaity sugeests its age is 2E1 Gyr.," From a simple linear interpolation of the DUSTY00 model values, the measured radius of WASP-30b and its $\sigma$ uncertainty suggests its age is $2 \pm 1$ Gyr."
 The measured radius is mconsistent with the DUSTY00 model for a 0.5-Cvr BD at the 5.7-0 level aud inconsistent with a 5-Civr DD at the lo level. but cousisteut for a αν BD at he 1.2-¢ level.," The measured radius is inconsistent with the DUSTY00 model for a 0.5-Gyr BD at the $\sigma$ level and inconsistent with a 5-Gyr BD at the $\sigma$ level, but consistent for a 1-Gyr BD at the $\sigma$ level."
 This agreciment would be slightly better if the DUSTY00 models took account of radiation., This agreement would be slightly better if the DUSTY00 models took account of irradiation.
" Baraffeetal.(2003). find that the inmadiation of a dust-ree atmosphere. at the level of radiation experieuccd w WASP-30b, results in radii larger by fora yplanct. aud laveer bv for a »plauet."," \citet{2003A&A...402..701B} find that the irradiation of a dust-free atmosphere, at the level of irradiation experienced by WASP-30b, results in radii larger by for a planet, and larger by for a planet."
" However. they note that the evolution of dust-yee atimospheres are more affected by mnmradiation than are dusty atmospheres. and WASP-30b is considerably nore Massive,"," However, they note that the evolution of dust-free atmospheres are more affected by irradiation than are dusty atmospheres, and WASP-30b is considerably more massive."
 WASP-30's lithium abundance favours a young svsteni age of 50.600 Avy. though lithium is a poor age indicator or an F8 star. aud an age of up to 2 Cyr is not ruled out.," WASP-30's lithium abundance favours a young system age of 50–600 Myr, though lithium is a poor age indicator for an F8 star, and an age of up to 2 Gyr is not ruled out."
 The apparent rotational svuchrouisation of the host star aces a lower Bt of 200 Myr (though the svstei may aave been born svuchronised) aud a possible companion star sugeests an older age of 2.3 Car., The apparent rotational synchronisation of the host star places a lower limit of 200 Myr (though the system may have been born synchronised) and a possible companion star suggests an older age of $~$ 2.3 Gyr.
 Given the BD's nieasured radius. the DUSTY00 BD model indicates au age of B41 Cor.," Given the BD's measured radius, the DUSTY00 BD model indicates an age of $2 \pm 1$ Gyr."
 Taken together. we sugeest that au age of 12 Cir is most likely.," Taken together, we suggest that an age of 1–2 Gyr is most likely."
" The discovery of WASP-30b heralds the first ""nieanbieuous observational determunation of the mass-radius relation (AIRR) iu the BD reeime. aud so we lave added BDs to white dwarfs aud neutron stars in the list of «quautunidonminated objects with radius deteruiuations."," The discovery of WASP-30b heralds the first unambiguous observational determination of the mass-radius relation (MRR) in the BD regime, and so we have added BDs to white dwarfs and neutron stars in the list of quantum-dominated objects with radius determinations."
 Chabrier&Baraffe(2000) perfomed the first quantitative theoretical calculation of the MIRR in the substellar aud low-mass star domain. predicting a muni in the ARR at high BD masses (sce Section 3.1 ofthat paper).," \citet{2000ARA&A..38..337C} perfomed the first quantitative theoretical calculation of the MRR in the substellar and low-mass star domain, predicting a minimum in the MRR at high BD masses (see Section 3.1 ofthat paper)."
 The location of WASP-30) in the MRR mun is consistent with the quantitative prediction of Chabrier&Barafte(2000).. thas confirming the theory.," The location of WASP-30b in the MRR minimum is consistent with the quantitative prediction of \citet{2000ARA&A..38..337C}, thus confirming the theory."
 Thus far. we know of two other higlhi-1uass brown dwarts that trausit stars: CoRoT-15b (63Mj:Bouchyctal.2010) aud LIIS 6315 C (63Mj:Johusouetal.2010.Johnson2010.privatecommunication ἐν.," Thus far, we know of two other high-mass brown dwarfs that transit stars: CoRoT-15b \citep[63 \Mjup;][]{2010arXiv1010.0179B} and LHS 6343 C \citep[63 \Mjup;][Johnson 2010, private communication ]{2010arXiv1008.4141J}."
 The radius of CoRoT-15b (1.12!v Πρ) is uncertain aud the age of the system is currently uucoustraimed., The radius of CoRoT-15b $1.12^{+0.30}_{-0.15}$ ) is uncertain and the age of the system is currently unconstrained.
 The faintuess of the host star (V.—16) makes improving this situation difficult., The faintness of the host star $V \sim 16$ ) makes improving this situation difficult.
 LIS 6313 C was found to transit oue member of an M-dsvarf binary svsteni using Ixepler photometry (KIC10002261:e.g.Doruckietal. 2010).., LHS 6343 C was found to transit one member of an M-dwarf binary system using Kepler photometry \citep[KIC 10002261; e.g.][]{2010Sci...327..977B}. .
 It initially seemed that LIS 6315 ο πας larger than predicted for a DD of its mass aud age (Jolhusonetal.2010)., It initially seemed that LHS 6343 C was larger than predicted for a BD of its mass and age \citep{2010arXiv1008.4141J}.
 However. after a re-evaluation of the treatiment of the third liebt in the system. it secus to be consistent (Jolson 2010. πανετο conunuuniication).," However, after a re-evaluation of the treatment of the third light in the system, it seems to be consistent (Johnson 2010, private communication)."
 Chabrieretal.2000} predict that BDs with 05 ddo not burn lithium. those with AL>0.06 yburn essentiallv all their lithium by an age of 0.5 Cor. and those with au intermediate mass (AL=0.055 AM.) mr half them lithium by 0.5 Cir. two-thirds by 1 Cx. and three-quarters by 5 Cyr.," \citet{2000ApJ...542..464C} predict that BDs with $M < 0.05$ do not burn lithium, those with $M > 0.06$ burn essentially all their lithium by an age of 0.5 Gyr, and those with an intermediate mass $M = 0.055$ ) burn half their lithium by 0.5 Gyr, two-thirds by 1 Gyr, and three-quarters by 5 Gyr."
 With a mass of 0.058210008 and an age of 12 Cyr. WASP-30hb is likely to have »nrned nost of its supply of lithinu.," With a mass of $0.0582 \pm 0.0008$ and an age of 1–2 Gyr, WASP-30b is likely to have burned most of its supply of lithium."